{ "1207/1207.5954_arXiv.txt": { "abstract": "{ We have determined the angular diameters of two metal-poor stars, \\object{HD\\,122563} and \\object{\\gmb}, using CHARA and Palomar Testbed Interferometer observations. For the giant star HD\\,122563, we derive an angular diameter $\\theta_{\\rm 3D} = 0.940 \\pm 0.011$ milliarcseconds (mas) using limb-darkening from 3D convection simulations and for the dwarf star Gmb\\,1830 (HD\\,103095) we obtain a 1D limb-darkened angular diameter $\\theta_{\\rm 1D} = 0.679 \\pm 0.007$ mas. Coupling the angular diameters with photometry yields effective temperatures with precisions better than 55 K (\\teff\\ = 4598 \\pmm\\ 41 K and 4818 \\pmm\\ 54 K --- for the giant and the dwarf star, respectively). Including their distances results in very well-determined luminosities and radii ($L = 230 \\pm 6$ \\lsol, $R = 23.9 \\pm 1.9$ \\rsol\\ and $L = 0.213 \\pm 0.002$ \\lsol, $R = 0.664 \\pm 0.015$ \\rsol, respectively). We used the CESAM2k stellar structure and evolution code in order to produce models that fit the observational data. We found values of the mixing-length parameter $\\alpha$ (which describes 1D convection) that depend on the mass of the star. The masses were determined from the models with precisions of $<$3\\% and with the well-measured radii excellent constraints on the surface gravity are obtained ($\\log g = 1.60 \\pm 0.04, 4.59\\pm 0.02$ dex, respectively). The very small errors on both $\\log g$ and \\teff\\ provide stringent constraints for spectroscopic analyses given the sensitivity of abundances to both of these values. The precise determination of \\teff\\ for the two stars brings into question the photometric scales for metal-poor stars.} {} {}{}{} ", "introduction": "Metal-poor stars are some of the oldest stars in the Galaxy and thus reflect the chemical composition of Galactic matter at the early stages of Galactic evolution. The determination of accurate {\\it observed} fundamental properties, and in particular their location in the Hertzsprung-Russell (HR) diagram, is a key requirement if we aim to constrain the {\\it unobservable} properties such as mass, age, and initial helium content by using stellar models. Among the most controversial {\\it observed} parameter is the effective temperature ($T_{\\rm eff}$) which can vary by more than 200 K for metal-poor stars from one method to another (see the PASTEL catalogue, \\citealt{soubiran10}). In particular, local thermodynamic equilibrium (LTE) is usually assumed and non-LTE (NLTE) effects must be included in spectroscopic analyses especially for metal-poor stars where these effects are enhanced \\citep{thev99,andrie10,merle11} and this leads to even more discrepancy between literature values. One solution is to measure the angular diameter and convert this to \\teff\\ to provide a {\\it direct} determination. The large majority of metal-poor stars belong to the halo or the old disk of the Galaxy which means that their apparent magnitude and or angular diameters are extremely small and difficult to measure. However, some instruments, in particular those on the CHARA array \\citep{tenbrummelaar05} are very capable of working at short wavelengths on long baselines to obtain the required angular resolution. Among the most exciting possible targets with CHARA working in the $K$ band are HD 122563 (=~\\object{HR 5270}, \\object{HIP 68594}, m$_V$ = 6.19 mag) and \\gmb\\ (=~\\object{HD\\,103095}, \\object{LHS 44}, \\object{HIP 57939}, m$_V$ = 6.45 mag) whose mean metallicities $[Z/X]_s$\\footnote{$[Z/X] = \\log Z/X_{\\rm star} - \\log Z/X_{\\odot}$ and Z/X$_{\\odot} = 0.0245$, see Sect.~\\ref{sec:models}} are $\\sim$--2.3 dex and --1.3 dex, respectively (see discussion in Sect.~\\ref{sec:fparamsobs}), where $Z$ and $X$ denote the metallicity and hydrogen (absolute) mass fraction in the star and the subscript refers to the observed surface value. HD 122563, a standard example of a very metal-poor field giant \\citep{wallerstein63, wolffram72}, has been extensively studied and presents similarities with metal-poor giants found in globular clusters. \\gmb\\ is a metal-poor halo dwarf star recognized as exhibiting depleted Li \\citep{deli94,king97} when compared to the mean value of halo dwarf stars \\citep{spite93, ryan05}. It is also the nearest halo dwarf and has an excellent parallax measurement. Combining interferometric measurements of these stars with other already measured old moderately metal-poor stars, such as $\\mu$ Cas ([$Z/X$]$_s$~=~-0.5 dex, \\citealt{boyajian08}), offers an excellent opportunity to constrain the \\teff\\ scale of metal-poor stars over a wide range of metallicities with possible implications for \\teff\\ calibrations of globular cluster stars. % In Table \\ref{tab:teffs} we summarize some of the most recent determinations of the atmospheric properties of both targets. Note that HD~122563 and \\gmb\\ have also been defined as benchmark stars for the Gaia mission under the SAM\\footnote{\\url{www.anst.uu.se/ulhei450/GaiaSAM/}} working group. Not only are temperature scales for metal-poor stars controversial, but stellar structure and evolution models often predict higher \\teff\\ than those observed for these stars (see e.g. Fig.~2 of \\citealt{leb00}). The difficulty encountered when trying to match evolutionary tracks to the observational data not only severely inhibits the determination of any fundamental properties but any chance of improving or testing the physics in the models is also limited. Considering the difficulties mentioned above, in this paper we aim to determine accurate fundamental properties of HD\\,122563 and \\gmb\\ based on interferometric observations (Sect.~\\ref{sec:observations}). In Sect.~\\ref{sec:diameters} we present our analysis of the observations to determine the angular diameters of both stars. We then determine the {\\it observed} values of \\teff, luminosity $L$, and radius $R$, and subsequently use stellar models to constrain the {\\it unobservable} properties of mass $M$, initial metal and helium content $Z_{\\rm_i},Y_{\\rm i}$, mixing-length parameter $\\alpha$ and age (Sect.~\\ref{sec:fparams}). We also predict their global asteroseismic properties in order to determine if such observations could further constrain the models. \\begin{table} \\begin{center} \\caption{Most recent photometric and spectroscopic determinations of atmospheric parameters for the target stars {\\label{tab:teffs}}} \\begin{tabular}{lllllllll} \\hline\\hline \\multicolumn{4}{l}{HD\\,122563} & \\multicolumn{4}{l}{\\gmb}\\\\ \\teff&\\logg&[Fe/H]&p/s$^a$& \\teff&\\logg&[Fe/H]&p/s\\\\ (K) & (dex) & (dex) & & (K) & (dex) & (dex) & \\\\ \\hline $4795^b$& \\dotfill & \\dotfill & p & 5129$^b$ & \\dotfill&\\dotfill&p\\\\ 4598$^c$& \\dotfill & \\dotfill & p& 5011$^c$& \\dotfill&\\dotfill&p\\\\ 4572$^d$& \\dotfill & \\dotfill & p&5054$^e$ &\\dotfill&\\dotfill&p\\\\ 4600$^f$& 1.50 & -2.53 & s&5250$^g$&5.00&-1.26&s\\\\ 4570$^h$& 1.10 & -2.42 & s&5070$^i$&4.69&-1.35&s\\\\ \\hline\\hline \\end{tabular} \\end{center} Notes. $^a$p/s = photometric/spectroscopic determination. $^b$\\citet{gonbon09} $^c$\\citet{rammel05} $^d$\\citet{aam99} $^e$\\citet{black98} $^f$NLTE analysis by \\citet{mas08} $^g$\\citet{luck06} $^h$\\citet{mis01} $^i$\\citet{gehren06} \\end{table} ", "conclusions": "We have determined the $T_{\\rm eff}$, $L$, and $R$ of HD 122563 and \\gmb\\ by using $K$ band interferometric measurements (Table~\\ref{tab:observ}) and 3D/1D limb-darkening for the giant/dwarf. We find angular diameters of $\\theta_{\\rm 3D} = 0.940 \\pm 0.011$ mas and $\\theta_{\\rm 1D} = 0.679 \\pm 0.015$ mas for HD\\,122563 and \\gmb, respectively, and these convert into \\teff\\ = 4598 \\pmm\\ 41 K for HD\\,122563 and \\teff\\ = 4818 \\pmm\\ 54 K for \\gmb. These new precision temperatures increase the well-known difficulty of fitting the error boxes of these two metal-poor stars with evolutionary tracks. Using the CESAM2k stellar structure and evolution code we found that we could match models to the data by using values of the mixing length (the parameter $\\alpha$) very different from that of the Sun. We found values of $\\alpha$ = 0.68 and 1.31 for the 0.63 \\msol\\ dwarf star and the 0.86 \\msol\\ giant, respectively. The order of these values seems consistent with recent model analyses \\citep{yildiz06,ker0861cyg}. We found that different equations of state lead to qualitatively but not quantitively different model parameters for the dwarf star but not for the giant. The initial helium content comes out similar to the big-bang value, the deduced masses are low and their ages are high, consistent with expected values for metal-poor halo stars (see Table~\\ref{tab:params}). The masses are determined with a few percent precision and coupling these with the radii yields well-constrained values of \\logg. For the giant star we found \\logg\\ = $1.60 \\pm 0.04$ somewhat higher than the typical values (1.1 - 1.5) adopted by spectroscopic analyses according to the PASTEL catalogue \\citep{soubiran10} and for the dwarf star we obtain \\logg\\ = 4.59 \\pmm\\ 0.02 dex. \\citet{barbuy03} determined the O abundance of HD\\,122563 assuming two different (both justified) values of \\logg, and they concluded that their resulting [O/Fe] = +0.7 abundance seemed most consistent when they adopt the {\\it Hipparcos}\\footnote{We note that with the new Hipparcos parallaxes the deduced \\logg\\~=~1.6.} \\logg\\ = 1.5 and not the value determined from ionization equilibrium of Fe, \\logg\\ = 1.1, a result due possibly to NLTE effects. This work supports their O determination. With both \\logg\\ and \\teff\\ now very precisely known, these provide very important inputs for any spectroscopic analyses, especially for the determination of neutron-capture element abundances which can constrain models of nucleosynthesis. Finally, we have also predicted the asteroseismic signatures \\mlsep\\ and \\numax\\ for the two stars and we showed that determinations of these quantities for the dwarf star are possible using ground-based observations. For the giant, however, we would require very long time series in order to resolve the frequency content of the oscillations, and this would only be possible with space-borne instruments. The asteroseismic data would provide very important constraints because it would allow us to determine the mass with better precision (using the radius from this work), and thus the initial helium abundance." }, "1207/1207.0576_arXiv.txt": { "abstract": "We describe the latest release of AtomDB, version 2.0.2, a database of atomic data and a plasma modeling code with a focus on X-ray astronomy. This release includes several major updates to the fundamental atomic structure and process data held within AtomDB, incorporating new ionization balance data, state-selective recombination data, and updated collisional excitation data for many ions, including the iron L-shell ions from Fe$^{+16}$ to Fe$^{+23}$ and all of the hydrogen- and helium-like sequences. We also describe some of the effects that these changes have on calculated emission and diagnostic line ratios, such as changes in the temperature implied by the He-like G-ratios of up to a factor of 2. ", "introduction": "X-ray spectra from astrophysical sources reveal both the constituent elements of those objects and the physics occurring within them. Successfully extracting this information from spectra requires both a model of the plasma and a large collection of data detailing the various atomic processes occurring within that plasma. The recognition of the importance of the dielectronic recombination (DR) process for plasmas, even at high temperatures \\citep{Burgess1964}, led to huge strides in modeling collisionally ionized, optically-thin astrophysical plasmas. The \\cite{Cox1969} collection was one of the earliest attempts to collect atomic data for these models. Since then there have been a steady stream of refinements to both the collisional plasma models and the atomic data which underpin them. Different evolutions of such models have included those of \\cite{Cox1971}, \\cite{Mewe1972}, \\cite{Landini1972}, \\cite{Raymond1977}, and \\cite{Brickhouse1995}. Both the quantity and quality of relevant atomic data continue to grow: increases in computational power available allow for improved calculations, while advances in experimental methods and equipment allow more accurate measurements of a wider range of quantities. In modern analysis of collisionally ionized X-ray astrophysical spectra there are three widely-used atomic databases: SPEX v2.0 \\citep{1996uxsa.conf..411K}, CHIANTI v7 \\citep{Landi2012} and AtomDB v1.3.1 \\citep{2001ApJ...556L..91S}. Each database has a slightly different focus: CHIANTI's main focus is on the EUV wavelengths for analyzing solar spectra, while SPEX and AtomDB focus on the X-ray ranges. Each of these databases continues to undergo periodic review and is updated as newer data become available; this paper introduces the release of AtomDB v2.0.2, describing the new data and improvements, how they have been implemented and what is planned for future releases. ", "conclusions": "We have presented the latest version of AtomDB, the first major update since 2001. Nearly every single piece of data in the database has been updated, with many new ions added. This is the result of a comprehensive evaluation of the previous data, and assessment of its replacement, addressing many of the known issues in the previous version. New recombination data significantly alter the ionization balance, and therefore the emissivities, of many lines; new collisional data for H- and He-like ions make significant improvements to common temperature diagnostics. The new data are now available on-line at www.atomdb.org, and are also available through spectral fitting packages such as XSPEC \\citep{Arnaud1996}, ISIS \\citep{Houck2000} and Sherpa \\citep{Freeman2001}.\\footnote{After the release of v2.0.0, errors were discovered affecting the radiative recombination continuum and the autoionization rates. This led to the release of v2.0.1 and v2.0.2 of AtomDB. All data in this paper refer to the corrected v2.0.2 data.} There are several improvements already planned for the next release of AtomDB. We will include final state resolved DR rates for the remaining ions in the database for which the data exist. We will largely target non-equilibrium ionization plasmas, with new inner-shell excitation data, as well as fluorescence line data. Work is also ongoing to document and release the APEC collisional ionization code, which will allow users to generate the non-equilibrium higher density plasma models. We wish to thank Paola Testa for discussion involving the line diagnostic issues. The authors gratefully acknowledge funding from NASA ADP grant \\#NNX09AC71G." }, "1207/1207.2059_arXiv.txt": { "abstract": "Bressert et al. recently showed that the surface density distribution of low-mass, young stellar objects (YSOs) in the solar neighbourhood is approximately lognormal. The authors conclude that the star formation process is hierarchical and that only a small fraction of stars form in dense star clusters. Here, we show that the peak and the width of the density distribution is also what follows if all stars form in bound clusters which are not significantly affected by the presence of gas and expand by two-body relaxation. The peak of the surface density distribution is simply obtained from the typical ages (few Myr) and cluster membership number (few hundred) typifying nearby star forming regions. This result depends weakly on initial cluster sizes, provided that they are sufficiently dense (initial half mass radius of $\\lesssim$ 0.3 pc) for dynamical evolution to be important at an age of a few Myr. We conclude that the degeneracy of the YSO surface density distribution complicates its use as a diagnostic of the stellar formation environment. ", "introduction": "\\label{sec:intro} We do not know what fraction of stars form in dense star clusters. Part of the problem is that there is no agreed definitions of what `dense' means and what a star cluster is \\citep*{2010ARA&A..48..431P}. In an attempt to shed light on this situation \\citet{2010MNRAS.409L..54B} studied a sample of young stellar objects (YSOs) in the solar neighbourhood. They calculate the surface density $\\sig$ around each YSO by finding the distance to the 7$^{\\rm th}$ nearest nearest neighbour, $d_7$, such that $\\sig=6/(\\pi d_7^2)$. They find that the distribution of $\\sig$ is roughly lognormal with a peak at about $22\\,\\pc^{-2}$ and a dispersion of 0.85. Because YSOs are very young (of order 1 Myr) they conclude that this distribution reflects the density distribution at the moment of star formation. They concluded that stars form in a broad and smooth spectrum of surface densities and that only a small fraction of the YSOs form in dense clusters. In this paper we argue that the observed peak in the surface density distribution of young stars (at around $20$ stars per square parsec) is an expected outcome from a wide range of initial clustering configurations. In Section \\ref{sec:peak} we argue that for `typical' cluster scales in star forming regions \\citep[$N \\simeq 100$,][]{2003ARA&A..41...57L} and young ages (about 1 Myr) of YSOs such a surface density represents the outcome of dynamical evolution (i.e. two-body relaxation) from a variety of plausible initial conditions. In Section \\ref{sec:model} we flesh out this argument by considering the factors that broaden this distribution (the range of surface densities in a given cluster together with a realistic spectrum of cluster membership number). In Section \\ref{sec:conclusions} we present a discussion and conclude that the observed surface density distribution is exactly what one expects if the majority of stars are born in clusters; the situation is however highly degenerate so that it is not possible to use this distribution to place unique constraints on the initial conditions. ", "conclusions": "\\label{sec:conclusions} We showed that the distribution of surface densities of YSOs presented by \\citet{2010MNRAS.409L..54B}, that was used to argue that only a small fraction of the stars form in dense clusters, can be reproduced by an extremely simple model in which all stars form in dense star clusters. In this model the location of the peak of the distribution is a measure of the age of the cluster population and the typical membership number of young clusters. Our model requires the clusters to form with higher densities than they have at the present day, which is a constraint on the physics of the star formation process which can be tested with future facilities, such as the Atacama Large Millimeter/submillimeter Array (ALMA). There are several physical effects that we have not included, such as the details of the spatial distribution of YSO on the sky and the presence of gas. Although it is beyond the scope of this work to present realistic models that include all these effects, it may be interesting to discuss the omission of gas in our modelling in a bit more detail. This is motivated by the observation that on the plane of the sky YSOs trace the molecular gas they form from and the total gas mass in star forming regions can be 10 or 20 times the total mass in YSOs \\citep[e.g.][]{2012arXiv1204.3552K}, with average gas surface (mass) densities three or four times the surface (mass) density of YSOs \\citep{2009ApJS..184...18G}. A star cluster embedded in an external potential (e.g. gas) will have a higher velocity dispersion, which slows down the two-body relaxation that causes the expansion we discuss here. This is because $\\trh$ depends on the stellar velocity dispersion $\\sigs$ and the stellar density $\\rho$ as $\\trh \\propto \\sigma^3/\\rhos$ \\citep{1971ApJ...164..399S}. To estimate the contribution to $\\sigs$ of the gas we assume two spherically symmetric distributions, with the same functional form for the density profile. Using $\\rs$ and $\\rg$ for the half-mass radii of the stellar and gaseous component, respectively, and $\\ms$ and $\\mg$ for the corresponding total masses, the velocity dispersion of the stars can then be expressed in the stellar and gas parameters as \\citep{1969ApJ...158L.139S} \\begin{equation} \\sigs^2 \\propto \\frac{G\\ms}{\\rs} \\left(1+\\frac{\\mg}{\\ms}\\frac{\\rs^3}{\\rg^3}\\right). \\label{eq:sig2} \\end{equation} The first term on the right-hand side is due to self-gravity of the stars and the second term is the contribution of the gas. If the gas and stellar distribution have the same half-mass radius then we find $\\sigs^2 \\propto (\\ms+\\mg)/\\rs$, which is what is usually assumed for models that consider the removal of natal gas from an embedded cluster. Introducing $\\eta = \\mg/\\ms$ and $\\mu = \\rg/\\rs$ we can compare the relaxation time-scale of a gas embedded system to that of a system containing only stars (i.e. what is assumed here) \\begin{equation} \\trh\\mbox{(stars and gas)} = \\left(1+\\frac{\\eta}{\\mu^3}\\right)^{3/2} \\trh\\mbox{(stars)}. \\end{equation} From this we see that $\\trh$ of a gas embedded system is longer if a significant amount of gas is present ($\\eta\\gtrsim1$) which has a comparable half-mass radius ($\\mu\\simeq1$). But the increase of $\\trh$ depends sensitively on the size of the gaseous system in which the stars evolve (a $\\mu^{-3}$ dependence), so the effect of an additional gas component becomes negligible if $\\mu\\gtrsim 3$. Note that $\\eta/\\mu^3$ is the ratio of the (volume) densities of the stars and the gas within their half-mass radii. This is not the same as the local (volume) density contrast. Estimates of $\\eta$ and (surface) density contrasts are available in some dense embedded clusters but it is not trivial to estimate the value of $\\mu$ from these observations. If $\\mu\\simeq1$ the effect of dynamical expansion can be overestimated by an order of magnitude. On the other hand, if $\\mu\\gtrsim3$, our assumption of gas free cluster evolution is justified. From SPH simulations it was found that gas dominated star forming regions are in fact gas-poor on the scale of the (sink) particles \\citep[i.e. $\\mu>>1$,][]{2012arXiv1205.1677M}. Arguments as to why this might happen can be found in \\citet{2012MNRAS.419..841K}. Similar results were obtained from adaptive mesh refinement (AMR) simulations \\citep{2012MNRAS.420.3264G}. The results of these numerical studies support our assumption of pure stellar dynamical evolution. We conclude that the distribution of surface densities can not be used as evidence that not all stars form in dense clusters. Similarly, the agreement between this model and the observations should not be construed as an argument that {\\it all} stars necessarily form in clusters. \\citet{2010MNRAS.409L..54B} argue that a (roughly) log-normal density distribution is evidence against a scenario in which stars form in distinct `clustered' and `distributed' star formation modes. Their argument is that the distribution would be bi-modal (or multi-modal) if there were distinct modes. However, if we interpret the model in Fig.~\\ref{fig:sigma} as a `clustered' star formation mode, we can add a `distributed` mode with a density of several tens of YSOs per $\\pc^2$ and some dispersion and the total distribution would still be unimodal. Our results support the suggestion of \\citet{2010MNRAS.409L..54B} that a (local) surface density threshold is not a useful tool to separate clusters from field stars, because in our model all stars are in clusters and there is a range of about four orders of magnitude in surface density." }, "1207/1207.5023_arXiv.txt": { "abstract": "We present a multi-wavelength analysis of the very fast X-ray transient \\maxi, which was detected by MAXI/GSC on 2011 November 11. The subsequent exponential decline of the X-ray flux was followed with \\textit{Swift} observations, all of which revealed spectra with low temperatures ($\\sim$100eV) indicating that \\maxi\\ is a new Supersoft Source (SSS). The \\textit{Swift} X-ray spectra near maximum show features around 0.8 keV that we interpret as possible absorption from \\oviii, and emission from O, Fe, and Ne lines. We obtained SAAO and ESO optical spectra of the counterpart early in the outburst and several weeks later. The early spectrum is dominated by strong Balmer and \\hei\\ emission, together with weaker \\heii\\ emission. The later spectrum reveals absorption features that indicate a B1/2IIIe spectral type, and all spectral features are at velocities consistent with the Small Magellanic Cloud. At this distance, it is a luminous SSS ($>10^{37}${\\lum}) but whose brief peak luminosity of $>10^{39}$\\lum in the 2--4 keV band makes it the brightest SSS yet seen at ``hard'' X-rays. We propose that \\maxi\\ is a Be--WD binary, and the first example to possibly enter ULX territory. The brief hard X-ray flash could possibly be a result of the interaction of the ejected nova shell with the B star wind in which the white dwarf (WD) is embedded. This makes \\maxi\\ only the third Be/WD system in the Magellanic Clouds, but it is by far the most luminous. The properties of \\maxi\\ give weight to previous suggestions that SSS in nearby galaxies are associated with early-type stellar systems. ", "introduction": "Supersoft sources (hereafter SSS) are a class of luminous X-ray sources, so-called because of their very soft ($kT\\rm{_{bb}}\\lessapprox$ 100~eV) X-ray spectrum, which can reach luminosities up to ${\\sim}10^{38}$\\lum). Initially discovered in the Magellanic Clouds, as a result of their low interstellar absorption in that direction \\citep{long79}, they have subsequently been found in the Milky Way and in nearby galaxies \\citep{d2003,d2004,kong2004,kong2005}. Key to understanding the nature of the SSS is that their effective blackbody radii are comparable to those of a white dwarf (WD). In fact, many hot WDs and pre-WDs have now been observed as SSS including several recent novae, symbiotic systems, and a planetary nebula, in all of which the WD nature of the compact object is well established (see the SSS catalog from \\cite{2000NewA....5..137G}). This led to \\cite{cbss} establishing what is now considered as the SSS paradigm, wherein they are WDs in close binary systems, accreting at very high rates, which leads to (quasi-)stable thermonuclear burning of hydrogen on the WD surface (see also \\cite{1994ApJ...426..692R,1996ApJ...470L..97H,1996LNP...472....3D}). Furthermore, this requires a massive WD for the most luminous of the class, and so SSS are considered to be strong candidates as progenitors of Type Ia supernovae in the single-degenerate scenario \\citep{d2010,compact}. SSS are therefore very important for enhancing our understanding of close binary interactions and the late stages of stellar evolution. It also appears that, under rare circumstances, SSS can exhibit ``ultraluminous'' (or ULX) levels, which are difficult to explain through a simple application of the steady nuclear burning model \\citep{2008ApJ...674L..73L}. While the \\textit{RXTE} monitoring programs of the Small Magellanic Cloud (SMC) have revealed an extensive population of Be X-ray pulsars (and hence are neutron star (NS) systems, see \\cite{2008ApJS..177..189G}), the detection of their WD cousins is hampered by the luminosity of the Be companion and the lack of continuous soft X-ray monitoring of these regions. To date, only two Be+WD systems (the first, XMMU J052016.0-692505, in the LMC \\citep{Kahabka06}; the second, XMMU J010147.5-715550, in the SMC \\citep{sssinsmc}) have been detected in the Magellanic Clouds. In late 2011, \\maxi, a new X-ray transient, detectable only at soft ($<$4~keV) energies, was discovered with the MAXI/GSC as a very brief ($<$90 minutes) X-ray flare, in the direction of the Wing of the SMC, but with poor ($\\sim0^{\\circ}.4$) location accuracy \\citep{atel3756}. This X-ray flare (denoted XRF 111111A) exhibited a very unusual spectrum in that all the X-ray flux was confined to the lowest GSC energy channel (2--4~keV), inferring a luminosity at these energies of $>6\\times10^{38}$\\lum \\citep{atel3756} at the distance of the SMC. This was substantially higher than any known Galactic/Magellanic Cloud SSS, and approaching the luminosities seen in some extragalactic SSS \\citep{2000NewA....5..137G}. A \\textit{Swift} X-ray and UV/optical observation \\citep{atel3758,atel3759} was performed within 10 hr of the MAXI trigger, revealing that \\maxi\\ was undergoing a supersoft phase of emission ($\\sim2\\times10^{37}$\\lum\\ (0.2--2~keV)), and with an accurate ($\\pm$3.6 arcsec) location of $\\mathrm{R.A.}=01$:59:25.6, $\\mathrm{decl.}=-$74:15:28, J2000.0). The X-ray spectra were compatible with a low-temperature blackbody ($kT_\\mathrm{bb}\\sim$100 eV) that is typical of the SSS. \\textit{Swift}/UVOT UV/optical images plus archival OGLE-IV \\textit{I} band photometry show that the optical counterpart is a bright (13th magnitude in the \\textit{U} band) blue star that increased in brightness by at least $\\sim$0.5 mag at the time of the MAXI/GSC X-ray flash, and declined on the same timescale as the \\textit{Swift}/XRT light curve. Ultraluminous X-ray sources (ULXs) are non-nuclear X-ray sources with $L_X{\\geq}10^{39}$\\lum, whose physical properties are still controversial \\citep{2004ApJS..154..519S}. Short-term variability indicates that they must be accreting compact objects, where the mass transfer rate is close to or exceeding the Eddington Limit. Once strong candidates for the long-sought intermediate-mass black holes (IMBHs), they are now considered likely to be extreme examples of stellar-mass black holes (see e.g., \\cite{ZampieriRoberts09}). It is thus of considerable current interest as to how accreting WDs can even come close to having properties that appear to overlap with the ULX. In this paper, we present a time-resolved, multi-wavelength follow-up of the \\maxi\\ X-ray flare and its subsequent decline over the next few months, using \\textit{Swift XRT/UVOT}, \\textit{Galaxy Evolution Explorer} \\textit{GALEX}, SAAO 1.9~m, ESO NTT, \\textit{Wide-field Infrared Survey Explorer} (\\textit{WISE}), and ATCA radio observations to reveal some of the detailed properties of this remarkable object. ", "conclusions": "By considering its WD-like radius inferred from the \\textit{Swift}/XRT X-ray spectra ($0.001$--$0.01$~$R_\\sun$) and its B1/2IIIe spectral type deduced from the ESO optical spectrum, we believe that \\maxi\\ is a member of the long-sought population of Be--WD binaries, and the first example to possibly enter ULX territory. To account for its properties (very short SSS phase of $\\leq$15 days) we deduce that \\maxi\\ is a heavy ($\\sim$1.35~$M_\\sun$) O--Ne WD which is slowly accreting material from the wind of its early-type companion until it undergoes unstable hydrogen burning on the WD surface, in what is essentially a classical nova explosion. Furthermore, with its location in the SMC, \\maxi\\ is a key example of the extreme behavior (apparently ultra-luminous X-ray flash followed by SSS phase) that a nova explosion can lead to if it is in the appropriate environment. It has pointed us in a new direction for understanding the QSS/SSS in external galaxies that have been proposed as IMBH associated with massive companions. If confirmed, it reveals an entirely new sub-class of SSS, showing once again that WDs are capable of mimicking BHs, and possibly even (some) transient ULXs." }, "1207/1207.5509_arXiv.txt": { "abstract": "In this contribution we present the first census of oxygen in star-forming galaxies in the local universe. We examine three samples of galaxies with metallicities and star formation rates at $z = 0.07, 0.8$ and $2.26$, including the SDSS and DEEP2 surveys. We infer the total mass of oxygen produced and mass of oxygen found in the gas-phase from our local SDSS sample. The star formation history is determined by requiring that galaxies evolve along the relation between stellar mass and star formation rate observed in our three samples. We show that the observed relation between stellar mass and star formation rate for our three samples is consistent with other samples in the literature. The mass-metallicity relation is well established for our three samples and from this we empirically determine the chemical evolution of star-forming galaxies. Thus, we are able to simultaneously constrain the star formation rates and metallicities of galaxies over cosmic time allowing us to estimate the mass of oxygen locked up in stars. Combining this work with independent measurements reported in the literature we conclude that the loss of oxygen from the interstellar medium of local star-forming galaxies is likely to be a ubiquitous process with the oxygen mass loss scaling (almost) linearly with stellar mass. We estimate the total baryonic mass loss and argue that only a small fraction of the baryons inferred from cosmological observations accrete onto galaxies. ", "introduction": "A complete theory of galaxy formation and evolution will have to be able to self-consistently account for, among other physical processes, the star formation and chemical evolution of galaxies. Our understanding of galaxy evolution is rooted in the currently accepted cosmological model in which large-scale structure in the universe traces out the cosmic web of dark matter and growth of the universe is accelerated by dark energy. In this theoretical framework, a hierarchical formation of galaxies is favored in which larger galaxies form as the dark matter halos within which they are embedded merge over time. It remains uncertain what epoch in cosmic history this is the dominant mode of growth. However, recent observations of strong correlations observed between fundamental galaxy parameters (e.g. mass, age, size, luminosity, baryonic content and angular momentum) have lead some to question the stochastic nature of the hierarchical formation scenario \\citep{Disney2008, Nair2010}. One possible resolution is that galaxies and groups of galaxies gather matter early on followed by quiescent, isolated evolution \\citep{Peebles2010}. The evolution of galaxies may be simpler than a hierarchical formation model suggests. A large number of studies have recently revealed that there exists a tight relation between stellar mass and star formation rates (SFRs) out to $z \\sim 2$ \\citep[among others]{Noeske2007a, Salim2007, Daddi2007, Elbaz2007, Pannella2009, Elbaz2011}. We refer to this as the MS relation. All these studies find the slope of the relation to be near unity and a $1\\sigma$ scatter of $\\lesssim0.3$ dex. The relation and its small scatter is taken as evidence that secular processes, such as gas accretion, are the dominant mechanism for star formation with mergers playing a minor role. In particular, \\citet{Noeske2007a} suggest that the presence of an MS relation with constant scatter at several epochs implies that star formation is gradually declining with galaxies spending 67\\% (95\\%) of their star formation lifetime within a factor of $\\sim$2 (4) of their average SFR. Several studies have applied the observational constraints imposed by the MS relation and its evolution to uncover star formation histories of galaxies. \\citet{Noeske2007b} show that their model of ``staged\" galaxy evolution accounts for the observed relation. In their model, less massive galaxies have later onset of initial star formation with longer timescales of exponential decay. Similar models result if star-forming galaxies are assumed to lie on the MS relation at all epochs. Several studies have focused on this simpler approach of continuity of star formation along MS relation. \\citet{Conroy2009b} combine this approach with abundance matching to dark matter halos, concluding that mergers play a minor role in mass growth of galaxies. Using this approach, \\citet{Peng2010} are able to explain the shape and evolution of the observed stellar mass function for star-forming galaxies. \\citet{Papovich2011} apply this technique to understand the gas accretion process at high redshifts. \\citet{Leitner2011} use this technique to show that gas recycling is sufficient to fuel the observed star formation in the local universe and \\citet{Leitner2012} argue that most star-forming galaxies in the local universe formed at $1 \\Delta M_b$. In the top panel of Figure \\ref{fig:db} we plot the total baryonic mass loss as a function of stellar mass for galaxies in the local universe. The black and red curves are estimates adopting an enriched outflow (Equation \\ref{eq:db}) and uniform wind model (Equation \\ref{eq:dbupper}), respectively. In the bottom panel of Figure \\ref{fig:db} we plot the logarithm of $\\delta_b$ which we define as \\begin{equation} \\delta_b = \\frac{\\Delta M_b}{M_g + M_\\ast}. \\end{equation} $\\delta_b$ is the ratio of the total baryonic mass loss to the baryonic mass of the galaxy, where the baryonic mass of the galaxy (gas + stellar). $\\delta_b$ represents the ratio of the mass of gas to be cycled in and out of the galaxy compared to the current baryonic mass of the galaxy. In the currently accepted cosmological model ($\\Lambda$CDM), the universal baryon fraction is precisely determined from the cosmic microwave background, the observed baryon acoustic oscillations and the Hubble constant. The universal baryon and dark matter density revealed by these observations is given by $\\Omega_b = 0.0456 \\pm0.0016$ and $\\Omega_c = 0.227 \\pm 0.014$, respectively \\citep{Komatsu2011}. Studies attempting to account for the baryons find that only a fraction of the expected baryons are observed in the low-redshift universe \\citep{Fukugita1998, Fukugita2004, Nicastro2005, Sommer-Larsen2006, Shull2011}. Only about a tenth of the baryons are found in the stars and gas of galaxies \\citep{Bell2003a}. While the Ly$\\alpha$ forest at low redshifts can account for another $\\sim30\\%$ \\citep{Penton2004, Sembach2004}, the majority of baryons are still missing. Some cosmological simulations favor the warm-hot intergalactic medium (WHIM) as the repository of the missing baryons \\citep[e.g][]{Cen1999, Dave2001, Cen2006, Oppenheimer2011}. However, observations of the hot gas at $10^5 - 10^7$K comprising the WHIM remain tentative and no compelling evidence for the detection of this phase yet exists \\citep{Bregman2007}. The distribution of the WHIM material is unknown and hot halos of massive galaxies are considered as a possible reservoir for substantial fraction of the missing baryons \\citep{Cen2006, Tang2009, Kim2009}, though this remains controversial \\citep{Anderson2010}. An open question is what fraction of the missing baryons were ejected from galaxies through feedback processes and what fraction never accreted in the first place. We can compare our estimates of the total baryonic mass that can be associated with galaxies (baryonic mass plus the baryonic mass loss) with the expected baryon content of galaxies and their halos inferred from cosmological estimates. The stellar-to-halo mass (SHM) relation parameterizes the relationship between stellar mass of galaxies and the dark matter halos in which they reside. \\citet{Moster2010} develop a statistical approach whereby halos and subhalos are populated in an $N$-body simulations with the requirement that the observed stellar mass function be reproduced. They parameterize the SHM as \\begin{equation} \\frac{M_\\ast}{M_h} = 2 \\left(\\frac{M_\\ast}{M_h}\\right)_0 \\left[ \\frac{M_h}{M_1}^{-\\beta} + \\frac{M_h}{M_1}^{-\\gamma} \\right]^{-1}. \\end{equation} Here $M_h$ is the halo mass and $\\left(\\frac{M_\\ast}{M_h}\\right)_0$, $M_1$, $\\beta$ and $\\gamma$ are free parameters. The relation evolves with redshift and the parameters are given by \\begin{eqnarray*} \\mathrm{log} M_1 (z) & =& 1.07 \\cdot(1+z)^{0.019} \\nonumber \\\\ \\left(\\frac{M_\\ast}{M_h}\\right)_0 (z) &=& 0.0282 \\cdot(1+z)^{-0.72} \\nonumber \\\\ \\gamma (z) &=& 0.556 \\cdot (1+z)^{-0.26} \\nonumber \\\\ \\end{eqnarray*} and \\begin{equation} \\beta (z) = 1.06 + 0.17 z. \\nonumber \\end{equation} We can estimate the expected baryon content of galaxies from the universal ratio of baryonic to dark matter given by $f_{bc} = \\Omega_b/\\Omega_c = 0.201 \\pm 0.014$. \\begin{figure} \\begin{center} \\includegraphics[width=\\columnwidth]{f16.eps} \\end{center} \\caption{The dashed cyan curve is the total baryonic (gas + stellar) mass of galaxies in the local universe. The dot-dashed blue line is the inferred baryon content from cosmological fraction and is given by $f_{bc} \\, M_h$. The solid black and red lines are the total baryonic mass plus the baryon mass loss for an enriched wind model (Equation \\ref{eq:db}) and uniform wind model (Equation \\ref{eq:dbupper}), respectively. In our estimate we have adopted the oxygen deficit given by the solid curve in Figure \\ref{fig:do}d.} \\label{fig:mb} \\end{figure} Figure \\ref{fig:mb} demonstrates the missing baryon problem for galaxies. The dashed cyan line is the current total baryonic (gas + stellar) mass content of galaxies plotted as a function of stellar mass. The dot-dashed blue line is the expected total baryonic mass content of galaxies assuming the universal baryonic to dark matter ratio and the SHM relation of \\citet{Moster2010} and is given by $M_b = f_{bc} \\, M_h$. The solid black and red curves are the sum of the baryonic mass and the baryonic mass loss using the two estimates given by Equations \\ref{eq:db} and \\ref{eq:dbupper}. The solid black and red curves can be interpreted as estimates of the baryon content associated with galaxies and are the amount of baryons currently found in local star-forming galaxies plus what was once in the galaxies but has since been cycled out through outflows. The baryon content associated with galaxies is still substantially lower than the inferred content from cosmology. The straightforward interpretation of Figure \\ref{fig:mb} is that the missing baryons were never accreted onto local star-forming galaxies and unless a large reservoir of gas is found in the halos of star-forming galaxies it is likely that the missing baryons reside in the IGM. Using observationally motivated constraints for the mass and radii of hot halos, \\citet{Anderson2010} come to a similar conclusion. \\begin{figure} \\begin{center} \\includegraphics[width=\\columnwidth]{f17.eps} \\end{center} \\caption{The effective mass loading parameter given in Equation \\ref{eq:ml} plotted against stellar mass. The red and black curves give the effective mass loading factor for our upper and lower limit estimates of the total baryon mass loss, respectively.} \\label{fig:ml} \\end{figure} From our estimate of the baryon mass loss we can derive an effective mass loading factor, $\\eta$, which is given by \\begin{equation} \\eta = \\frac{\\Delta M_b}{M_\\ast/(1-R)}. \\label{eq:ml} \\end{equation} The quantity $\\Delta M_b$ is the total amount of baryonic mass loss and $M_\\ast/(1-R)$ represents the \\emph{total} amount of star formation. In the literature the instantaneous mass loading factor is defined as $\\eta_i = \\dot{M}_w/\\dot{M_\\ast}$. The instantaneous mass loading factor is the mass loss rate divided by the SFR. The effective mass loading as we have defined it in Equation \\ref{eq:ml} is the instantaneous mass loading factor, $\\eta_i$, averaged over the star formation history. In Figure \\ref{fig:ml} we plot the effective mass loading factor as a function of stellar mass. Here, we have adopted $R=0.5$. Assuming a smaller value for $R$ would lower the estimate. Even adopting our upper limit estimates of baryon mass loss, the effective mass loading for star-forming galaxies in the local universe is $<1$. \\subsection{Future Prospects} A complete and self-consistent theory of galaxy evolution will require a detailed account of galaxy growth and chemical evolution along with physical mechanisms governing these processes. In this contribution we present empirical models attempting to self-consistently integrate chemical evolution and galaxy growth. Our self-consistent approach is complementary to cosmological simulations and semi-analytical models and the self-consistent census approach presented here should be compared with those models. Both theoretical and observational advances are crucial to constraining the model results developed in this study. The method of analysis used in this study provides useful tests for consistency of a diverse set of observations and theories related to chemical properties of galaxies. Here we address future prospects for improvement. One of the greatest outstanding astrophysical problems is the large uncertainty in the absolute calibration of the nebular abundances scale. High S/N observations of large sample of HII regions covering a broad range of physical parameters will be extremely important in statistically establishing and testing diagnostics. Such observations will be crucial in developing empirical calibrations relying on recombination lines which are thought to be less susceptible to effects of temperature variations associated with auroral lines \\citep{Esteban2002, Peimbert2005, Bresolin2007}. A complementary approach will be to use the recently developed wide-field integral field spectrographs to observe nearby HII regions in order to understand the discrepancy between empirically and theoretically calibrated strong-line methods. A well calibrated diagnostic applied to large data sets investigating the MZ relation at higher redshifts and to lower stellar masses will provide important constraints for our understanding of chemical evolution and our census of oxygen in the stellar and gas-phase. Theoretical models incorporating stellar rotation and mass loss into computations of nucleosynthesis for stars at all masses are yet to be developed \\citep{Romano2010}. Current models of stellar nucleosynthesis are not yet able to explain the full diversity of chemical abundance patterns observed. While the oxygen yields are much better constrained by models owing to its primary origin, there is still a factor of $\\sim2$ discrepancy. A single, self-consistent model of nucleosynthesis able to reproduce the abundance ratios observed in the Milky Way and other nearby galaxies will likely alleviate some of the discrepancy, thus constraining the total oxygen production in star-forming galaxies. Part of the uncertainty in chemical evolution models rests on the assumption of a particular IMF \\citep{Romano2005}. Adopting a constant IMF is problematic for galaxies, where even a constant stellar IMF, could lead to variations in the integrated galactic IMF \\citep{Kroupa2003} and studies of the IMF from integrated measurements of star-forming galaxies indirectly indicate variations \\citep{Hoversten2008, Lee2009, Meurer2009}. A varying IMF would have important implications for chemical evolution models \\citep[e.g.][]{Romano2005, Koppen2007, Calura2010}. Despite the mounting evidence, no direct evidence of IMF variations in star-forming galaxies is currently available leading to a lack of consensus on the constancy of the IMF. Distinguishing whether a universal IMF or IGIMF formulation provides a better description of large scale star formation will be important for chemical evolution studies of galaxies. Large statistical studies of the integrated properties of galaxies observed over cosmic time, in particular observations of the FUV and FIR properties which are now becoming available, will be useful in establishing any systematic variations of the IMF. In our study, tighter constraints on the IMF will alleviate much of the uncertainty in the total amount of oxygen produced and the return fraction. Total gas masses are required to measure the absolute content of metals within galaxies from observed relative abundances. In the local universe, large surveys of atomic and molecular gas have been conducted \\citep{Helfer2003, Walter2008, Leroy2009, Saintonge2011}. Molecular gas at higher redshifts has also been detected in star-forming galaxies \\citep{Daddi2008, Daddi2010, Tacconi2010}. The next generation of radio and sub-mm observatories such as the Atacama Large Millimeter/submillimeter Array and the Square Kilometer Array will revolutionize the study of cold gas in the universe allowing us to probe larger samples to far greater depths. Currently, we are only able to estimate the total oxygen in the gas-phase of local star-forming galaxies. Measurements of cold gas in nearby and distant star-forming galaxies taken together with a well calibrated gas-phase metallicity diagnostic will allow us to track the total mass, not just relative abundance, of oxygen in the gas-phase of star-forming galaxies over cosmic time. The structure, content and chemical composition of outflowing gas is highly uncertain. Understanding the physical properties of galactic winds is crucial for estimating the total mass of gas outflowing from galaxies. Observations reveal that galaxy scale outflows have a complicated multiphase structure \\citep[see][]{Veilleux2005} and observations suggest that structure and composition vary in each of the phases \\citep[e.g.][]{Tripp2011}. Multi-wavelength observations are required. The cold gas component can be observed in absorption lines of neutral and low ionization state metals and the current generation of integral field spectrographs on 8-10m class telescopes and wide-field integral field spectrographs on 4m class telescopes will provide census of galactic scale outflows in local galaxies and presence of galactic winds in the distant universe. Hotter phases can be observed using space-based UV spectrometers such as the \\emph{Cosmic Origins Spectrograph} (COS) onboard Hubble. A crucial but inaccessible phase is the so-called wind fluid which drives the stellar winds. A hard x-ray telescope with high spatial resolution and sensitivity is required to observe this phase. For the foreseeable future, astronomers will likely have to rely on detailed theoretical models and indirect observations to understand the wind fluid. COS will also provide important insight into the circum-galactic and warm-hot ionized mediums, both thought to be important repositories of baryons. Studies of these regions will begin to reveal the baryonic content of the hot haloes and intergalactic medium surrounding galaxies. Observations of these regions along with a census of baryons associated with galaxies will be crucial to potentially identifying the large fraction of missing baryons in the local universe and resolving this long-standing problem \\citep{Bregman2007}. A benchmark for these studies is to resolve the oxygen deficit in star-forming galaxies presented here. In this study we have applied the best theoretical and observational constraints available in undertaking a census of oxygen in star-forming galaxies. We are unable to draw any strong quantitative conclusions from the models developed owing to the large uncertainties associated with both our adopted model parameters and observational inputs. Nonetheless, this study represents one of the first attempts to self-consistently account for the stellar mass growth and chemical evolution of galaxies. As theoretical and observational advances allow for ever greater constraints, we hope the approach taken in this study will prove to be an important ingredient in testing and developing a fully self-consistent theory of galaxy evolution. In this study, we have consistently incorporated chemical evolution in the framework of stellar mass growth in order to conduct a census of oxygen in local star-forming galaxies. We are able to estimate the total oxygen production from total amount of star formation inferred from the current stellar mass. The mass of oxygen in the gas-phase is constrained by the observed MZ relation in the local universe and the relation between gas fraction and stellar mass. The most difficult to constrain observationally is the amount of oxygen locked up stars. Our empirical models self-consistently incorporate stellar mass growth and chemical evolution, thus allowing us to track the metallicity of the gas from which stars are formed over cosmological timescales and giving us empirical constraints on the oxygen mass locked up in stars. The main results of this study are given below. \\begin{enumerate} \\item{We conduct a census of oxygen and show that the amount of oxygen in the stellar and gas phase of galaxies does not fully account for the total amount of oxygen produced. We conclude that the most straight-forward interpretation of the oxygen deficit is that oxygen has been expelled from galaxies by outflows. Our results establish the need for ejective feedback in normal star-forming galaxies.} \\item{We compare our oxygen deficit with the observed lower limit of oxygen found in the CGM of star-forming galaxies and conclude that oxygen mass loss is a ubiquitous process in star-forming galaxies. Furthermore, the oxygen deficit in our preferred model scales with stellar mass and we predict that either more massive galaxies should be found to contain a greater mass of oxygen in their halos or that oxygen escapes the galaxy potential well altogether and therefore massive galaxies contribute for IGM enrichement.} \\item{We estimate the total amount of mass lost from the ISM of star-forming galaxies and find that it is a small fraction of the total baryon content expected from the cosmological baryon density. We conclude that only a small fraction of the total baryons in the universe ever cycled through star-forming galaxies.} \\end{enumerate} Our empirical model provides an important test of self-consistency for many physical processes governing galaxy evolution. Future theoretical and observational advances will provide ever increasing constraints on the census of oxygen in star-forming galaxies and our models provide important benchmarks with which to compare theory and observation." }, "1207/1207.0606_arXiv.txt": { "abstract": "The pseudoscalar-photon mixing in presence of large scale magnetic field induces polarization in light from distant cosmological sources. We study the effect of these pseudoscalars or axion like particles (ALPs) on Cosmic Microwave Background Radiation (CMBR) and constrain the product of mixing strength $g_{\\phi}$ times background magnetic field $B$. The background magnetic field has been assumed to be primordial and we assume large scale correlations with the correlation length of 1 Mpc. We use WMAP seven year foreground reduced polarization and temperature data to constrain pseudoscalar-photon mixing parameter. We look for different mass limits of the pseudoscalars and find $g_{\\phi}B\\le 1.6\\times10^{-13} GeV^{-1} nG$ with ALPs of mass $10^{-10} eV$ and $g_{\\phi}B\\le3.4\\times10^{-15} GeV^{-1} nG$ for ultra light ALPs of mass $10^{-15} eV$. ", "introduction": "\\label{one} The pseudoscalar-photon mixing and its effects on distant cosmological sources have been studied in literature \\cite{Harari:1992,Mohanty:1993,Das:2001,Kar:2002,Kar:2002cqg,Csaki:2002,Csaki:2002prl,Grossman:2002,Jain:2002vx, Song:2006,Mirizzi:2005,Raffelt:2008,Gnedin:2007,Mirizzi:2007,Finelli:2009,Agarwal:2012,Ostman:2005,Lai:2006, Hooper:2007,Hochmuth:2007,Chelouche:2009}. These hypothetical axion like particles (ALPs), arise naturally as pseudo-Goldstone bosons in theories with spontaneously broken global symmetries \\cite{Peccei:1977prl,Peccei:1977prd,McKay:1977,Weinberg:1978,Wilczek:1978, McKay:1979,Kim:1979,Dine:1981,Kim:1987}. ALPs have an interaction vertex with two photons and hence in an external magnetic field ALPs can convert into a photon and vice versa \\cite{Clarke:1982,Sikivie:1983,Sikivie:1985,Sikivie:1988,Maiani:1986,Raffelt:1988,Carlson:1994,Bradley:2003, Das:2004qka,Das:2004ee,Ganguly:2006,Ganguly:2009}. Although this effect is very small, it becomes significant at cosmological scales and leads to many interesting signatures on electromagnetic radiation. This pseudoscalar-photon mixing phenomena causes changes in intensity as well as polarization in radiation from distant sources \\cite{Harari:1992,Mohanty:1993,Das:2001,Kar:2002,Kar:2002cqg,Csaki:2002,Csaki:2002prl,Grossman:2002,Jain:2002vx, Song:2006,Mirizzi:2005,Raffelt:2008,Gnedin:2007,Mirizzi:2007,Finelli:2009,Agarwal:2012,Ostman:2005,Lai:2006, Hooper:2007,Hochmuth:2007,Chelouche:2009}. The contribution of this effect has been investigated for CMBR \\cite{Mirizzi:2005,Agarwal:2008ac}, radio \\cite{Harari:1992,Jain:2002vx,Jain:1998r,Ralston:2004} and optical \\cite{Agarwal:2012,Agarwal:2011,Hutsemekers:1998,Hutsemekers:2001fv,Hutsemekers:2005iz,Hutsemekers:2008} sources. Various experiments are looking for these pseudoscalars and providing limit on the coupling constant $g_{\\phi}$ and their masses $m_{\\phi}$\\cite{Mohanty:1993,Raffelt:2008,Dicus:1978,Dearborn:1986,Raffelt:1987,Raffelt:1988prl,Turner:1988, Janka:1996,Keil:1997,Brockway:1996,Grifols:1996,Raffelt:1999,Rosenberg:2000,Horns:2012,Zioutas:2005, Lamoreaux:2006,Yao:2006,Jaeckel:2007,Andriamonje:2007,Robilliard:2007,Zavattini:2008,Rubbia:2008}. In the present paper we study the effect of pseudoscalar-photon mixing on CMBR multipoles. We show, using WMAP observations that this leads to a new constraint on the product of magnetic field $B$ and the pseudoscalar-photon coupling $g_{\\phi}$. We consider the background as a large number of correlated magnetic field domains and do a complete $3D$-simulation to calculate the Stokes parameters for CMBR. The origin of background magnetic field is considered as primordial \\cite{Subramanian:2003sh,Seshadri:2005aa, Seshadri:2009sy,Jedamzik:1998,Subramanian:1998} and we assume a smooth variation of the magnetic field over the scale of 1 Mpc. The magnetic field correlations are assumed to obey a power law with spectral index $n_B$. The details for the background magnetic field model are discussed in Sec.\\ref{sc:PMF}. As we do simulation over a very large distances (redshift 1000), we choose domain size around 16 Mpc. The strength of magnetic field in each domain is assumed to be order of nG \\cite{Csaki:2002prl,Grossman:2002,Mirizzi:2005}. The initial pseudoscalar density is assumed to be zero or negligible as compared to photon density as assumed by most authors \\cite{Das:2004ee,Agarwal:2008ac,Agarwal:2012,Agarwal:2011}. We made this assumption as the pseudoscalars are likely to decouple from cosmic plasma at very early times. After pseudoscalar decoupling, photon density would be enhanced by many processes such as QCD phase transition, $e^- e^+$ annihilation etc. It may not even be in equilibrium after inflation. Hence, it is reasonable to assume pseudoscalar density as negligible as compared to photon density. We compare our result with the WMAP 7-year data and constrain the coupling parameter $g_{\\phi}$ times $B$. The limit presented in the paper is bound to certain assumptions. We list all of them as follows:\\\\ 1)The background magnetic field follows a simple cosmological evolution.\\\\ 2)We have assumed a definite value for the spectral index $n_{B}=-2.37$ which correspond to the best fit of matter and CMBR power spectrum\\cite{Yamazaki:2010nf}. However we also determine its dependence on $n_{B}$. \\\\ 3)CMBR is assumed to be unpolarized initially and the initial density for pseudoscalars is zero.\\\\ The paper is organized as follows. In Sec.\\ref{sc:mixing} we briefly review the pseudoscalar-photon mixing in presence of plasma and uniform magnetic field in a flat expanding universe. In Sec.\\ref{sc:PMF} we model the background magnetic field, which is correlated in real space and discuss the numerical method for generating the $3D$ magnetic field. In Sec.\\ref{sc:sim} we present our simulation result and compare with the WMAP observations. Finally, in Sec.\\ref{sc:dis} we conclude and compare our results with available literatures. ", "conclusions": "\\label{sc:dis} We have done full $3D$ simulation of pseudoscalar-photon mixing for CMBR at a very high resolution over the full sky. A comparison with WMAP observation results in a new and more stringent limit on the factor $g_{\\phi}B$. It depends on the pseudoscalar mass and we simulate the limit on the factor $g_{\\phi}B$ for two different masses of pseudoscalars. Recently\\cite{Horns:2012}, a bound on factor $g_{\\phi}B$ as $g_{\\phi}B\\le10^{-11}GeV^{-1}nG$ has been derived from ultraviolet photon polarization emerging from active galactic nuclei. Here the derived limit corresponds to ultra light ALPs($m_{\\phi}\\le10^{-15}eV$). In Ref.\\cite{Mirizzi:2005,Mirizzi:2009} the limits on $g_{\\phi}B$ has been studied through CMBR spectral distortion, giving $g_{\\phi}B\\le10^{-13}\\sim 10^{-11}GeV^{-1}nG$ for ALPs masses between $10^{-15}eV$ and $10^{-4}eV$. The pseudoscalar-photon mixing may also contribute to the dimming of Type Ia supernovae\\cite{Csaki:2002,Csaki:2002prl,Avgoustidis:2010}. The phenomenon fixes $g_{\\phi}B$ to $~10^{-11}GeV^{-1}nG$ for a axion of mass $10^{-16}eV$\\cite{Csaki:2002prl}. We may constrain $g_{\\phi}$ form our bound on $g_{\\phi}B$. However the constrain on $g_{\\phi}$ is subject to uncertainties in the background magnetic field. Assuming the background magnetic field $B_0$ as $1nG$, our results bound $g_{\\phi}\\le1.6\\times10^{-13}GeV^{-1}$ and $g_{\\phi}\\le3.4\\times10^{-15}GeV^{-1}$ for the ALPs of $10^{-10}eV$ and $10^{-15}eV$ respectively. Our limits can be compared with the direct experimental limits from SN1987A , which is $g_{\\phi}\\le10^{-11}GeV$\\cite{Brockway:1996} and $g_{\\phi}\\le3\\times10^{-12}GeV$\\cite{Grifols:1996} for very light ALPs ($\\le10^{-9}eV$). We also recall the results from CAST\\cite{Andriamonje:2007,Zioutas:2005}, $g_{\\phi}\\le8.8\\times10^{-11}GeV$ for the ALPs of $0.02eV$, which of course is not for the ultralight ALPs and can not be directly compared with our results. We conclude that the CMBR multipole anisotropy imposes a stringent constraint on the pseudoscalar-photon coupling. We have obtained the lowest value of $g_{\\phi}B$ as compared to available literatures." }, "1207/1207.4213_arXiv.txt": { "abstract": "The initial conditions of massive star and star cluster formation are expected to be cold, dense and high column density regions of the interstellar medium, which can reveal themselves via near, mid and even far-infrared absorption as Infrared Dark Clouds (IRDCs). Elucidating the dynamical state of IRDCs thus constrains theoretical models of these complex processes. In particular, it is important to assess whether IRDCs have reached virial equilibrium, where the internal pressure balances that due to the self-gravitating weight of the cloud plus the pressure of the external environmental. We study this question for the filamentary IRDC G035.39-00.33 by deriving mass from combined NIR \\& MIR extinction maps and velocity dispersion from $\\ceto$ (1-0) \\& (2-1) line emission. In contrast to our previous moderately super-virial results based on $\\cothree$ emission and MIR-only extinction mapping, with improved mass measurements we now find that the filament is consistent with being in virial equilibrium, at least in its central parsec-wide region where $\\sim 1000\\:M_\\odot$ snakes along several parsecs. This equilbrium state does not require large-scale net support or confinement by magnetic fields. ", "introduction": "Identified by obscuration of the mid-infrared (MIR) (i.e. $\\sim 10\\:{\\rm \\mu m}$) Galactic background, Infrared Dark Clouds (IRDCs) are likely to be representative of the initial conditions of massive star and star cluster formation, since their high mass surface densities ($\\Sigma \\gtrsim 0.1\\gcc$) and densities ($n_{\\rm H} \\gtrsim 10^{4} {\\rm cm}^{3}$) are similar to regions with such star formation activity (e.g. Teyssier et al. 2002; Rathborne et al. 2006; Tan 2007; Butler \\& Tan 2009, hereafter BT09; Zhang et al. 2009; Ragan et al. 2009; Butler \\& Tan 2012, hereafter BT12). The kinematics of IRDCs can be measured via their molecular line emission to determine if they are gravitationally bound and/or in virial equilibrium. Hernandez \\& Tan (2011, hereafter HT11) used $\\thco$(1-0) emission from the Galactic Ring Survey (Jackson et al. 2006) to measure velocity dispersions in two filamentary IRDCs (F \\& H in the BT09 sample). $\\Sigma$ was estimated by averaging two methods: (1) MIR extinction (MIREX) mapping (BT09), given an assumed MIR opacity per gas mass and a foreground correction based on an analytic model of Galactic hot dust emission; (2) $\\thco$ emission, given an assumed abundance of this isotopologue. Cloud mass was then calculated assuming near kinematic distances to the IRDCs. A filamentary virial analysis following Fiege \\& Pudritz (2000, hereafter FP00) was performed in several orthogonal strips along the filaments. HT11 concluded surface pressure terms are dynamically important, suggesting that the filaments have not yet reached virial equilibrium. Here we revisit this analysis for one of the filaments, G035.30-00.33 (H in the BT09/BT12 sample), which is 2.9~kpc distant with a few thousand solar masses of material spread over about 4~pc (projected) at its northern end. Although there are several 24~$\\rm \\mu m$ sources seen towards the filament (Carey et al. 2009), which are likely to be embedded protostars, most of the region appears MIR dark and starless. For our kinematic measurements, we use IRAM 30m observations of $\\ceto$ (1-0) and (2-1). The results of these and other molecular line observations are being presented in a series of papers on the formation and evolution of this filamentary IRDC. Paper I (Jim\\'enez-Serra et al. 2010) presented maps of SiO, CO, $\\thco$ and $\\ceto$. Widespread SiO emission was observed, perhaps suggesting the presence of large-scale shocks that may have been involved in forming the filament. Paper II (Hernandez et al. 2011) compared the $\\ceto$ and BT12 MIREX maps, showing CO suffers widespread gas phase depletion in the IRDC by factors of up to $f_D^\\prime \\simeq 5$, i.e. just 1 in 5 CO molecules remain in the gas phase. Paper III (this {\\it Letter}) performs a filamentary virial analysis of the IRDC. Paper IV (Henshaw et al. 2012) studies the detailed kinematics of the filament and its surroundings, including analysis of the dense gas tracer $\\rm N_2H^+$. ", "conclusions": "\\label{S:discussion} There is some evidence that the IRDC H filament may have formed recently from converging flows of molecular gas. J\\'imenez-Serra et al. (2010, Paper I) reported widespread, more than parsec-scale SiO emission from the filament, which may have resulted from large-scale shocks with speeds of at least several km/s. Henshaw et al. (2012, Paper IV) studied the kinematics of the region and found evidence for the formation of the main filament via the merging flow of surrounding filaments observed in $\\ceto$ (with typical densities of $n_{\\rm H,flow}\\sim 10^3\\:{\\rm cm^{-3}}$). They find that the line of sight relative velocity between each component of the merging flow and the main filament is $v_{\\rm flow}\\sim 3\\:{\\rm km\\:s^{-1}}$. If the main filament has formed relatively recently, is this consistent with its observed state of near virial equilibrium? It should take at least a signal crossing time, $t_{s,f} = 2R_f/\\sigma_f$ for a region of the filament to settle into virial equilibrium. These times are $\\sim 0.8$~Myr for the Inner Filament (Table~\\ref{tab:1}). To form the filament from two converging flows takes a time $t_{{\\rm form},f}= 0.236 (m_{f}/100 M_\\odot\\:{\\rm pc}^{-1})/[(R_f/{\\rm pc})(v_{\\rm flow}/3\\:{\\rm km\\:s^{-1}}) (n_{{\\rm H,flow}}/10^3\\:{\\rm cm^{-3}})]$~Myr. Applying this to the average properties of the Inner Filament ($R_f=0.432$~pc, $m_f = 259\\:M_\\odot$), yields $t_{{\\rm form},f}= 1.4$~Myr (see also Paper IV). Thus, even in the scenario of a recently formed filament, enough time should have elapsed for it to have settled into virial equilibrium. The widespread CO depletion observed in Paper II and here (Figure 1e), also constrains the age of the filament. The CO depletion time is $t_{D,f}=0.16 (n_{{\\rm H},f}/10^4{\\rm cm^{-3}})^{-1}$~Myr, assuming a sticking probability of order unity (Tielens \\& Allamandola 1987). Table~\\ref{tab:1} lists $t_{D,f}$ for the different strips. For the Inner Filament, these are relatively short, $\\sim0.1$~Myr, which thus provides a lower limit for its age. The implications of virial equilibrium of self-gravitating IRDC filaments are profound. If this result for G035.39-00.33 applies more generally to IRDCs, then it indicates that the {\\it initial conditions} for star formation, including the cores that form massive stars and the clumps that form star clusters, are created from environments where approximate pressure equilibrium has been established. The value of this pressure is set by the self-gravitating weight of the larger scale cloud, i.e. the IRDC, which dominates over the pressure of its surrounding environment (unlike for most clumps in GMCs: e.g., Bertoldi \\& McKee 1992; Kainulainen et al. 2011b). This would confirm a basic assumption of the initial conditions of the Core Accretion model of massive star formation of McKee \\& Tan (2002, 2003) and is also expected under the scenario of Equilibrium Star Cluster Formation (Tan, Krumholz \\& McKee 2006). The fact that our results differ from those of HT11, highlights the importance of improved estimates of masses and surface pressures of IRDCs. These effects may help explain other reported discrepancies between dynamical (i.e. virial) masses and true masses (e.g. Battersby et al. 2010). Measurement of large scale magnetic field strengths and geometries in IRDCs would help to further constrain the models. \\begin{figure*}[!tb] \\begin{center}$ \\begin{array}{c} \\includegraphics[width=4.5in,angle=0]{Fig2.eps} \\end{array}$ \\end{center} \\caption{ Ratio of surface to internal pressure, $P_e/P_f$, versus ratio of mass per unit length to virial mass per unit length, $m_f/m_{{\\rm vir},f}$, for strips 1 to 4 of IRDC H. Red line and points labelled $\\#_o$ show the ``Outer Filament''. Magenta line and points labelled $\\#_i$ show the ``Inner Filament''. The error bars in the lower-left show typical uncertainties. The smooth curves show the conditions satisfied by equation~(\\ref{cylvireqn}) for confining magnetic fields with ${\\cal M}_f/|W_f|<0$ (dotted lines), supporting magnetic fields with ${\\cal M}_f/|W_f|>0$ (dashed lines), and no magnetic fields, i.e. ${\\cal M}_f/|W_f|=0$ (long-dashed line). The results for the Inner Filament indicate a dynamical state consistent with magnetically-neutral virial equilibrium. } \\label{fig:fiege} \\end{figure*}" }, "1207/1207.1436_arXiv.txt": { "abstract": "{The spatial distribution and variability of Fe-K$\\alpha$ emission from molecular clouds in the Galactic Centre region may provide an important key to the understanding of the recent history of Sgr A*. A very plausible interpretation is that this variability represents an echo in the reflected radiation from the clouds of a past episode of high activity in Sgr A*.} {We examine the temporal and spectral properties of nine Fe-K$\\alpha$ bright molecular clouds within about 30 pc of Sgr A*, in order to understand and constrain the primary energising source of the Fe fluorescence.} {The variability of the Fe-K$\\alpha$ line at 6.4-keV was investigated by spectrally fitting the data derived from the EPIC MOS cameras, after subtracting a modelled background. We have also studied the reflection imprints in time-averaged pn-spectra of each cloud, by measuring the equivalent width (EW) of the 6.4 keV line and the optical depth of the Fe-K absorption edge at 7.1 keV.} {Significant Fe-K$\\alpha$ variability was detected, with a spatial and temporal pattern consistent with that reported in previous studies. The main breakthrough that sets our paper apart from earlier contributions on this topic is the direct measurement of the column density and the Fe abundance of the MCs in our sample. We used the EW measurements to infer the average Fe abundance within the clouds to be 1.6$\\pm$0.1 times solar. The cloud column densities derived from the spectral analysis were typically of the order of 10$^{23}$ cm$^{-2}$, which is significantly higher than previous estimates. This in turn has a significant impact on the inferred geometry and time delays within the cloud system. } { Past X-ray activity of Sgr A* is the most likely source of ionisation within the molecular clouds in the innermost 30 pc of the Galaxy. In this scenario, the X-ray luminosity required to excite these reflection nebulae is of the order of 10$^{37}$-10$^{38}$ erg s$^{-1}$, significantly lower than that estimated for the Sgr B2 molecular cloud. Moreover, the inferred Sgr A* lightcurve over the past 150 years shows a long-term downwards trend punctuated by occasional counter-trend brightening episodes of at least 5 years duration. Finally, we found that contributions to the Fe fluorescence by X-ray transient binaries and cosmic-ray bombardment are very likely, and suggest possible ways to study this latter phenomenon in the near future.} ", "introduction": "The Galactic Centre (hereafter GC) region is a unique environment within the local Universe, which provides many tests of our understanding of fundamental issues in astrophysics. The region hosts the nearest Super Massive Black Hole (SMBH), Sgr A*, with a mass of 4$\\times$10$^{6}$M$_{\\odot}$ \\citep[][]{2002Natur.419..694S, 2008ApJ...689.1044G}. It is also a region in which high energy phenomena abound. For example, in X-rays the Sgr B2 and Sgr C molecular complexes, located at projected distances of 90 pc and 70 pc from Sgr A*, shine brightly through 6.4-keV Fe-K$\\alpha$ line emission \\citep[][]{1996PASJ...48..249K, 2000ApJ...534..283M, 2009PASJ...61S.233N}, consistent with the prediction of \\citet{1993ApJ...407..606S}. The physical mechanism responsible for the Fe-K$\\alpha$ emission from these molecular clouds near the GC is the fluorescence of cold, neutral or near-neutral matter irradiated by high energy particles or X-ray photons. So far several hypotheses have been proposed as to the nature and origin of the primary source of this irradiation. The possibilities include: the X-ray reflection nebulae model \\citep[XRN,][]{1998MNRAS.297.1279S}, heating by low-energy cosmic-rays (CR) \\citep[][]{2002ApJ...568L.121Y}, shock mechanisms \\citep[][]{1997AIPC..410.1027Y}, and electron bombardments \\citep[][]{2003AN....324...73P}. A recent suggestion is that subrelativistic protons can be created via accretion of stellar debris onto the central black hole, thus explaining both the observed X-ray continuum and the 6.4 keV line emission from the GC \\citep[][]{2009PASJ...61..901D}. Other particle-like candidates can be CR electrons originating in supernova events \\citep[][]{2000ApJ...543..733V}. If the fluorescence observed in molecular complexes is the result of irradiation by X-ray photons, a luminous localised source of X-rays must be invoked to power the observed emission, since the diffuse hot plasma which permeates the central regions of the Galactic plane produces a factor ten less photons than is required to account for the observed Fe-K$\\alpha$ flux. It has been estimated that, depending on its position relatively to the Sgr B2 and Sgr C molecular clouds, this source should have a 2-10 keV luminosity of the order of 10$^{39}$erg~s$^{-1}$. Presently no persistent sources in the GC region have such a high luminosity; however a past transient outburst in a source, which has now returned to a relatively quiescent state, might well match the requirement. Sgr A* itself is arguably the best candidate, since although its activity is currently rather weak \\citep[the brightest flare measured with \\textit{XMM-Newton} has a luminosity of the order of a few 10$^{35}$ erg~s$^{-1}$,][]{2003A&A...407L..17P}, the SMBH could have undergone a period of high-state activity in the past \\citep[][]{1996PASJ...48..249K}. In fact, the projected light travel time to Sgr B2 implies an outburst roughly 75-150 years ago, according to a recent estimate \\citep[][] {2010ApJ...719..143T}. Since there are a number of filaments nearer to Sgr A* which emit at 6.4-keV, the same argument would require Sgr A* to also have flared more recently, say within the last 100 years, depending on the relative position of the clouds along the line of sight. One method of distinguishing between the point source and the particle hypothesis is to investigate the lightcurve of the filaments which emit the 6.4-keV line. If we are dealing with a transient source as an engine of the fluorescent emission, one might expect the temporal evolution of the observed line flux to follow the light curve of the outburst, subject to blurring arising from the spread in light-travel delays inherent in the source to cloud geometry. Such effects have been found in the Sgr B2 molecular cloud where the brightest peak of the Fe-K$\\alpha$ emission observed in 1995 by ASCA \\citep[][]{1996PASJ...48..249K,2000ApJ...534..283M} maintained the same brightness level when remeasured in 2000 with Chandra \\citep[][]{2001ApJ...558..687M} but then, some five years later when observed by Suzaku, showed a marked decline to approximately half the peak value \\citep[][]{2008PASJ...60S.201K,2009PASJ...61S.241I}. Similarly, \\citet[][]{2009PASJ...61S.233N} discovered variability of the Fe-K$\\alpha$ flux from Sgr C on the basis of recent Sukaku observations. Moreover, \\citet[][]{2007ApJ...656L..69M} have reported the evolution in intensity and morphology of the 4-8 keV continuum emission in two filamentary regions located close to Sgr A*. The latter changes occurred on parsec scales (in projection), which requires a brightening/fading of the illuminating source over a 2-3 year period, with an inferred 2--8 keV luminosity of at least 10$^{37}$ erg~s$^{-1}$. Recently, \\citet[][]{2010ApJ...714..732P} showed that the 6.4 keV line flux from the molecular filaments with 15 arcmin of Sgr A* exhibit a complex pattern of variability. If these molecular clouds/knots have a particular distribution along the line of sight, then it is possible to argue that they were all energised by the same outburst on Sgr A*, consistent with the XRN scenario (\\citealt{2010ApJ...714..732P}). However, given the observational complexities, it is very possible that this is not the complete story; indeed, very recently \\citet[][]{2011A&A...530A..38C} studied the Fe-K$\\alpha$ line emission from the MCs in the Arches cluster region (about 20 pc in projection from Sgr A*), showing that the XRN/Sgr A* scenario can hardly describe the spectral and temporal properties of those clouds. The above results open once more the question whether Sgr A* has exhibited AGN activity in the past and, if so, what is the exact nature of that activity. It is in this context that we reassess in this paper, the morphology, variability and spectral properties of the 6.4-keV emitting clouds within 15\\arcmin of Sgr A*, using the extensive set of \\textit{XMM-Newton} observations targeted at this region. Our goal is to investigate both the past role of Sgr A* (or other transient sources) in illuminating the GC molecular clouds and also to seek evidence for a contribution from alternative mechanisms, such as CR bombardment, in the excitation of the Fe fluorescence which characterises the GC environment. Throughout this work, the distance to the GC has been taken to be 8 kpc \\citep[][]{2009ApJ...692.1075G}. \\onltab{1} { \\begin{table*} \\caption{Specifications for the selected OBSIDs: MODE/FILTER combination used for each of the pointings, and GTI compared to the total exposure for each instrument. F=Full Frame MODE; E=Extended Full Frame MODE; T=Thick filter; M=Medium Filter. OBSID 0506291201 has PN in Timing MODE.} \\label{obs_table} \\centering \\begin{tabular}{|c|c|ccc|ccc|} \\hline & & & Instrument specifics & & & GTI \\& Exposure (ks) & \\\\ \\hline OBSID & Obs Date & PN & MOS1 & MOS2 & PN & MOS1 & MOS2 \\\\ & yyyy-mm-dd & mode/filter & mode/filter & mode/filter & GTI/exp & GTI/exp & GTI/exp \\\\ \\hline \\hline 0111350101 & 2002-02-26 & F/T & F/M & F/M & 38.590/40.030 & 42.262/52.105 & 41.700/52.120 \\\\ 0202670501 & 2004-03-28 & E/M & F/M & F/M & 13.320/101.170 & 33.070/107.784 & 30.049/108.572 \\\\ 0202670601 & 2004-03-30 & E/M & F/M & F/M & 25.680/112.204 & 32.841/120.863 & 35.390/122.521 \\\\ 0202670701 & 2004-08-31 & F/M & F/M & F/M & 59.400/127.470 & 80.640/132.469 & 84.180/132.502 \\\\ 0202670801 & 2004-09-02 & F/M & F/M & F/M & 69.360/130.951 & 94.774/132.997 & 98.757/133.036 \\\\ 0402430301 & 2007-04-01 & F/M & F/M & F/M & 61.465/101.319 & 61.002/93.947 & 62.987/94.022 \\\\ 0402430401 & 2007-04-03 & F/M & F/M & F/M & 48.862/93.594 & 40.372/97.566 & 41.317/96.461 \\\\ 0402430701 & 2007-03-30 & F/M & F/M & F/M & 32.337/32.338 & 26.720/33.912 & 27.685/33.917 \\\\ 0505670101 & 2008-03-23 & F/M & F/M & F/M & 74.216/96.601 & 73.662/97.787 & 74.027/97.787 \\\\ 0554750401 & 2009-04-01 & F/M & F/M & F/M & 30.114/38.034 & 32.567/39.614 & 33.802/39.619 \\\\ 0554750501 & 2009-04-03 & F/M & F/M & F/M & 36.374/42.434 & 41.376/44.016 & 41.318/44.018 \\\\ 0554750601 & 2009-04-05 & F/M & F/M & F/M & 28.697/32.837 & 37.076/38.816 & 36.840/38.818\\\\ \\hline \\end{tabular} \\end{table*} } \\begin{figure}[!Ht] \\begin{center} \\includegraphics[width=0.4\\textwidth,angle=-90]{19544fig1.eps} \\end{center} \\caption{\\footnotesize {Soft proton flare filtering of the MOS1 dataset from OBSID 0202670701. \\textit{Upper panel}: 2.5-12 keV count rate histogram. The blue lines mark the region used in the Gaussian fit, the green line represents the best fit Gaussian and the red lines show the bounds used to filter the data. \\textit{Mid panel}: 2.5-8.5 keV lightcurve of the IN FOV region. \\textit{Lower panel}: 2.5-8.5 keV lightcurve of the corner data. In both the Mid and the Lower panels, the points coloured green in the light-curves correspond to the selected GTI intervals (see text).}} \\label{espfilt} \\end{figure} ", "conclusions": "\\label{Sect:conclusions} We have studied the X-ray properties of selected molecular clouds and filaments to the east of Sgr A* at a projected distance of $\\approx$8-30 pc. \\begin{itemize} \\item{\\textit{Fe-K$\\alpha$ topology}: we have demonstrated that the Fe-K$\\alpha$ emission delineates both compact and more diffuse structures. We have studied the temporal and spectral properties of 9 clouds; three of which are newly studied (regions C, DS1 and DS2).} \\item{\\textit{Fe-K$\\alpha$ variability}: significant variability is seen from some of the molecular clouds, although the pattern is by no means a simple one. Our Fe-K$\\alpha$ flux and variability measurements agree reasonably well with previous published results, although in the case of one cloud (B1) we measured a decrease of the Fe-K$\\alpha$ line flux, in contrast with what found previously.} \\item{\\textit{Fe-K$\\alpha$ surface brightness}: we measure the surface brightness of the 6.4-keV line to be of the same order of magnitude in all the molecular filaments we have studied. Typically the observed variability only involves a factor of two change. In the Sgr A*/XRN scenario, the surface brightness of the filaments might be expected to decline with distance from Sgr A*. This issue is resolved by our finding that the X-ray luminosity of Sgr A* has declined substantial over the last 150 years.} \\item{\\textit{Fe-K$\\alpha$ EW and Fe-K edge measurements}: we measured a high value of the EW of the 6.4-keV line from all the MCs in our sample, consistent with an origin of the bulk of the fluorescence in photoionisation. Moreover, all the clouds show an absorption feature at the Fe-K edge energy of 7.1 keV.} \\item{\\textit{Spectral hardness of the reflected continuum}: the spectral shape of the reflected/ionising X-ray continuum emission associated with the Fe fluorescence is found to be $\\Gamma \\approx$1.9, a result compatible with the XRN/Sgr A* scenario, since this is a value rather typical of that found in AGN.} \\item{\\textit{Fe abundance}: we use the measured EW values of the Fe-K$\\alpha$ line in our sample of MCs to calculate the average Fe abundance (relative to solar). We show that Z$\\rm{_{Fe}}$=1.6, a result which is consistent with the general finding that a higher than solar metallicity characterises all the ISM phases in the GC region.} \\item{\\textit{Column density through the MCs}: we have employed a relatively direct method to calculate the hydrogen column density, N$\\rm{_{H_{C}}}$, within the MCs. This is based on the joint modelling of the low-energy absorption and the absorption at the Fe-K edge imprinted on the incident continuum in its passage through the cloud (both up to and after the point of scattering). This approach benefits from prior knowledge of the relatively iron abundance, Z$\\rm{_{Fe}}$, in the cloud. The column densities so calculated are close to 10$^{23}$ cm$^{-2}$ and, in most cases, significantly higher than those inferred in previous studies, which made use of radio molecular line measurements.} \\item{\\textit{The past X-ray activity of Sgr A*}: based on our study of the 6.4-keV line emission properties of the MCs in the inner 30 pc of the GC region, we outline a model for the X-ray emission of Sgr A* encompassing the last $\\sim$150 years. Over this period the X-ray luminosity has declined from an apparent peak of $\\sim$10$^{39}$ erg~s$^{-1}$ roughly 150 years ago, to 10$^{37-38}$ erg~s$^{-1}$ perhaps 100 years ago, down to typically 10$^{33-35}$ erg~s$^{-1}$ at the present time. This decreasing long-term trend has, however, been punctuated by counter-trend episodes of brightening by factors of a few over timescales in excess of $\\sim$ 5 years.} \\item{\\textit{Cloud C - an XRN energised by a transient source?}: we have found that the 6.4-keV line flux variability measured from cloud C could also be associated with the transient X-ray source XMMU J174554.4-285456. In this scenario, the Fe fluorescence in this complex is composed of a steady non-zero level, possibly produced by the interaction of CR with the ambient gas plus a superimposed contribution due to the localised variable X-ray irradiation. If this picture is confirmed, this is the second XRN found in the GC region which has been illuminated by a transient X-ray source \\citep[][]{2011A&A...530A..38C}.} \\item{\\textit{Low surface brightness 6.4-keV line emission}: we have found a good correlation between the TeV emission and the Fe-K$\\alpha$ line emission observed on arcminute scales in the inner CMZ. Specifically we measured diffuse, low surface brightness emission to the east of Sgr A*, which seems to permeate the whole region between the Sgr A* and the giant Radio Arc. A pedestal (\\textit{i.e.} constant) fluorescence component might in fact be present in all the lightcurves which show a varying Fe-K$\\alpha$ line flux. We have also shown that this diffuse low surface brightness 6.4-keV line emission correlates quite well with the contours of the brightest features seen in TeV $\\gamma$-rays. This suggests that CRs may contribute to the ionisation and fluorescence of cold gas in this region.} \\item{\\textit{CR contribution}: we presented arguments in favour of a CR contribution to the total Fe fluorescence in the GC MCs. In particular, LECR may be responsible for the constant of 6.4-keV line emission observed in several structures (including the Bridge), as well as in other regions of the CMZ, like the Arches cluster, the Radio Arc and Sgr C. We suggest that further study of the Fe-K$\\alpha$ surface brightness versus N$\\rm{_{H_{C}}}$ relation for GC clouds will help reveal the contribution of CR to the Fe-K fluorescence observed in the GC region.} \\end{itemize} \\noindent We are far from unveiling the full mystery of the Fe fluorescence observed in the GC region. It will be of great importance in the future to continue to monitor the Fe-K$\\alpha$ line emission from the MCs in the CMZ so as to better characterise the past activity of Sgr A*. Future X-ray observations can also complement studies in different energy domains, from the radio band (continuum and lines) up to the TeV regime, which aim to fully chart the influence of CR particles on the GC environment." }, "1207/1207.3928_arXiv.txt": { "abstract": "We describe a simple step-by-step guide to qualitative interpretation of galaxy spectra (Fig.~\\ref{decision}). Rather than an alternative to existing automated tools, it is put forward as an instrument for quick-look analysis, and for gaining physical insight when interpreting the outputs provided by automated tools. Though the recipe is of general application, it was developed for understanding the nature of the Automatic Spectroscopic K-means based (ASK) template spectra. They resulted from the classification of all the galaxy spectra in the Sloan Digital Sky Survey data release 7 (SDSS-DR7), thus being a comprehensive representation of the galaxy spectra in the local universe. Using the recipe, we give a description of the properties of the gas and the stars that characterize the ASK classes, from those corresponding to passively evolving galaxies, to HII galaxies undergoing a galaxy-wide starburst. The qualitative analysis is found to be in excellent agreement with quantitative analyses of the same spectra. We compare the mean ages of the stellar populations with those inferred using the code {\\sc starlight}.% We also examine the estimated gas-phase metallicity with the metallicities obtained using electron-temperature based methods. A number of byproducts follow from the analysis. There is a tight correlation between the age of the stellar population and the metallicity of the gas, which is stronger than the correlations between galaxy mass and stellar age, and galaxy mass and gas metallicity. The galaxy spectra are known to follow a 1-dimensional sequence, and we identify the luminosity-weighted mean stellar age as the affine parameter that describes the sequence. All ASK classes happen to have a significant fraction of old stars, although spectrum-wise they are outshined by the youngest populations. Old stars are metal rich or metal poor depending on whether they reside in passive galaxies or in star-forming galaxies. ", "introduction": "There are several automated tools for inferring the properties of the stellar populations contributing to the integrated galaxy spectra. The list includes {\\sc moped} \\citep{2004MNRAS.355..764P}, {\\sc starlight} \\citep{2005MNRAS.358..363C}, {\\sc steckmap} \\citep{2006MNRAS.365...74O}, {\\sc vespa} \\citep{2007MNRAS.381.1252T}, or {\\sc ulyss} \\citep{2009A&A...501.1269K}, as well as the use of line indices like the Lick indices \\citep{1994ApJS...94..687W}. Similarly, there are semi-automatic procedures to deduce the properties of the gas \\citep[e.g.,][]{1995PASP..107..896S,2006IAUS..234..439J,pynebular}, including the so-called strong-line ratio methods \\citep[e.g.,][]{1979MNRAS.189...95P,2000MNRAS.312..130D,2002MNRAS.330...69D,2005A&A...437..849S}. These tools are (and will be) fundamental for understanding the galaxy formation and evolution, but the blind use of the codes results quite unsatisfactory from a physical stand point. One obtains a precise quantitative description of the stellar populations contributing to the integrated spectra, but ignores the reason why the code has chosen them rather than other potential alternatives. The educated eye of an astronomer is often far more telling from a physical point of view. Unfortunately, the know-how of qualitatively interpreting a spectrum is learned after a long experience of working in the field. The information on which particular spectral feature informs of which particular physical property is scattered among a large number of technical publications, difficult to identify and to deal with for a newcomer. This paper aims at providing a step-by-step guide to qualitative interpretation of galaxy spectra. Moreover, it will be compared with up-to-date numerical techniques to show that both qualitative and quantitative results are in excellent agreement. The work was originally planned as a mere academic exercise to understand the nature of the classes resulting from the k-means classification of all the galaxy spectra in the Sloan Digital Sky Survey data release 7 \\citep[SDSS-DR7, ][]{2010ApJ...714..487S}. We wanted to translate the spectral shapes into physical units like stellar ages and metallicities, so that this information can be used to tailor class-based searches \\citep[e.g.,][]{2012A&A...540A.136A},% or when interpreting spectra \\citep[e.g.,][]{sanchezjanssen2012}. However, the exercise is of interest beyond the original scope. The simple decision tree we use is suitable to characterize any galaxy spectrum. We know of its generality because it allows to separate and characterize the 28 Automated Spectroscopic K-means-based (ASK) classes \\citep{2010ApJ...714..487S} which, by construction, are proxies that condense the properties of the some one-million SDSS spectra \\citep[][]{2002AJ....123..485S,2009ApJS..182..543A}. The ASK class characterization represents a significant part of the paper, that are discussed in detail as an illustration of the procedure. As we stress above, our qualitative analysis may have several other applications, e.g., (1) to gain physical insight when interpreting quantitative Star Formation Histories (SFHs) derived from modern automated tools, (2) for quick-look galaxy classification (not only in the local universe, but also at moderate-high redshifts, since the Hubble expansion shifts the UV-visible spectrum to the near IR), (3) for interpreting noisy spectra where eyeball inspection is often better than detailed inversion, (4) as reference for identifying unusual galaxies, or (5) for educational purposes to develop physical intuition. The paper is organized as follows. Section~\\ref{ask_class} introduces the ASK spectral classification of galaxy spectra whose templates serve as reference point. Section~\\ref{list_features} lists and discusses spectral features commonly used when interpreting galaxy spectra. They are employed to set up the recipe introduced in Sect.~\\ref{decision_tree}, which is abridged in a schematic shown in Fig.~\\ref{decision}. The recipe (or algorithm) is used in Sect.~\\ref{qualitative_classes} to disclose the physical properties of all the ASK classes. The results of such qualitative analysis are compared with state-of-the-art quantitative analyses in Sects.~\\ref{starlight} and \\ref{quantitative_lines} -- Sect.~\\ref{starlight} deals with the comparison of stellar components, whereas Sect.~\\ref{quantitative_lines} refers to the gas components. Section~\\ref{additional_results} discusses several additional properties of the ASK templates, whereas Sect.~\\ref{conclusions} summarizes the proposed qualitative analysis. ", "conclusions": "As argued in the introduction, we have % sophisticated computer codes for inferring the properties of the stellar populations contributing to the observed galaxy spectra. Similarly, tools are available for qualitative diagnostics of the physical properties of the galaxy gas. They have been developed by specialist groups, and then kindly offered to a much broader community. Technicalities often complicate the interpretation of the results, therefore, there is a natural tendency to apply these sophisticated tools in black-box fashion, which turns out to be quite unsatisfactory for a physical stand point. One obtains a detailed description of the stars and gas producing the observed galaxy spectra, but overlooks the reasons why the computer code has preferred them rather than other alternatives. We provide a simple step-by-step guide to qualitative interpretation of galaxy spectra. It is not precise, and has not been planed as an alternative to the existing tools. However, it allows a quick-look that yields the main properties of the spectra in a intuitive fashion. This may be of interest in various applications, e.g., to provide physical insight when using sophisticated tools, or to interpret noisy spectra. Moreover, the results of the qualitative analysis agree with those inferred using up-to-date computer codes. The step-by-step guide is described in Sect.~\\ref{decision_tree}, and it has been summarized as a simple questionary in Fig.~\\ref{decision}. Emission and absorption lines are analyzed separately, which give rise to a classification with one entry for the gas and another for the stars. (In real galaxies, however, the properties of gas and stars are tightly correlated; see Sect.~\\ref{additional_results}.) The analysis has been systematically applied to the set of ASK template spectra that resulted from the classification of all galaxy spectra in SDSS-DR7 (see Sect.~\\ref{ask_class}). Their physical properties are summarized in Table~\\ref{table_summary}. With the caveats pointed out in Sect.~\\ref{additional_results}, the ASK classes represent a comprehensive set of galaxy spectra, that go all the way from passively evolving red galaxies (e.g., ASK~0) to HII galaxies, dominated by massive newborn stars having no absorption lines (e.g., ASK 15). Since it works for this set, the analysis should work for most galaxies. The qualitative analysis is found to be in excellent agreement with quantitative numerical codes. We show how the index for stellar-age (SAI) follows an almost one-to-one correlation with the mean stellar age assigned by the code {\\sc starlight} (Fig.~\\ref{age_vs_age}). Similarly, we found how the proxy for gas metallicity is in good agreement with the (oxygen) metallicity inferred by applying the direct method to the emission lines of the ASK templates (Fig.~\\ref{pplike}). The ASK templates are freely available (see footnote \\#~\\ref{my_foot}) and, together with their physical properties listed in Table~\\ref{table_summary}, they can be used as benchmarks so that any other galaxy spectrum can be analyzed by reference to them." }, "1207/1207.6350_arXiv.txt": { "abstract": "Whereas current cosmological observations suggest that the universe is dominated by a positive cosmological constant ($\\Lambda > 0$), the AdS/CFT correspondence tells us that the case $\\Lambda<0$ is still worthy of consideration. In this paper we study the McVittie solution with $\\Lambda<0$. Following a related study, the solution is understood here by way of a systematic construction of conformal diagrams based on detailed numerical integrations of the null geodesic equations. As in the pure Robertson - Walker case, we find that $\\Lambda<0$ ensures collapse to a Big Crunch, a feature which completely dominates the global structure. ", "introduction": "Recently \\cite{lake}, a detailed study of the McVittie solution \\cite{mcvittie} was carried out for a non-negative cosmological constant ($\\Lambda \\geq 0$). The McVittie solution has been known for many years, but it continues to attract interest \\cite{newmcvittie}. Even though it is now widely believed that the universe is dominated by a positive cosmological constant, the remarkable AdS/CFT correspondence \\cite{Maldacena} presents a strong argument that the case $\\Lambda< 0$ should also be examined. Following \\cite{lake} we systematically construct a global view of the McVittie solution with $\\Lambda< 0$ based on numerical integrations of the null geodesics. What results is a situation very distinct from the case $\\Lambda \\geq 0$: the global structure is completely dominated by a collapse to a Big Crunch, just as in the pure Robertson - Walker case. ", "conclusions": "Motivated by the AdS/CFT correspondence, we have examined the McVittie solution with a negative cosmological constant $\\Lambda < 0$. A detailed construction of the global structure has been given for the case of a background of dust. We have found that the situation is very distinct from the cases $\\Lambda \\geq 0$ \\cite{lake}. As in the pure Robertson - Walker case, we find that $\\Lambda<0$ ensures collapse to a Big Crunch, a feature which completely dominates the global structure." }, "1207/1207.6399_arXiv.txt": { "abstract": "We have searched for prompt radio emission from nine Gamma Ray Bursts (GRBs) with a 12~m telescope at 1.4~GHz, with a time resolution of $64~\\unit{\\mu s}$ to 1~s. We detected single dispersed radio pulses with significances $>6 \\sigma$ in the few minutes following two GRBs. The dispersion measures of both pulses are well in excess of the expected Galactic values, and the implied rate is incompatible with known sources of single dispersed pulses. The arrival times of both pulses also coincide with breaks in the GRB X-ray light curves. A null trial and statistical arguments rule out random fluctuations as the origin of these pulses with $>95\\%$ and $\\sim 97\\%$ confidence, respectively, although a simple population argument supports a GRB origin with confidence of only 2\\%. We caution that we cannot rule out RFI as the origin of these pulses. If the single pulses are not related to the GRBs we set an upper limit on the flux density of radio pulses emitted between 200 to 1800~s after a GRB of $1.27 w^{-1/2} \\unit{Jy}$, where $6.4 \\times 10^{-5} \\unit{s} < w < 32 \\times 10^{-3} \\unit{s}$ is the pulse width. We set a limit of less than 760~Jy for long timescale ($>1$~s) variations. These limits are some of the most constraining at high time resolution and GHz frequencies in the early stages of the GRB phenomenon. ", "introduction": "When processing archival data from the Parkes 64~m telescope, \\citet{lorimer2007bmr} detected a single 30~Jy burst with a spectral index of -4 ($S \\propto \\nu ^{\\alpha}$) with $\\sim 5 \\unit{ms}$ duration at a dispersion measure (DM)of $375 \\unit{pc~cm^{-3}}$. The dispersion measure was much larger than the Galactic contribution in the direction \\citep{Cordes02} implying an extragalactic origin. From this detection, \\citet{lorimer2007bmr} derived a brightness temperature around $\\sim 10^{34}~\\unit{K}$ and an event rate of $3.8 \\times 10^{-4} \\unit{hr^{-1} deg^{-2}}$ (based on a sample of 1). \\citet{lorimer2007bmr} proposed that these parameters were broadly compatible with a Gamma Ray Burst (GRB) origin, but noted that no mechanisms had been discussed that could produce such a burst. \\citet{Keane11} also report a detection with the same telescope of a somewhat weaker burst (4~Jy) with higher DM ($745 \\unit{pc~cm^{-3}}$). To date, a GRB origin for these short-timescale, GHz radio pulses has yet to be observationally tested. Furthermore, some doubt has been cast on the astronomical origin of these bursts \\citep{BurkeSpolaor11lg,Kocz12}, and the state of published theoretical mechanisms with the observed timescales and frequencies has not progressed. There are some, as yet unobserved mechanisms that can produce longer timescale radio emission at lower frequencies( $\\sim 100~\\unit{MHz}$) both for for the collapsar model for long GRBs \\citep{Usov00, Sagiv02, Inoue04, Moortgat05} and short GRBs \\citep{Lipunov96, Hansen01, Pshirkov10, Shibata11}, and scattering effects may limit the observability of short-timescale bursts in certain circumstances \\citep{Macquart07,Lyubarsky08}. Nonetheless, the rewards for detecting prompt emission from GRBs are great. Many aspects of the explosion physics can be probed by measurements of the prompt radio emission, such as the jet opening angle, and Lorentz factor \\citep{Macquart07}, the density and distance of any scattering material \\citep{Lyubarsky08}, and the structure of the fireball magnetic field \\citep{Sagiv04}. Additionally, a short radio pulse from an extragalactic source is expected to undergo dispersion as it propagates through the intergalactic medium (IGM). Such dispersion would not only provide direct evidence for the existence of the majority of the baryons in the Universe \\citep{Ginzburg73}, but, for bursts of sufficiently high redshift, would also differentiate between different models of cosmic reionization history \\citep{Inoue04}. There have been a number of unsuccessful searches for prompt emission from GRBs, although no searches were sensitive to the burst of \\citet{lorimer2007bmr}. Most searches have been non-directed, in which a large fraction of the sky was monitored, with the hope of a GRB occurring somewhere in this region. Early results at 151 and 408~MHz with an integration time of 300~ms detected some pulses within $\\pm 10 \\unit{min}$ of the gamma-ray trigger, but did not confidently associate any with GRBs \\citep{Cortiglioni81,Inzani82}. A survey at 843~MHz, which was sensitive to pulses between 0.001~ms and 800~ms, also made no definitive detections \\citep{Amy89}, although it was not clear if any GRBs were present in the field of view during the observations. More recently, \\citet{Katz03} performed an all-sky survey at 611~MHz, with a time resolution of 125~ms and a flux density detection threshold of 27~kJy. This search detected $\\sim 4 \\times 10^6$ bursts in 18 months, but rejected 99.9\\% as RFI and identified the remaining bursts with solar activity. In a similar vein, \\citet{Lazio10} detected no transients above 500~Jy for pulse widths of about 300~s at 73.8~MHz. A number of other surveys have begun but are yet to publish results \\citep{Balsano98, Morales05}. The only report of automatic follow-up at radio frequencies is that performed by \\citet{Koranyi95} and \\citet{Dessenne96} at 151~MHz with a time resolution of 1.5~s. Based on observations of two GRBs, \\citet{Dessenne96} report upper limits on any radio emission of 16--73~Jy between 5 hrs before, and 2 hrs after the GRB. A search for the evaporation of primordial black holes at 3~GHz with a time resolution of $2\\mu$s also failed to detect any radio emission \\citep{Osullivan78}. The lack of radio detections of previous surveys may be due to their low operating frequencies, low time resolution, insufficient sensitivity, and low sensitivity to the GRB rate. While low frequency observations have the advantage of a large predicted radio luminosity due to steep spectral index, which is predicted by some models of prompt radio emission (e.g. \\citep{Sagiv02}), and has been observed in one case \\citep{lorimer2007bmr}, some low frequency effects make detecting short duration pulses more difficult. For example, scatter broadening, which substantially reduces the detectability of radio pulses \\citep{Cordes03}, and sky temperature are worse at low frequencies. Low time resolution reduces the detectability of short duration radio pulses, including those required to avoid the brightness temperature constraint from induced scattering \\citep{Lyubarsky08}. In particular, the directed searches of \\citet{Koranyi95} and \\citet{Dessenne96} used a relatively low time resolution of 1.5~s. The sensitivities of most of the above surveys are low, and not approaching the flux levels required to detect the burst of \\citet{lorimer2007bmr}, while more sensitive blind experiments have not had the field of view and on-sky time to obtain a GRB in-beam \\citep{Wayth12, BurkeSpolaor11,Deneva09}. In this paper we describe a survey to detect prompt radio emission from GRBs at 1.4~GHz, to test whether the bursts reported by \\citet{lorimer2007bmr} and \\citet{Keane11} have a GRB origin. We used a single 12~m dish that slews automatically to the GRB coordinates, based on the gamma-ray position, and observes the position with high time resolution. Our aim was twofold: to attempt to detect any GHz radio emission within the first few minutes of the GRB, and to gain experience in automating radio follow-up for a potential future experiment. ", "conclusions": "We have searched for prompt radio emission from gamma-ray bursts at 1.4~GHz, using a robotic telescope and a pulsar backend. Our telescope was typically on source within 200~s of the gamma-ray trigger. We detected single dispersed pulses following two GRBs at significances $>6 \\sigma$. Simple statistical arguments, and a null trial based on randomizing channels on existing data, rule out random fluctuations as the origin of these pulses at $95\\%$ and $\\sim 97\\%$, respectively. The arrival times of both pulses also coincide with breaks in the X-ray light curves. While high DM and absence of adjacent RFI lend credence to an astronomical origin, weak impulsive RFI and atmospheric origins of these pulses remain a distinct possibility. A simple population arguments suggests a GRB origin for these pulses of only 2\\%. If the radio pulses are associated with the corresponding GRBs, they could be related to changes in the central engine, in particular the delayed formation of a black hole due to spin-down of a rotationally-supported magnetar. The non-detection of radio pulses by previous surveys of \\citep{Koranyi95, Dessenne96, Katz03} can be explained by insufficient sensitivity of those surveys, a somewhat flatter spectrum than measured for the Lorimer burst, or the possibility that only a fraction of GRBs emit radio pulses. If the single pulse is not related to the GRB, we set an upper limit on the flux density of radio pulses between 200 to 1800~s after GRB trigger of $1.27 w^{-1/2} \\unit{Jy}$, where $6.4 \\times 10^{-5} \\unit{s} < w < 32 \\times 10^{-3} \\unit{s}$ is the pulse width. This limit is substantially better than the all-sky 27~kJy limit of \\citet{Katz03} at 611~MHz, although not as constraining as the limits on two GRBs by \\citet{Dessenne96}. We have detected no convincing repeating dispersed candidates. We also detect no candidates on timescales $>1 \\unit{s}$, but our experiment was not primarily designed for such detections. Nonetheless, we set an upper limit of a change of 760~Jy on any long-duration emission ($>1$~s) between 200 to 1800~s from our GRB triggers. The detection of single dispersed pulses in this experiment is intriguing. Clearly the next step is to determine whether these pulses are related to their GRBs, for which the key problem is ruling out RFI, statistical fluctuations and other equipment-related sources as the origin of these pulses. The simplest future experiment to rule out these origins is to use a coincidence detection, by employing the same telescope setup at two widely separated sites. The fact that some 25\\% of GRBs may be accompanied by a radio pulse, even with our relatively poor sensitivity, suggests that sensitivity is not the key factor in this experiment. Therefore, similar dishes, feeds and backends can be used. More important parameters of this experiment are the short on-source time (preferably $<200$~s), and a wide separation between antennas. A simultaneous detection of a single pulse at two widely separated sites, even at $6 \\sigma$, would almost certainly rule out RFI and statistical fluctuations, and render atmospheric effects a very remote possibility. If such a detection were to be made, the future would be wide open to probe the astrophysical and cosmological implications of these phenomena." }, "1207/1207.6485_arXiv.txt": { "abstract": "{Transition disks are believed to be the final stages of protoplanetary disks, during which a forming planetary system or photoevaporation processes open a gap in the inner disk, drastically changing the disk structure. From theoretical arguments it is expected that dust growth, fragmentation and radial drift are strongly influenced by gas disk structure, and pressure bumps in disks have been suggested as key features that may allow grains to converge and grow efficiently.} {We want to study how the presence of a large planet in a disk influences the growth and radial distribution of dust grains, and how observable properties are linked to the mass of the planet.} {We combine two-dimensional hydrodynamical disk simulations of disk-planet interactions with state-of-the-art coagulation/fragmentation models to simulate the evolution of dust in a disk which has a gap created by a massive planet. We compute images at different wavelengths and illustrate our results using the example of the transition disk LkCa15.} {The gap opened by a planet and the long-range interaction between the planet and the outer disk create a single large pressure bump outside the planetary orbit. Millimeter-sized particles form and accumulate at the pressure maximum and naturally produce ring-shaped sub-millimeter emission that is long-lived because radial drift no longer depletes the large grain population of the disk. For large planet masses around 9~$M_{\\mathrm{Jup}}$, the pressure maximum and, therefore, the ring of millimeter particles is located at distances that can be more than twice the star-planet separation, creating a large spatial separation between the gas inner edge of the outer disk and the peak millimeter emission. Smaller grains do get closer to the gap and we predict how the surface brightness varies at different wavelengths.} {} ", "introduction": "\\label{sec1} Circumstellar disks are the birthsites of planets. The physical conditions and the evolution of these disks control the planet formation mechanisms. An important goal is to provide theoretical models and observational constraints to understand the various stages in the evolution of gas and dust in the disk. Decrease of mass accretion rate \\citep{sicilia10,fedele10}, and near-infrared excess with time \\citep{hernandez07, andrews11b} indicate that disks have a range of lifetimes from 1 to 10~Myr. With the advent of powerful infrared space telescopes such as \\textit{Spitzer}, a new class of objects has been identified, called the transition disks \\citep[e.g.,][]{espaillat10}. Their spectral energy distribution (SED) and direct sub-millimeter (mm) imaging suggest that warm dust in the inner disk is strongly depleted compared to the outer disk. The small number of transition disks \\citep{muzerolle10} suggests an inside-out evolution that occurs rapidly. Various mechanisms have been proposed so far to explain the inner disk clearing: photoevaporation winds \\citep[e.g.][]{owen11}, grain growth \\citep{klahr97, tanaka05} and dynamical interactions with companions \\citep{lin79}. Transition disks are therefore excellent laboratories for planet formation models. The clearing of a gap by a companion or planet, from a simplistic point of view, depends on the competition between the viscous torque from the disk and gravitational torques from the planet. For a laminar disk, a 1~$M_\\mathrm{Jup}$ planet can clear a gap or hole. The recent discovery of a companion inside a massive disk in T Cha \\citep{huelamo11} supports the scenario of a dynamical clearing, at least for this object. However, models show that a single planet seems unlikely to be capable of creating the observed large holes, which require multiple systems \\citep{zhu11, dodson11}. Interestingly, for most of these transition disks the inner cavity is not empty. They still present relatively high accretion rates ($\\thicksim 10^{-8}M_{\\odot}yr^{-1}$; see e.g.,\\cite{calvet05, espaillat07}) which implies that the inner cavity is not completely empty and that some gas flows through the gap. To allow mass flowing, a limit for the planet mass can be inferred, depending on the disk viscosity \\citep{lubow06}. In addition to the gas, some transition disks also present a strong near-infrared excess, indicating the presence of dust close to the star. \\cite{rice06} studied the filtration of dust in the gap, considering a fixed size for the dust particles, and concluded that increasing the planet mass from e.g. 0.5~$M_\\mathrm{Jup}$ to 5~$M_\\mathrm{Jup}$, the maximum particle size that sweeps into the gap decreases from $\\sim10~\\mu m$ to a few tenths of a micrometer. Transition disks are potentially interesting laboratories to study processes related to the impact of planet formation on the disk. One of the most stubborn problems in planet formation is the so-called ``meter-size barrier\". A one meter size object at 1~AU drifts towards the central star in timescales shorter than the growth timescales, impeding it to grow \\citep[see e.g.,][]{brauer08, birnstiel10}. In addition, high relative velocities lead to numerous fragmentation collisions converting a large object into smaller dust particles. The same physical process happens to the millimeter-size particles that are observed in the outer regions of the disk \\citep[e.g.,][]{wilner00, ricci10, guilloteau11}. One possible solution to prevent the rapid inward drift and trap dust particles, is to consider pressure bumps \\citep{klahr97, fromang05, brauer08, johansen09, pinilla11}. A long-lived positive pressure gradient can lead dust particles to move outwards, causing an accumulation of dust at the location of the pressure maximum. A large pressure bump is expected in protoplanetary disks as a consequence of the presence of a massive planet in a disk. In fact, when a planet carves a gap in a disk, the gas surface density shows a significant depletion, resulting in a large pressure bump at the outer edge of the gap. The dust material is trapped and piles up in this local pressure maximum where the gas motion is exactly Keplerian, and as a result there is no frictional drag between the gas and the particles. For that reason, not only do the particles not drift anymore, they also do not experience the potentially damaging high-velocity collisions due to relative radial and azimuthal drift. Under those circumstances, growth to larger-than-usual sizes is expected, possibly even a breakthrough that leads to overcome the growth barrier. However, if turbulence is still strong enough, particles may fragment due to their relative turbulent velocities. This scenario is only possible if we assume the presence of a planet, formed by another mechanism than dust agglomeration. In this paper, our goal is to test the idea that the outer edge of a planetary gap is a particle trap. We consider the dust evolution in a disk, where the gas density profile is determined by its interaction with a massive planet. We then explore the case of LkCa~15, a transition disk \\citep{mulders10}. The disk has been intensively observed at millimeter wavelengths \\citep{pietu06, andrews11b} with a maximum angular resolution corresponding to 28~AU \\citep{isella12}. The dust continuum images show a ring-like structure from $\\sim$~42 to $\\sim$120~AU, which is best fitted by a flat surface density profile. In addition, a $\\thicksim~$6~-~15$~M_{\\mathrm{Jup}}$ planet in circular orbit and located at $15.7~\\pm2.1$~AU, was claimed \\citep{kraus11}, but not confirmed yet. We describe the numerical simulations of planet-disk interactions that are considered for this work, as well as the coagulation/fragmentation model of the dust evolution in Sect.~\\ref{sec2}. Section~\\ref{sec3} describes the results of the numerical simulations. In addition, we present observation predictions in Sect.~\\ref{sect_obs}. Section~\\ref{sec4} is a discussion of the main results of our work and Sect.~\\ref{sec5} is a summary. ", "conclusions": "\\label{sec5} The sub-Keplerian radial velocity of the gas in protoplanetary disks makes millimeter-size particles in the outer regions of the disk exposed to a rapid inward drift, implying that they migrate towards the star on timescales shorter than a million years. Therefore, in planet formation, rapid inward drift is one of the main issues with models to form planetesimals. The idea of the presence of pressure bumps in protoplanetary disks has been proposed as a solution to stop the rapid inward drift. With the presence of a massive planet in a disk, a pressure bump is created in the outer edge of the cleared gap. Particles may experience a positive pressure gradient and as a result, do not drift anymore and not experience the high-velocity collisions due to relative radial and azimuthal drift. Nevertheless, particles can still fragment due to turbulence motion, and the resulting micron-size particles are less easily trapped and they can finally drift inwards. Some transition disks reveal gaps that can result from the presence of a massive planet, making them potentially interesting laboratories for studying dust growth under the favorable circumstances of having a significant pressure bump. In this paper, we combine two-dimensional hydrodynamical simulations for the gas with one-dimensional coagulation/fragmentation dust evolution models, to study how dust evolves in a disk which gas density has been disturbed by a massive planet. We investigate the influence of the disk geometry, turbulence and planet mass on the potential trapping of particles. The disk geometry does not have a significant effect on the gap opening process. For a 1~$M_\\mathrm{Jup}$ planet, there is an important influence of the $\\alpha$ turbulence parameter on the depth of the gap and has important consequences for the trapping of particles. Unlike the case of $\\alpha=10^{-2}$, with $\\alpha=10^{-3}$, the particles are trapped at a distance of $\\sim~0.5~r_p$ from the planet position $r_p$. While the gas gap has a radius of $\\sim~0.35~r_p$ (or 5 Hill radii), the dust is retained at a further distance. Without considering turbulence mixing for the dust dynamics would produce trapping of particles for certain sizes regardless of $\\alpha$ value provided that the pressure gradient is positive. In the case of 9~$M_\\mathrm{Jup}$ planet, the gap depth is independent of the turbulence and as a consequence the pressure gradient behaves similarly for all $\\alpha$ values. The dust particles are trapped at a distance of 1.4~$r_p$ from the planet orbit $r_p$. The planet mass strongly influences the location where the dust is retained and the width of the dust bump. Observations at millimeter wavelengths reveal some transition disks with very wide gaps. We show that the location where the dust piles up does not coincide with the gap outer edge in the gaseous disk. This mismatch suggests that multiple planets may not necessarily be needed to account for the observed large opacity holes. We reproduce the main observed features of the disk around LkCa15, in which a companion was recently detected. Reproducing asymmetric features in the disk, such as the ones found in HD135344B \\citep{brown09} is subject of future work, as they will be well detectable with ALMA." }, "1207/1207.4952_arXiv.txt": { "abstract": "{ The search for the sources of cosmic rays is a three-fold assault, using charged cosmic rays, gamma rays and neutrinos. The first conceptual ideas to detect high energy neutrinos date back to the late fifties. The long evolution towards detectors with a realistic discovery potential started in the seventies and eighties, with the pioneering works in the Pacific Ocean close to Hawaii and in Lake Baikal in Siberia. But only now, half a century after the first concepts, such a detector is in operation: IceCube at the South Pole. We do not yet know whether with IceCube we will indeed detect extraterrestrial high energy neutrinos or whether this will remain the privilege of next generation telescopes. But whatever the answer will be: the path to the present detectors was a remarkable journey. This review sketches its main milestones. } % ", "introduction": "The year 2012 marks the hundredth anniversary of the detection of cosmic rays by Viktor Hess \\cite{Hess-1912}. As we know today, cosmic rays consist of protons and nuclei of heavier elements; electrons contribute only on the percent level. Since cosmic rays are electrically charged, they are deflected by cosmic magnetic fields on their way to Earth. Precise pointing -- i.e. astronomy -- is only possible with electrically neutral, stable particles: electromagnetic waves (i.e. gamma rays at the energies under consideration) and neutrinos. High energy neutrinos, with energies much beyond a GeV, must be emitted as a by-product of collisions of charged cosmic rays with matter. Actually, only neutrinos provide incontrovertible evidence for acceleration of hadrons since gamma rays may also evolve from inverse Compton scattering of accelerated electrons and other electromagnetic processes. Since neutrinos can escape much denser celestial environments than light, they can be tracers of processes which stay hidden to traditional and gamma ray astronomy. At the same time, however, their extremely low reaction cross section makes their detection a challenge ($\\sigma_{\\nu\\,p} \\sim E_{\\nu} \\times 10^{-38}$\\,cm$^2$, with $E_{\\nu}$ in GeV). Neutrino astronomy is reality already now in the {\\it low-energy} sector, where the detection of neutrinos from the Sun and the Supernova SN\\,1987A has been accomplished and was honored by the 2002 Nobel Prize for physics. Figure~\\ref{all-nu} shows a compilation of the spectra of dominant natural and artificial neutrino fluxes. \\begin{figure}[ht] \\begin{center} \\includegraphics[width=10cm]{figs/all-nu-spectrum-mod.png} \\caption {Measured and expected fluxes of natural and reactor neutrinos (see text for explanations). The energy range from keV to several GeV is the domain of underground detectors. The region from tens of GeV to about 100 PeV, with its much smaller fluxes, is addressed by Cherenkov light detectors underwater and in ice. The highest energies are only accessible with huge detector volumes and methods described in section \\ref{new-methods}.} \\label{all-nu} \\end{center} \\end{figure} The range of $\\mu$eV and meV is that of cosmological (or \"relic\") neutrinos, i.e. the 1.9 Kelvin neutrino counterpart to the 2.7 Kelvin cosmic microwave background. No practicable idea exists on how to detect these neutrinos, since their reaction cross section as well as the energy of the recoil products from their interactions are frustratingly small. The keV-MeV range is populated by neutrinos from the Sun, from supernovae, from nuclear reactors and the from the interior of the Earth. Neutrinos from a nuclear reactor first recorded in 1956 by Clyde Cowan and Frederick Reines \\cite{Reines-1956} mark the discovery of neutrinos. It was acknowledged with the 1995 Nobel Prize for physics. Solar neutrinos have been measured first by Ray Davis in 1968 \\cite{Davis-1968} in the Homestake mine in USA. The apparent deficit of neutrinos observed by Davis -- the long-standing \"solar neutrino puzzle\" -- could eventually be explained by neutrino oscillations which transform a large part of the original solar electron neutrinos to muon and tau neutrinos; with respect to these, the detector of Davis and many of its successors were blind. Supernova neutrinos from the supernova 1987A in the Large Magellanic Cloud have been recorded at February 23, 1987 by three detectors: Kamiokande in Japan, IMB in the USA (both water Cherenkov detectors) and the Baksan scintillation detector in Russia. The Nobel Prize for physics 2002 was awarded to Masatoshi Koshiba (spokesman of the Kamioka collaboration) and Ray Davis, for \"pioneering contributions to astrophysics, in particular the detection of neutrinos from the Sun and a supernova\". Neutrinos from radioactive decay processes in the interior of the Earth (\"geo-\" or \"terrestrial\" neutrinos) have been identified only recently \\cite{Kamland-2005,Borexino-2010}. Next on the energy scale come \"atmospheric neutrinos'' created in cosmic ray interactions in the Earth's atmosphere. They have been detected in 1965 and will be in the focus of section \\ref{sec-concepts}. The highest energies are the domain of neutrinos from sources like supernova remnants, Gamma Ray Bursts or Active Galactic Nuclei (marked AGN in the figure) or from interactions of ultra-energetic protons with the 2.7\\,K cosmic microwave background (marked \"cosmogenic\") \\cite{Berezinsky-Zatsepin}. These cosmic neutrinos will hopefully be detected by neutrino telescopes in this decade, even though predictions for their fluxes are uncertain by orders of magnitude in many cases. This review is about neutrinos related to cosmic rays. First ideas to detect extraterrestrial high energy neutrinos data back to the end of the fifties, i.e. we look back to a journey of more than fifty years. I will focus to the first four decades and keep the developments of the last decade comparatively short. I refer the reader to the 2011 review of the field \\cite{Katz-Spiering-2011} for more detailed information on actual results and plans for future detectors. ", "conclusions": "" }, "1207/1207.6957_arXiv.txt": { "abstract": "{} {We aim at characterizing the inward transition from convective to radiative energy transport at the base of the convective envelope of the solar-like oscillator \\object{HD~52265} recently observed by the {\\small CoRoT} satellite.} {We investigated the origin of one specific feature found in the \\object{HD~52265} frequency spectrum. We modelled the star to derive the internal structure and the oscillation frequencies that best match the observations and used a seismic indicator sensitive to the properties of the base of the envelope convection zone.} { The seismic indicators clearly reveal that to best represent the observed properties of \\object{HD~52265}, models must include penetrative convection below the outer convective envelope. The penetrative distance is estimated to be $\\sim0.95 H_P$, which corresponds to an extent over a distance representing $6.0$ per cents of the total stellar radius, significantly larger than what is found for the Sun. The inner boundary of the extra-mixing region is found at $0.800\\pm0.004\\ R$ where $R=1.3\\ R_\\odot$ is the stellar radius.} {These results contribute to the tachocline characterization in stars other than the Sun.} ", "introduction": "Low-mass main-sequence (MS) stars have convective envelopes in which the chemical elements are mixed on short time scales and - at least in the deeper parts of the envelope - the energy transport by fluid elements can be treated as an adiabatic process. In the standard description, convective zones (CZ) are regions where the Schwarzschild criterion is fulfilled. Their boundaries lie at the border where the adiabatic temperature gradient % equals the radiative one. % It is expected, however, that fluid elements penetrate into the adjacent radiative zone due to their inertia \\citep[see e.g.][]{1991A&A...252..179Z}. In the Sun, the transition region (tachocline) is believed to be the site where the dynamo originates. Penetrative convection (PC) corresponds to efficient convective heat transport - and material mixing - by downward flows that establish a close to adiabatic temperature stratification below which the downward plumes are no longer able to modify the temperature stratification that remains close to radiative. The extent of penetrative convection and/or overshoot in stars cannot be derived from first principles and is still largely unknown. A crucial point therefore is to find observational signatures of penetrative convection in low-mass MS stars. This can be achieved by means of asteroseismology. The abrupt change of energy transport from a convective to a radiative regime is visible in the sound speed profile and impacts the oscillations frequencies as well as some characteristic frequency spacings, which then show an oscillatory (periodic) behaviour \\citep{1990LNP...367..283G,1994MNRAS.268..880R}. Owing to the importance of the tachocline region, helioseismic studies have attempted to measure the extent of the PC in the Sun \\citep{1993ASPC...40...60B, 1994A&A...283..247M}. \\citet{2011MNRAS.414.1158C} have recently found that convective envelope overshoot is necessary over an estimated extent of $0.37 H_P$ ($H_P$ is the pressure scale height). The high-quality solar data and wide available range of values of mode degrees $\\ell$ allowed \\citeauthor{2011MNRAS.414.1158C} to also show that the transition between the convective and the radiative stratification must be smooth, intermediate between a classical - ballistic - overshoot formulation and a no-overshoot one. Solar-like oscillations have been identified in many stars first from the ground, then from space by the {\\small CoRoT} \\citep{2002ESASP.485...17B} and Kepler \\citep{2010cosp...38.2513K} high-precision photometry missions, leading to an accuracy in frequency measurements of a few tenths $\\mu$Hz. Theoretical studies of the effects on the oscillations frequencies of PC at the base of a CZ have been conducted using low-degree modes, the only ones expected to be detectable in stars \\citep[see e.g.][ and references therein]{2009A&A...493..185R}. The period of the oscillatory component is found to be related to the location of the discontinuity inside the star and its amplitude to depend on the height of the discontinuity. One then expects that the amplitude of the oscillatory signal grows with an increasing extent of PC that causes a larger jump from the adiabatic temperature gradient to the radiative one below. Moreover, for a larger PC extent, the discontinuity is located deeper inside the star and the period of the oscillation is expected to be longer \\citep{1993ASPC...42..173R}. \\citet{2011A&A...530A..97B} analysed the {\\small CoRoT} oscillation spectrum of the star \\object{HD~52265} (\\object{HIP~33719}), a high-metallicity G0V star hosting an exoplanet. They found a typical p-mode solar-like spectrum, and identified 28 reliable low-degree p-modes of degrees $\\ell=0, 1, 2$ and order $n$ in the range $14-24$ (see their Table\\ 4). The frequencies $\\nu_{n, \\ell}$ are in the range $1500-2550\\ \\mu$Hz with a frequency at maximum amplitude $\\nu_\\mathrm{max}=2090\\pm 20\\mu$Hz. The error on each frequency is a few tenths of $\\mu$Hz. Such a high quality data set enables to probe the interior structure of the star. Here we report on one evidence for penetrative convection below the upper convective region of \\object{HD~52265} as indicated by its internal structure modelling. ", "conclusions": "We have detected an oscillatory signal with a significant amplitude in the variation of the separations $rr_{01/10}(n)$ with frequency for the {\\small CoRoT} star \\object{HD~52265}. A comparison with the same signal arising from appropriate stellar models shows that it cannot be reproduced unless penetrative convection is included at the base of the outer convective envelope. This is the first time that such a feature is firmly detected in a star other than the Sun. A best fit of the signal provides a measure of the extent of the mixed region below the CZ, in terms of the proxy $d_\\mathrm{ov}=0.95\\pm0.08~H_p$. The periodic signal for \\object{HD~52265} is more pronounced than for the Sun \\citep{ 2011MNRAS.414.1158C}, with a longer period, and so is the measured extent of PC in terms of $H_p$ and normalized radius. We point out that \\object{HD~52265} is similar to the Sun in all aspects except for the higher metallicity. Therefore, to understand the amplitude difference between the two stars, it is important to investigate the impact of metallicity on the structure and dynamics of the tachocline. {Progress is expected in a near future when seismic missions will provide similar observations for many stars spanning a wide metallicity range.} With the above results, we provided an additional clue that the physical description of the turbulent convection must be improved in stars, particularly at interfaces with radiative regions. Future modelling of the dynamics in the region of the tachocline will have to comply with our findings as well as with the results previously obtained for the Sun." }, "1207/1207.2165_arXiv.txt": { "abstract": "We investigate the change in stellar magnetic topology across the fully-convective boundary and its effects on coronal properties. We consider both the magnitude of the open flux that influences angular momentum loss in the stellar wind and X-ray emission measure. We use reconstructed maps of the radial magnetic field at the stellar surface and the potential-field source surface method to extrapolate a 3D coronal magnetic field for a sample of early-to-mid M dwarfs. During the magnetic reconstruction process it is possible to force a solution towards field geometries that are symmetric or antisymmetric about the equator but we demonstrate that this has only a modest impact on the coronal tracers mentioned above. We find that the dipole component of the field, which governs the large-scale structure, becomes increasingly strong as the stellar mass decreases, while the magnitude of the open (wind-bearing) magnetic flux is proportional to the magnitude of the reconstructed magnetic flux. By assuming a hydrostatic and isothermal corona we calculate X-ray emission measures (in magnitude and rotational modulation) for each star and, using observed stellar densities as a constraint, we reproduce the observed X-ray saturation at $Ro \\le 0.1$. We find that X-ray rotational modulation is not a good indicator of magnetic structure as it shows no trend with Rossby number but can be useful in discriminating between different assumptions on the field geometry. ", "introduction": "Stellar coronal magnetic activity is predominantly investigated in the X-ray and radio bands. A linear relationship between the X-ray luminosity, $L_X$, and the radio luminosity, $L_R$, established by \\citet{Guedel_X-ray_Microwave_1993}, pointed to a close connection between the hot coronal plasma responsible for the thermal X-ray emission and the free electrons causing the non-thermal radio emission, \\begin{equation} LogL_{X} \\approx LogL_{R} + 15.5 . \\label{eq.Lx-Lr} \\end{equation} While this linear relationship holds for several classes of active, main sequence star, between types F and early-M, and is also verified for T Tauri stars and solar flares, the assumption that the $L_X - L_R$ relationship would continue into the realm of the very low-mass stars proved to be short lived; the detection of radio emission from an M9 dwarf, LP944-20, defied this relation by almost four orders of magnitude \\citep{Berger_Discovery_2001}. Since the discovery of radio emission from LP944-20, many more observations have been carried out on low-mass stars to reveal that they also show the same radio phenomenon \\citep{Berger_flaring_2002,Berger_Magnetic_2005,Burgasser_Quiescent_2005,Hallinan_RadioObservations_2006,McLean_RadioSurvey_2011}. Radio observations are important diagnostic tools as they can help determine the magnetic field configuration \\citep{Hallinan_ECMI_2008} and the nature of plasma emitting regions. X-ray observations on the other hand help provide a deeper understanding of the generation of the magnetic field. Originating principally in the magnetically confined plasma of the hot stellar corona - at temperatures greater than $10^{6}K$ \\citep{Guedel_XrayReview_2004}- X-ray emissions prove useful for determining stellar parameters, such as the extent of the stellar corona, which also aids in establishing the structure and evolution of the magnetic field. X-ray emission has shown to be prevalent throughout the main sequence, with the maximum value of X-ray luminosity observed for each spectral type decreasing with bolometric luminosity, (e.g. \\citealt{Barbera_LxLbol_1993}) - for spectral type down to early-M. K and M stars are the most numerous in the Galaxy, and coupled with their high levels of X-ray emission, they are prime candidates for investigating the dependence of X-ray emission on physical parameters such as mass, rotation rate, and more recently, magnetic topology. So far, the X-ray luminosity has been shown to correlate well with either rotational velocity or Rossby number (the ratio of the stellar rotation period, P, to the convective turnover time, $\\tau_{c}$) (e.g.,\\citealt{Mangeney_XrayRotation_1984,Marilli_Chromospherid-coronal_1984,Schmidt_EinsteinXray_1985, Fleming_RelationXrayRotation_1989}); in general, $L_{X}/L_{bol}$ increases and then saturates ($L_{X}/L_{bol} \\approx10^{-3}$; \\citep{Delfosse_LxLbolM7_1998}) with increasing rotation rate. This behaviour is attributed to coronal saturation \\citep{Vilhu_Chromospheric_1987,Stauffer_radial_1994}; however, it is not clear whether saturation is a reflection of the dynamo itself or the total filling of the stellar surface with active regions \\citep{Vilhu_Magnetic_1984}.\\\\ For a small sample of stars later than M5, the $L_{X}/L_{bol}$ ratio was also shown to decline once again at the lowest Rossby numbers within the sample stars (e.g. \\citealt{Golub_Quiescent_1983,Bookbinder_PhDThesis_1985,Rosner_X-ray_1985,Jeffries_coronalSaturation_2011}) and \\citet{Berger_radio_2006} found that the value drops to $L_{X}/L_{bol} \\approx10^{-3.5}$ at spectral type M7. This indicates that outwith flaring, the emission from the thermal plasma cannot exceed a certain fraction of the total stellar flux. This \\textit{super-saturation} has been attributed to negative dynamo feedback \\citep{Kitchatinov_NegativeFeedback_1994}, lack of coverage of active regions \\citep{Stepien_Supersaturation_2001}, or centrifugal stripping of the corona \\citep{Jardie_CoronalStripping_2004}. It is important to note that there is a change in internal structure between the early-M dwarfs (dM) and mid-M dwarfs. The former have a solar-like internal structure: a turbulent, electrically conducting convective envelope, surrounding a radiative core. The interface between these two zones i.e. \\textit{the tachocline}, is the site where amplification of the magnetic field is believed to take place \\citep{Spiegel_tachocline_1992}. The relative size of the radiative core dramatically drops with decreasing temperature for early M dwarfs, and at spectral type M4 ($\\approx 0.35M_\\odot$), the stars are fully convective (e.g. \\citealt{Chabrier_Structure_1997}). This change in internal structure coincides with a change in magnetic topology (\\citealt{Donati_EarlyM_2008,Morin_MidM_2008,Morin_LateM_2010} hereafter D08, M08 and M10, respectively). In the higher mass stars, $M \\ge 0.45M_\\odot$, the field configuration is more complex, i.e. a non-axisymmetric poloidal field with a strong toroidal component, in comparison to the nature of the field structure at spectral types later than M4, which is mainly axisymmetric and poloidal. As for magnetic activity, it is crucial to look at coronal properties such as emission measure, especially for fast rotators, to explore whether saturation and super-saturation occur at a fixed rotation period or Rossby number. This could potentially help in determining the physical mechanism that is responsible for the saturation, i.e. an effect of the dynamo efficiency, or a consequence of the fast rotation rates exhibited by these stars. Stellar models (e.g. \\citealt{Gilliland_Chromospheric_1986,Kim_Rossby_1996}) suggest that the convective turnover time is longer in lower mass stars; if one considers the Sun and an M-dwarf with an equivalent rotation period, then the M dwarf has a smaller Rossby number and yet is more active in terms of $L_{X}/L_{bol}$. Although coronal saturation occurs at $L_{X}/L_{bol} \\approx10^{-3}$ for spectral type G,K and M, the rotation period at which it sets in is larger for the lower masses. Therefore, this results in a Rossby number of 0.1 being the point where coronal saturation sets in \\citep{Patten_Evolution_1996,Pizzolato_ActivityRotation_2003,Jeffries_coronalSaturation_2011}. However, \\citet{Jeffries_coronalSaturation_2011} show that the super-saturation is more clear as a function of rotation period within each spectral type, rather than Rossby number, with the effect occurring at periods of $P \\le 0.3$days for K-dwarfs and $P \\le 0.2$days for M dwarfs. Magnetic activity and rotation are well correlated \\citep{Stauffer_Rotation-Activity_1997} and both are strongly dependent on stellar age for main-sequence solar-like stars \\citep{Skumanich_AgeActivityRotation_1972}, \\begin{equation} \\Omega_{0} \\propto t^{-1/2} , \\label{eq.age-rotation} \\end{equation} where $\\Omega_{0}$ is stellar angular velocity. During their lifetime, rotational evolution of stars can be governed by disc braking, pre-main-sequence contraction i.e. \\textit{spin-up}, and/or magnetic winds. From observations of young open clusters (e.g.\\citealt{Irwin_rotation_2009}), it is clear that rapidly rotating dM stars are a common occurrence; however, in older open clusters, this number is reduced. \\citet{Scholz_RotationPeriods_2011} find a mass-dependent exponential rotational braking law \\begin{equation} P \\propto exp[t/\\tau] , \\label{eq.rotational-braking} \\end{equation} where towards lower masses the spin-down timescale $\\tau$ increases, with $\\tau \\approx 0.5$Gyr for $0.3M_{\\odot}$ and $\\tau > 1$Gyr for $0.1M_{\\odot}$. This indicates that at an age of 600Myr stars of $0.3M_{\\odot}$ are almost exclusively fast rotators. As well as X-ray emission, another coronal tracer of magnetic activity is radio emission. The relation that correlates the X-ray luminosity with the radio luminosty (eqn. ~\\ref{eq.Lx-Lr}), suggests that thermal X-ray emission, assumed to be due to hot coronal plasma, and non-thermal radio emission, possibly generated by the electron cyclotron maser instability \\citep{Melrose_Dulk_EMC_1982} or gyrosynchrotron, are both consequences of the magnetic field. Throughout the subtypes of dM stars, the radio luminosity remains approximately constant, while $L_{X}$ tracks the bolometric luminosity (i.e. $L_{X}/L_{bol} \\approx10^{-3}$). The sharp deviation from the $L_{X}-L_{R}$ relation, first noted by \\citet{Berger_Discovery_2001}, does not present itself until beyond spectral subtype M7 \\citep{Berger_Basri_2010}, where the relation evolves from $L_{R}/L_{X} \\approx 10^{-15.5}$ to $\\approx 10^{-11.5}$. This would indicate that the deviation must not directly relate to the transition to full convection. Along with this increase in radio luminosity, it is still unclear as to why the emission is present on these stars during one set of observations and then absent during the next (e.g. \\citealt{Berger_Basri_2010}). In order to investigate the role of field topology on the coronal structure and emission properties of M dwarfs, we use the observed surface magnetic field maps of their coronae. From this we can predict the X-ray emission measure. ", "conclusions": "\\label{sec.Summary} We have used reconstructed maps of the radial magnetic field at the stellar surface for a sample of early-to-mid M dwarfs to extrapolate their 3D coronal magnetic field (using the PFSS method). We have investigated the topology of the large-scale magnetic field at the stellar surface and the structure of the extrapolated 3D corona. By assuming a hydrostatic, isothermal corona, we have modelled the density structure within the corona and hence determined the X-ray emission measure. We have focussed in particular on variations with Rossby number. We find the following: 1. As the Rossby number decreases, the polar field strength of the dipole component of the field increases and then appears to saturate. Stars with low Rossby numbers have strong, mainly dipolar fields. 2. A similar variation with Rossby number is seen in both the total (unsigned) surface magnetic flux and the flux of open field (which can carry the stellar wind). The role of the topology of the large-scale field is apparent when we calculate the magnitude of the open flux. This is significantly less than would be predicted if all of the surface magnetic flux were contained in a purely dipole field. The contribution of the higher multipoles therefore reduces the open flux and may also significantly influence the angular momentum loss rate, which for a Weber-Davies model scales as the square of the open flux. Both the strength and also the topology of the large-scale field are therefore important in angular momentum loss. 3. As is observed, a rise and then saturation of the X-ray emission measure with decreasing Rossby number is also found. The stellar coronae are compact, with most of the emission originating from regions below approximately 1.5 stellar radii. Our sample contains a large range of both the inclinations of the stellar rotation axis and also the tilt of the magnetic axis. As a result, there is a large spread in the values of the rotational modulation of the X-ray emission and no clear trend with Rossby number can be detected. For low-mass stars, the observed variation in X-ray emission with Rossby number results naturally from the observed variation in the surface magnetic field. We note that we choose the parameter $\\kappa$ that scales the surface pressure in such a way that we reproduce typical X-ray fluxes and we do not model the ionisation fraction sometimes invoked in the atmospheres of later spectral types. While there is a range of magnetic topologies within our sample, the spread of values for the rotational modulation of the X-ray emission is too great for it to be a useful indicator of the field structure. Although not a good indicator of field structure when considering a large sample of stars with different inclination, it could be useful when considering a single object. For example, checking the symmetry and providing independent confirmation of the predominant mode (e.g. dipole versus quadrupole), or observing a magnetic cycle with a change from predominantly quadrupolar to predominantly dipolar as on the Sun \\citep{Sanderson_Sun_2003}. We find that both the strength of the field and its geometry, however, affect the magnetic flux that is open (wind-bearing) and which therefore allows the star to lose mass and angular momentum. The magnitude of this open flux is significantly reduced by departures from a purely dipolar field. This suggests that simple scalings for angular momentum losses based on dipolar field geometries may not be sufficient to explain the angular momentum evolution of low mass stars. The high values of open flux for stars with the lowest Rossby numbers may indicate that they have significant angular momentum loss rates. \\begin{table*} \\caption{Data for stellar sample of early-to-mid M dwarfs. Mass, radius, rotation period, inclination and, where available, $B_{V}$, the average large-scale magnetic field derived from spectropolarmetric measurements are provided by \\citet{Donati_EarlyM_2008,Morin_MidM_2008}. $B_{I}$, is the average magnteic field (i.e. small + large-scale field) derived from unpolarised spectroscopy, supplied by \\textit{a)} \\citet{Reiners_Basri_MagneticTopology_2009}, \\textit{b)} \\citet{Saar_Recent_1996}, \\textit{c)} \\citet{ReinersBasri_FirstDirect_2007}, \\textit{d)} \\citet{JohnsKrull_Valenti_2000}. Rossby number are from \\citet{Donati_EarlyM_2008,Morin_MidM_2008} and were computed from empirical $\\tau_{c}$ suited to the stellar mass from \\citet{Kiraga_Stepien_MDwarfs_2007}. $\\beta$, the estimated angle between the rotation and magnetic axis, are from this paper, along with the predicted values for emission measure (both magnitude and rotational modulation) and coronal density. \\label{tab.stellardata}} \\centering \\begin{tabular}{cccccccccccccc} \\hline Star & Sp Type & Mass ($M_\\odot$) & Radius ($R_\\odot$) & P (days) &Ro($10^{-2}$) &i($^{\\circ}$) &$\\beta_{M}$ ($^{\\circ}$) & $B_{V}$ (kG) &$B_{I}$ (kG)&LogEM ($cm^{-3}$)&Rot Mod &Log$\\overline n_{e}$ ($cm^{-3}$) \\\\ \\hline \\hline GJ 182 & M0.5 & 0.75 & 0.82 & 4.35 & 7.44 & 60 & 41.1 & 0.172 & $2.5^{a}$ & 50.33&12.18&8.62\\\\ DT Vir & M0.5 & 0.59 & 0.53 & 2.85 & 9.2 & 60 & 83.6 & 0.149 & $3.0^{b}$ &50.91&2.82& 9.26\\\\ DS Leo & M0 & 0.58 & 0.52 & 14.0 & 43.8 & 60 & 41.1 & 0.087 & - & 48.36&9.93&7.98\\\\ GJ 49 & M1.5 & 0.57 & 0.51 & 18.6 & 56.4 & 45 & 10.9 & 0.027& - & 46.83&18.02&7.05\\\\ OT Ser & M1.5 & 0.55 & 0.49 & 3.40 & 9.7 & 45 & 12.1 & 0.123 & - &50.66& 48.83&9.22\\\\ CE Boo & M2.5 & 0.48 & 0.43 & 14.7 & 35.0 & 45 & 7.4 & 0.10 & $1.8^{a}$ & 49.47&2.17&8.52\\\\ AD Leo & M3 &0.42 & 0.38 &2.3399 & 4.7 & 20 & 4.5& 0.19& $2.9^{c}$&50.22&3.27&8.95\\\\ EQ Peg A & M3.5 & 0.39 & 0.35 & 1.061 & 2.0 & 60 & 25.9 & - & - & 51.42&51.02&9.66\\\\ EV Lac & M3.5 & 0.32 & 0.30 & 4.3715 & 6.8 & 60 & 45.8 & 0.53 & $3.9^{d}$ &52.02&12.05&10.14\\\\ YZ CMi & M4.5 & 0.31 & 0.29 & 2.7758 & 4.2 & 60 & 24.8 & 0.50 & $\\ge 3.9^{c}$ &50.21&9.89& 10.27\\\\ V374 Peg & M4 & 0.28 & 0.32 & 0.44565 & 0.6 & 70 & 9.3 & - & - &52.62&26.49&10.26\\\\ EQ Peg B & M4.5 & 0.25 & 0.25 & 0.404 & 0.5 & 60 & 6.5 & -& - & 51.11&14.05&9.64\\\\ \\\\ \\hline \\end{tabular} \\end{table*}" }, "1207/1207.0483_arXiv.txt": { "abstract": "{Intriguing features in the angular distribution of the cosmic microwave background (CMB), such as the north-south asymmetry, were reported in the one- and three-year Wilkinson Microwave Anisotropy Probe (WMAP) data and should be studied in detail. We investigate some of these asymmetries in the CMB temperature angular distribution considering the $\\Lambda$CDM model in the three, five and seven year WMAP data.} {We aim to analyze the four quadrants of the internal linear combination (ILC) CMB maps using three different Galactic cuts: the WMAP KQ85 mask, a $|b|<10^\\circ$ Galactic cut, and the WMAP KQ85 mask $+$ $|b|<10^\\circ$ Galactic cut. } {We used the two-point angular correlation function (TPCF) in the WMAP maps for each of their quadrants. The same procedure was performed for 1000 Monte Carlo (MC) simulations that were produced using the WMAP team $\\Lambda$CDM best-fit power spectrum. In addition, we changed the quadrupole and octopole amplitudes obtained from the $\\Lambda$CDM model spectrum. We changed this to fit the quadrupole and octopole amplitudes to their observable values from the WMAP data. We repeated the analysis for the 1000 simulations of this modified $\\Lambda$CDM model, hereafter M$\\Lambda$CDM. } {Our analysis showed asymmetries between the southeastern quadrant (SEQ) and the other quadrants (southwestern quadrant (SWQ), northeastern quadrant (NEQ) and northwestern quadrant (NWQ)). Over all WMAP ILC maps, the probability for the occurrence of the SEQ-NEQ, SEQ-SWQ and SEQ-NWQ asymmetries varies from 0.1\\% (SEQ-NEQ) to 8.5\\% (SEQ-SWQ) using the KQ85 mask and the KQ85 mask $+$ $|b|<10^\\circ$ Galactic cut, respectively. We also calculated the probabilities for the M$\\Lambda$CDM using only the KQ85 mask and found no significant differences in the results. Moreover, the cold spot region located in the SEQ quadrant was covered with masks of 5,10 and 15 degrees radius and again the results remained unchanged. Furthermore, this analysis was repeated for random regions in the SEQ quadrant with a 15-degree mask and the SEQ quadrant still remained asymmetric with respect to the other quadrants of the CMB map.} {We found an excess of power in the TPCF at scales $>$ 100 degrees in the SEQ with respect to the other quadrants that is independent of the Galactic cut used. Moreover, we tested a possible relation between the Cold Spot and the SEQ excess of power and found no evidence for it. Finally, we could not find any specific region within the SEQ that might be considered responsible for the quadrant asymmetry.} ", "introduction": "After the cosmic microwave background (CMB) discovery by A. Penzias e R. Wilson \\citep{penzias1965}, several experiments were developed to characterize this radiation, leading to high-precision observations that raised the status of the CMB, which is since considered to be one of the main pillars of the $\\Lambda$CDM model. Particularly, the CMB cosmological fluctuations first detected by the Cosmic Background Explorer (COBE) satellite data \\citep{1992smoot} set up a new era in cosmology studies and paved the way to what today is called precision cosmology. However, detailed studies of the angular distribution of temperature fluctuations showed unexpected results if analyzed within the framework of the so-called cosmological concordance model ($\\Lambda$CDM model). These peculiar features in the CMB angular distribution were found for the first time in the COBE data and attracted much interest since then. A quadrupole amplitude smaller than that expected according to the $\\Lambda$CDM model was reported by the COBE team \\citep{1992smoot} and was confirmed by all WMAP data releases \\citep{2003bennett.1, 2007hinshaw,2009hinshaw,2011jarosik}. Other anomalies were found in the WMAP data that were not expected according to the $\\Lambda$CDM model either, such as the alignment between the quadrupole and octopole (e.g. \\citep{2004bielewicz,2004schwarz,2004copi,2004deoliveiracosta,2005bielewicz,2005land,2006copi,2006abramo,2010frommert,2010gruppuso}), the low quadrupole and octopole amplitudes (e.g., \\citep{2003mukherjee,2010ayata,2010cayon,2010cruz}), the north-south asymmetry (e.g., \\citep{2004eriksen,2004hansen,2004eriksen.2,2004hansen.2,2005donoghue,2009hoftuft,2010paci,2010pietrobon,2010vielva}), the anomalous alignment of the CMB features toward the Ecliptic poles (e.g., \\citep{2006wiaux,2007vielva}), and the cold spot (e.g., \\citep{2004vielva,2005cruz,2007cruz,2010vielva.2}). Recently, \\citet{2011aluri} analyzed in detail the signature of parity asymmetry first found by \\citet{2010kim} in the WMAP best- fit temperature power spectrum, confirming this asymmetry on a 3-$\\sigma$ level. \\citet{2011aluri} also concluded that their result is not due to residual foregrounds or to foreground cleaning. In addition, the preferred direction of the parity asymmetry coincides with the CMB kinematic dipole, showing that it may somehow be related to the quadrupole-octopole alignment \\citep{2011naselsky}. On the other hand, \\citet{2011bennett} reviewed the anomalies reported in the CMB temperature fluctuations and claimed that they are not statistically significant and, for this reason, do not in disagree with the $\\Lambda$CDM concordance model. Nevertheless, in this work we report an asymmetry in the WMAP temperature anisotropy data appearing in the two-point angular correlation function (TPCF) at scales above 100 degrees. In Section 2 we present our method to prepare the MC sky map simulations to confront them with real data and describe our TPCF method. In Section 3 we show our results. Finally, in Section 4 we present the discussion, followed by the conclusions in Section 5. ", "conclusions": "We found a significant asymmetry between the SEQ and the other quadrants by considering the temperature WMAP ILC maps. We calculated the probability of occurrence for this asymmetry using MC simulations, and 1 out of 1000 simulations for the $\\Lambda$CDM model corresponded to the SEQ-NEQ asymmetry found in the WILC7 using the KQ85 mask. We also showed that different Galactic cuts do not influence the result in a significant way, leading us to believe that this effect is not caused by the asymmetric mask. Moreover, the use of KQ85y7 preserves the same asymmetries (SEQ-NEQ, SEQ-SWQ and SEQ-NWQ), as expected. Considering all Galactic cuts and maps used, the highest probability of having an asymmetry is 8.5\\%, for the WILC5 with the KQ85 mask + data clipping for $|b|<10^\\circ$. The possibility that the asymmetries described in this work were related to the reported lack of power in the quadrupole and octopole was tested. We constructed simulations based on a modified $\\Lambda$CDM model, adjusting the amplitudes of the quadrupole and the octopole to their observational WMAP values. No explicit relation between the quadrant asymmetries and the low amplitude of the first two non-zero multipoles was found. Furthermore, we found no evidence of a relationship between the cold spot region and the SEQ excess of power, as pointed out by \\cite{2009bernui}, who suggested that the Cold Spot is responsible for 60\\% of the Southern Hemisphere power. Masking this region and some other regions in the SEQ does not change the TPCF in a significant way, leading us to conclude that the asymmetries between the SEQ and the other quadrants are not related to any specific region in the SEQ. We conclude that the excess of power found in the SEQ is likely related to the north-south asymmetry, in which the South Hemisphere presents more power than the Northern one (see \\citep{2004eriksen,2004hansen} and \\citep{2010paci,2010pietrobon,2010vielva} for recent discussions on the topic). Our results support the claims that there is indeed a north-south asymmetry and show that the excess of power occurs in the SEQ. Additional investigation is needed to find a better explanation for the north-south asymmetry. An explanation for these asymmetries is still missing. They could be primordial or caused by residual foregrounds or systematic effects. The upcoming CMB data from the Planck satellite when analyzed with more accurate foreground removal techniques will enable us to study in more detail the CMB anomalies reported in the literature. Finally, in addition to the SEQ excess of power, we can notice a lack of correlation in the SWQ, NWQ and NEQ in the TPCF in scales between 20 and 100 degrees for all Galactic cuts in the present work (see Figures \\ref{TPCF-NWQ-NEQ}-\\ref{TPCF-SWQ-SEQ-cut}). A lack of correlation in the TPCF was already reported by \\cite{2007copi} in scales above 60 degrees for a full sky analysis using the Kp0 mask in the first and third year of WMAP data. We would like to thank Paolo Cabella for useful discussions. We also acknowledge the use of HEALPix packages, of the Legacy Archive for Microwave Background Data analysis (LAMBDA). L. Santos thanks CAPES-Brazil for financial support. T. Villela acknowledges CNPq support through grant 308113/2010-1. C. A. Wuensche acknowledges CNPq grant 308202/2010-4." }, "1207/1207.0817_arXiv.txt": { "abstract": "We characterize the near-IR sky background from 308 observations with the FIRE spectrograph at Magellan. A subset of 105 observations selected to minimize lunar and thermal effects gives a continuous, median spectrum from 0.83 to 2.5 $\\mu$m which we present in electronic form. The data are used to characterize the broadband continuum emission between atmospheric OH features and correlate its properties with observing conditions such as lunar angle and time of night. We find that the moon contributes significantly to the inter-line continuum in the $Y$ and $J$ bands whereas the observed $H$ band continuum is dominated by the blended Lorentzian wings of multiple OH line profiles even at $R=6000$. Lunar effects may be mitigated in $Y$ and $J$ through careful scheduling of observations, but the most ambitious near-IR programs will benefit from allocation during dark observing time if those observations are not limited by read noise. In $Y$ and $J$ our measured continuum exceeds space-based average estimates of the Zodiacal light, but it is not readily identified with known terrestrial foregrounds. If further measurements confirm such a fundamental background it would impact requirements for OH-suppressed instruments operating in this regime. ", "introduction": "The development of low-noise HgCdTe focal plane arrays has motivated the construction of a new generation of medium-resolution, near-infrared spectrometers for ground-based telescopes \\citep{sim08,triplespec,xshooter}. These instruments resolve the well-known forest of sky emission line features induced by hydroxyl (OH) ions and other atmospheric molecules. For faint object spectroscopy, instrument sensitivity is therefore limited by the inter-line sky continuum, which is a superposition of terrestrial, astronomical, and instrumental backgrounds. Reliable estimates of the inter-line IR continuum are critical for the design of successful observations and instruments. Yet calibrated measurements remain sparse in the literature precisely because the strong foreground line emission makes such a calibration very challenging. As we shall demonstrate, the contrast ratio between narrow emission peaks and the broadband continuum approaches 5 magnitudes in the $J$ band and 7 magnitudes in $H$. Attempts to measure the inter-line background accordingly require certain characteristics of instrumentation and data processing. First, the spectral resolution must be sufficiently high to separate lines cleanly. Second, the detector must exhibit low dark current to not overwhelm the sky signal with thermal shot noise. Read noise should also be low, but it may be mitigated by averaging many exposures. Finally, and perhaps most importantly, attention must be given to the reduction of scattered light from optical surfaces in the instrument. The FIRE spectrograph at Magellan was designed to address these criteria while capturing both line and continuum emission of the sky across the $Y$, $J$, $H$, and $K$ bands simultaneously in each exposure. As part of its data processing pipeline, a precise model of the sky from 0.8 to 2.5 $\\mu$m is generated for each science exposure and saved to disk. In this paper, we gather these data and present calibrated sky brightness measurements taken from 308 deep exposures obtained across several observing runs since the commissioning of FIRE in March 2010. Section 2 describes the data processing techniques and instrumental considerations used to arrive at the continuum measurement. In Section 3, we examine correlations between the sky brightness and common observational conditions, including lunar phase, moon-object angle, and local time, and we present sky measurements under dark conditions which minimize their influence. Finally, in Section 4, we discuss the implications of our findings for the construction of new instrumentation, planning observations, and allocation of bright versus dark telescope time for near-IR spectroscopic observations. ", "conclusions": "We have presented a composite spectrum of the near-infrared night sky obtained by stacking 308 exposures obtained with Magellan/FIRE on 23 nights over seven observing runs. Careful attention was given to correcting electronic artifacts from the detector and scattered light in the instrument. Absolute flux calibration was taken from hot spectrophotometric standards, although the systematic error in the instrument efficiency is uncertain at the $\\sim 15\\%$ level because of slit losses. Analysis of a high signal-to-noise arc lamp composite indicates a small amount of OH line flux is scattered into broad Lorentzian wings at the level of $\\sim 10^{-4}$ times the peak line intensity. We explored measurements of FIRE's inter-line sky continuum and compared to previous estimates with the following main results: \\begin{enumerate} \\item{The mean of the inter-line continuum falls at $Y_{AB}=20.05\\pm0.04$, $J_{AB}=19.55\\pm0.03$, and $H_{AB}=18.80\\pm0.02$ (stat.) $\\pm 0.2$ (sys.) mag arcsec$^{-2}$. This is consistent with what is reported in \\citet{mai93} for the $H$ band.} \\item{The $H$ band background correlates most strongly with OH intensity since the $H$ band ``continuum'' is largely a superposition of scattered light from OH emission features. OH intensity decreases after sunset, giving a moderate correlation between $H$ continuum flux and time of night, but thermal emission also contributes. Our $H$ band measurement and that of \\citet{mai93} may not achieve the true $H$ continuum, which we extrapolate to being 19.4 AB mag arcsec$^{-2}$.} \\item{The $Y$ and $J$ band inter-line continuum fluxes correlate with moon-object angle and moon elevation, which are degenerate in our data set. Observations obtained at $\\lesssim 30$ degrees from a gibbous to full moon exhibit backgrounds 2.5 times brighter on average than those from a dark sky. This trend is much weaker in $H$ and not seen in $K$.} \\item{Under dark conditions, the $Y$ and $J$ continua are still higher than the predicted signal from scattered OH emission, and our measurements seem to detect a true broadband background. At $R\\sim 6000$, 60-80\\% of spectral bins achieve this background level. If real, it exceeds the average level of Zodiacal light measured by HST/NICMOS, yet it is not easily explained by known terrestrial foregrounds. } \\item{The most challenging programs in $Y$ and $J$ will benefit from obtaining telescope allocations during the dark half of the lunar cycle {\\em provided} that instruments are sufficiently sensitive (or exposures are sufficiently long) for sky noise to dominate the higher read noise of CMOS/HgCdTe detectors over CCDs.} \\end{enumerate}" }, "1207/1207.3814_arXiv.txt": { "abstract": "We investigate the means by which cold gas can accrete onto Milky Way mass galaxies from a hot corona of gas, using a new smoothed particle hydrodynamics code, `SPHS'. We find that the `cold clumps' seen in many classic SPH simulations in the literature are not present in our SPHS simulations. Instead, cold gas condenses from the halo along filaments that form at the intersection of supernovae-driven bubbles from previous phases of star formation. This {\\it positive feedback} feeds cold gas to the galactic disc directly, fuelling further star formation. The resulting galaxies in the SPH and SPHS simulations differ greatly in their morphology, gas phase diagrams, and stellar content. We show that the classic SPH cold clumps owe to a numerical thermal instability caused by an inability for cold gas to mix in the hot halo. The improved treatment of mixing in SPHS suppresses this instability leading to a dramatically different physical outcome. In our highest resolution SPHS simulation, we find that the cold filaments break up into bound clumps that form stars. The filaments are overdense by a factor of 10-100 compared to the surrounding gas, suggesting that the fragmentation results from a physical non-linear instability driven by the overdensity. This `fragmenting filament' mode of disc growth has important implications for galaxy formation, in particular the role of star formation in bringing cold gas into disc galaxies. ", "introduction": "The star formation rate (SFR) of the Universe has been rapidly falling since a redshift of $\\sim 2-3$ \\citep[e.g.,][]{LillyEtal1996, MadauEtal1998, HippeleinEtal2003}, dominated by the most massive galaxies and galaxy groups. By contrast, the Milky Way -- similarly to other disc galaxies -- has been forming stars at a near-continuous rate for the past $\\sim 8$\\,Gyr since $z \\sim 1$ \\citep[e.g.,][]{NohScalo1990, Rocha-PintoEtal2000}. The observed SFR of $\\sim 1-3 \\msun$ yr$^{-1}$ is hard to explain given the amount of cold gas present in the disc today; the Milky Way must have continuously accreted cold gas at a rate of $\\simgt 1\\,\\msun$ yr$^{-1}$ over this time \\citep{FraternaliTomassetti2012}. Merger-driven accretion appears to account for just $\\sim 0.1 \\msun $ yr$^{-1}$ \\citep{SancisiEtal2008}, and there is to date no evidence of a low redshift `cold flow' accretion mode \\citep[e.g.][]{StewartEtal2011}. This has led to two main solutions in the literature. The first is recycling of gas from existing stars in the disc through stellar winds \\citep[e.g.,][]{Roberts1963, Sandage1986, KennicuttEtal1994}. \\cite{LeitnerKravtsov2011} estimate that this recycled gas could contribute at least half of the global SFR for a galaxy of Milky Way mass at low redshift ($z \\simlt 0.5$). The second is accretion from a massive hot halo, or corona, of gas surrounding the Galaxy. Such a hot halo has not yet been directly observed, but several indirect lines of evidence exist: observations of X-ray emitting gas \\citep{GuptaEtal2012}; absorption along sight lines to quasars \\citep[e.g.,][]{WilliamsEtal2005, FangEtal2006, KacprzakEtal2008}; pulsar dispersion measures \\citep[e.g.,][]{2010ApJ...714..320A, GaenslerEtal2008}; a significant Galactic baryon deficiency when compared to the universal baryon fraction \\citep[e.g.,][]{FukugitaPeebles2006, NicastroEtal2008}; and evidence of ram pressure stripping of the Magellanic stream \\citep{MastropietroEtal2005} and other nearby dwarf galaxies \\citep{GrcevichPutman2009}. These studies give a hot halo mass of $> 5 \\times 10^9 - 1.5 \\times 10^{10} \\msun$ assuming a \\cite{NavarroFrenkWhite1996} (NFW) profile or $\\simgt 4 \\times 10^{10}$ assuming a flattened power-law profile \\citep{AndersonBregman2011}. Similar results are seen in other disc galaxies \\citep[e.g.,][]{2005RSPTA.363.2693R, MoEtal2005, SancisiEtal2008, AndersonBregman2011}. While it is likely that hot gaseous coronae surround disc galaxies, it is not clear \\emph{how} the gas cools and condenses out of these coronae and onto the disc to form stars. One particular mode of cold gas supply that has been seen in a number of numerical simulations \\citep{Sommer-Larsen2006, KaufmannEtal2006, KaufmannEtal2007, KaufmannEtal2009, PutmanEtal2009} but was initially proposed by \\cite{Nulsen1986}, is that of direct cooling from the halo via thermal instability. However, doubt has been cast on this picture by \\cite{MalagoliEtal1987} \\& \\cite{BinneyEtal2009} who show that hot haloes are linearly stable to density perturbations, making direct cooling unlikely. \\cite{JoungEtal2012} extend this treatment to the non-linear regime with dedicated numerical simulations, finding that while a runaway process of cooling and collapse is possible, it can only occur for overdensities of $\\simgt 10-20$ with respect to the local background density. These results suggest that if gas is to cool from the hot coronae then some mechanism is required to \\emph{excite} a thermal instability. \\cite{MarinacciEtal2011} suggest a mechanism whereby a galactic fountain seeds metal-rich gas into the metal-poor hot haloes, giving rise to a thermal instability that causes cold gas to rain down onto the disc in the form of $\\sim 10^5 \\msun$ clouds \\citep[see also][]{FraternaliBinney2008}. This model provides an excellent fit to both the kinematics and spatial distribution of warm HI gas in the Milky Way \\citep{MarascoEtal2012}, although it cannot account for the high-velocity clouds (HVCs) \\citep[e.g.,][]{SembachEtal2003, TrippEtal2003, CollinsEtal2005}. Alternative models include cooling stripped gas from dwarf galaxies or warm clouds at the disc-corona interface \\citep[e.g.,][]{PutmanEtal2009, HeitschPutman2009, Peek2009}. The `direct cooling' mode mentioned above is seen regularly in `classic'\\footnote{for a careful definition of `classic' SPH see Section \\ref{sec:classic}.} smoothed particle hydrodynamics (SPH) simulations \\citep[e.g.,][]{Sommer-Larsen2006, KaufmannEtal2006, KaufmannEtal2007, KaufmannEtal2009, PutmanEtal2009}, where a thermal instability leads to a sudden and widespread condensation of gas from the halo in the form of cold, dense clumps. If correct, no special mechanism would be required to explain the continued star formation of disc galaxies over the past $\\sim 8$\\,Gyrs. However, `classic' SPH is known to exhibit an artificial surface tension that inhibits mixing of different gaseous phases, leading to poor performance on a variety of hydrodynamic test problems \\citep{AgertzEtal2007, 2008JCoPh.22710040P, 2008MNRAS.387..427W, ReadEtal2010}. In recent work, some of the present authors have developed a new `flavour' of SPH -- SPHS -- that solves these problems, giving excellent performance and numerical convergence on a wide range of test problems \\citep{2012MNRAS.tmp.2941R}. In this paper, we use SPHS to revisit the problem of thermal instabilities in hot gaseous coronae. Our primary goals are to determine whether the instabilities seen in the classic SPH simulations are physical or numerical; and under what circumstances cold gas can condense out of a hot halo to fuel star formation in disc galaxies. The former goal is further motivated by the absence of such cold clumps in the hot haloes of galaxy formation simulations using adaptive mesh refinement (AMR) codes \\citep[e.g.,][]{AgertzEtal2009} or the recent moving Voronoi mesh code AREPO \\citep{VogelsbergerEtal2011}. This paper is organised as follows. In Section \\ref{sec:method}, we briefly review the SPHS algorithm and `classic' SPH. We describe our initial conditions, and present our implementation of radiative cooling, star formation and stellar feedback, and our treatment of a central supermassive black hole (SMBH). In Sections \\ref{sec:results} \\& \\ref{sec:fullpower} we present our results, which we discuss in Section \\ref{sec:discussion}. Finally, in Section \\ref{sec:conclusions}, we present our conclusions. ", "conclusions": "\\label{sec:conclusions} We have presented simulations of a cooling gaseous halo in a Milky Way mass galaxy, using this particular problem to perform the first scientific investigation with a new hydrodynamics code, SPHS. We have compared the results obtained (at identical spatial resolution) with that of the standard (`classic') SPH method employed in the literature, finding significant differences in the mode of gas cooling in the halo and subsequently the mode of disc feeding. The puzzle of the many cold clumps seen forming from the halo in many SPH simulations of galaxy formation is attributed to a numerical inability driven by unresolved mixing of different gas phases. We demonstrate both with our full simulations (Section \\ref{sec:differences}) and with a more idealised test (Section \\ref{sec:numericalclump}) that the removal of pressure blips in an otherwise smooth flow prevents the formation of the clumps, leading instead to the formation of cold filaments that feed the disc. The resulting galaxies in the SPH and SPHS simulations differ greatly in their morphology, gas phase diagrams, and stellar and gaseous disc/bulge ratio. We have explored in more detail the mode of disc feeding seen in our SPHS simulations, going to higher resolution and employing a kernel that gives improved force accuracy. We find a new way of bringing cold gas to the galactic disc; namely, the fragmentation and collapse through non-linear thermal instability of filament(s) formed at the intersection of supernovae-driven bubbles. The feeding rate of cold gas ($T \\simlt 10^4$ K) to the disc is found to be approximately a solar mass per year, which suggests this is a promising model for fuelling late-time star formation in real spiral galaxies. We emphasise that our focus in this paper was on understanding what drives thermal instabilities in hot halo gas, and in particular performing a comparison between the `classic' SPH and the SPHS numerical methods. Our numerical experiments are idealised and do not present a complete picture of galaxy formation. Nonetheless, by employing a hydrodynamics method that resolves the mixing of different gas phases, we find a novel mode of cold gas accretion and disc growth that may be very relevant for galaxy formation." }, "1207/1207.3025_arXiv.txt": { "abstract": "We combine for the first time all available information about the spectral shape and morphology of the radio halo of the Coma cluster with the recent $\\gamma$-ray upper limits obtained by the Fermi-LAT and with the magnetic field strength derived from Faraday rotation measures. We explore the possibility that the radio emission is due to synchrotron emission of secondary electrons. First we investigate the case of pure secondary models that are merely based on the mechanism of continuous injection of secondary electrons via proton-proton collisions in the intra-cluster medium. We use the observed spatial distribution of the halo's radio brightness to constrain the amount of cosmic ray protons and their spatial distribution in the cluster that are required by the model. Under the canonical assumption that the spectrum of cosmic rays is a power-law in momentum and that the spectrum of secondaries is stationary, we find that the combination of the steep spectrum of cosmic ray protons necessary to explain the spectrum of the halo and the very broad spatial distribution (and large energy density) of cosmic rays result in a $\\gamma$--ray emission in excess of present limits, unless the cluster magnetic field is relatively large. However this large magnetic field required to not violate present $\\gamma$--ray limits appears inconsistent with that derived from recent Faraday rotation measures. Second we investigate more complex models in which the cosmic rays confined diffusively in the Coma cluster and their secondary electrons are all reaccelerated by MHD turbulence. We show that under these conditions it is possible to explain the radio spectrum and morphology of the radio halo and to predict $\\gamma$-ray fluxes in agreement with the Fermi-LAT upper limits without tension with present constraints on the cluster magnetic field. Reacceleration of secondary particles also requires a very broad cosmic ray spatial profile, much flatter than that of the intracluster medium, at least provided that both the turbulent and magnetic field energy densities scale with that of the intracluster medium. However, this requirement can be easily alleviated if we assume that a small amount of (additional) seed primary electrons are reaccelerated in the cluster's external regions, or if we adopt flatter scalings of the turbulent and magnetic field energy densities with distance from the cluster center. ", "introduction": "Galaxy clusters host several potential accelerators of cosmic ray (CR) electrons and protons, from ordinary galaxies to active galaxies (AGN) and cosmological shock waves, driven in the intracluster medium (ICM) during the process of hierarchical cluster formation (see Blasi et al. 2007 for a review). The long lifetime of CR protons (or nuclei) in the ICM and the large geometrical size of the magnetized region of galaxy clusters make them efficient storage rooms for the hadronic component of CRs produced within their volume (V\\\"olk et al. 1996, Berezinsky et al. 1997, Ensslin et al. 1997). The accumulation of CRs inside clusters over cosmological times leads to the assumption that an appreciable amount of energy may be stored in the ICM in the form of non-thermal particles. If this energy is sufficiently high, the flux of $\\gamma$ radiation induced by the production and decay of neutral pions may reach potentially detectable levels, thereby providing us with a powerful diagnostic tool of the CR energy content of clusters (Colafrancesco \\& Blasi 1998, Blasi \\& Colafrancesco 1999, V\\\"olk \\& Atoyan 1999, Miniati 2003, Pfrommer \\& En\\ss lin 2004, Wolfe et al. 2008). So far only upper limits to the $\\gamma$-ray emission from galaxy clusters have been obtained (Reimer et al.~2003; Perkins et al. 2006, Aharonian et al. 2009a,b; Aleksic et al. 2010; Ackermann et al. 2010) \\footnote{See however Han et al. 2011 for Virgo}. These upper limits, together with several constraints from complementary approaches based on radio observations lead us to conclude that CR protons contribute less than a few percent of the energy of the ICM, at least in the central Mpc--size region (Reimer et al. 2004, Brunetti et al. 2007, 2008, Brown et al. 2011, Aleksic et al. 2011). CR electrons are very well traced in the ICM through their radio emission which appears in the form of diffuse (Mpc scale) synchrotron {\\it giant radio halos} from the cluster X-ray emitting regions, and {\\it relics}, typically in the clusters' peripheral regions (see Ferrari et al. 2008, Venturi 2011, for recent reviews on observations). \\noindent Giant radio halos are the most spectacular and best studied non-thermal large scale phenomena in the universe. They appear in about $1/3$ of the most massive galaxy clusters (e.g. Giovannini et al. 1999, Kempner \\& Sarazin 2001, Cassano et al. 2008), in a rather clear connection with dynamically disturbed systems, while ``off-state'' clusters (those with no evidence of diffuse emission) are generally more relaxed (Cassano et al. 2010a and references therein). The connection between cluster mergers and radio halos suggests that such emission traces the hierarchical cluster assembly and probes the dissipation of gravitational energy during the dark matter-driven mergers that lead to the formation of clusters. The physical mechanisms responsible for the generation and evolution of radio halos are still a matter of debate, but two main lines of thought have been developed throughout the years. \\noindent One is based on the idea that seed electrons may be re-accelerated by turbulence produced during merger events (Brunetti et al. 2001, Petrosian 2001). In this class of models the $\\gamma$-ray emission is predicted to be rather low, though it may be substantial under the hypothesis that the electron seeds are secondaries produced in inelastic collisions of a subdominant hadronic CR component in the ICM (Brunetti \\& Blasi 2005; Brunetti \\& Lazarian 2011). \\noindent The second line of thought is based on the idea of clusters as storage rooms of CR protons: the radio halos may be generated as a result of synchrotron emission of secondary electrons and positrons from pp collisions (Dennison 1980, Blasi \\& Colafrancesco 1999, Pfrommer \\& En\\ss lin 2004). On one hand this idea serves as a solution to the problem of the large spatial dimensions of the radio emitting region, larger than the typical loss length of electrons: secondary electrons and positrons are produced {\\it in situ} in inelastic CR collisions and radiation is produced near the production region. On the other hand, secondary models do not explain in a {\\it natural} way the observed association between cluster mergers and giant radio halos, in that CRs accumulate inside the ICM on cosmological time scales and not in direct connection with acceleration events such as those associated with mergers\\footnote{see however Ensslin et al. 2011, where diffusion/transport of CRs is studied under peculiar conditions}. One proposal is that during merger events the cluster magnetic field becomes larger than in the quiescent state, thereby turning the halo on (Kushnir et al. 2009, Keshet \\& Loeb 2010), although studies based on Rotation Measures of clusters' and background radio sources disfavour this scenario (see Bonafede et al.~2011a and ref. therein). In fact, some pieces of observations put tension on a {\\it pure} hadronic origin of radio halos. These include the steepening in the spectrum (or the very steep spectrum) of several halos (Schlickeiser et al. 1987, Thierbach et al. 2003, Reimer et al. 2004, Brunetti et al. 2008, Dallacasa et al. 2009, Giovannini et al. 2009, Macario et al. 2010, van Veeren et al. 2011) and the very large spatial extent of several halos (or their flat radio-brightness distribution) (Brunetti 2004, Murgia et al. 2009, Donnert et al. 2010a, Brown \\& Rudnick 2011); in all cases observations would imply that the energy budget of CR protons is uncomfortably large, at least assuming that clusters are magnetised at $\\sim \\mu$G level, consistent with the results of RM (see Bonafede et al.~2010 and references therein). \\noindent The most distinct prediction of models of radio halos that are based on secondary particles is the production of $\\gamma$--rays (e.g., Blasi \\& Colafrancesco 1999, Sarazin 2004, Pfrommer \\& En\\ss lin 2004, Brunetti 2009, Brunetti \\& Lazarian 2011, En\\ss lin et al. 2011). Nowadays there is agreement on the fact that the abundance of secondaries required to fit the spectrum of (at least some) radio halos should produce $\\gamma$--ray emission detectable with the sensitivity of the Fermi-LAT, assuming low/medium level of cluster magnetic field (eg. Marchegiani et al 2007, Pfrommer 2008, Brunetti 2009, Jeltema \\& Profumo 2011). Thus the combination of the available information on the spectrum and morphology of radio halos in conjunction with the limits on their $\\gamma$-ray emission provides a powerful tool to gain insights into the origin of the halo emission. \\noindent Following this pathway, in this paper we concentrate on the case of the Coma cluster, which represents a prototypical example of giant radio halos, with a wealth of data available on its spectrum and morphology. The incoming upper limits on the Coma $\\gamma$-ray emission with the Fermi-LAT telescope are invaluable in imposing stringent limits on the amount of cosmic rays that can be stored in the intracluster medium and serve as sources of secondary electrons and positrons. In particular in this paper we combine for the first time all available information about the spectral shape and morphology of the radio halo of the Coma cluster with the recent $\\gamma$-ray upper limits and magnetic field strengths derived from Faraday rotation measures. We show that the requirement of reproducing the properties of the Coma radio halo in the context of a pure secondary electron model leads to a large CR energetics and fluxes of $\\gamma$-rays which results in a tension with the existing upper limits from the Fermi-LAT (Ackermann et al.~2010). This situation is readily alleviated in the case of a large cluster magnetic field, however the magnetic fields required by the model are appreciably larger than those inferred from recent Faraday rotation measures. \\noindent This tension disappears when including the effect of turbulent reacceleration in combination with the process of injection of secondary particles. We adopt a physically motivated picture in which secondary products of cosmic ray interactions are reaccelerated by MHD turbulence during clusters mergers. In this sense we value the physical insight behind the concept of cosmic ray confinement and the production of secondary electrons, but we do not assume that these particles are ``directly'' responsible for the formation of the radio halo. We find that even reacceleration models of this type require a broad spatial distribution of the parent cosmic rays, but the expected $\\gamma$-ray fluxes are well consistent with the Fermi-LAT upper limits. The paper is organized as follows: we discuss the hadronic model of radio emission in \\S \\ref{sec:hadro} and the reacceleration model in \\S \\ref{sec:reacc}. A critical discussion of our results is provided in \\S \\ref{sec:disc}. A $\\Lambda$CDM cosmology ($H_{o}=70\\,\\rm km\\,\\rm s^{-1}\\,\\rm Mpc^{-1}$, $\\Omega_{m}=0.3$, $\\Omega_{\\Lambda}=0.7$) is adopted throughout the paper. ", "conclusions": "\\label{sec:disc} We presented a combined analysis of the spectrum and morphology of the giant radio halo in the Coma cluster, the available measurements of cluster magnetic fields from RM and the upper limits on the $\\gamma$-ray emission from this cluster as recently obtained by the Fermi-LAT telescope (Ackermann et al. 2010)\\footnote{see also Ando \\& Nagai (2012) and Han et al. (2012)}, in the context of models based on secondary electrons: we concentrated on a pure secondary electron model and on a scenario where secondaries are reaccelerated by MHD turbulence during mergers. \\subsection{Hadronic models} The pure secondary electron model is based on the concept of CR effective confinement in the ICM (V\\\"olk et al. 1996, Berezinsky et al. 1997) and consists of explaining the giant radio halo emission as the result of synchrotron emission of secondary electrons (and positrons) from inelastic collisions of cosmic ray protons with gas in the ICM. These collisions result in the production and decay of charged and neutral pions. The former lead to production of secondary electrons, while the latter provide a channel of continuous production of gamma radiation. This model has been widely implemented also in cosmological simulations. These simulations, that include, to some extent, CR physics and the acceleration of CRs at cosmological shocks provided a picture of the radio to $\\gamma$--ray properties of galaxy clusters, under the assumption that the emitting particles (electrons) are generated through pp collisions in the ICM or accelerated at shocks (e.g., Pfrommer et al. 2008 and references therein). The first simulations of this kind predicted that clusters would be potentially detectable in $\\gamma$-rays with present day $\\gamma$--ray telescopes (Miniati 2003, Pfrommer 2008). The most important assumption in these simulations is in the efficiency of particle acceleration at weak shocks that is poorly known (see Gabici \\& Blasi 2003, 2004 and Kang et al.~2007 for a critical view). More recently, numerical simulations of large-scale structure formation made an attempt to reconcile their results with the lack of detection of galaxy clusters in the $\\gamma$-ray band (Pinzke \\& Pfrommer 2010, see also Aleksic et al. 2010,11). Although these simulations provide expectations (still) consistent with the Fermi-LAT upper limits, it is worth mentioning that they do not allow to reproduce the observed properties of radio halos if we assume that halos originate via secondary emission, in particular their (very broad) spatial extent (Donnert et al. 2010a,b; Brown \\& Rudnick 2011). In this respect, for the sake of completeness, in Fig. \\ref{fig:numerical} we show a comparison between the brightness profile of the Coma radio halo and expectations based on CRs distributions from Pfrommer et al. (2008) and from Pinzke \\& Pfrommer (2010) simulations. This can be made by using the semi-analytical prescription of the spatial distribution of CR in galaxy clusters, as derived from high resolution numerical simulations of clusters (Pinzke \\& Pfrommer 2010), and using parameters of the Coma cluster to calculate the production rate of secondaries and the synchrotron emissivity. According to these simulations the ratio of the CR and thermal pressure in typical non-CC clusters is quasi constant up to a distance $R/R_{vir} \\sim 0.2$ and increases by a factor 2-3 up to the virial radius (see Fig.14 in Pinzke \\& Pfrommer 2010); such an increase is however mainly driven by the temperature decrement in the ICM at large distances. \\noindent In reality, the microphysics of the ICM and of CRs in galaxy clusters is very complicated and unfortunately beyond the capabilities of present simulations. For this reason in our paper we have carried out a analysis that does not depend on the way CR protons are generated in the ICM and on the complex processes of CR transport and reacceleration that can take place in the cluster volume. Rather than modeling these complex processes we indeed derived constraints on the spectral and spatial distributions of CR protons directly from the observed spectrum and morphology of the Coma halo, using parameters of the ICM from the X-ray observations (\\S 2.2). \\noindent We derived the azimuthally averaged brightness profile of the Coma halo and used it in order to obtain constraints on the spatial profile of the magnetic field strength and of CRs density as functions of the distance from the cluster center. Our profile is obtained using the deep WSRT observations from Brown \\& Rudnick (2011) and allows us to obtain unprecedented constraints on the non-thermal cluster properties on 3--3.5 $r_c$ scales. We find that fitting the volume averaged radio spectrum of the Coma halo and its azimuthal brightness distribution requires a rather steep spectrum of CR protons, $s\\sim 2.6$, in the ICM with a spatial distribution much flatter than the spatial distribution of the thermal gas. This leads to an exceedingly large energy content in the form of CRs confined in the ICM and correspondingly large $\\gamma$-ray emission, exceeding the Fermi-LAT upper limits. Moreover the spectral steepening observed at high frequency can only marginally be explained for the cases with steeper slope ($s \\geq 3$) and the combined effect of the SZ decrement. This case is however the one that violates the Fermi-LAT upper limits more clearly. \\noindent If we neglect the spectral steepening in our analysis, the model can potentially explain radio observations. However the Fermi-LAT $\\gamma$--ray limits set a lower limit to the strength of the cluster magnetic field that is in disagreement with a recent analysis of Faraday rotation measures (Bonafede et al. 2010); the more so for steeper spectra of the parent CR protons. \\noindent {\\it Pure} secondary electron models of the Coma halo are thus disfavoured when all available observational data are considered. In this respect any further improvement of upper limits over the next several years, or even a $\\gamma$-ray detection of the Coma cluster (i.e. with flux 2-3 times smaller the Ackermann et al. limits) will conclusively establish the incompatibility of a hadronic origin of the Coma radio halo with the radio (including RM) and $\\gamma$-ray observations. \\begin{figure} \\begin{center} { \\includegraphics[width=0.79\\textwidth]{syn_brigh_SB_S1p6_NEWNEW_PP.ps}} \\end{center} \\caption{Azimuthal averaged brightness profile of the Coma halo (as in Fig.~1) compared with expectations based on numerical simulations that include the acceleration of CR protons and the generation of secondary electrons in the ICM. Points show the expected synchrotron profile from secondary electrons in the massive cluster gs72 from Pfrommer et al. 2008 (circles mark the case of radiative simulations). The solid line (black) show the expectations based on the semi-analytic model of CR protons in Coma based on numerical simulations (Pinzke \\& Pfrommer 2010). A magnetic field profile $B(r)^2 \\propto \\epsilon_{ICM}$ and $B_0 = 5 \\mu$G are assumed.} \\label{fig:numerical} \\end{figure} \\subsection{A comment on the most relevant assumptions} \\noindent A discussion of the assumptions adopted in our analysis is in order. Conclusions derived above are based on canonical assumptions used to model non-thermal emission in galaxy clusters. The most notable assumption is that the CR spectrum in the cluster volume is a power law in momentum, $N(p) \\propto p^{-s}$, and that the spectrum of secondary particles emitting in the radio band can be calculated under stationary conditions. We constrained the value of the slope $s$ from the spectrum of the radio halo in the frequency range $\\sim$0.1-1 GHz (Fig. \\ref{fig:hadroComa}) and the $\\gamma$-ray emission of the cluster is calculated by using such values of $s$. In principle, the shape of the spectrum of CR protons might also change with location, moving away from cluster core toward the periphery, although there is no strong indication that this may be happening from the total radio spectrum of the Coma halo\\footnote{Indeed a mixture of constributions from different power-law spectra results in a ``concave'' shape of the total synchrotron spectrum.}. In this case the expected $\\gamma$--ray emission is reduced if the spectrum of CR protons is flatter in the regions where the majority of $\\gamma$--rays are produced, at distances 2--3.5 $r_c$ from the center. This however would result in a halo spectrum that is flatter in the external regions, in contrast with present observations that suggest a ``radial spectral steepening'' (Giovannini et al 1993, Deiss et al 1997, this is also evident from the comparison between the synchrotron brightness distributions of the halo at 350 and 1400 MHz, Fig. 1). \\noindent The situation may change if the spectrum of CRs becomes harder at low enough energies. In this case one may expect that the flux of low energy $\\gamma$-rays (from decays of neutral pions) is also reduced allowing for some room for pure hadronic models to accomodate current observational constraints. Present radio data limit the possibility of a flattening in the radio spectrum to frequencies $\\nu \\leq 100$ MHz, leading to a limit to the energy where a possible break occurs in the spectrum of the primary CR protons, $E_b \\leq 10-20$ GeV. Under these conditions we note that a substantial reduction of the $\\gamma$-ray luminosity due to $\\pi^0$--decay in the 1-10 GeV band, necessary to make hadronic models still consistent with --at least-- the limits derived from the first 18 month of FERMI-LAT data (Ackermann et al. 2010) (Fig. \\ref{fig:gamma}), would require a prominent CR spectral break, by $\\Delta s \\ge 0.5$. The two points that disfavour pure hadronic models for the radio halo are the steepening of the halo spectrum observed at higher frequencies and the inconsistency between the properties of the cluster magnetic field constrained by Faraday RM and by Fermi-LAT limits under the assumption of a hadronic origin of the halo. The latter point however might simply indicate that RM and synchrotron emission trace different magnetic fields. As already mentioned in Sect. 2.3 the simplest way to have RM and synchrotron emission sample different fields is to assume that some of the RM come from regions adjacent to and influenced by the radio sources, in this case however the magnetic field in the ICM as inferred from the analysis of RM would likely be biased to higher values of the magnetic field (Rudnick \\& Blundell 2003) thus making the inconsistency between magnetic fields even stronger. On the other hand, if one assumes that RM originates entirely in the ICM, RM and synchrotron emission may sample different fields in the case of a highly inhomogeneous fields, because the synchrotron emissivity depends non-linearly on the magnetic field and its fluctuations. Positive fluctuations however increase also the radiative losses of particles (assuming they vary on time scale sufficiently long, $> 10^7-10^8$yrs\\footnote{The minimum RM scale of about 2 kpc for Coma radio galaxies derived by Bonafede et al 2010, and the typical fractional polarization of about 10\\%, imply that any ICM fluctuation on smaller scales, that could be very intermittent, must be weak.}) and this is expected to partially quench the expected boosting of the synchrotron emissivity of electrons generated in regions with medium/high fields. Under (at least quasi--) stationary conditions the emissivity reads: \\begin{equation} _{Volume} = \\propto {\\hat B}^{1+\\alpha} {{ ( 1 + (\\delta B/{\\hat B})^2)^{(1+\\alpha)/2} } \\over{ {\\hat B}^2 + \\delta B^2 + B^2_{cmb} }} \\, , \\label{jsyn2} \\end{equation} \\noindent where we introduced a magnetic field made of a large scale component ${\\hat B}$ and a turbulent component $\\delta B$, such that $<\\bf{ {\\hat B}} + \\bf{\\delta B}>_{Volume}= {\\bf {\\hat B}}$ and $<\\bf{ {\\hat B}} + \\bf{\\delta B}>^2_{Volume}= {\\hat B}^2 + \\delta B^2$. Under these assumptions we estimate that the ratio of $\\gamma$--ray and radio cluster luminosities decreases by (only) a factor $\\sim 1.7$, compared to results using the formalism in Sect.~2.1, if we assume $\\delta B^2 \\sim {\\hat B}^2$, with $\\delta B^2$ anchored to the best fit value from RM studies (normalisation and spatial profile). On the other hand, RMs suggest that $\\delta B\\gg B$, in which case Eq.\\ref{jsyn2} becomes equivalent to Eq.\\ref{jsyn} in Sect.~2.1\\footnote{ In the analysis by Bonafede et al. the magnetic field is parameterised with a power spectrum $B^2 = 8\\pi \\int P_B(k) dk$ between a maximum and minimum scale and its properties are constrained by comparing observations with simulated Faraday Rotation maps derived for different parameters}. In this sense the question of whether Faraday RM and synchrotron emission measure the same magnetic field only lies in the capability of RM studies to provide a good description of the ICM magnetic field. Future radio telescopes, including LOFAR and SKA, will greatly improve the sensitivity to RMs allowing to detect and use many tens of background radio sources per clusters to sample magnetic fields along many lines of sight, thus removing potential biases in present studies. \\noindent Finally, for the sake of completeness we mention that in principle, spatial diffusion of particles in inhomogeneous fields may also affect the value of the ratio synchrotron/$\\gamma$-ray luminosity in hadronic models, if the diffusion time necessary to cover the spatial scales on which field inhomogeneities occur is smaller than both the life-time of particles and the time-scale of magnetic field (local) variations. This however depends on the details of the magnetic field and diffusion model. \\subsection{Beyond the pure hadronic model: turbulent reacceleration of secondaries} In \\S \\ref{sec:reacc} we went beyond the pure hadronic model and discussed the case of the reacceleration model were secondary particles are reaccelerated by MHD turbulence. We find that this model allows to obtain a good description of the radio spectrum and morphology of the Coma cluster and at the same time they may easily be compatible with the Fermi-LAT upper limits on the $\\gamma$-ray emission from this cluster and with the measured magnetic field strength as inferred from RMs. \\noindent In the simplified assumption that both the ratios of turbulent and thermal energy density, $\\epsilon_{tur}/\\epsilon_{ICM}$, and of magnetic field and thermal energy density , $\\epsilon_{B}/\\epsilon_{ICM}$, are constant in the radio emitting volume, the brightness profile of the Coma halo leads us to infer that the spatial distribution of CRs in the cluster should be very broad, flat (or slightly increasing in the external regions) on the halo size-scale. This is because the particular reacceleration model adopted in this paper faces drawbacks that are in part similar to those of a pure hadronic model for the origin of the halo. This is simply because the seed electrons for reacceleration are generated by proton-proton collisions in the ICM. A flat spatial distribution of CRs is a fairly strong requirement, although possible support for a rather flat distribution of CRs, significantly broader than that of the thermal ICM, comes from very recent numerical cosmological simulations that include CRs accelerated at shocks (Vazza et al.~2012). Present radio data do not constrain the energy density of CR protons on scales larger than the halo size-scale. However if we speculate that the flat spatial distribution of CR protons extends on larger scales the resulting energy budget of CRs in the cluster would be significantly larger than that required by our modeling. In general several, effects may contribute to mitigate this situation. First, numerical simulations show that the turbulent energy density (and the ratio $\\epsilon_{tur}/\\epsilon_{ICM}$) increases outside the cores of simulated clusters (Vazza et al. 2009, 11; Iapichino \\& Niemeyer 2008, Iapichino et al. 2011), implying a synchrotron brightness profile potentially flatter than that calculated under the assumption of a constant ratio $\\epsilon_{tur}/\\epsilon_{ICM}$. Second, we limit our analysis to the case $B^2 \\propto \\epsilon_{ICM}$ ($\\eta =0.5$), that gives the best fit to Faraday RM. On the other hand for smaller values of $\\eta$ the spatial distribution of CR protons that is required by the model to match the observed synchrotron brightness profile of the halo is significantly steeper with radius, although still flatter than the distribution of the thermal energy density of the cluster. Finally we note that the situation is greatly mitigated as soon as one would relax the assumption of having only secondary electrons in the ICM. Indeed primary electrons accelerated at shocks, or during the activity of cluster galaxies and AGN, can survive a substantial fraction of the Hubble time in the cluster outskirts (e.g. Sarazin 1999). These primaries provide a natural population of seed electrons to reaccelerate in a turbulent ICM (Brunetti et al 2001, Petrosian 2001) and may significantly contribute to the synchroton emission in the external regions of giant radio halos. If a substantial contribution to the radio halo emission comes from the reacceleration of primary electrons, the expected $\\gamma$--ray emission from the Coma cluster becomes even smaller than that calculated in \\S \\ref{sec:reacc}. \\noindent Ongoing and future observations in the deep non-thermal (high and very-high energy) regime of the Coma cluster will therefore provide precious information on its non-thermal content." }, "1207/1207.1485_arXiv.txt": { "abstract": "We test models for the evolution of neutron star (NS) magnetic fields (B). Our model for the evolution of the NS spin is taken from an analysis of pulsar timing noise presented by Hobbs et al. (2010). We first test the standard model of a pulsar's magnetosphere in which B does not change with time and magnetic dipole radiation is assumed to dominate the pulsar's spin-down. We find this model fails to predict both the magnitudes and signs of the second derivatives of the spin frequencies ($\\ddot{\\nu}$). We then construct a phenomenological model of the evolution of $B$, which contains a long term decay (LTD) modulated by short term oscillations (STO); a pulsar's spin is thus modified by its B-evolution. We find that an exponential LTD is not favored by the observed statistical properties of $\\ddot{\\nu}$ for young pulsars and fails to explain the fact that $\\ddot{\\nu}$ is negative for roughly half of the old pulsars. A simple power-law LTD can explain all the observed statistical properties of $\\ddot{\\nu}$. Finally we discuss some physical implications of our results to models of the $B$-decay of NSs and suggest reliable determination of the true ages of many young NSs is needed, in order to constrain further the physical mechanisms of their $B$-decay. Our model can be further tested with the measured evolutions of $\\dot{\\nu}$ and $\\ddot{\\nu}$ for an individual pulsar; the decay index, oscillation amplitude and period can also be determined this way for the pulsar. ", "introduction": "% Many studies on the possible magnetic field decay of neutron stars (NSs) have been done previously, based on the observed statistics of their periods ($P$) and period derivatives ($\\dot{P}$). Some of these studies relied on a population synthesis of pulsars (Holt \\& Ramaty 1970; Bhattacharya et al. 1992; Han 1997; Regimbau \\& de Freitas Pacheco 2001; Tauris \\& Konar 2001; Gonthier et al. 2002; Guseinov et al. 2004; Aguilera et al. 2008; Popov et al. 2010); however, no firm conclusion can be drawn on if and how the magnetic fields of NSs decay (e.g., Harding \\& Lai 2006; Ridley \\& Lorimer 2010; Lorimer 2011). Alternatively some other studies used the spin-down or characteristics ages, $\\tau_{\\rm c}=(P-P_{0})/2\\dot{P}$ ($P_{0}$ is the initial spin period of the given pulsar), or $\\tau_{\\rm c}=P/2\\dot{P}$ if $P\\gg P_{0}$, as indicators of the true ages of NSs and found evidence for their dipole magnetic field decay (Pacini 1969; Ostriker \\& Gunn 1969; Gunn \\& Ostriker 1970). However, as we have shown recently (Zhang \\& Xie 2011), their spin-down ages are normally significantly larger than the ages of the supernova remnants physically associated with them, which in principle should be the unbiased age indicators of these NSs (Lyne 1975; Geppert et al., 1999; Ruderman 2005). It has been shown that this age mismatch can be understood if the magnetic fields of these NSs decay significantly over their life times, since the magnetic field decay alters the spin-down rate of a NS significantly, and thus cause the observed age mismatch (Lyne 1975; Geppert et al. 1999; Ruderman 2005; Zhang \\& Xie 2011). Pulsars are generally very stable natural clocks with observed steady pulses. However significant timing irregularities, i.e., unpredicted times of arrival of pulses, exist for most pulsars studied so far (see Hobbs et al. 2010, for an extensive reviews of many previous studies on timing irregularities of pulsars). The timing irregularities of the first type are ``glitches\", i.e., sudden increases in spin rate followed by a period of relaxation; it has been found that the timing irregularities of young pulsars with $\\tau_{\\rm c}<10^5$ years are dominated by the recovery from previous glitch events (Hobbs et al. 2010). In many cases the NS recovers to the spin rate prior to the glitch event and thus the glitch event can be removed from the data satisfactorily without causing significant residuals over model predictions. However in some cases, glitches can cause permanent changes to both $P$ and $\\dot P$ of the NS, which cannot always be removed from the data satisfactorily, for instance, if the glitch has occurred before the start of the data taking and/or long term recoveries from glitch events cannot be modeled satisfactorily. These changes can be modeled as a permanent increase of the surface dipole magnetic field of the NS; consequently some of these NSs may grow their surface magnetic field strength gradually this way and eventually have surface magnetic field strength comparable to that of magnetars over a life time of 10$^{5-6}$ years (Lin \\& Zhang 2004). These are the first studies that linked the timing irregularities of pulsars with the long term evolution of magnetic fields of NSs. The timing irregularities of the second type have time scales of years to decades and thus are normally referred to as timing noise (Hobbs et al. 2004; 2010). Hobbs et al. (2004, 2010) carried out so far the most extensive study of the long term timing irregularities of 366 pulsars. Besides ruling out some timing noise models in terms of observational imperfections, random walks, and planetary companions, Hobbs et al. (2004, 2010) concluded that the observed $\\ddot{\\nu}$ values for the majority of pulsars are not caused by magnetic dipole radiation or by any other systematic loss of rotational energy, but are dominated by the amount of timing noise present in the residuals and the data span. Some of their main results that are modeled and interpreted in this work are: (1) All young pulsars have $\\ddot{\\nu} > 0$; (2) Approximately half of the older pulsars have $\\ddot{\\nu} > 0$ and the other half have $\\ddot{\\nu} < 0$; (3) Quasi-oscillatory timing residuals are observed for many pulsars; and (4) The value of $\\ddot{\\nu}$ measured depends upon the data span and the exact choice of data processed. Proper understanding of these important results can lead to better understanding of the internal structures and the magnetospheres of NSs. Since the spin evolution of pulsars cannot be caused by any systematic loss of rotational energy (Hobbs et al. 2004; 2010), it is natural to introduce some oscillatory parameters. Indeed, some of the observed periodicities have been suggested as being caused by free-precession of the NS (Stairs et al. 2000). Some strong quasi-periodicities have been identified recently as due to abrupt changes of the magnetospheric regulation of NSs, perhaps due to varied particle emissions (Lyne et al. 2010). It was suggested that such varied magnetospheric particle emissions is also responsible for the observed long term timing noise (Liu et al. 2011). Alternatively it has been suggested that the propagation of Tkachenko waves in a NS may modulate its moment of inertia, and thus produce the observed periodic or quasi-periodic timing noise (Tkachenko 1966; Ruderman 1970; Haskell 2011). Because all of the above processes can only cause periodic or quasi-period modulations to the observed spin of NSs, all these modulations can be modeled phenomenologically as some sort of oscillations of the observed surface magnetic field strengths of NSs. It is thus reasonable to attribute the observed wide-spread quasi-oscillatory timing residuals to the oscillations of the magnetic fields of NSs. As shown in this work, the oscillating magnetic fields can also naturally explain the fact that approximately half of the older pulsars have $\\ddot{\\nu} > 0$ and the other half have $\\ddot{\\nu} < 0$, when the long term magnetic field decay does not dominate the spin evolutions of pulsars. This is the main advantage of this model, and thus justifies invoking oscillating magnetic fields for modeling the long term spin evolutions of pulsars. In this work we attempt to use the known statistical properties of the spin evolution of the pulsars to test some phenomenological models for the evolution of the dipole magnetic field of a NS. In sections 2-5, we test the following three models: (1) the standard magnetic dipole radiation model with constant magnetic field, (2)** long-term exponential decay with an oscillatory component and (3) long-term power-law decay modulated with oscillations. We find the last model is the only one compatible with the data, with the power-law index $\\alpha>0.5$ and the oscillation amplitude of around $5\\times 10^{-5}$ for an oscillation period of several years. Finally we discuss the physical implication of our results, summarize our results and make conclusions and discussions, including predictions of the model for further studies. All timing data of pulsars used in this work are taken from Hobbs et al. (2010). ", "conclusions": "In this work we have tested models of magnetic field evolution of NSs with the statistical properties of their timing noise reported by Hobbs et al. (2010). In all models the magnetic dipole radiation is assumed to dominate the spin-down of pulsars; therefore different models of their magnetic field evolution lead to different properties of their spin-down. Our main results in this work are summarized as follows: (1) We tested the standard model of pulsar's magnetosphere with constant magnetic field. We found that this model under-predicts the second derivatives of their spin frequencies ($\\ddot{\\nu}$) by several orders of magnitudes for most pulsars (Fig.~(\\ref{Fig:1})). This model also fails to predict $\\ddot{\\nu}<0$, which is observed in nearly half of the old pulsars. The standard model of pulsar's magnetosphere with constant magnetic field is thus ruled out, and $\\ddot{\\nu}$ is found to be dominated by the varying magnetic field. (2) Statistically about half of the pulsars in this sample show increasing or decreasing magnetic fields. The magnetic field strength of most pulsars are found to either increase or decrease with time scales significantly shorter than their characteristic ages (Fig.~(\\ref{Fig:2})). (3) We constructed a phenomenological model of the evolution of the magnetic fields of NSs, which is made of a long term decay modulated by short term oscillations. (4) A simple exponential decay is not favored by the apparent correlation of the inferred decay constant with their characteristic ages for young pulsars (Fig.~(\\ref{Fig:3})(b)) and fails to explain the fact that $\\ddot{\\nu}$ is negative for roughly half of old pulsars. (5) A simple power-law decay, however, can explain all the observed statistical properties of $\\ddot{\\nu}$ (Figs.~(\\ref{Fig:5} and \\ref{Fig:6})) and infer a narrow distribution of an oscillation amplitude of about $5\\times 10^{-5}$ with a preferred decay index of about unity (Figs.~(\\ref{Fig:7} and \\ref{Fig:8})). We thus conclude that long term magnetic field and short term oscillations are ubiquitous in all NSs and that the data support the model consisting of a power-law decay modulated by short term oscillations for the evolution of their magnetic fields. However we can only constrain the decay index to be $\\alpha\\geq0.5$ (Figs.~(\\ref{Fig:4}) and (\\ref{Fig:8})(b)). Our results are consistent with the models of magnetic field decay we discussed, in which ohmic dissipation, Hall effect, and ambipolar diffusion are possible effects causing magnetic field decay. However, the uncertainty in the decay index $\\alpha$ makes it difficult to constrain further or distinguish between these effects. The problem in determining $\\alpha$ is the unavailability of the true ages of pulsars. With the true age $t$ available for some pulsars, $\\alpha$ can be determined directly with Eq.~(\\ref{ddot_p1}); $f=0$ can be safely assumed for young pulsars. Unfortunately only several pulsars in the sample of Hobbs et al (2010) have the ages of their supernova remnants available, thus preventing a statistically meaningful determination of $\\alpha$ this way. We thus suggest that reliable determination of the true ages of a large sample of young NSs is needed, in order to constrain further the physical mechanisms of the magnetic field decay of NSs. This can also provide an independent test of our model. The main prediction of this model is that the spin evolution of each pulsar should follow, \\begin{equation}\\label{ddot_p2} \\ddot{\\nu}(t)\\simeq 2\\dot{\\nu}(t)(-\\frac{\\alpha}{t}+\\sum f_i\\cos(\\phi_i+2\\pi\\frac{t}{T_i})), \\end{equation} where both $\\ddot{\\nu}(t)$ and $\\dot{\\nu}(t)$ should be time-resolved, rather than the averaged value conventionally reported over the whole data span. Hobbs et al. (2010) have shown that both the value and sign of the averaged $\\ddot{\\nu}$ depend on the time span used, which can be reproduced with our model presented here (Zhang \\& Yi 2012). However a clear-cut and robust test of our model is to compare the prediction of the above equation. As discussed above, the first or the second term in the above equation may be ignored for old millisecond pulsars or young pulsars, respectively. The current on-going programs of intensive long term precise timing monitoring of pulsars should produce the required data for validating our model. In doing so, both the decay index and the oscillating properties of each individual pulsar can be determined, which in turn can provide important insights into the physics of NSs." }, "1207/1207.0204_arXiv.txt": { "abstract": "Identifications of quasars at intermediate redshifts ($2.2 3.5$) (Fan et al. 2001a,b; Richards et al. 2002). However, in the redshift range $2.22.2$ quasars to study cosmic baryon acoustic oscillation (BAO) (White 2003; McDonald \\& Eisenstein 2007) and using these quasars to construct the accurate luminosity function to study quasar evolution in the mid-redshift universe (Wolf et al. 2003; Jiang et al. 2006), we need to explore other more efficient ways to identify the $2.22.2$ quasar density in S82 than that based on optical color selected quasars (Palanque-Delabrouille et al. 2011; Ross et al. 2012). However, only for very limited sky areas the multi-epoch observational data have been publicly available, so at present the variability methods can not be broadly used for selecting quasars over a large sky area. Another possible way for separating $z>2.2$ quasars from stars is to utilize their near-IR colors. Due to the different radiative mechanisms of stars and quasars, the continuum emission from stars usually has a blackbody-like spectrum and decreases more rapidly from optical to near-IR wavelengths than that of quasars, which usually display a power-law spectra over a broad range of wavelength plus thermal emissions from accretion disks and dusts. This leads to obvious color differences in near-IR band between stars and quasars, even though their optical spectra are similar. Because of this difference, a K-band excess technique has been proposed for identifying quasars at $z>2.2$ (e.g. Warren, Hewett \\& Foltz 2000; Croom, Warren \\& Glazebrook 2001; Sharp et al. 2002; Hewett et al. 2006; Chiu et al. 2007; Maddox et al. 2008,2012; Smail et al. 2008; Wu \\& Jia 2010). Based on a sample of 8498 quasars and a sample of 8996 stars complied from the photometric data in the ugriz bands of SDSS and YJHK bands of UKIRT InfraRed Deep Sky Surveys (UKIDSS\\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).}), Wu \\& Jia (2010) proposed an efficient empirical criterion, (i.e. $\\rm Y-K>0.46(g-z)+0.82$, where YK magnitudes are Vega magnitudes and $\\rm gz$ magnitudes are AB magnitudes) for selecting $z<4$ quasars. A check with the VLA-FIRST (Becker, White \\& Helfand 1995) radio-detected SDSS quasars, which are thought to be free of color selection bias (see McGreer, Helfand \\& White (2009), however), also proved that with this Y-K/g-z criterion they can achieve the completeness higher than 95\\% for these radio-detected quasars at $z<3.5$, which seems to be difficult when using the SDSS optical color selection criteria alone where two dips around $z\\sim 2.7$ and $z\\sim 3.4$ obviously exist (Richards et al. 2002,2006; Schneider et al. 2007,2010). Recently, Peth, Ross \\& Schneider (2011) extended the study of Wu \\& Jia (2010) to a larger sample of 130,000 SDSS-UKIDSS selected quasar candidates and re-examined the near-IR/optical colors of them. Although by combining the variability and optical/near-IR color we may achieve the maximum efficiency in identifying $2.22.2$ quasars (Wu et al. 2010a,b; Wu et al. 2011), more efforts are still needed to check whether using our Y-K/g-z criterion can help us to discover more quasars at $2.22.2$ quasars. The recently released Wide-field Infrared Survey Explorer (WISE) all-sky data (Wright et al. 2010) also provided abundant photometric data in the near(middle)-IR bands, which will be very helpful for quasar selections (Wu et al. 2012; Stern et al. 2012; Edelson \\& Malkan 2012; Yan et al. 2013) Fortunately, several ongoing optical and near-IR photometric sky surveys will also provide us further oppotunities to apply our optical/near-IR color selections of quasars to larger and deeper fields. In addition to SDSS III (Eisenstein et al. 2011), which has taken 2,500 deg$^2$ further imaging in the south galactic cap, the SkyMapper (Keller et al. 2007) and Dark Energy Survey (DES; The Dark Energy Survey Collaboration 2005) will also present the multi-band optical photometry in 20,000/5,000 deg$^2$ of the southern sky, with the magnitude limit of 22/24 mag in $i$-band, respectively. The Visible and Infrared Survey Telescope for Astronomy (VISTA; Arnaboldi et al. 2007) is carrying out the VISTA Hemisphere Survey (VHS) in the near-IR YJHK bands for 20,000 deg$^2$ of the southern sky with a magnitude limit at K=20.0, which is about five and two magnitude deeper than the Two Micron ALL Sky Survey (2MASS; Skrutskie et al. 2006) and UKIDSS/LAS limits (Lawrence et al. 2007), respectively. Therefore, the optical and near-IR photometric data obtained with these ongoing surveys will provide us a large database for quasar selections. Needless to say, the ongoing Panoramic Survey Telescope \\& Rapid Response System (Pan-STARRS; Kaiser et al. 2002) and the future Large Synoptic Survey Telescope (LSST; Ivezic et al. 2008) will also provide us with multi-epoch photometry in multi-bands covering a large area of the sky, which will undoubtedly help us to construct a much larger sample of quasars based on both optical/near-IR colors and variability features. On the other hand, the spectroscopic observations are still crucial to determine the quasar nature and redshifts for the quasar candidates selected from the optical/near-IR colors. The ongoing SDSS-III/BOSS is expected to obtain the spectra of 150,000 quasars at $2.23$ at \\mbox{95\\%~CL~\\cite{ Smith:2011es, Hamann:2011ge, Archidiacono:2011gq, Hamann:2011hu, Nollett:2011aa}.} Given these hints for extra energy density in the early universe and bearing in mind the expectation that the Planck experiment will soon provide a significantly improved measurement of $\\Neff$ at photon decoupling (e.g.~see~\\cite{Galli:2010it}), we consider it a pertinent time to consider models in which $\\Neff$ is increased. The conventional way is to introduce extra `dark' radiation, which leads to the same value of $\\Neff$ at BBN and photon decoupling (see~e.g.~\\cite{Nakayama:2010vs,Feng:2011uf,Dreiner:2011fp}). While it is usually assumed that $\\Neff$ does not change between BBN and photon decoupling, intriguingly, the experimental data are consistent with a slightly larger value of $\\Neff$ at photon decoupling. This increase in $\\Neff$ after BBN can be achieved through decays to dark radiation during or after BBN \\mbox{(see~e.g.~\\cite{Ichikawa:2007jv, Fischler:2010xz, Hasenkamp:2011em, Menestrina:2011mz,Hooper:2011aj, Bjaelde:2012wi})} or from primordial gravitational waves~\\cite{Smith:2006nka}. In this paper we explore another scenario that has received less attention in the literature, which predicts a value of $\\Neff$ greater than three, and moreover, predicts that $\\Neff$ at photon decoupling is larger than the value during BBN. As we will show in section~\\ref{Section:Current}, $\\Neff$ scales as $(T_{\\nu}/T_{\\gamma})^4$ so increasing the neutrino-to-photon temperature ratio leads to an increase in $\\Neff$. A relative increase in $T_{\\nu}/T_{\\gamma}$ can be achieved by reducing the photon temperature relative to the neutrino temperature or by increasing the neutrino temperature relative to the photon temperature. A decrease in $T_{\\gamma}$ can be obtained through the production of hidden photons \\cite{Jaeckel:2008fi,Foot:2011ve} while a new light mediator may lead to a change in either $T_{\\gamma}$ or~$T_{\\nu}$~\\cite{Blennow:2012de}. In this paper we follow the example of the standard cosmological model to increase $T_{\\nu}$: in the standard cosmological model, the photons are hotter than the neutrinos because the transfer of entropy from the electrons and positrons to the photons (when they become non-relativistic) happens after the neutrinos decouple from the electromagnetic plasma at $T_{\\rm{D}}\\approx2.3$~MeV~\\cite{Enqvist:1991gx}. In an analogous way, we show that an additional `generic' particle $\\chi$ that remains in thermal equilibrium solely with the neutrinos after they decouple from the electromagnetic plasma and until it is non-relativistic, reheats the neutrinos relative to the electromagnetic plasma and as a result, leads to a higher value for $\\Neff$. Requiring that the neutrino reheating, which happens when $\\chi$ becomes non-relativistic and transfers its entropy to the neutrinos, happens after neutrino decoupling implies that $\\chi$ must have a mass $m_{\\chi}\\lesssim {\\rm{few}}\\cdot T_{\\rm{D}}\\sim10$ MeV. The goal of this paper is to investigate the current and future constraints on $\\chi$ from the inferred values of $\\Neff$ from BBN and the CMB . The only condition that we impose on the `generic' particle $\\chi$ is that the neutrino-$\\chi$ interaction rate is sufficiently high that they remain in thermal equilibrium until $\\chi$ is non-relativistic. Therefore, the constraints we derive may apply to, for instance, light dark matter particles that obtain their abundance from thermal freeze out (see e.g.~\\cite{Boehm:2000gq, Boehm:2002yz, Boehm:2006mi,Farzan:2009ji}) or light mediators that have large couplings to neutrinos (see e.g.~\\cite{Beacom:2004yd, Hannestad:2004qu}). The parameter that determines $\\Neff$ is $m_{\\chi}$ as this dictates the additional energy density, so it is this parameter that we will constrain. During BBN, the photon temperature is similar to $m_{\\chi}$ so $\\chi$ makes a direct contribution to $\\Neff$. In addition, there may also be an indirect contribution from the increase in $T_{\\nu}/T_{\\gamma}$. At photon decoupling, we assume $\\chi$ is non-relativistic so its direct contribution to $\\Neff$ is Boltzmann suppressed; the increase in $\\Neff$ arises solely from the increase in $T_{\\nu}/T_{\\gamma}$. As we show in section~\\ref{Section:Current}, this difference in the origin of the extra energy density leads to a larger value of $\\Neff$ at photon decoupling. This paper is organised as follows. In section~\\ref{Section:Current} we discuss the impact of light particles in thermal equilibrium with neutrinos on the energy density and neutrino-to-photon temperature ratio (some results are fully derived in an Appendix). Using these results, we first find the effect of $\\chi$ on the abundance of primordial nuclei produced during BBN. This has previously been studied in~\\cite{Kolb:1986nf, Serpico:2004nm}. Here, we independently calculate the primordial abundance of $\\He$ and $\\DHy$ as a function of $m_{\\chi}$ and compare with recent experimental measurements. Our calculation uses a more recent BBN code than \\cite{Kolb:1986nf, Serpico:2004nm} and includes the most recent values of the neutron lifetime and baryon density. We then calculate the value of $\\Neff$ as a function of $m_{\\chi}$ at photon decoupling and compare with the experimental values inferred by ACT and SPT. Turning to consider experimental results from the near future, in section~\\ref{Section:Planck} we forecast the constraints that will soon be placed on $\\chi$ from Planck's measurement of $\\Neff$. Finally, we conclude in section~\\ref{Section:Conclusion}. ", "conclusions": "\\label{Section:Conclusion} Motivated by discrepancies in the determination of $\\Neff$ from BBN and the CMB, we have considered the impact of a new light particle which remains in thermal equilibrium with neutrinos until it becomes non-relativistic, which we assume occurs before photon decoupling. Such a particle leads to extra energy density in the early universe, either directly or through its effect on the ratio of the neutrino-to-photon temperature. To quantify the impact of such particles on BBN, we updated the analysis of \\cite{Kolb:1986nf, Serpico:2004nm} through modifications of the $\\PArthENoPE$ code to take into account the two effects of light particles on the energy density mentioned above. These particles lead to an increased value of $\\Neff$ and bring the $\\He$ abundance into better agreement with recent observations (see the upper panel of figure~\\ref{fig:YpDH}). We also showed (in the lower panel of figure~\\ref{fig:YpDH}) that the $\\rm{D}$ abundance is compatible with recent measurements. We then considered the effect of such particles at the time of formation of the CMB. At this epoch, the increase in $\\Neff$ comes solely from the increase in the neutrino-to-photon temperature ratio, as $\\chi$ is non-relativistic. As we showed in figure~\\ref{fig:NeffCMB}, this brings the value of $\\Neff$ into better agreement with the values reported by CMB experiments. In general, we found that the value of $\\Neff$ at the time of CMB formation is larger than the value at BBN. This is demonstrated explicitly in figure~\\ref{fig:Neffcomp}. Fixing $Y_p$ to the value inferred from $\\Neff$ at photon decoupling, as is often done, is not applicable when setting constraints on this scenario. As the current experimental constraints on our scenario are still relatively weak, we forecast the sensitivity of the Planck satellite to the effect of these particles on the CMB. We find that Planck is highly sensitive to the effects of such particles. If the value of $\\Neff$ derived from Planck agrees with that of standard cosmology with three neutrinos, we showed in the upper panel of figure~\\ref{fig:Planck} that it will rule out (for example) the existence of a light Majorana fermion with a mass of less than 7.4 MeV which couples to the neutrinos. The lower panel of figure~\\ref{fig:Planck} indicates that it will be difficult to distinguish between our scenario and the standard case of dark radiation (in which $\\Neff$ is the same at BBN and photon decoupling) at much more than $1\\sigma$ when only data from astrophysical measurements of $\\He$ and Planck is considered. On the other hand, considering only particles in thermal equilibrium with neutrinos, we demonstrated in the lower panel of figure~\\ref{fig:Planck} that as well as providing significantly improved constraints on $m_{\\chi}$, in some regions of parameter space, Planck is even able to discriminate between a real scalar and a complex scalar or Majorana fermion." }, "1207/1207.6491_arXiv.txt": { "abstract": "{From a dynamical analysis of the orbital elements of transneptunian objects (TNOs), Ragozzine \\& Brown reported a list of candidate members of the first collisional family found among this population, associated with (136\\,108) Haumea (a.k.a. 2003~EL$_{61}$).} {We aim to distinguish the true members of the Haumea collisional family from interlopers. We search for water ice on their surfaces, which is a common characteristic of the known family members. The properties of the confirmed family are used to constrain the formation mechanism of Haumea, its satellites, and its family.} {Optical and near-infrared photometry is used to identify water ice. We use in particular the $CH_4$ filter of the Hawk-I instrument at the European Southern Observatory Very Large Telescope as a short $H$-band ($H_S$), the $(J - H_S)$ colour being a sensitive measure of the water ice absorption band at 1.6 $\\mu m$.} {Continuing our previous study headed by Snodgrass, we report colours for 8 candidate family members, including near-infrared colours for 5. We confirm one object as a genuine member of the collisional family (2003~UZ$_{117}$), and reject 5 others. The lack of infrared data for the two remaining objects prevent any conclusion from being drawn. The total number of rejected members is therefore 17. The 11 confirmed members represent only a third of the 36 candidates.} {The origin of Haumea's family is likely to be related to an impact event. However, a scenario explaining all the peculiarities of Haumea itself and its family remains elusive.} ", "introduction": "\\indent The dwarf planet (136\\,108) Haumea \\citep{2005-IAUC-8577-Santos-Sanz} is among the largest objects found in the Kuiper Belt \\citep{2006-ApJ-639-Rabinowitz, 2008-SSBN-3-Stansberry}, together with Pluto, Eris, and Makemake. It is a highly unusual body with the following characteristics: \\begin{enumerate} \\item It has a very elongated cigar-like shape \\citep{2006-ApJ-639-Rabinowitz, 2010-AA-518-Lellouch}. \\item It is a fast rotator \\citep[$P_{rot} \\sim 3.9$\\,h,][]{2006-ApJ-639-Rabinowitz}. \\item It has two non-coplanar satellites \\citep{2006-ApJ-639-Brown, 2009-AJ-137-Ragozzine, 2011-AA-528-Dumas}. \\item It is the largest member of a dynamical family \\citep{2007-Nature-446-Brown, 2007-AJ-134-Ragozzine}, whose velocity dispersion is surprisingly small \\citep{2009-ApJ-700-Schlichting, 2010-ApJ-714-Leinhardt}. \\item Its surface composition is dominated by water ice \\citep{2007-AJ-133-Tegler, 2007-ApJ-655-Trujillo, 2007-AA-466-Merlin, 2009-AA-496-Pinilla-Alonso, 2011-AA-528-Dumas}, yet it has a high density of 2.5-3.3\\,g\\,cm$^{-3}$ \\citep{2006-ApJ-639-Rabinowitz}. \\item It surface has a hemispherical colour heterogeneity, with a dark red ``spot'' on one side \\citep{2008-AJ-135-Lacerda, 2009-AJ-137-Lacerda}. \\end{enumerate} \\indent \\citet{2007-Nature-446-Brown} proposed that Haumea suffered a giant collision that ejected a large fraction of its ice mantle, which formed both the two satellites and the dynamical family and left Haumea with rapid rotation. A number of theoretical studies have since looked at the family formation in more detail (see Sect.~\\ref{sec: discussion}). \\\\ \\indent A characterisation of the candidate members \\citep[35 bodies listed by][including Haumea itself]{2007-AJ-134-Ragozzine} however showed that only 10 bodies out of 24 studied share their surface properties with Haumea \\citep{2010-AA-511-Snodgrass}, and can thus be considered genuine family members. Moreover, these confirmed family members cluster in the orbital elements space \\citep[see Fig.~4 in][]{2010-AA-511-Snodgrass}, and the highest velocity found was $\\sim$123\\,m\\,s$^{-1}$ (for 1995 SM$_{55}$).\\\\ \\indent We report on follow-up observations to \\citet{2010-AA-511-Snodgrass} of 8 additional candidate members of Haumea's family. We describe our observations in Sect.~\\ref{sec: obs}, the colour measurements in Sect.~\\ref{sec: colours}, the lightcurve analysis and density estimates in Sect.~\\ref{sec: dens}, and we discuss in Sect.~\\ref{sec: discussion} the family memberships of the candidates and the implication of these for the characteristics of the family. ", "conclusions": "\\indent We have presented optical and near-infrared colours for 8 of the 36 candidate members of Haumea's collisional family \\citep{2009-AJ-137-Ragozzine}, in addition to the 22 objects we already reported \\citep{2010-AA-511-Snodgrass}. We confirmed the presence of water ice on the surface of 2003~UZ$_{117}$, confirming its link with Haumea, and rejected 5 other candidates \\citep[following our prediction that most of the remaining objects would be interlopers,][]{2010-AA-511-Snodgrass}.\\\\ \\indent Of the 36 family member candidates including Haumea, only 11 (30\\%) have been confirmed on the basis of their surface properties, and a total of 17 have been rejected (47\\%). All the confirmed members are tightly clustered in orbital elements, the largest velocity dispersion remaining 123.3\\,m s$^{-1}$ (for 1995~SM$_{55}$). These fragments, together with the two satellites of Haumea, Hi`iaka and Namaka, account for about 1.5\\% of the mass of Haumea. \\\\ \\indent The current observational constraints on the family formation can be summarised as: \\begin{enumerate} \\item A highly clustered group of bodies with unique spectral signatures. \\item An elongated and fast-rotating largest group member. \\item A velocity dispersion and total mass lower than expected for a catastrophic collision with a parent body of Haumea's size, but a size distribution consistent with a collision. \\end{enumerate} Various models have been proposed to match these unusual constraints, although so far none of these match the full set of constraints." }, "1207/1207.4341_arXiv.txt": { "abstract": "The properties of the short, energetic bursts recently observed from the $\\gamma$-ray binary {\\lsi}, are typical of those showed by high magnetic field neutron stars, and thus provide a strong indication in favor of a neutron star being the compact object in the system. Here, we discuss the transitions among the states accessible to a neutron star in a system like {\\lsi}, such as the ejector, propeller and accretor phases, depending on the NS spin period, magnetic field and rate of mass captured. We show how the observed bolometric luminosity ($\\simgt {\\rm few}\\times10^{35}$ erg s$^{-1}$), and its broad-band spectral distribution, indicate that the compact object is most probably close to the transition between working as an ejector all along its orbit, and being powered by the propeller effect when it is close to the orbit periastron, in a so-called {\\it flip-flop} state. By assessing the torques acting onto the compact object in the various states, we follow the spin evolution of the system, evaluating the time spent by the system in each of them. Even taking into account the constraint set by the observed $\\gamma$-ray luminosity, we found that the total age of the system is compatible with being $\\approx$ 5--10 kyr, comparable to the typical spin-down ages of high-field neutron stars. The results obtained are discussed in the context of the various evolutionary stages expected for a neutron star with a high mass companion. ", "introduction": "{\\lsi} is one the few high-mass X-ray binaries (HMXB) discovered so far to emit the largest part of their luminosity at high energies \\citep{hermsen1977,gregory1978,albert2006}, being therefore a member of the class of $\\gamma$-ray binaries. Variability of its emission, at the timescale set by the $\\sim$26.5 d orbital period, has been found at almost all wavelengths, e.g., \\citet{albert2008,abdo2009,torres2010,zhang2010}. The companion star is a massive B0Ve star, with a mass between 10 and 15 $M_{\\odot}$, in an eccentric 26.5 d orbit \\citep{casares2005}. For the nature of the compact object in $\\gamma$-ray binaries, models involving an accreting black hole launching a relativistic jet \\citep[the microquasar scenario; see, e.g.][and references therein]{boschramon2009} and a rotation-powered neutron star (NS in the following) emitting a relativistic wind of particles \\citep[see, e.g.][]{maraschi1981,dubus2006,sierpowska2008}, have been proposed. The presence of a NS in {\\lsi} would be definitely proven by the observation of pulsations, but deep searches in the radio \\citep{mcswain2011} and X-ray band \\citep{rea2010} were not successful, so far. This is not surprising, since free-free absorption easily washes out the pulses in the radio band, while the upper limit of $\\approx10\\%$ (3 $\\sigma$ confidence level) on the pulsed fraction in X-rays could well be larger than the actual pulsed fraction of the source. However, in the past few years, a couple of energetic ($\\approx 10^{37}$ erg s$^{-1}$), short ($\\simlt 0.3$ s) bursts were detected by the {\\it Swift}-Burst Alert Telescope from a region of a few arc minutes of radius, compatible with the position of {\\lsi} \\citep[][see Table~\\ref{tab:bursts} for their observed properties]{depasquale2008,barthelmy2008,burrows2012}. The properties of the two bursts are typical of those observed in magnetars, namely NSs for which emission is believed to be powered by their strong magnetic energy. It is probable that the bursts were emitted by \\lsi\\ itself. Otherwise, we would be witnessing the unlikely alignment, within a couple or arcmin, of a gamma-ray binary (a population of objects for which a handful members are known) with a magnetar-like burst-emitting object (for which we know 20 sources). If the \\lsi\\ origin is accepted, any model of its multi-wavelength emission should thus provide an explanation of such bursts. Under the common assumption of pulsars emitting their rotational energy via magnetic dipolar losses, the NS surface dipolar magnetic field can be estimated from the observed period and period derivative \\citep{pacini1967,gold1969}. For the known magnetars, this usually ranges from $\\sim5\\times10^{13}$ to $2\\times10^{15}$\\,G; recently, however, two sources with a lower field, $\\simgt 7\\times10^{12}$ G, were discovered, the emission of which is still believed to be powered by non-dipolar components of the magnetic field \\citep{rea10,turolla2011,rea2012}. About 20 magnetars are known to date, all being isolated pulsars with periods ranging from 0.3--12\\,s, usually young spin-down ages ranging between 0.7--230 kyr (again with the two exceptions reported above which are also much older systems), and X-ray luminosities of the order of $10^{33-35}~$erg~s$^{-1}$ (see \\citealt{mereghetti2008,rea11} for recent reviews). Magnetars, historically divided into the two subclasses of Anomalous X-ray Pulsars (AXPs) and the Soft Gamma Repeaters (SGRs), display a large variety of bursts and flares, with properties at variance with those observed from other compact objects such as accreting NSs or BHs. Magnetars bursts can be empirically divided in three main classes (although there is probably a continuum among them): the short bursts ($\\sim0.01-1$\\,s; $10^{37-40}~$erg~s$^{-1}$), the intermediate bursts ($\\sim5-50$\\,s; $10^{40-42}~$erg~s$^{-1}$) and the giant flares ($\\sim100-500$\\,s; $10^{43-47}~$erg~s$^{-1}$). \\citet{torres2012} have started to study how a high-field NS could cope with the multi-wavelength phenomenology of {\\lsi}. In their scenario the NS would behave as a usual rotation-powered pulsar only when far from the companion star, whereas close to periastron, the increased pressure exerted by the matter of the Be equatorial disk would rather overcome the pulsar pressure, quenching the rotation-powered pulsar behavior. Though, accretion of the matter captured would be inhibited by the quick rotation the NS, which would then act as a propeller \\citep{illarionov1975}. Such alternation between ejector and propeller states along the orbit, which we refer to as a {\\it flip-flop} state, was originally proposed by \\citet{gnusareva1985} for NS in close-binary orbits of high eccentricity, and already applied to the case of {\\lsi} by \\citet{campana1995} and \\citet{zamanov1995}. In this paper we delve further into this scenario, estimating the interval of spin periods at which a NS in {\\lsi} is expected to behave either as an ejector or a propeller, and the duration of each of the states experienced by a NS in an eccentric binary system, as it spins down during its initial evolution. By taking into account the constraints set on the parameters of the system by the observed $\\gamma$-ray luminosity, we also estimate the relative likelihood of observing the assumed NS in one of the different states. Results are discussed, comparing the case of an assumed high-field NS in {\\lsi} to possibly related systems, such as rotation powered sources in eccentric binary systems, as well as very long period HMXBs, thought to have host a magnetar in their early evolutionary stages. \\begin{deluxetable}{lrr} \\tablecaption{Bursts observed by {\\it Swift}-BAT from {\\lsi}\\label{tab:bursts}} \\tablehead{ & \\colhead{Burst No.\\Rmnum{1}\\tablenotemark{(a)}} & \\colhead{Burst No.\\Rmnum{2}\\tablenotemark{(b)}} } \\startdata Date & 2008 Sep 10 & 2012 Feb 5 \\\\ Position uncertainty & 2.1' & 3' \\\\ Ang. sep. & 0.60' &1.07' \\\\ T$_{100}$ (s) & 0.31 & 0.044 \\\\ Fluence ($10^{-8}$ erg cm$^{-2}$) & $1.4\\pm0.6$ & $0.58\\pm0.14$ \\\\ $\\Gamma$ & $2.0\\pm0.3$ & $3.9\\pm0.4$ \\\\ Luminosity ($10^{37}$ erg s$^{-1}$) & 2.1 & 6.3 \\enddata \\tablenotetext{(a)}{\\citet{torres2012}} \\tablenotetext{(b)}{From \\citet{burrows2012}, see also {http://gcn.gsfc.nasa.gov/notices\\_s/513505/BA/}} \\tablecomments{The positional uncertainty is given at a 90$\\%$ confidence level, including also systematic uncertainties. The angular separation is calculated with respect to the position of the optical counterpart. The $T_{100}$ duration and the fluences are estimated in the 15--50 keV band. Burst spectra were fitted by a power law with index $\\Gamma$. The average luminosity is estimated by assuming a distance of 2 kpc \\citep{frail1991}. } \\end{deluxetable} ", "conclusions": "\\label{sec:discussion} According to the conventional picture \\citep[see, e.g.][]{lipunov1992}, a NS in a binary system evolves through ejector, propeller and accretor states, as it spins down. While on long timescales the transition between these states is mainly set by the evolution of the spin period of the NS, on shorter timescales a key role is played by the rate at which the compact object captures the mass lost by the companion star. In particular, the large values of the ratio between the maximum and minimum rate of mass captured by the NS achieved if the orbit is eccentric, may induce state transitions along an orbital cycle, a {\\it flip-flop} state. The range of mass capture rates spanned through an orbital cycle is even larger if the non degenerate star is of Be class, losing mass also through a dense equatorial disc which is transversed by the NS when it is close to periastron. The existence of systems alternating states on the timescale set by the orbital period is therefore a natural consequence of the way a NS evolves. While the observation of two magnetar-like bursts from a few-arcmin region compatible with the position of {\\lsi} provided a strong indication in favor of a NS nature of the compact object in the system, it is unlikely that an accreting NS is hosted by {\\lsi}. Accreting NS in Be-HMXB show in fact X-ray pulsations at a period larger than few seconds. Moreover, their X-ray energy continuum spectrum is described by a power law with an exponential cutoff at 10--30 keV, on which cyclotron and/or iron emission features are generally superimposed \\citep[see][and references therein]{reig2011}. No X-ray pulsation has been detected so far from {\\lsi} \\citep{rea2010}, while its X-ray spectrum is featureless and does not show a cut-off in the X-ray energy band \\citep{chernyakova2006,zhang2010}. On the other hand, the luminosity observed from the system is of the same order of the spin-down power liberated by the NS when its spin period is close to the critical value marking the transition from the ejector to the {\\it flip-flop} phase. An important clue to estimate the likelihood of finding a NS in {\\lsi} in such a phase follows from the assessment of the time spent by the system in such a state. We showed how such a timescale crucially depends on the details of the assumed propeller torque. By assuming the rotating NS to release to the incoming matter the energy needed to unbind it (see Eq.~\\ref{is}), or the energy effectively stored in the NS rotation (Eq.~\\ref{ghosh}), largely different estimates of the evolutionary timescales are obtained. This is essentially because the rotational energy of the NS when spinning close to the ejector/propeller transition largely exceeds the gravitational energy possessed by the in-falling matter, evaluated at the large magnetospheric radius implied by a strong magnetic field. Among the propeller torques considered here, the weaker yields a total duration of the {\\it flip-flop} phase which is much larger ($\\approx 280$ kyr) than the one implied by the stronger torque ($\\approx 25$ kyr). Even considering the stronger propeller torque, which is favored if the system is powered by the propeller effect when the NS is close to periastron, the expected total duration of the {\\it flip-flop} phase [see Eq.~(\\ref{eq:flipfloptime})] is larger by a factor of $\\approx 4$ with respect to the timescale spent by the object in the pure ejector state: \\begin{equation} \\frac{t_{ff}}{t_{ej}}\\approx 4 \\:b_1^{-0.24}\\: (\\dot{m}_1^{max})^{0.27} \\label{eq:ratio} \\end{equation} where we have made explicit only the dependence of Eq.~(\\ref{eq:emtime}) and (\\ref{eq:flipfloptime}) on the NS magnetic field, and on the maximum mass accretion rate. It is then reasonable to find the system in the {\\it flip-flop} state. Moreover, even if it is taken into account that the system must be relatively young to emit a spin-down luminosity $\\simgt{\\rm few}\\times10^{35}$ erg s$^{-1}$, the time spent by the system in the {\\it flip-flop} phase is of few kyr, comparable to the total duration of the previous ejector phase, and yielding an age of the system $\\approx 5$--$10$ kyr, of the order of typical spin-down ages of magnetars. On the other hand, a spin-down luminosity of the order of $\\approx10^{37}$ erg s$^{-1}$, would imply a smaller age for the system, with the NS emitting as an ejector all along the orbit. \\begin{deluxetable}{lccccc} \\tablecaption{Ante-diluvian systems and {\\lsi}\\label{tab:antediluvian}} \\tablehead{ \\colhead{Name} & \\colhead{$P_S$(s)} & \\colhead{$P_{orb}$(d)} & \\colhead{$e$} & \\colhead{$B_1$(G)}\\tablenotemark{(a)} & $M_2$ ($M_{\\odot}$) } \\startdata J1740-3052 & 0.57 & 231.0 & 0.57 & 3.9$\\times10^{12}$ & 11.0--15.8 \\\\ J1638-4725 & 0.76 & 1941 & 0.95 & 1.9$\\times10^{12}$ & 5.8--8.1 \\\\ J0045-7319& 0.93 & 51.1 & 0.81 & 2.1$\\times10^{12}$ & 3.9--5.3 \\\\ B1259-63 & 0.048 & 1237 & 0.87 & 3.3$\\times10^{11}$ & 3.1--4.1 \\\\ & & & & & \\\\ {\\lsi}& \\nodata & 26.5 & 0.63 & \\nodata & 10--15 \\enddata \\tablenotetext{(a)}{The magnetic field of pulsars is determined as $3.2\\times10^{19}\\:(P\\dot{P})^{1/2}$ G.} \\tablerefs{ \\citet{mcconnell1991, johnston1992,kaspi1996a,stairs2001, wang2004,casares2005,lorimer2006}; we acknowledge the use of the ATNF Pulsar Catalogue, http://www.atnf.csiro.au/people/pulsar/psrcat/, \\citet{manchester2005}} \\end{deluxetable} Either it lying in the ejector or in the {\\it flip-flop} state, the presence of a young NS in {\\lsi} would make the system closely related to so-called {\\it ante-diluvian systems} \\citep{vandenheuvel2001}, rotation powered pulsars orbiting a high mass companion in a eccentric orbit, considered as the progenitors of HMXB. The properties of the four sources of this class discovered so far are listed in Table~\\ref{tab:antediluvian}, together with those of {\\lsi}. We note that B1259-63 is also one of the brightest $\\gamma$-ray binaries known. All these NS have a superficial magnetic field in the range $3\\times10^{11}$--$4\\times10^{12}$ G, as derived from the observed spin down rate. Indeed, it was early proposed that some of these systems could be found in a propeller state when the NS was close to periastron. However, the rate at which the companion star would have to lose mass in order to overcome the pulsar pressure should be much larger than expected \\citep[][and references therein]{campana1995,ghosh1995a,tavani1997} or observed \\citep{kaspi1996b}. Despite a proper evaluation of the ratio defined by Eq.~(\\ref{eq:ratio}) for the known {\\it ante-diluvian} systems is far from the scope of this paper, it can be noted how the likelihood of observing a system in the pure ejector state increases when the maximum mass capture rate decreases, as it is the case of at least three out of the four known {\\it ante-diluvian} systems, lying in a much larger orbit with respect to that of {\\lsi}. If confirmed by the discovery of pulsations, {\\lsi} would be the first magnetar to be discovered in a binary system. The large luminosity variation shown by super fast X-ray transients on short timescales led \\citet{bozzo2008} to argue how a magnetar-like magnetic field could represent an efficient gating mechanism to make these systems to rapidly switch between propeller and accreting states. That a number of accreting HMXB should have host in the past a NS with a field in the magnetar-range has also been claimed on the basis of their very large spin periods (e.g. the cases of 2S 0114+650 with a period of 2.7 hr, \\citealt{li1999}, and 4U 2206+54 with a period of 5560 s, \\citealt{finger2010,ikhsanov2010,reig2012}). The spin period at which a system eventually starts accreting mass increases in fact with the strength of the magnetic field (see Eq.~[\\ref{eq:pmax}] and the similar expression derived by \\citealt{shakura2012} for the equilibrium period of NS in the settling regime), essentially because the value of the magnetic field sets the strength of the propeller torques that are responsible for the NS spin down. In this context the case of the Be/X-ray binary in the Small Magellanic Cloud, SXP 1062, with a period of 1062 s, and with an estimated age of 16 kyr \\citep{haberl2012} is noteworthy. \\citet{popov2012} showed how such a short age strongly points to the presence of a NS with a large initial magnetic field in the system, $\\sim 10^{14}$ G, which could have spun the NS down very efficiently before the start of the accretion phase. A comparison of the age proposed for SXP 1062 with the timescales of the different propeller mechanism listed in Table~\\ref{tab}, indicates how the faster and stronger expression for the propeller torque is probably closer to the torque effectively experienced by a NS in that system during the propeller stage. How a {\\it strong} propeller torque, of the order of that given by Eq.~(\\ref{ghosh}), should be possibly preferred in describing the evolution of systems like {\\lsi}, is also indicated by the ratio between the ejector and {\\it flip-flop} timescales defined by Eq.~(\\ref{eq:ratio}). If the weaker propeller torques were in place, a NS would spend a much larger time in the {\\it flip-flop} state than in a pure ejector state, and this is not indicated by the number of systems observed in the latter state (4) with respect to the single possible case of {\\lsi}. We finally note that, contrary to the case of SXP 1062, no association with a supernova remnant could be been made for {\\lsi} \\citep{frail1987}. This is not entirely surprising for a source with an estimated age between 10 and 20 kyr, since such an association can be found only for a fraction 0.64 and 0.55 of the radio pulsars with such ages estimated from their electromagnetic spin down, respectively (ATNF pulsar database, \\citealt{manchester2005}). Furthermore, no other Be X-ray binary has been observed embedded in a SNR, despite the relatively young age and large number of observed systems. This is most probably due to the large wind of the two progenitor massive stars, which has swept away most of the material around the binary, resulting in an under-luminous (hence undetectable) SNR after the explosion." }, "1207/1207.1108_arXiv.txt": { "abstract": "It is strongly believed that Andromeda's double nucleus signals a disk of stars revolving around its central super-massive black hole on eccentric Keplerian orbits with nearly aligned apsides. A self-consistent stellar dynamical origin for such apparently long-lived alignment has so far been lacking, with indications that cluster self-gravity is capable of sustaining such lopsided configurations if and when stimulated by external perturbations. Here, we present results of N-body simulations which show unstable counter-rotating stellar clusters around super-massive black holes saturating into uniformly precessing lopsided nuclei. The double nucleus in our featured experiment decomposes naturally into a thick eccentric disk of apo-apse aligned stars which is embedded in a lighter triaxial cluster. The eccentric disk reproduces key features of Keplerian disk models of Andromeda's double nucleus; the triaxial cluster has a distinctive kinematic signature which is evident in HST observations of Andromeda's double nucleus, and has been difficult to reproduce with Keplerian disks alone. Our simulations demonstrate how the combination of eccentric disk and triaxial cluster arises naturally when a star cluster accreted over a pre-existing and counter-rotating disk of stars, drives disk and cluster into a mutually destabilizing dance. Such accretion events are inherent to standard galaxy formation scenarios. They are here shown to double stellar black hole nuclei as they feed them. ", "introduction": "Over the years, the nucleus of the Andromeda galaxy (M31) went from being asymmetric, to doubling, then tripling. The asymmetry was first noted in the balloon-born Stratoscope observatory \\citep{Light1974}. It was photometrically resolved into a double nucleus by Hubble Space Telescope (HST) observations \\citep{Lauer1993}; asymmetry hence turned into lopsidedness, with a luminous feature (referred to as P1) shining a few parsecs away from a dimmer one (referred to as P2), which is closer to the center of the host bulge. The resolved double nucleus had already been suspected to host a Super-Massive Black Hole (SMBH), located somewhere close to P2, with a neighborhood that was known to shine brighter than the rest of the nucleus in the ultraviolet \\citep{Dressler1988, Kormendy1988}. Detailed HST spectroscopy added a third component (known as P3), a disk of young massive stars, whose size and rotation speed rules out viable alternatives to a central SMBH of $\\sim 10^8$ solar masses {\\it (\\SM) } \\citep{Bender2005}. In the currently favored model of M31's lopsided double nucleus \\citep{Tremaine1995, Peiris2003}, the brighter peak P1 \\citep{Lauer1993} is thought to coincide with the common apo-centric region of an eccentric disc of stars revolving on nearly apse-aligned Keplerian ellipses, around M31's central super-massive black hole (hereafter SMBH). Similar eccentric discs are thought to underlie lopsided nuclei detected around SMBHs in the centers of nearby galaxies \\citep{Lauer1996, Lauer2005, Gultekin2011}. That such kinematic configurations can be sustained in self-gravitating disks around SMBHs was clarified in a series of dynamical studies. Indeed, investigation of modes of hot nearly-Keplerian disks indicated that slow (m=1, lopsided) modes can be stably excited \\citep{Tremaine2001}. This conclusion was corroborated by results of planar N-body simulations of disks around SMBHs which showed that, when started in asymmetric conditions, disks can relax into long-lived uniformly precessing lopsided configurations \\citep{Bacon2001, Jacobs2001}. Subsequently, razor-thin, self-gravitating, lopsided equilibria were constructed, yielding encouraging agreement with photometric and kinematic observations of M31's nucleus \\citep{Salow2001, Sambhus2002}. These and related studies left open questions about how such modes are excited, how they ultimately saturate into global lopsided configurations, and whether constructed equilibria were stable or not. The early suggestion \\citep{Tremaine1995} that an initially circular disk could become eccentric under the influence of dynamical friction with the host bulge has not been thoroughly explored. A recently proposed scenario \\citep{Hopkins2010a, Hopkins2010b} has eccentric disks forming in the notoriously complex -and poorly understood \\citep{Silk2011}- environment which couples star formation and its feedback to gas accretion in the early stages of SMBH growth.\\\\ \\indent In this work, we opt for a minimalist and dynamically self-consistent route to three-dimensional, lopsided nuclei, out of counter-rotating (CR hereafter) collisionless stellar distributions dominated by SMBHs. Such CR distributions are prone to violent m=1 instabilities in the presence of moderate counter-rotation \\citep{Jog2009, Touma2002}, hence the suggestion that the lopsided structure in Andromeda's nucleus may have been triggered by the accretion of a retrograde globular cluster in a pre-existing disk of stars \\citep{Sambhus2002}. With the intention of exploring the outcomes of such accretion events, we performed a series of N-body simulations of CR stellar distributions evolving in the sphere of influence of an SMBH. Below, we report on results which show unstable CR distributions evolving into stable, uniformly precessing lopsided nuclei. We follow the instability in our featured simulation from multiple complementary perspectives. We then probe the ensuing lopsided nucleus, and show that in addition to displaying all the key observational signatures which Keplerian disks are meant to model, it recovers asymmetries in the tail of Adromeda's line of sight velocity distribution \\citep{Bender2005} which a thick lopsided disk alone is unable to reproduce\\citep{Peiris2003}. \\begin{figure} % \\includegraphics[scale=0.65]{panel1.eps} \\caption{The eccentricity-inclination dynamics of the particles up to 1 Myr is shown. In row {\\bf a}, we follow a top view of the prograde particles, from initial axisymmetric annular configuration, through a phase of growth of lopsidedness, with growing mean eccentricity (0.41 around 1 Myrs), and alignment of apsides. Row {\\bf b} shows a top view of retrograde particles over the same period; they experience greater increase in mean eccentricity (around 0.73 by 0.72 Myrs), and show signs of apse alignment by 0.5 Myrs, before dispersing (by 1 Myrs) into a disky blob. A side view of retrograde particles in row {\\bf c} shows how the in-plane dispersal seen in row {\\bf b} reflects in projection a burst of out-plane dynamics which puffs the initially retrograde disk into a triaxial cluster of stars. The eccentricity-inclination dynamics of the full cluster is captured in row {\\bf d}, with prograde in black, and retrograde in blue; the mean eccentricity of both populations grows as the instability develops; by 1 Myrs, the mean inclination of the prograde particles has hardly changed, while the highly eccentric retrograde population, is widely dispersed in inclination.} \\label{fig:part} \\end{figure} \\section[]{Counter-Rotating Instability: From Linear Growth to Saturation} Our N-body simulations were performed with the parallel version of \\gadget~\\citep{Springel2005}, an oct-tree code, which is popular with the cosmological structure formation community. We pushed this tool to the limit of extreme mass ratios, one for which it was not specifically designed, one in which it did the desired job, albeit with stringent force accuracy, and time stepping criteria. With M31 in mind, all experiments have an SMBH with a mass of \\msmbh=10$^{8}~$\\SM and a main disk component (prograde in our convention) with a tenth of SMBH mass \\citep{Bender2005, Peiris2003}. An exhaustive exploration, with sufficient realism, of a scenario in which a CR cluster ($10^5-10^6$ \\SM) decays under dynamical friction to nuclear regions \\citep{tremaine1975}, then gets disrupted \\citep{Kim2003, Fujii2010} as it couples to the massive nuclear stellar disk, would have been forbidding to pursue in detail with our current computational resources. Instead we opted for an extensive exploration of a realistic configuration which is known to be unstable \\citep{Touma2002, Sridhar2010}, and in which the disrupted CR cluster is modeled by overlaying the prograde disk with a retrograde disk of 1/10 its mass. Details on initial conditions for this experiment, along with \\gadget~simulation parameters, are described in Appendix \\ref{sec:supp-methods}; variations on the fiducial experiment are discussed in Appendix \\ref{sec:crExplored}. \\subsection{Dissecting the Instability} CR configurations are known to be (linearly) unstable \\citep{Touma2002, Sridhar2010}. We are here concerned with the (non-linear) saturation of this instability and shall first map its unfolding with complementary views of particle dynamics (Fig. \\ref{fig:part}). A top view of stars, in both prograde and retrograde populations (hereafter PP and RP respectively), shows them developing lopsidedness by coupling growing eccentricity to apse-alignment. The more massive PP maintains its coherent, apse-aligned precession (Fig. \\ref{fig:part},{\\bf (a)}) whereas the lighter RP dissolves gradually into a disky blob after 0.5 Myrs (Fig. \\ref{fig:part},{\\bf (b)}). Viewed from the side, the gradual in-plane dispersal of the RP occurs along with a dramatic excitation of out-of-plane motion; by 1 Myrs the RP unfolds into a triaxial structure (Fig. \\ref{fig:part},{\\bf (c)}). Scatter plots of particle eccentricities and inclinations (Fig. \\ref{fig:part},{\\bf (d)}) reveal how, by 0.72 Myrs, the now highly eccentric RP is well on its way to complete triaxial disruption. Tucked in this emerging cluster, the PP, which has absorbed much of the RP's (negative) angular momentum, heats up slightly in inclination as it consolidates its eccentric lump, then adjusts slowly to whatever little angular momentum is left in the RP. A look at the mean eccentricity and inclination of both populations (Fig. \\ref{fig:statvstime} {\\bf a, b}) confirms these observations, as it reveals three distinct phases: {\\bf Phase-I} of near coplanar growth of the mean eccentricity of both populations which lasts till about 0.5 Myrs; {\\bf Phase-II} of continued mean eccentricity growth, but now coupled to growth in the mean inclination of the RP, both of which reach a maximum around 0.72 Myr; {\\bf Phase-III}, with cycles of increasingly small amplitude leading to a near-steady regime. As emphasized in the caption of Fig. \\ref{fig:statvstime}, transitions in mean eccentricity and inclination are neatly imprinted on the modal structure of the projected density and the pattern speed of the dominant m=1 mode, as the full cluster transitions from a thin axisymmetric initial state, to a saturated, lopsided, and slowly precessing mode (Fig. \\ref{fig:statvstime}, {\\bf c, d}, and Fig. \\ref{fig:dens}). \\begin{figure} \\centering \\includegraphics[scale=1.0]{panel2.eps} \\caption{The linearly unstable, counter-rotating configuration saturates into a lopsided uniformly precessing disk. Panels {\\bf a} and {\\bf b} show the mean eccentricity $<$e$>$ and inclination $<$I$>$ of PP (black) and RP (blue). $<$e$>$ grows, at near constant $<$I$>$, till about 0.5 Myrs, at which point $<$I$>$ takes off, with both $<$e$>$ and $<$I$>$ peaking around 0.72 Myrs. Cycles of increasingly smaller amplitude follow this initial growth phase, leading to a saturated eccentric prograde disk which is shrouded by a triaxial halo. In panel {\\bf c}, the power of the four dominant modes in projected density is displayed; an initially axisymmetric configuration (m=0) lends way to a dominant m=1 mode, along with non-negligible m=2 and m=3 contributions. Triaxial dispersal of the RP is responsible for the slight growth of the m=0 amplitude, and the related slight decrease in the m=1 mode amplitude after 0.72 Myrs. In panel {\\bf d}, we follow the precession rate of the m=1 mode in time: initial large amplitude oscillations reflect back and forth libration of prograde and retrograde m=1 excitations about the uniform precession state; these oscillations die out with the dispersal of the retrograde bunch, leaving an m=1 mode which precesses at a near constant rate of 19 km s$^{-1}$ pc$^{-1}$.} \\label{fig:statvstime} \\end{figure} Saturation of the m=1 mode is evidently correlated with the dispersal of the initially retrograde population into a triaxial cluster of high eccentricity orbits. The dispersal appears to be associated with collective dynamics along eccentricity-inclination cycles. These cycles are excited when a population of stars (the retrograde population) develops sufficient eccentricity (through a CR instability) in the presence of a dominant eccentric perturber (the more massive eccentric and precessing prograde disk) \\cite{Touma2009}. One can presumably describe the inception of these cycles in a generalized Kozai-Lidov framework \\citep{naoz2011}, but then one is left with the harder task of accounting for the dispersal of the population on these cycles. We gained valuable insight into dispersal then saturation by modeling a closely related process in a 2D analog of the 3D cluster (Kazandjian \\& Touma, in preparation). The model in question permits a blow by blow account of the mutual sculpting of planar prograde and retrograde populations in terms of capture in, then escape from a drifting, trapping region in phase-space \\citep{sridhar1997}. It can naturally explain the near-coplanar growth and apse alignment in phase {\\bf I} of 3D dynamics. Furthermore, it clarifies the collisionless damping process at work in the decaying cycles of phase {\\bf III}. How the eccentricity growth of phase {\\bf I} paves the way for inclination growth in phase {\\bf II} cannot be captured in a planar analogue and shall await a more sophisticated treatment of the 3D deployment of unstable CR clusters. ", "conclusions": "Andromeda's kinematics shows a marked asymmetry between the tails of its LOSVD which is clearly unaccounted for in Keplerian disk models, and appears to be cleanly resolved by combining an eccentric disk with a triaxial cluster of highly eccentric orbits. To be sure, thorough modeling is called for before we can rigorously quantify the improvements of a disk-cluster combination over the thick eccentric disk model of \\cite{Peiris2003}. Still, we find it remarkable that this combination achieves, with minimal probing, the level of qualitative and quantitative agreement described above, and to see it emerge as a natural end product of our process can only strengthen the case for CR stimuli of double nuclei \\citep{Lauer2005}. That case is made stronger when one learns that a rich variety of CR configurations are equally prone to developing lopsidedness (details in Appendix \\ref{sec:crExplored}.). Experiments with CR point masses suggest that a cluster, initially on a circular trajectory, can drive a coplanar disk unstable provided its orbital radius is less than a critical radius (which is roughly equal to twice the mean radius of the disk); clusters on highly inclined circular trajectories can still drive a CR disk into a lopsided state; on the other hand, clusters with too large an initial eccentricity will end up with little negative angular momentum to drive a coplanar disk lopsided. The CR perturbation envisaged in this work is delivered by a CR cluster which spirals deep into the nucleus by dynamical friction with bulge stars \\citep{tremaine1975} before it disrupts in the combined tidal field of bulge and SMBH \\citep{Quillen2003}. The orbital radius at which the cluster disrupts depends critically on the cluster's core density. For a cluster migrating on a near-circular trajectory into M31's nucleus, we estimate \\citep{Quillen2003} that a core density $\\geq$ 5 $\\times$ 10$^5$ \\SM pc$^{-3}$ is sufficient for the cluster to cross past the critical radius for instability. Mass segregation and/or intermediate-mass black hole formation \\citep{Kim2003, Fujii2010} can amplify core density significantly, thus further delaying disruption till the cluster migrates into the overlapping configurations explored here. Such configurations may also arise when a cluster penetrates the nucleus on an eccentric orbit, and disrupts upon close encounter with the SMBH. CR clusters will likely approach a pre-existing nuclear disk on eccentric, and inclined trajectories; they may thus find themselves on the eccentricity-inclination cycles observed above, then disperse as they drive the disk in a saturated lopsided configuration. Given on one hand the strong likelihood of counter-rotating excitations in galactic centers \\citep{Jog2009}, and on the other the robustness and efficiency of the proposed mechanism, double nuclei are likely to prove ubiquitous in stellar clusters dominated by super-massive black holes \\citep{Lauer2005}. More generally (and irrespective of whether or not a given lopsided nucleus results from a CR instability), the proposed mechanism can be deployed to customize triaxial equilibrium configurations with which to model observed kinematics of stellar black hole nuclei \\citep{alexander2005PhR} (the Milky Way's included) and improve estimates of the mass and feeding rate \\footnote{We note in passing that, as it drives a substantial fraction of stars to near radial orbits, the CR instability populates the loss cone of the SMBH \\citep{Magorrian99}. Repeated CR accretion events will thus enhance the feeding rate of the central SMBH as they sculpt stellar distributions in its sphere of influence.} of the black hole within \\citep{Magorrian98, kormendy2004}. Whether exploring realistic CR scenarios or tailoring triaxial equilibrium models, extensive numerical simulations of CR stellar systems with state of the art solvers \\citep{Fujii2010, Touma2009} will surely contribute invaluable insights into the dynamics and structure of stellar black hole nuclei." }, "1207/1207.0424_arXiv.txt": { "abstract": "{We present optical and NIR spectroscopic observations of U\\,Sco 2010 outburst. From the analysis of lines profiles we identify a broad and a narrow component and show that the latter originates from the reforming accretion disk. We show that the accretion resumes shortly after the outburst, on day +8, roughly when the super-soft (SSS) X-ray phase starts. Consequently U\\,Sco SSS phase is fueled (in part or fully) by accretion and should not be used to estimate $m_{\\mathrm{rem}}$, the mass of accreted material which has not been ejected during the outburst. In addition, most of the He emission lines, and the He\\,{\\sc ii} lies in particular, form in the accretion flow/disk within the binary and are optically thick, thus preventing an accurate abundance determination. A late spectrum taken in quiescence and during eclipse shows Ca\\,{\\sc ii}\\,H\\&K, the G-band and Mg\\,{\\sc i}b absorption from the secondary star. However, no other significant secondary star features have been observed at longer wavelengths and in the NIR band. } {} {} {} {} ", "introduction": "U\\,Sco is a well know recurrent nova (RN - see Bode and Evans 2008, and/or Warner 1995 for review) which has been observed in outburst in 1863, 1906, 1917, 1936, 1945, 1969, 1979, 1987, 1999 and 2010 (Schaefer et al. 2010). It is also the prototype of the RNe subclass, which is characterized by 1) very fast outburst decline and evolution, 2) orbital periods of the order of one day, which point to an evolved secondary star, and 3) no development of forbidden emission lines (i.e. a nebular spectrum) from the ejecta. The interest in RNe is linked to the fact that they are possible candidates for supernova (SN) type Ia progenitors, due to their massive white dwarf. In addition U\\,Sco has been a highly debated object for the nature and composition of its donor star, which are still quite uncertain: past works have often reported anomalously high He abundances and concluded that the donor is possibly a He rich and evolved star (Barlow et al. 1981, Williams et al. 1981, Evans et al. 2001, but see also Iijima 2002, Anupama and Dewangan 2000, and Maxwell et al. 2012); while multicolor photometry and spectral analysis have identified it with a G0 (Hanes 1985) or K2 type star (Anupama and Dewangan 2000). The U\\,Sco 2010 outburst had an extensive follow-up both from space and ground, and from X-ray to NIR wavelengths (Schaefer et al. 2010, Munari et al. 2010, Schaefer et al. 2011, Banerjee et al. 2010, Yamanaka et al. 2010, Diaz et al. 2010, Kafka and Williams 2010, Manousakis et al. 2010, Ness et al. 2012, Orio et al. 2010, Osborne et al. 2010). The observations of the 2010 outburst have already provided two important results, revising our current understanding of this RN: i) U\\,Sco develops a nebular spectrum (provided it is followed long enough after the outburst) as shown by Diaz et al. (2010) and Mason (2011); and ii) U\\,Sco hosts an ONe white dwarf (Mason 2011), whose ultimate fate, in the case of mass accretion, is to collapse in a neutron star rather than explode as SN-Ia (e.g. Nomoto and Kondo 1991). In this paper we present a series of spectra collected across the U\\,Sco 2010 outburst with FEROS and X-Shooter and reveal new unexpected results about the outburst evolution and the binary system. ", "conclusions": "The analysis of the data presented here shows the following results \\begin{itemize} \\item In U\\,Sco, the mass transfer from the secondary star recovers $\\sim$8-10 days after maximum at every outburst. The narrow emission lines from H, He\\,{\\sc i} and He\\,{\\sc ii} lines arise from the optically thick gas of the reforming accretion disk \\item The SSS phase (plateau phase in the light curve) is at least in part fueled by accretion and not just by residual activity on the surface of the white dwarf. \\item The ejecta do not produce He\\,{\\sc ii} emission lines. Only weak He\\,{\\sc i} emission lines form in the ejecta and they are primarily blended with transitions from low ionization elements, thus preventing a reliable abundance calculation. The narrow He\\,{\\sc i} and He\\,{\\sc ii} emission lines from the disk, being optically thick, do not allow abundance calculation either. \\item The secondary star spectral type remains ambiguous though constrained between a F3 and a G sub-giant. Higher signal to noise data taken during the eclipses are needed in order to better identify the donor spectral type and ascertain which absorption features show emission components due to a Wilson-Bappu effect. \\item U\\,Sco spectral characteristics during the post outburst phases are remarkably similar to those of SSS object and LMXBs and might deserve further attention. \\end{itemize}" }, "1207/1207.2421_arXiv.txt": { "abstract": "We present a new model for the formation of stellar halos in dwarf galaxies. We demonstrate that the stars and star clusters that form naturally in the inner regions of dwarfs are expected to migrate from the gas rich, star forming centre to join the stellar spheroid. For dwarf galaxies, this process could be the dominant source of halo stars. The effect is caused by stellar feedback-driven bulk motions of dense gas which, by causing potential fluctuations in the inner regions of the halo, couple to all collisionless components. This effect has been demonstrated to generate cores in otherwise cuspy cold dark matter profiles and is particularly effective in dwarf galaxy haloes. It can build a stellar spheroid with larger ages and lower metallicities at greater radii without requiring an outside-in formation model. Globular cluster-type star clusters can be created in the galactic ISM and then migrate to the spheroid on 100\\thinspace Myr timescales. Once outside the inner regions they are less susceptible to tidal disruption and are thus long lived; clusters on wider orbits may be easily unbound from the dwarf to join the halo of a larger galaxy during a merger. A simulated dwarf galaxy ($\\text{M}_{\\rm vir}\\simeq10^{9}\\text{\\thinspace M}_{\\odot}$ at $z=5$) is used to examine this gravitational coupling to dark matter and stars. ", "introduction": "\\label{sec:intro} \\indent Dwarf galaxies are the predominant star forming objects in the early universe and dwarf spheroidals, in particular, are fossil remnants of this era \\citep{dekel1986}. Normal star formation (post-Population III) occurs in the densest gas accumulating in the centres of galactic potential wells. In this case, we might expect dwarf spheroids to be simple, highly concentrated, star piles. In contrast, the stars in observed dwarfs are diffuse and many lack a conspicuous nucleus. Further, dwarfs at or above the luminosity of Fornax have their own globular cluster systems \\citep{mateo1998}. If these galaxies were the first objects large enough to have a high-pressure ISM in their centres, capable of forming large clusters, we need to explain how such clusters could end up orbiting at substantial radii with a distribution similar to that of the overall light \\citep{miller2009}. Radial age and metallicity gradients are also observed \\citep{mcconnachie2012}, suggesting an outside-in formation scenario reminiscent of the ``monolithic collapse'' model \\citep[hereafter ELS]{eggen1962}.\\\\ \\indent In this paper, we explore how features present in the old stellar populations of dwarf galaxies can occur naturally in contemporary cosmological models through star formation and feedback in these galaxies. Young dwarf galaxies have a high gas content and form stars vigorously. In prior work, \\citet{mashchenko2008} were able to show that stellar feedback in a simulated dwarf galaxy will drive bulk gas motions that couple gravitationally to all matter near the centre of the dwarf. As discussed in \\S\\ref{fb}, this mechanism has been shown to act in actively star forming galaxies at a range of masses and is believed to be generic. The process pumps energy into the orbits of all material passing near the centre, transforming an initial dark matter cusp into a broad core, consistent with observations. Here we study the evolution of the stellar content, which is formed self-consistently in those simulations. Orbit pumping also operates on stars, the other key collisionless component of galaxies, to grow stellar spheroids from the inside out, as well as place massive star clusters on large radial orbits.\\\\ \\indent It is widely understood that the $\\Lambda$CDM cosmology predicts the hierarchical assembly of galaxies: dwarf proto-galaxies interact and merge into larger galaxies, contrary to the ELS model. \\citet{searle1978} and \\cite{zinn1980} refined this model by invoking a late in-fall of old stars that would contribute to both the stellar halo and its globular cluster population. Subsequent work has focused on reconciling this picture for the formation of the Galactic stellar halo with the standard hierarchical framework \\citep[for a recent review see][]{helmi2008}. Chemical enrichment models combined with descriptions of a Milky Way (MW)-type merger history \\citep[e.g.][]{robertson2005,bullock2005,delucia2008} and cosmological simulations of MW-type galaxies \\citep[e.g.][]{zolotov2010} can be made to match the abundance patterns of the stellar halo \\citep[e.g.][]{carollo2007,carollo2010,dejong2010}. A general conclusion is that dwarf progenitors play a major role in building the MW halo, owing to their high rates of star formation at early times and their ability to retain supernova-enriched gas.\\\\ \\indent However, MW-scale simulations poorly resolve dwarf galaxies which thus readily disintegrate and contribute their entire stellar contents to the halo. This conclusion is a direct consequence of low numerical resolution and is at odds with how star formation would be expected to occur in dwarfs. A closer understanding of star formation in dwarf galaxies is needed to establish how those stars are produced and how readily they can contribute to the observed Galactic stellar populations and their radial variations.\\\\ \\indent In \\S\\ref{fb} we discuss the stellar redistribution mechanism and how it operates. In \\S\\ref{stars} we explore the impact this has on the formation of the stellar spheroid in dwarfs including stellar systems such as globular clusters. We also discuss implications for dwarfs contributing their stars to the spheroids of larger galaxies. ", "conclusions": "We have presented a new framework for understanding the formation of the stellar spheroid in dwarf galaxies: All stars form in the nuclear regions and are then redistributed to eventually occupy the entire halo. The redistribution mechanism relies on strong fluctuations in the baryon-dominated central gravitational potential that are associated with stellar feedback as first demonstrated by \\cite{mashchenko2008}. These fluctuations irreversibly affect the orbits and hence distributions of the collisionless components: dark matter, stars and star clusters. The key implications are:\\\\[10pt] \\noindent $\\bullet\\ $ This process directly affects dwarf galaxies. In these galaxies a mild gradient with radius of increasing age and decreasing metallicity would be created as older stars achieve the largest orbits. Orbital redistribution stops when vigorous star formation ceases.\\\\[10pt] \\noindent $\\bullet\\ $ The central density of stars stays fairly constant as new stars form to replace those migrating outwards.\\\\[10pt] \\noindent $\\bullet\\ $ Globular cluster-like star clusters form in the ISM (and thus have no associated dark matter) and migrate outward over several orbital periods.\\\\[10pt] \\noindent $\\bullet\\ $ The star clusters may form multiple generations of stars from enriched gas readily available in the nuclear regions. They will lose access to new gas as their orbits become larger.\\\\[10pt] \\noindent $\\bullet\\ $ Continuous creation and outward migration of stars and globular clusters avoids the formation of a super-nucleus at the centre of most dwarf galaxies.\\\\[10pt] \\noindent $\\bullet\\ $ Larger clusters become protected against tidal destruction as their orbits grow and the dwarf's dark-matter core becomes flattened.\\\\[10pt] \\noindent $\\bullet\\ $ Mergers and tidal stripping will deposit these loosely bound stars and clusters into the halo of later generations of larger galaxies.\\\\[10pt] \\noindent $\\bullet\\ $ Large star clusters formed in dwarf galaxies at high redshift, rather than in dark matter mini-halos, could be the primary source of Globular Clusters in all galaxies." }, "1207/1207.0754_arXiv.txt": { "abstract": "We presented the first particle based, Lagrangian code that can follow the collisional/accretional/dynamical evolution of a large number of km-sized planetesimals through the entire growth process to become planets. We refer to it as the {\\it Lagrangian Integrator for Planetary Accretion and Dynamics} or {\\it LIPAD}. LIPAD is built on top of SyMBA, which is a symplectic $N$-body integrator \\citep{{Duncan:1998aj116:2067}}. In order to handle the very large number of planetesimals required by planet formation simulations, we introduce the concept of a {\\it tracer} particle. Each tracer is intended to represent a large number of disk particles on roughly the same orbit and size as one another, and is characterized by three numbers: the physical radius, the bulk density, and the total mass of the disk particles represented by the tracer. We developed statistical algorithms that follow the dynamical and collisional evolution of the tracers due to the presence of one another. The tracers mainly dynamically interact with the larger objects ({\\it planetary embryos)} in the normal $N$-body way. LIPAD's greatest strength is that it can accurately model the wholesale redistribution of planetesimals due to gravitational interaction with the embryos, which has recently been shown to significantly affect the growth rate of planetary embryos \\citep{Levison:2010AJ:139:1297}. We verify the code via a comprehensive set of tests which compare our results with those of Eulerian and/or direct $N$-body codes. ", "introduction": "\\label{sec:intro} The construction of a comprehensive, end-to-end model of the accumulation of the terrestrial planets and giant planet cores has been an elusive goal for planetary scientists because of the huge dynamic range inherent in the problem. The region of the proto-planetary disk from which the planets formed originally contained something like $10^{\\sim\\!14}$ objects with radii {\\color{MyChange}perhaps as small as $\\sim\\!100\\,$m \\citep{Weidenschilling:2011Icar:214:671} or as large as $1000\\,$km \\citep{Johansen:2011EM&P:108:39} depending on the planetesimals formation model}. These objects grew into the planets via a process that includes both complex collisional (both accumulation and fragmentation) and dynamical evolution. In addition, at different stages of this process, the action occurred on very different temporal and physical scales, making the construction of comprehensive models very difficult. Take, for an example, the formation of terrestrial planets. Studies have shown that {\\color{MyChange} once the first macroscopic planetesimals have formed (which is a field of study in itself),} solid body growth can occur in three distinct stages. In the first stage, planetesimals grow by so-called {\\it runaway} accretion \\citep{Wetherill:1989Icar:77:330, Greenberg:1978Icar:35:1}. During this stage, the largest objects do not affect the dynamical state of the rest of the disk and so an object's mass accretion rate scales as $M^{4/3}$. As a result, the largest bodies grow the fastest --- mainly by feeding off of much smaller objects. \\cite{Ida:1993Icar:106:210} showed that runaway accretion ends when the growing planets are only roughly $100\\times$ their original mass. Because this stage requires the study of hundreds of billions of objects, the codes used to study it employ Eulerian statistical algorithms which divide the problem into a multidimensional grid, usually 2-dimensional in heliocentric distance and size \\citep{Wetherill:1989Icar:77:330, Spaute:1991Icar:92:147, Kenyon:1999aj118:1101, Kenyon:2001aj:121:538, Morbidelli:2009Icar:204:558, Bromley:2011apJ:731:101} which evolves the total mass, and RMS eccentricity and inclination in each bin. These codes usually accurately follow the detailed collisional/fragmentational evolution of system, while using relatively simple, semi-analytic equations to evolve the dynamics. These dynamical equations are appropriate in this stage because the dynamics are local and well behaved --- there is little dynamical mixing and the surface density of the system remains smooth. In the middle stage, the largest bodies become big enough to gravitationally ``stir their own soup'' of planetesimals \\citep{Ida:1993Icar:106:210, Kokubo:1998Icar:131:172, Kokubo:2000Icar:143:15, Thommes:2003Icar:161:431, Chambers:2006ApJ:652L:133}, and thus the mass accretion rate of the largest bodies scales as $M^{2/3}$. In this phase, the largest few objects at any given time are of comparable mass. As the system evolves, the mass of the system is concentrated into an ever-decreasing number of bodies, known as {\\it planetary embryos}, of increasing masses and separations. This stage ends at a given location in the disk when the surface density of the local ``oligarchs'' becomes similar to that of the planetesimals {\\color{MyChange}\\citep{Kenyon:2006aj:131:1837}}. This occurs when the largest bodies reach roughly half their so-called {\\it isolation mass}, which is the mass they would have if they had consumed all planetesimals within their gravitational reach. In the terrestrial planet region, typical disk models produce isolation masses of only about Mars mass --- thus a third, very violent, phase must take place in which these bodies' orbits cross and they collide to form Earth- and Venus-mass bodies. The middle- and late-stages have mainly been studied with direct $N$-body simulations \\citep[][for example]{Chambers:1998Icar:136:304, Agnor:1999Icar:172:219, Chambers:2001Icar:152:205, OBrien:2006Icar:184:39, Raymond:2009Icar:203:644}. The $N$-body codes accurately follow the dynamical evolution of the system, which is necessary because there is much mixing and chaotic behavior. Studies of these stages are required to represent the large number of planetesimals remaining in the system by a smaller number of more massive {\\it tracer} particles in order to make the problem computationally tractable. In addition, they assume that when two bodies collide, they merge with 100\\% efficiency; there is no fragmentation. There have been a couple of attempts at constructing an end-to-end simulation of planet formation that started with a population of small planetesimals and built a complete planetary system \\citep[][for example]{Spaute:1991Icar:92:147, Weidenschilling:1997Icar:128:429, Kenyon:2006aj:131:1837, Bromley:2011apJ:731:101}. These have employed codes that graft an $N$-body algorithm onto Eulerian statistical code. The dynamics of the growing planetary embryos are handled correctly by the $N$-body algorithm, the accretion/fragmentation of the planetesimals are handled by the Eulerian code, and the interaction between the two populations are handled via analytical expressions (for example, applying dynamical friction to the embryos by the planetesimals). The embryos can affect the eccentricities and inclinations of the planetesimals, but not their surface density distribution. The last point above is likely to be a serious limitation of these algorithms. In \\citet[][hereafter LTD10]{Levison:2010AJ:139:1297} we showed that the growth rate of planetary embryos is strongly effected by the wholesale redistribution of planetesimals due to gravitational interaction with the embryos, themselves. In particular, we found that growth can stop if a gap opens around a embryo. In addition, the embryos can migrate as a result of gravitational scattering of the nearby planetesimals \\citep[see also][]{Fernandez:1984Icar:58:109, Hahn:1999AJ:117:3041, Ida:2000AJ:534:428, Levison:2007prpl.conf:669, Kirsh:2009Icar:199:197}. This so-called {\\it planetesimal driven migration} can significantly enhance growth \\citep[LTD10;][]{Minton:2012Icar:sub}. Unfortunately, this result calls into question the bulk of the models of the early stages of planet formation because they rely on algorithms that do not take this process into account. We realized that in order to adequately incorporate our results into full planet formation simulations would require a totally new, Lagrangian approach to the problem. Fortunately, we also realized that the code used in LTD10 supplied us with a basic structure in which to develop this algorithm. Here we report on the first particle-based Lagrangian code that can follow the dynamical/collisional/accretional evolution of a large number of km-sized planetesimals through the entire growth process to become planets. We call this code {\\it LIPAD} for Lagrangian Integrator for Planetary Accretion and Dynamics. In \\S{\\ref{sec:code}}, we describe the code in detail. In \\S{\\ref{sec:tests}}, we present a comprehensive set of tests and show that LIPAD represents the behavior of a systems containing a large number of planetesimals better than Eulerian codes. Finally, our conclusions are presented in \\S{\\ref{sec:conl}}. ", "conclusions": "\\label{sec:conl} We presented the details of the first particle based (i.e$.$ Lagrangian) code that can follow the collisional/accretional/dynamical evolution of a large number of km-sized planetesimals through the entire growth process to become planets. We refer to it as the {\\it Lagrangian Integrator for Planetary Accretion and Dynamics} or {\\it LIPAD}. LIPAD is built on top of SyMBA, which is a symplectic $N$-body integrator \\citep{{Duncan:1998aj116:2067}}. In order to handle the very large number of planetesimals required by planet formation simulations, we introduce four types of particles in LIPAD: \\begin{enumerate} \\item{} {\\it Tracers:} These objects are intended to represent a large number of planetesimals on roughly the same orbit and size as one another. Each tracer is characterized by three numbers: the physical radius $s$, the bulk density $\\rho$, and the total mass of the disk particles represented by the tracer, $m_{\\rm tr}$. As a result, each tracer represents $N_{\\rm tr}\\!=\\!m_{\\rm tr}/{{4\\over 3}\\pi\\rho s^3}$ planetesimals. They gravitationally interact with each other through Monte Carlo routines that include viscous stirring, dynamical fraction, and collisional damping (\\S{\\ref{ssec:tracer}}). They are gravitationally stirred by the larger objects (i.e$.$ full- and sub-embryos, see immediately below) via the $N$-body routines. \\item{} {\\it Full-Embryos:} These are the most massive objects in LIPAD. They interact with all classes of particles through the direct summation of individual forces already present in the SyMBA code. SyMBA routines also monitor whether physical collisions occur. The algorithm that LIPAD uses to handle these collisions is described in \\S{\\ref{sssec:EmCol}}. \\item{} {\\it Sub-Embryos:} These objects interact with full-embryos and each other through SyMBA routines. However, the only dynamical effect that the tracers have on them is through dynamical friction and planetesimal-driven migration routines (\\S{\\ref{ssec:embryos}}). Collisions are handled in the same way as those of the full-embryos. \\item{} {\\it Dust Tracers:} These are tracers that can no longer fragment. The user can set the code so that these objects do not interact with the other tracers. However, they always interact with the embryos via SyMBA's $N$-body routines. The user also has the option to apply Poynting-Robertson drag. \\end{enumerate} Perhaps LIPAD's greatest strength is that it can accurately model the wholesale redistribution of planetesimals due to gravitational interaction with the embryos, which has recently been shown to significantly affect the growth rate of planetary embryos, themselves (LTD10). This redistribution controls growth in two ways. First, it can open gaps around the embryos thereby effectively stopping accretion. Additionally, the embryos can migrate as a result of gravitational scattering of the near-by planetesimals, which can enhance growth. On the minus side, LIPAD struggles with being able to accurately resolve the side-distribution of collisional tails. However, we show that this does not affect embryo growth rates. We have carefully verified and tested LIPAD. In \\S{\\ref{sec:tests}}, we present experiments that independently exercise all of LIPAD's abilities. We find that it out performs Eulerian statistical algorithms previously used to study this problem. Of particular note, LIPAD's viscous stirring routines are particularly accurate. Our Lagrangian approach has an advantage over most previous attempts to study planet formation because, rather than using analytical expressions to estimate the global evolution of the system, our code mimics the important micro-physics (i$.$e., local accelerations and individual collisions) and lets the global system evolve naturally. Therefore, there are fewer assumptions made, and the interactions between different mechanisms are handled more realistically. {\\color{MyChange} LIPAD has many free parameters. In addition, we have not discussed any issues concerning the convergence of the code. Of course, each problem is different and so a user is going to be required to investigate these issues on their own. Thus, we decided that the best approach is to present an illustrative example of a production run we are currently undertaking with LIPAD. In particular, we are doing a series of simulations of terrestrial planet formation in the region between 0.7 and $1.5\\,$AU. This region of the planetary system is populated with $2.9\\,M_\\oplus$ of planetesimals whose initial radii varied from 10 to $50\\,$km depending on the run. The surface density of the planetesimals initially scale as $r^{-1.5}$. We represent these planetesimals with $12,000$ tracers of $1.4 \\times 10^{24}\\,$g (roughly 50\\% more massive than Ceres). This implies that tracers transition to sub-embryos when they have a radius of only $\\sim\\!450\\,$km. The transition from sub-embryo to embryo is set to $2.8 \\times 10^{26}\\,$g, roughly 200 times a tracer mass. We use $N_{\\rm s-bin}=30$ spanning sizes from 1 to $450\\,$km, and have 166 annular bins stretching from 0.5 to $60\\,$AU. We set $s_{\\rm dust} = 30\\,\\mu$ and adopt BA99's values for high velocity rock in the calculations of $Q^*_D$. We embed this material in a minimum mass gas disk \\citep{Hayashi:1985prpl.conf:1100}, which we assume has an opacity of 2\\% the ISM value when in the atmospheres of the embryos. Type~I eccentricity damping is included with $c_e=1$, but we disable Type~I migration by setting $c_a=0$. The gas disk decays with a lifetime of $2\\,$My. The timestep for the $N$-body code is set to 0.025 years, while the one for the statistical code is 3 times longer. We will present an analysis of these calculations in an upcoming paper. }" }, "1207/1207.2084_arXiv.txt": { "abstract": "The GREGOR Fabry-P\\'erot Interferometer (GFPI) is one of three first-light instruments of the German 1.5-meter GREGOR solar telescope at the Observatorio del Teide, Tenerife, Spain. The GFPI allows fast narrow-band imaging and post-factum image restoration. The retrieved physical parameters will be a fundamental building block for understanding the dynamic Sun and its magnetic field at spatial scales down to 50~km on the solar surface. The GFPI is a tunable dual-etalon system in a collimated mounting. It is designed for spectropolarimetric observations over the wavelength range from 530--860~nm with a theoretical spectral resolution of ${\\cal R}\\approx 250,000$. The GFPI is equipped with a full-Stokes polarimeter. Large-format, high-cadence CCD detectors with powerful computer hard- and software enable the scanning of spectral lines in time spans equivalent to the evolution time of solar features. The field-of-view of $50\\arcsec \\times 38\\arcsec$ covers a significant fraction of the typical area of active regions. We present the main characteristics of the GFPI including advanced and automated calibration and observing procedures. We discuss improvements in the optical design of the instrument and show first observational results. Finally, we lay out first concrete ideas for the integration of a second FPI, the Blue Imaging Solar Spectrometer, which will explore the blue spectral region below 530~nm. ", "introduction": "\\label{SEC01} Solar physics has made tremendous progress during recent years thanks to numerical simulations and high-resolution, spectropolarimetric observations with modern solar telescopes such as the Swedish Solar Telescope\\cite{2003SPIE.4853..341S}, the Solar Optical Telescope on board the Japanese HINODE satellite,\\cite{2007SoPh..243....3K} and the stratospheric Sunrise telescope.\\cite{2010ApJ...723L.127S} Taking the nature of sunspots as an example, many important new observational results have been found, e.g., details about the brightness of penumbral filaments, the Evershed flow, the dark-cored penumbral filaments, the net circular polarization, and the moving magnetic features in the sunspot moat. Telescopes with apertures of about 1.5~m such as the GREGOR solar telescope\\cite{2007msfa.conf...39V, 2010SPIE.7733E..18V, 2010AN....331..624V, 2012arXiv1202.4289S} or the New Solar Telescope\\cite{2006SPIE.6267E..10D, 2010SPIE.7733E..93C} will help to discriminate among competing sunspot models and to explain the energy balance of sunspots. New results on the emergence, evolution, and disappearance of magnetic flux at smallest scales can also be expected. However, these 1.5-meter telescopes are just the precursors of the next-generation solar telescopes, i.e., the Advanced Technology Solar Telescope\\cite{2008SPIE.7012E..16W} and the European Solar Telescope,\\cite{2010SPIE.7733E..15C} which will finally be able to resolve the fundamental scales of the solar photosphere, namely, the photon mean free path and the pressure scale height. Fabry-P\\'erot interferometers (FPIs) have certainly gained importance in solar physics during the last decades, because they deliver high spatial and spectral resolution, and a growing number of such instruments is in operation at various telescopes. Although most of the instruments have been initially designed only for spectroscopy, most of them have now been upgraded to provide full-Stokes polarimetry.\\cite{2011A&A...533A..21P, 2011SoPh..268...57M, 2010SPIE.7733E..14R} The Universit\\\"ats-Sternwarte G\\\"ottingen developed an imaging spectrometer for the German Vacuum Tower Telescope (VTT) in the early 1990s. This instrument used a universal birefringent filter (UBF) as an order-sorting filter for a narrow-band FPI mounted in the collimated light beam.\\cite{1992A&A...257..817B} The spectrometer was later equipped with a Stokes-$V$ polarimeter and the UBF was replaced by a second etalon in 2000.\\cite{1995A&A...304L...1V} A fundamental renewal of the G\\\"ottingen FPI during the first half of 2005 was the starting point of the development of a new FPI for the 1.5-meter GREGOR solar telescope.\\cite{2006A&A...451.1151P} New narrow-band etalons and new large-format, high-cadence CCD detectors were integrated into the instrument, accompanied by powerful computer hard- and software. From 2006 to 2007, the optical design for the GREGOR Fabry-P\\'erot Interferometer (GFPI) was developed, the necessary optical elements were purchased, and the opto-mechanical mounts were manufactured.\\cite{2007msfa.conf...45P} An upgrade to full-Stokes spectropolarimetry followed in 2008.\\cite{2008A&A...480..265B, 2009IAUS..259..665B, 2011ASPC..437..351B} In 2009, the Leibniz-Institut f\\\"ur Astrophysik Potsdam took over the scientific responsibility for the GFPI, and the instrument was finally installed at the GREGOR solar telescope.\\cite{2010SPIE.7735E.217D} During the commissioning phase in 2011, three computer-controlled translation stages (two filter sliders and one mirror stage) were integrated and the software was prepared for TCP/IP communication with external devices according to the Device Communication Protocol (DCP).\\cite{2011arXiv1111.5509P} This permits automated observing and calibration procedures and facilitates easy operations during observing runs. ", "conclusions": "" }, "1207/1207.3683_arXiv.txt": { "abstract": "{The progenitors of many Type II supernovae have been observationally identified but the search for Type Ibc supernova (SN Ibc) progenitors has thus far been unsuccessful, despite the expectation that they are luminous Wolf-Rayet (WR) stars. } {We investigate how the evolution of massive helium stars affects their visual appearances, and discuss the implications for the detectability of SN Ibc progenitors. } {Evolutionary models of massive helium stars are analysed and their properties compared to Galactic WR stars. } {Massive WR stars that rapidly lose their helium envelopes through stellar-wind mass-loss end their lives when their effective temperatures -- related to their hydrostatic surfaces -- exceed about 150kK. At their pre-supernova stage, their surface properties resemble those of hot Galactic WR stars of WO sub-type. These are visually faint with narrow-band visual magnitudes $M_v = -1.5 \\cdots -2.5$, despite their high bolometric luminosities ($\\log L/L_\\odot = 5.6 \\cdots 5.7$), compared to the bulk of Galactic WR stars ($M_v < -4$). In contrast, relatively low-mass helium stars that retain a thick helium envelope appear fairly bright in optical bands, depending on the final masses and the history of the envelope expansion during the late evolutionary stages.} {We conclude that SNe Ibc observations have so far not provided strong constraints on progenitor bolometric luminosities and masses, even with the deepest searches. We also argue that Ic progenitors are more challenging to identify than Ib progenitors in any optical images. } ", "introduction": "\\label{sect:intro} Direct detections of the progenitor stars of supernovae discovered in the nearby Universe provide one of the most stringent constraints on stellar evolution theory. Since the unambiguous identification of the progenitor of the supernova 1987A, observers have identified progenitor stars of numerous Type IIP supernovae and several Type IIb and IIn supernovae in pre-supernova images \\citep[see][for a review]{Smartt09}. However, despite their relevance to Galactic chemical evolution, the progenitor stars of supernovae (SNe) Ibc remain as yet elusive. For decades, there had been a widely held belief that SNe Ibc progenitors are massive Wolf-Rayet (WR) stars, formed either through stellar-wind mass-loss or Roche-lobe overflow in a binary system. Also the identification of broad SN Ic features in the vicinity of several long GRBs \\citep{Galama98, Hjorth03, Stanek03} for which the progenitors were also suggested to be massive WR stars \\citep{Woosley93} added further evidence to this assertion. An alternative to the massive WR scenario is that of lower mass binaries \\citep[e.g.][]{Podsiadlowski92, Wellstein99, Eldridge08, Yoon10}. Although many previous searches for SNe Ibc progenitors were hampered by insufficient detection limits, WR stars with $\\log L/L_\\odot \\gsim 5.3$ could have been identified as the progenitors of the Type Ib supernova 2000ds~\\citep{Maund05a}, and the Type Ic supernovae 2004gt~\\citep{Maund05b} and 2002ap~\\citep{Crockett07}. On the other hand, the most recent stellar evolution models predict that WR stars originating from massive single stars with initial masses higher than about 30~\\Msun{} have bolometric luminosities of $\\log L/L_\\odot \\gsim 5.4$ at the pre-supernova stage~\\citep{Meynet03, Crockett07, Georgy12}. New detection limits, e.g., on SNe 2002ap, have mostly been interpreted as evidence against the massive WR scenario, but in favour of progenitors of lower mass binaries instead and largely dismissing the option of massive WR stars \\citep[e.g.,][]{Crockett07, Smartt09, Yoon10}. However, in these interpretations the crucial final stages of massive-star evolution are not properly accounted for, as we argue in the following. In this paper, we present the seemingly counter-intuitive idea that the more massive -- and bolometrically more luminous -- single WR progenitors are more challenging to detect than low-mass binaries. ", "conclusions": "Our study has shown that the masses of SNe Ibc progenitor stars are not linearly correlated with their optical brightness and that the evolution of the surface properties of massive helium stars during their final evolutionary stages should be carefully investigated. The non-detection of a SN Ibc progenitor even with a good detection limit does not necessarily imply that its progenitor is a relatively low-mass helium star, and vice versa. This should be properly taken into account in future observational efforts to directly identify SNe Ibc progenitors." }, "1207/1207.1686_arXiv.txt": { "abstract": "Within a density-dependent relativistic mean-field model using in-medium meson-hadron coupling constants and meson masses, we explore effects of in-medium hyperon interactions on properties of neutron stars. It is found that the hyperonic constituents in large-mass neutron stars can not be simply ruled out, while the recently measured mass of the millisecond pulsar J1614-2230 can constrain significantly the in-medium hyperon interactions. Moreover, effects of nuclear symmetry energy on hyperonization in neutron stars are also discussed. ", "introduction": "The in-medium hyperon interactions play an important role in determining properties of hypernuclei and the hyperonization in neutron stars. Conversely, observed properties of hypernuclei and neutron stars can be used to constrain the in-medium hyperon interactions. For instance, it has been shown that the properties of hypernuclei are indeed very useful in extracting in-medium hyperon potentials at subsaturation densities~\\citep{Fr2007}. In bulk matter, such as neutron stars, hyperons can be produced by virtue of strong interactions. They can actually become important constituents of neutron stars and thus have important effects on astrophysical observations (for a review, see~\\citep{Gl2001}). In fact, it is well known that the hyperonization can reduce the maximum mass of neutron stars as much as 3/4$M_\\odot$~\\citep{Gl1985,Gl1991,Ji2006}. Currently, a number of phenomenological models, considering only the minimum compositions of nucleons and leptons using interaction parameters that are well calibrated by using terrestrial nuclear laboratory data, can produce the maximum mass of neutron stars around or below 2$M_\\odot$ ~\\citep{da02,Pi2007}. Interestingly, several neutron stars with large masses around $2M_\\odot$ have been observed~\\citep{Ni05,Ni08,Oz06} recently. In particular, the $2M_\\odot$ pulsar J1614-2230 was measured rather accurately through the Shapiro delay~\\citep{De10}. Since properties of neutron stars are determined by the nuclear equation of state (EOS) and the hyperonization reduces significantly the maximum mass of neutron stars, it has been stated that these observations seem to rule out almost all currently proposed hyperon EOS. Recent evidence for this can also be found in the work of the Brueckner approach~\\citep{Sch2011}. Though most of the hyperon EOSs give lighter neutron stars, there were actually a few endeavors in the past to stiffen the EOS either by invoking strong repulsions for hyperons or pushing upwards the onset density of hyperons, leading to heavy neutron stars involving hyperons~\\citep{Ho01,Ta02,St07,De08}. Recently, by virtue of nonlinear self-interacting terms involving a vector meson with hidden strangeness, Bednarek et al. obtained a stiff hyperon EOS with which the large mass of the PSR J1614-2230 can be produced~\\citep{Be2011}. Usually, the SU(6) relations are imposed to constrain the meson-hyperon coupling constants. Such SU(6) relations were recently reexamined by Weissenborn et al. and an arbitrary breaking of such relations can also result in the stiffening of hyperon EOS and the uplift of the maximum mass of neutron stars with the inclusion of hyperons~\\citep{We11}. On the other hand, the quark deconfinement may occur in the medium as the spatial overlap of nucleons becomes sufficient to dissolve the boundary of color singlets with the increase of density. Of course, the quantitative understanding of such a color deconfinement in cold medium is still model dependent. Typical effective QCD models include the NJL-like models, e.g., see~\\citep{Kl2007,Ip2008,Pa2008,Bo2012} and widely used MIT bag models, e.g., see~\\citep{Pr1997,Al05,We11b}, as well as Schwinger-Dyson approaches~\\citep{Li2011}. With the Maxwell or Gibbs constructions for the hadron-quark phase coexistence, the resulting quark EOS may give rise to two possible types of stars: strange stars which are totally made of absolutely stable strange matter~\\citep{Gl2000} and hybrid stars with a quark core and hadron out-layer. In order to be consistent with the recent observation of the 2$M_\\odot$ pulsar, the strong coupling and/or color superconductivity were shown to be necessary~\\citep{Xu03,Al05,Kl2007,Ip2008,Pa2008,We11b,Bo2012}. While the compositions of hybrid stars are rather model dependent, it is interesting to mention that Yasutake et al. suggested the hyperon suppression with the MIT model using a density-dependent bag constant~\\citep{Ya2011}. Noticing the recent investigations on the consequences of various quark EOS, we examine in this work the consistency of the hyperon EOS with the recent observation of the PSR J1614-2230. Despite the impressive progress made in recent decades in constraining the nuclear EOS using both astrophysical observations and nuclear reaction data, see, e.g.~\\citep{You99,da02,Li08}, many uncertainties still remain. The in-medium hyperon interactions are among the most uncertain ingredients of neutron star models. We shall thus seek in-medium hyperon interactions that can produce the observed maximum mass of neutron stars. In view of the fact that some microscopic theories, such as the Brueckner approach, are still having difficulties to obtain the $2M_\\odot$ of hyperonized neutron stars, here we resort to the phenomenological models developed in refs.~\\citep{Ji2007,Ji2007b} to analyze effects of various in-medium hyperon interactions on properties of neutron stars. ", "conclusions": "Since the parameters for the nucleonic sector of the SLC and SLCd are clearly given in~\\citep{Ji2007b}, we just list in table~\\ref{t:t1} the parameters for the hyperonic sector. We see in table~\\ref{t:t1} that in the same case all parameters but $g_{\\rho\\Xi}^0$ are the same for the models SLC and SLCd. This is because the unique difference between the models SLC and SLCd is that the latter has a softer symmetry energy than the former. Since the neutron star properties are rather insensitive to the coupling parameters of $\\Sigma$ and $\\Xi$ hyperons owing to their small fractions in the core of neutron stars, for simplicity we take $f_{\\sigma \\Sigma}=f_{\\sigma \\Xi}=f_{\\sigma \\Lambda}$ for the SC in the calculation. For the similar reason, we prefer the choice $X_{\\omega \\Lambda}=X_{\\omega \\Sigma}=X_{\\omega \\Xi}$ in the UC calculation, unless otherwise denoted. For the $\\rho$ meson in the UC, the usual relation $X_{\\rho \\Sigma}=2X_{\\rho \\Xi}=2$ is used. In both the SC and UC, the $g_{\\rho\\Lambda}^*$ is zero. Note that all parameters used in this work for hyperons meet the relation (\\ref{eqhpot}). As a well-known consequence, the emergence of the hyperon degree of freedom results in the softening of the EOS. While the models SLC and SLCd that were constructed based on the BR scaling, the softening turns out to be too appreciable at high densities to stabilize the neutron star in the UC with relatively small $X_{\\omega\\Lambda}$. As shown in the upper panel of Fig.~\\ref{fmass}, this is related to the rapid decrease of the nucleon effective mass, corresponding to an increasingly large scalar field that provides the attraction. As known before~\\citep{Ji2007}, the vector coupling constant is decisive to generate a stiff EOS at high densities. Thus, the stiffening of the EOS can follow from increasing the parameter $X_{\\omega\\Lambda}$. With larger $X_{\\omega\\Lambda}$, for instance, $X_{\\omega\\Lambda}=0.9$, the EOS is stiffened to recover the stability of neutron stars. For the SC, the EOS can be stiffened by increasing the parameter $f_{\\sigma Y}$, since the latter results in the decrease of the scalar coupling constant, as shown in Fig.~\\ref{fcoup}. Generally, the EOS obtained with the SC is much stiffer than that with the UC. Meanwhile, the accelerating decrease of the baryon effective mass due to the inclusion of hyperons can be greatly suppressed in the SC, as shown in the lower panel of Fig.~\\ref{fmass}. In Fig.~\\ref{fmass}, the $\\Lambda$ and $\\Xi$ hyperon effective masses are also displayed. The upward shift of hyperon masses at high densities in the lower panel of Fig.~\\ref{fmass} is due to the decrease of the source term (namely, the hyperon density) of the strange mesons, also see below. The hyperon effective mass is much larger than the nucleon one. This would justify the use of the different in-medium interactions for hyperons and nucleons in the SC. \\begin{figure} \\epsscale{1.0} \\plotone{nmss5.eps} \\caption{(Color online) Baryon effective masses as a function of density in the UC and SC with the SLC. Three curves (solid, dashed and dotted) of the UC in the upper panel are calculated with $X_{\\omega\\Xi}=0.6$ and $X_{\\omega\\Lambda}=2/3,0.8$ and 0.9, respectively. In the lower panel, three curves (solid, dashed and dotted) of the SC are obtained with $f_{\\sigma\\Lambda}=f_{\\sigma\\Xi}= 0.7, 0.8$ and 0.9, respectively. \\label{fmass}} \\end{figure} \\begin{figure*}[thb] \\epsscale{1.5} \\plotone{fsl31.eps} \\caption{(Color online) Particle fractions in the UC and SC with the model SLC as a function of density. In the UC (left panel), three curves for each particle, e.g., denoted by the number 1, 2, and 3, are calculated with $X_{\\omega\\Lambda}=2/3$, 0.8 and 0.9, respectively, while $X_{\\omega\\Xi}=0.9$. The same number in the left panel denotes the results obtained with the same parameters. In the SC (right panel), three curves for each particle, for instance, in a rising order of the $\\Lambda$ hyperon appearance, are obtained with $f_{\\sigma\\Lambda}=f_{\\sigma\\Xi}= 0.7, 0.8$ and 0.9, respectively. \\label{ffrac}} \\end{figure*} The dropping of the baryon effective mass in chemically equilibrated matter associates tightly with the beginning density and fractions of hyperons. In Fig.~\\ref{ffrac}, we display the particle fractions as a function of density. We see that the fractions of hyperons are rather sensitive to the variation of the in-medium interactions. More noticeably, quite large differences between the left and right panels can be observed. For instance, in the UC (left panel), the interval between the $\\Lambda$ and $\\Xi^-$ beginning densities depends sensitively on the ratios of coupling constants, although the relation (\\ref{eqhpot}) is always satisfied, while in the SC such a sensitivity does not exist. Moreover, the emergence of the $\\Sigma$ hyperon depends on the model interaction. The $\\Sigma$ hyperon does not appear in the SC. Remarkably, with the increase of density, the hyperon fractions in the SC tend downwards till to disappear after reaching the maximum, as shown in the right panel of Fig.~\\ref{ffrac}. The occurrence of this phenomenon is due to the fact that the vector meson-hyperon coupling constant $g_{\\omega Y}^*$ has a weaker density dependence [see Eq.(\\ref{eqphiY})] than the meson-nucleon coupling constant $g_{\\omega N}^*$. At high densities, $g_{\\omega Y}^*$ exceeds $g_{\\omega N}^*$, and so do the vector potentials. The chemical equilibrium thus makes the hyperon Fermi momenta and fractions lower. As an application in studying properties of neutron stars, this actually results in the exclusion of hyperons in the core of neutron stars and accordingly the re-stiffening of the EOS at high densities. In addition to our scheme, we note that there are other attempts to decrease the number of hyperons in neutron stars. For instance, different couplings of a new boson to hyperons and nucleons were proposed to decrease the number of hyperons~\\citep{Kr2009}. Including the nonlinear self-interactions involving a vector meson with hidden strangeness, Bednarek et al. found that the onset density of hyperons can be as high as $3\\rho_0$ accompanied by smaller fractions of hyperons~\\citep{Be2011}. In~\\citep{Ta02,Ts2009}, however, hyperons were found to appear above $4\\rho_0$ and thus played a rather limited role in the EOS of neutron star matter. In this work, we find that the onset density of hyperons in the SC can vary upwards within the region $2.2-3\\rho_0$ as the ratio of vector meson coupling $g_{\\omega\\Lambda}^{0}/g_{\\omega N}^0$ increases from 0.2 to 0.8 (In table~\\ref{t:t1}, this ratio is 2/3), while the hyperon fractions get suppressed significantly at larger onset densities. A large rise of this density up to $4.5\\rho_0$ can be obtained. However, the binding relation (\\ref{eqhpot}) should then be reduced to $U^{(N)}_\\Lambda(\\rho_0)=0$ MeV. In this case, the hyperon fraction becomes so small that the hyperonic constituents become unimportant. On the quark level, it is interesting to see that the mechanism of small numbers of strange quarks in hybrid stars was explored within a specific quark model~\\citep{Bu2004}. We mention that the baryon fractions in neutron stars are also sensitive to the symmetry energy. In~\\citep{Ji2006}, it was illustrated that the onset density for hyperons increases moderately with the softening of the symmetry energy. For the same reason, the onset density for hyperons with the SLCd is about 0.2$\\rho_0$ larger than those with the SLC. To save space herein, we do not display particle fractions with the SLCd in a figure similar to Fig.~\\ref{ffrac}. \\begin{figure}[thb] \\epsscale{1.0} \\plotone{peden31.eps} \\caption{EOS of isospin-asymmetric matter with models SLC and SLCd. Curves are obtained with parameters listed in Table~\\ref{t:t1}. The arrow indicates the critical point to quark matter. The result including the muon but without hyperons is also depicted for comparisons. \\label{fpeden}} \\end{figure} Although the emergence of hyperons is a cause for softening the nuclear EOS, the specific behavior relies indeed on the in-medium interactions, as clearly shown in Fig.~\\ref{fpeden}. The softening persists at high densities for the UC. However, the softening is succeeded by a stiffening in the SC where hyperons feel a different in-medium interaction from nucleons. Looking back to the right panel of Fig.~\\ref{ffrac}, we see that the stiffening of the EOS occurs with the suppression of hyperon fractions. As the hyperon vanishes, the EOS returns to the normal EOS without hyperons. As shown in Fig.~\\ref{fpeden}, the EOS evolves to become stiffer than the normal one with increasing density. The neutron star matter thus transits to the normal isospin-asymmetric matter prior to the vanishing of hyperons. This eventually leaves a limited density window allowing the existence of hyperons in neutron stars. Interestingly, we find that the influence of the hyperonization in the SLCd model falls prominently, as compared to the SLC model. This is attributed to the larger onset density of hyperons with the SLCd due to the softening of the symmetry energy, as mentioned above. In the UC, since the nucleon effective mass vanishes at certain critical density, we need to consider the EOS beyond the critical density. Though the relationship between the chiral restoration and deconfinement occurrence is still discussed, it is nevertheless a convenient and usual way to neglect the distinction between them. Beyond the critical density, we thus adopt a quark matter EOS described by the MIT model with the appropriate bag parameter [$(179.5 MeV)^4$] to connect smoothly the hadron matter and quark matter EOSs. The connection to the quark matter further softens the EOS, as shown in Fig.~\\ref{fpeden}. \\begin{figure}[thb] \\epsscale{1.0} \\plotone{msr222.eps} \\caption{(Color online) Mass-radius relation of neutron stars in the UC and SC with the SLC and SLCd. In the SC, two curves are obtained with $f_{\\sigma Y}=0.8$ and 0.9, respectively. The hatched areas give the probability distributions with $1\\sigma$ (red) and $2\\sigma$ (green) confidence limits for $r_{ph}\\gg R$ summarized in~\\citep{St2010}. \\label{fmsr}} \\end{figure} We now turn to the consequences of hyperonizations on properties of neutron stars. In particular, it is interesting to see whether our model with hyperonization can give properties of neutron stars that are compatible with the recent observation of the millisecond pulsar J1614-2230. The mass of this pulsar was accurately determined to be $1.976\\pm0.04 M_\\odot$~\\citep{De10}. It was concluded that such a large mass can rule out almost all currently proposed hyperon equations of state. The mass-radius relation of neutron stars is obtained from solving the standard TOV equation with the nuclear EOS being specified above. For the low-density crust, we adopt the EOS in~\\citep{Ba1971,Ii1997}. It is seen in Fig.~\\ref{fmsr} that all EOS's of the SC can produce neutron stars with the maximum mass around $2M_\\odot$ and with the corresponding radius ranging from $9.3$ to $9.5$ km. Since the maximum mass of neutron stars is more sensitive to the EOS at high densities, it is understandable that the EOS including hyperonization at intermediate densities can lead to masses compatible with that of the pulsar J1614-2230. Thus, the EOS with hyperonization can not be simply ruled out. Of course, the hyperon fraction in neutron stars with the SC is rather limited, as compared to that with the UC. In addition, a subtle factor affecting the hyperonization is the density dependence of nuclear symmetry energy. Due to the softening of the symmetry energy in the SLCd, the hyperonization in the SC becomes unimportant for the EOS of neutron star matter, see Fig.~\\ref{fpeden}, which equivalently expels most hyperons in neutron stars. Consequently, the mass-radius relation with the SLCd is almost independent of the parameter $f_{\\sigma Y}$ in the SC, as shown in Fig.~\\ref{fmsr}. For the UC, it looks undoubtedly that the EOS is ruled out by the recent observation~\\citep{De10}, since the maximum mass of neutron stars in this case just sprints to $1.4M_\\odot$ with a much compacter size than the canonical neutron star without hyperons. It is worth adding here some discussions about the influence of the symmetry energy on the radii of neutron stars. It is now well established that the maximum mass of neutron stars is dominated by the high-density behavior of the EOS, while the radius is primarily determined by the slope of the symmetry energy at intermediate densities (1 -3$\\rho_0$) ~\\citep{La2001,La2004,St2005,Xu2009}. The large difference between the radii of low-mass neutron stars obtained with the SLC and SLCd can be attributed to the difference in the slope parameter $L$. As discussed in detail earlier in ref.~\\citep{Fat10}, the central density of an 18 - 20 km star is near $\\rho_0$, and the crust ends approximately at $(1/3-1/2)\\rho_0$. In this density range, the pressure is dominated by the symmetry energy and not the incompressibility of symmetric nuclear matter. Our results shown in Fig. 5 are consistent with those in ref.~\\citep{Fat10}. For the 1.4$M_\\odot$ neutron stars in the SC, we see however that the inclusion of hyperons reduces moderately the range of the neutron star radius. For instance, the radius ranges from 11.4km with the SLCd to 12.3km with the SLC for $f_{\\sigma Y}=0.9$, which are well situated in the domain extracted by Steiner et al.~\\citep{St2010}. With the emergence of hyperons, the sensitivity of the star radius to the symmetry energy is reduced clearly. This is because the $\\Lambda$-hyperon, being the dominant component of hyperons, is an isospin scalar. The suppression of the isovector potential $g_\\rho^*b_0$ due to the appearance of $\\Lambda$ hyperons~\\citep{Ji2006} results in the reduction of the pressure of asymmetric matter and thus the star radius, while the magnitude of the suppression depends on the specific values of the symmetry energy and hyperon fraction. As a result, the less sensitivity of the star radius to differences in the symmetry energy is observed in calculations with both the SC and UC. In Fig.~\\ref{fmsr}, we also include the constraints of mass-radius trajectories for $r_{ph}\\gg R$ obtained by Steiner et al. in~\\citep{St2010}. Our results in the SC are either within (for SLCd) or not very far off (for SLC) the constrained region, though the inclusion of hyperons seems to tilt the vertical trajectories. For low-mass neutron stars, the radii with the SLC are predicted to go beyond the optimal region extracted very recently in~\\citep{St2012b}, unless some other scenarios that allow a loose extension of the radii are invoked to extract the radius constraints~\\citep{Su2011,Zh2007}. However, the maximum mass of neutron stars is almost independent of the slope $L$ and the radii of low-mass neutron stars. In our cases, the maximum mass is only reduced by about 1\\% by softening the symmetry energy from the model SLC to SLCd. On the other hand, it is known that the measurements of the neutron star radii are far less precise than the mass measurements, see, e.g., refs.~\\citep{Su2011,Zh2007} and references therein. We note that an option of $r_{ph}= R$ can cause a visible slanting of the vertical trajectories~\\citep{St2010}. The agreement of our results can thus be better with the constraints obtained for $r_{pc}=R$. We note that there had been a few endeavors in the past to involve the hyperons in heavy neutron stars either by invoking strong repulsions for hyperons or pushing upwards the onset density of hyperons\\citep{Ho01,Ta02,St07,De08,Be2011,We11}. The main purpose of our work is to constrain the in-medium hyperon potentials using the 2 solar mass constraint of neutron stars. This indeed requires strong repulsions for hyperons. Due to different density-dependencies of nucleonic and hyperonic potentials in the SC, the hyperonic vector potential exceeds the nucleonic one, leading to a very significant suppression of hyperons at high densities. As a result, the hyperons would exist in a shell of neutron star core even with a small $L$ value. Besides the effect of hyperonization on the static properties, another consequence of the hyperonization is on the thermal evolution of neutron stars. For the SC in SLC, we see from Fig.~\\ref{fmsr} that $\\Lambda$ hyperons start to appear above $2.5\\rho_0$. At this central density the neutron star mass is about 1.2$M_\\odot$. For most observed neutron stars that have larger masses, it appears that the direct Urca (DU) process with nucleons \\citep{La1991} and/or hyperons \\citep{Pr1992} would occur since the proton fraction in the neutron star matter with hyperonization can be in excess of the DU threshold (14\\%) with the SC. According to the thermal evolution of observed neutron stars analyzed in \\citep{Pa2004,Ya2004}, the fast cooling with the DU processes seemed to be excluded in most neutron stars except the massive ones. Slower cooling is possible when the neutrino emissivity can be suppressed by the superfluity of constituent particles like nucleons and hyperons when the temperature falls below the critical temperature. Page et al. set a stringent requirement for the critical temperature of neutron superfluidity without which enhanced cooling from DU processes may be needed in at least half of the observed young cooling neutron stars~\\citep{Pa2009}. While the DU cooling involving nucleons only was regarded to be too fast, it was pointed out by Tsuruta et al. that the DU cooling with hyperons in neutron stars can be compatible with the observations, provided that the hyperon superfluidity is appropriately accounted for~\\citep{Ts2009}. In our models with the SC, the hyperonization can thus be favorable to make up the potential incompatibility in the thermal evolution by considering the hyperon superfluidity. On the other hand, the threshold mass of neutron stars which allow the DU process increases with the softening of the symmetry energy due to which the proton fraction exceeds the threshold value at larger densities. For instance, this mass is around 1.3$M_\\odot$ with the SC of the model SLCd, and the threshold density for the DU process is about 4$\\rho_0$, which is coincidently the onset density of hyperons in~\\citep{Ts2009}. Finally, we discuss some details concerning the in-medium interactions for hyperons. Constrained by the large mass of observed pulsars, it is favorable for us to select the in-medium interactions for hyperons in the SC. It is however interesting to find that the single-particle potentials for hyperons in the UC and SC almost overlap in a large density domain ranging from zero density to intermediate density. The significant departure appears only at high densities (roughly $\\ge 4\\rho_0$). Especially at lower densities, the SC and UC descriptions give rather limited differences in single-particle properties. This thus indicates that without compromising the success in describing properties of hypernuclei, one can constrain significantly the high-density hyperon interactions with the large mass of neutron stars. In addition, we have noticed that many efforts using various quark models with the postulate of strong interactions and/or color superconductivity can also lead to large-mass hybrid stars~\\citep{Al05,Kl2007,Ip2008,We11b,Bo2012} with stiff EOS's at high densities. Our models respecting chiral limits possess the similarly stiff EOS at high densities. On the other hand, neutron stars may be composed of more complicated constituents including quarks and meson condensates. It is useful to consider these non-baryonic degrees of freedom and their interplay with hyperons to constrain more quantitatively the in-medium interactions for hyperons. However, this is beyond the scope of the present work." }, "1207/1207.6773_arXiv.txt": { "abstract": "We study the prospects for studying line features in gamma-ray spectra with upcoming gamma-ray experiments, such as HESS-II, the Cherenkov Telescope Array (CTA), and the GAMMA-400 satellite. As an example we use the narrow feature at 130 GeV seen in public data from the Fermi-LAT satellite. We found that all three experiments should be able to confidently confirm or rule out the presence of this 130 GeV feature. If it is real, it should be confirmed with a confidence level higher than 5$\\sigma$. Assuming it to be a spectral signature of dark matter origin, GAMMA-400, thanks to a projected energy resolution of about 1.5 \\% at 100 GeV, should also be able to resolve both the $\\gamma\\gamma$ line and a corresponding $Z\\gamma$ or $H\\gamma$ feature, if the corresponding branching ratio is comparable to that into two photons. It will also allow to distinguish between a gamma-ray line and the similar feature resulting from internal bremsstrahlung photons. ", "introduction": "As the Large Hadron Collider (LHC) keeps accumulating data at high luminosity (and soon at full energy), hopes are high that it will help elucidating the nature of the particle making up around 23 \\% of the energy density of the universe, the {\\it dark matter} particle \\cite{bertone_book,reviews}. So far, no such new mass scale has been found, although the prediction from supersymmetric (SUSY) models that the lightest Higgs boson should weigh less than 130 GeV~\\cite{lhc_higgs}, which seems to be confirmed by the detection recently done at CERN's LHC, which gives a mass of the potential Higgs boson of around 125 GeV. As for dark matter candidates, only constraints on the parameter space of the most popular extensions of the Standard Model, in particular Supersymmetry, have been obtained \\cite{SUSYfits}, but even if such candidates were to be found, it will be hard to prove with LHC data only that they actually constitute most of the dark matter in the Universe, as the required lifetime of many times the age of the universe would seem impossible to verify in accelerator experiments \\cite{Bertone:2010at}. Fortunately, {\\it direct} and {\\it indirect} dark matter searches will provide complementary information, possibly allowing a precise identification of dark matter particles \\cite{Bertone:2010rv,Bertone:2011pq}. Direct detection by scattering of dark matter particles traversing the earth in ultra-pure counting experiments has historically been the most advanced technique, but indirect detection methods have recently received increased interest (see e.g. Ref. \\cite{bertone_book} for reviews). Indirect detection is based on the search for secondary photons, antimatter, and neutrinos produced by the annihilation or decay of dark matter particles. For $\\gamma$-rays coming from annihilations of dark matter particles in the halo, Fermi-LAT has very successfully delivered bounds that have started to probe into the parameter space of viable models, in line with pre-launch expectations \\cite{prelaunch}, in particular for dwarf spheroidal galaxies \\cite{fermidwarfs} and galaxy clusters \\cite{fermiclusters}. Recently, a possible hint of a dark matter signal in the form of a narrow {\\it spectral line} or an {\\it internal bremsstrahlung} (IB) feature, has been found in analyses of public data from the Fermi-LAT satellite detector \\cite{bringmann, weniger} (see also \\cite{tempel, su_fink}). The signal is too weak to claim a discovery, but being of a type and at an energy where there is no other known astrophysical explanation\\footnote{The very fine-tuned pulsar model from \\cite{aharonian_pulsar} can be disregarded since the signal is significantly extended.} it is important to further study this type of signature in independent experiments. We take in this paper, as an exercise, the existence of these recent indications for a line or an IB bump seriously, and we discuss how this effect, if real, would appear in a number of existing (Fermi-LAT \\cite{fermi-lat}, HESS-II \\cite{hess-2}) and planned (CTA \\cite{cta}, GAMMA-400 \\cite{gamma-400}) $\\gamma$-ray detectors. If the present indications of a line structure in the Fermi-LAT public data would disappear, our results should be useful for future indirect dark matter searches. In particular, we discuss the possibility of one or more associated lines coming from the $Z\\gamma$ and $H\\gamma$ (with $H$ the Higgs boson) annihilation channels in some models, and we investigate whether one can separate a line signal from other spectral features like IB emission. We will show that with the upcoming detector HESS-II, and the proposed Cherenkov Telescope Array (CTA) and the GAMMA-400 satellite these gamma-ray structures, if real, indeed will be confirmed with much higher confidence. The paper is organized as follows: in Sec. \\ref{sec:theory} we review the spectral features arising from dark matter annihilation and the tentative detection of a feature in Fermi data. In Section \\ref{sec:prospects} we discuss the prospects to detect such feature with future observatories focusing on improvements in energy resolution and effective area. In \\ref{sec:conclusions} we discuss our results and present our conclusions. ", "conclusions": "\\label{sec:conclusions} The detection of a sharp feature at an energy of 130 GeV in Fermi-LAT data has sparked the interest of the astroparticle community, since the presence of gamma-ray lines has long been considered a smoking-gun signature of new physics, possibly pointing to the annihilation of dark matter particles. Of course, future Fermi-LAT data will be very important: If the Fermi-LAT collaboration can exclude instrumental effects as the cause of the structure, it may well, in case upcoming data strengthens the feature, confidently establish discovery of the effect. In any case, future gamma-ray observatories would provide necessary independent confirmation and are expected to clarify the experimental situation, in view of their increased effective area or better angular resolution. In particular we focused here on three upcoming experiments: HESS-II, CTA and GAMMA-400. We summarize here the main results: \\begin{itemize} \\item We have calculated the sensitivity to gamma-ray lines for the three experiments, and we have shown that all of them will be able to confirm or rule out the presence of the 130 GeV line. In all cases, in fact, the feature found in Fermi-LAT data would be detectable with a significance higher than 5$\\sigma$. \\item We have assessed, for each experiment, the prospects for identifying the presence of additional lines, which would allow a better reconstruction of the particle properties of the annihilating dark matter particle. We found that only GAMMA-400, thanks to a claimed energy resolution of about 1.5~\\% at 100 GeV, will be able to separate a $\\gamma\\gamma$ line from a $Z\\gamma$ or $H\\gamma$, if the corresponding branching ratio is comparable to that into two photons, while HESS-II and CTA cannot separate them. \\item We investigate for which signal fluxes a signal arising from the annihilation into $\\gamma\\gamma$ can be distinguished from IB photons, and found that GAMMA-400 would be able to distinguish at a $95\\%$CL between a gamma-ray line and IB photons if the 130 GeV feature is real, and we have identified the broader region of the parameter space where a discrimination is possible. \\end{itemize} HESS-II will soon be operational and given the good performances foreseen for the instrument in hybrid-mode, we stress that it should offer a quick confirmation of the genuineness of the signal reported in \\cite{weniger} (our estimates are based on an exposure time of 50 hours, assuming intermediate zenith angles), since this could provide on a short timescale an independent observation with completely different background and systematic errors. As for CTA, the actual construction of the array should start in 2015, and the first data should realistically be available by 2018. In the case of GAMMA-400, the claimed improvement in energy and angular resolution over Fermi-LAT make it an invaluable tool for dark matter searches. We have demonstrated that it has an enormous potential in the detection and discrimination of lines, despite the smaller effective area compares to the NASA satellite, and we therefore strongly encourage this experimental effort. \\smallskip" }, "1207/1207.1962_arXiv.txt": { "abstract": "{} {We study the possible detection of and properties of very high-energy (VHE) \\gr\\ emission (in the energy band above 100~GeV) from high redshift sources.} { We report on the detection of VHE \\gr\\ flux from blazars with redshifts $z>0.5$. We use the data of Fermi telescope in the energy band above 100~GeV and identify significant sources via cross-correlation of arrival directions of individual VHE \\gr s with the positions of known Fermi sources. } {There are thirteen high-redshift sources detected in the VHE band by Fermi/LAT telescope. The present statistics of the Fermi signal from these sources is too low for a sensible study of the effects of suppression of the VHE flux by pair production through interactions with Extragalactic Background Light photons. We find that the detection of these sources with ground-based \\gr\\ telescopes would be challenging. However, several sources including BL Lacs PKS 0426-380 at $z=1.11$, KUV 00311-1938 at $z=0.61$, B3 1307+433 at $z=0.69$, PG 1246+586 at $z=0.84$, Ton 116 at $z=1.065$ as well as a flat-spectrum radio quasar 4C +55.17 at $z=0.89$ should be detectable by HESS-II, MAGIC-II and CTA. A high-statistics study of a much larger number of VHE \\gr\\ sources at cosmological distances would be possible with the proposed high-altitude Cherenkov telescope 5@5. } {} ", "introduction": "\\gr\\ emission from distant blazars is suppressed by the effect of pair production through interactions of these \\gr s with the low-energy photons forming the Extragalactic Background Light (EBL) \\citep{gould67,kneiske,stecker06,mazin07,franceschini08,gilmore}. This prevents observations of high-redshift sources using the technique of imaging the Cherenkov emission produced through \\gr\\ induced air showers, used by the ground-based Cherenkov telescopes HESS, MAGIC and VERITAS \\citep{aharonian_review}. Most of the VHE \\gr\\ loud blazars, detected by these telescopes, are situated in the local Universe, at (several) hundred megaparsec distances, in the redshift range $z\\sim 0-0.2$\\footnote{\\tt http://tevcat.uchicago.edu}. Only one source at redshift $z>0.5$, 3C 279, was possibly detected by the MAGIC telescope during a flaring activity \\citep{3C279_MAGIC}. Another source at redshift 0.444, 3C 66A, was detected by VERITAS \\citep{3c66a}. One more relatively high redshift source, PG 1553+113 (at $z>0.4$, reported by \\citet{danforth}) has been detected by MAGIC and VERITAS \\citep{PG1553_MAGIC,PG1553_Veritas}. Measurement of the effect of suppression of the VHE \\gr\\ flux from low-redshift blazars is commonly used for the estimation of the density of the EBL in the local Universe \\citep{kneiske,stecker06,mazin07,franceschini08,gilmore}. Constraints on the EBL density and spectral characteristics, extracted from the VHE blazar observations, are useful for understanding the cosmological evolution of galaxies. Observations of the suppressed VHE \\gr\\ emission from blazars at non-negligible redshifts can provide a measurement or constraint not only of the EBL density in the local Universe, but also new information on the cosmological evolution of the EBL, which is largely uncertain. Observations of high-redshift sources in the VHE band would also provide valuable information on the cosmological evolution of the VHE \\gr\\ loud blazars (and, more generally, \\gr\\ emitting active galactic nuclei, AGN), which is also highly uncertain. This information is important, in particular, for an understanding of the origin of Extragalactic \\gr\\ background \\citep{fermi_background,neronov_EGB}. Most of the VHE \\gr\\ emitting blazars at zero redshift belong to a High-energy peaked BL Lac (HBL) sub-class of the blazar population. Negative cosmological evolution of this sub-class was reported based on the observations in the visible and X-ray band \\citep{giommi}. Such puzzling evolution is, apparently, opposite to the general positive evolution of blazars and radio galaxies (BL Lac parent AGN population) with the cumulative power of the sources increasing as $(1+z)^k$, $k>0$ \\citep{hodge09,sadler07,smolcic09}. Independent verification of this hypothesis using \\gr\\ observations would be possible only if a significant amount of VHE emitting blazars could be detected in the redshift range $z\\sim 0.5-1$. Taking into account the importance of the study of the VHE \\gr\\ emission from distant AGN, we report here on the detection of VHE \\gr\\ signals from sources at redshifts $z>0.5$ by the Large Area Telescope (LAT) on board of Fermi satellite \\citep{atwood09}. The effective area of LAT, about 1~m$^2$, is several orders of magnitude smaller than that of the ground-based \\gr\\ telescopes, so that LAT has detected only one-to-several photons from the brightest extragalactic VHE \\gr\\ sources in $\\sim 4$~years of observations. However, an extremely low level of the background (including the residual cosmic ray background and Galactic / extragalactic \\gr\\ background) makes even a several-photon signal in the energy band above 100~GeV significant. This fact has been used by \\citet{ic310,100GeV_paper1} for the search of extragalactic VHE \\gr\\ sources via the analysis of clustering of VHE \\gr s on the sky, based on the method first proposed by \\citet{Gorbunov_method}. In this paper we use the correlation of arrival directions of VHE \\gr s detected by Fermi with the positions of high-redshift blazars to identify high-redshift sources of VHE \\gr s and to study their spectral characteristics. ", "conclusions": "The detection of VHE \\gr\\ emission at energies above 100~GeV from blazars with redshifts beyond $z=0.5$ carries important implications for the possibility of studying the cosmological evolution of the blazar population and the overall energy output from the galaxies, in the form of the EBL. \\begin{figure} \\includegraphics[height=\\linewidth]{Emax_vs_z} \\caption{Energies of VHE photons from BL Lacs and FSRQ as a function of redshift. Curves correspond to the optical depth $\\tau=1$ (black) and $\\tau=3$ (grey) with respect to absorption on the EBL. Solid line: EBL from {\\citet{franceschini08}}; dotted line: EBL from {\\citet{gilmore}}; short-dashed line: EBL from {\\citet{finke}}; long-dashed line: EBL from {\\citet{kneiske}}} \\label{fig:photons} \\end{figure} Fermi, being able to discover the VHE \\gr\\ signal from the high-redshift sources, does not have a capability of studying the details of the spectral characteristics of these high-redshift sources. At the same time, Fermi's detections of these sources provides a clear indication for the selection of the high-redshift targets for observations with existing and next-generation ground-based \\gr\\ telescopes. In the previous section we have compared the Fermi flux measurements in the $E>100$~GeV band with the sensitivity of the ground-based \\gr\\ telescopes. The results from such comparisons indicate that the flux levels of sources listed in table \\ref{tab:list} are already at or below the sensitivity limit of the current generation ground-based telescopes and are, in fact just at the sensitivity limit of the next generation facility CTA. This means that Table \\ref{tab:list} provides a more-or-less exhaustive list of the high-redshift sources accessible for detection using ground-based \\gr\\ telescopes. The situation with the limited capabilities for the detection of high-redshift sources from the ground could, nevertheless, change if a telescope system specially optimized for the reduction of the low-energy threshold, like 5@5 would be realized. In this case most of the sources listed in Table~\\ref{tab:list} would be detectable by 5@5 already at energies of about 10~GeV, with very high signal statistics. This would allow a high-quality study of the details of the spectral characteristics of the VHE \\gr\\ emission from high redshift sources. Lowering the energy threshold with 5@5 would open the possibility for detecting a much larger number of the high-redshift sources, compared to the just thirteen sources listed in Table~\\ref{tab:list}. Low statistics of the VHE \\gr\\ signal from the Fermi high redshift blazars does not allow a study of the effect of absorption of \\gr s through their interactions with EBL photons on a source-by-source basis. However, the significance of this effect for different models of cosmological evolution of EBL could be evaluated collectively for all thirteen sources from Table \\ref{tab:list}, following the method discussed by \\citet{abdo10}. Fig. \\ref{fig:photons} shows the energies of VHE \\gr s from distant blazars as a function of the source redshift. The data from the previous analysis at lower energies, reported by \\citet{abdo10}, are shown in grey in the same figure for comparison. The black/grey curves show the energies at which the optical depth with respect to pair production on the EBL is $\\tau=1$ and $\\tau=3$, respectively. This implies a flux suppression by a factor of $\\simeq 3$ and $\\simeq 20$. A naive expectation is that there should be no VHE photons from the sources in the upper part of the plot, much beyond the corresponding $\\tau=1$ or $\\tau=3$ curves. From Fig. \\ref{fig:photons} one can see that this trend of decreasing photon counts beyond the $\\tau=1$ curve holds in the redshift range $z=0.5-1$, with only one photon beyond the solid black curve corresponding to $\\tau=1$ for the \\citet{franceschini08} model of EBL evolution. At the same time, the trend seems to be broken for redshifts of about $z\\gtrsim 1$, where a large number of photons beyond the $\\tau=1$ curve appears. This large number of photons beyond the $\\tau=1$ curve is definitely not due to the larger source number in this redshift range. Table \\ref{tab:list} contains seven sources at $z<1$ and six sources at $z\\ge 1$. Thus, the larger statistics of photons beyond the $\\tau=1$ curve should be due to a different reason. One possibility is that the determination of redshifts of the $z\\ge 1$ sources listed in the table \\ref{tab:list} are not reliable and, in fact, the sources are at lower redshifts. In the previous section it is mentioned that this might be the case for B2 0912+29. Another possibility is that at least one or two photons in the $z>1$ part of the diagram might still be from the background and not from the high-redshift sources. This might be the case for the highest energy photon beyond $\\tau=3$ curves, which came from the direction of RGB J0250+172. This photon has a 0.8\\% probability of being from the background. If none of these possible explanations holds (this needs to be checked with more data and verification of the redshift measurements), then the over-abundance of the VHE photons beyond the $\\tau=1$ curves in the redshift range $z>1$ in Fig. \\ref{fig:photons} should be due a physical effect. One possibility is that the EBL level at high redshifts is lower, so that the Universe is more transparent to the VHE photons, than what is implied by the currently existing models. Otherwise, new propagation effects, such as origin of the \\gr\\ emission in the ultra-high-energy cosmic ray induced cascade \\citep{Essey2009,murase,kusenko} or even by new physics, like axions \\citep{axions}, should be considered. In the scenario of \\citet{Essey2009}, UHE protons with energies in the EeV range propagate cosmological distances and lose energy primarily through the proton pair production process. In this case secondary TeV gamma-rays are produced by such protons at distances 100-300 Mpc and easily reach the Earth. For this model to be valid, one needs a relatively low extra-galactic magnetic field with values 1-10 pG everywhere on the way of the protons, not only in the voids of large scale structure. Apparently, a support for such a scenario is found in a recent work of \\citet{Landt2012}, in which absorption lines at the redshift beyond $z=1$ were found in the spectra of two BL Lacs, PKS 0447-439 and PMN J0630-24. One of these BL Lacs has been recently observed in the VHE band by HESS {\\citep{zech}. Explanation of the detection of PKS 0447-439 in the VHE band within conventional models would be challenging if the source redshift is indeed $z>1$ \\citep{kusenko}. Thorough verification of the redshift measurement for both PKS 0447-439 and for the $z\\ge 1$ sources listed in the Table \\ref{tab:list} is extremely important for the clarificaiton of consistency / inconsistency of the VHE band detections of the high-redshift sources with the current understanding of the mechanisms of formation of the VHE \\gr\\ spectra of the sources and of the cosmological evolution of the EBL." }, "1207/1207.1709_arXiv.txt": { "abstract": "The Jeans analysis is often used to infer the total density of a system by relating the velocity moments of an observable tracer population to the underlying gravitational potential. This technique has recently been applied in the search for Dark Matter in objects such as dwarf spheroidal galaxies where the presence of Dark Matter is inferred via stellar velocities. A precise account of the density is needed to constrain the expected gamma ray flux from DM self-annihilation and to distinguish between cold and warm dark matter models. Unfortunately the traditional method of fitting the second order Jeans equation to the tracer dispersion suffers from an unbreakable degeneracy of solutions due to the unknown velocity anisotropy of the projected system. To tackle this degeneracy one can appeal to higher moments of the Jeans equation. By introducing an analog to the Binney anisotropy parameter at fourth order, $\\beta'$ we create a framework that encompasses all solutions to the fourth order Jeans equations rather than the restricted range imposed by the separable augmented density. The condition $\\beta' = f(\\beta)$ ensures that the degeneracy is lifted and we interpret the separable augmented density system as the order-independent case $\\beta'=\\beta$. For a generic choice of $\\beta'$ we present the line of sight projection of the fourth moment and how it could be incorporated into a joint likelihood analysis of the dispersion and kurtosis. The framework is then extended to all orders such that constraints may be placed to ensure a physically positive distribution function. Having presented the mathematical framework, we then use it to make preliminary analyses of simulated dwarf spheroidal data leading to interesting results which strongly motivate further study. ", "introduction": "The favoured $\\Lambda$CDM model of cosmology is consistent with a large invisible non-baryonic component of matter. To infer its existence astronomers thus look for the gravitational effect of its significant mass upon luminous tracer objects or for the observable products of DM annihilation and/or decay such as gamma rays \\citep{gunn78,stecker78} which have been used in recent searches \\citep[e.g][]{abdo} for dark matter. In both instances the density distribution of the system is critical with the Earth-incident flux of annihilation products dependent not only on model-dependent properties derived from particle physics \\citep[see e.g.][]{pieri2009} but also on the square of the density distribution of dark matter within the source. This is encoded in what is known as the \\textit{astrophysical J-factor} which can be written \\citep{walker2011}, \\begin{equation}\\label{jfac} J(\\theta_{\\text{int}}) = \\frac{4\\pi}{d^2}\\int^{\\theta_{\\text{int}}d}_{0}r^2\\rho^{2}_{\\text{dm}}(r)\\text{d}r \\end{equation} where $d$ is the distance to the source, $\\rho_{\\text{dm}}(r)$ is the local density of dark matter and $\\theta_{\\text{int}}$ is the integration angle which is related to a given solid angle of the source via $\\Delta\\Omega=2\\pi(1-\\cos\\theta_{\\text{int}})$. The quadratic dependence upon the density introduces a very significant and DM model independent contribution to the flux which makes the choice of astrophysical source critically \\citep{penarrubia} important for optimising DM searches. Though one might expect that the galactic center would provide the strongest signal, the strong and chaotic astrophysical backgrounds make it arguably less favourable than dwarf spheroidal galaxies (dSphs) of the local group which have been identified as having a large mass-to-light discrepancy \\citep{aaronson} suitable for dark matter searches \\citep{Lake, evans04}. This, in conjunction with their relative proximity to earth, makes them natural laboratories for DM and in recent years it has been possible to obtain samples of stellar positions and velocities \\citep[e.g][]{walkdat} that are large enough for statistical treatment. Since the typical angular resolution of gamma ray telescopes is larger than the angular size on the sky of dwarf spheroidal galaxies \\citep{abdo}, it turns out that the J-factor is not extremely sensitive to the distribution of dark matter in the core of dwarf spheroidals, although in the event of a signal being observed we would ideally want a better indication of the J-factors than the range of current estimates which vary over about an order of magnitude. Another reason why it is important to study the centre of dwarf spheroidals is to probe another aspect of dark matter, namely its primordial velocity. In the cold dark matter hypothesis the kinetic energy of dark matter at the start of structure formation is some very small fraction of its rest mass energy \\citep{stefan, green} and the smallest structures above this free streaming scale are the first to form. In models of hot dark matter \\citep[see e.g.][]{zeldovich} dark matter begins completely relativistic and the largest structures form first. A (rather finely tuned) compromise between these two extremes is the idea of warm dark matter \\citep{wdm} where dark matter is not created with highly relativistic velocities, but nevertheless with significant velocities, meaning that the normal growth pattern of cold dark matter proceeds only above a length scale related to the free streaming length corresponding to the initial velocity. This idea has been invoked to explain the lack of predicted satellites of the Milky Way \\citep{wherearethey} as well as some interpretations of tracer populations in dwarf spheroidals wherein it is argued that dark matter halos possess a significant core, possibly due to some inherent initial kinetic energy \\citep{gilmore}. At the same time, the importance of the role of baryons upon dark matter density in the core of halos is becoming increasingly clear \\citep{governato}. Whatever the underlying physics, it is clear that we would like to be able to interpret the stellar velocity dispersion in such objects more effectively. To infer the DM density from the kinematic data the Jeans analysis is used to relate the joint distribution of tracer stars' positions and velocities to the underlying potential of the system. Traditionally the second order Jeans equation \\citep{binney} is used to generate the velocity dispersions for a set of input parameters including the potential which is then fitted to the dSph data with a likelihood analysis by radially binning the line of sight velocities for the variance and assuming Gaussianity. It has long been known however \\citep{dejonghe87,merritt87} that this analysis may not be used to uniquely specify the potential for anisotropic systems for which the variances of the radial and tangential velocity components are not equal. As it is only possible to observe the projected quantities along the line of sight, the intrinsic dispersions of the system are convolved such that there is a degeneracy of indistinguishable solutions to the Jeans analysis. Indeed it has been shown that the observed line of sight dispersion can be generated by any given parameterisation of the anisotropy parameter \\citep{evans2009} thus leaving the potential almost completely unconstrained. This is the so-called Jeans degeneracy problem which is the main subject of this work. A discussion of the higher order moments is presented herein with a mathematical description of how they enter the Jeans analysis and what assumptions are required to ensure that the Jeans degeneracy is solved or at least partially lifted. This is then placed into the practical context of a joint likelihood analysis of the variance and kurtosis in dwarf spheroidal galaxies. We evaluate the contribution by \\cite{Lokas02} in establishing a model for the kurtosis that may be used to lift the degeneracy \\citep{Lokas05} and extend the method to general anisotropy as proposed by \\cite{an11a} with the separable augmented density system. To simplify the mathematical description of the higher order Jeans equations there has been much success in the literature since the advent of the augmented density formalism by \\cite{dejonghe86}. Whilst an application \\citep{dejonghe87,baes07} of this method has generally been limited to models \\citep{plummer,hernquist} with particularly simple potential-density pairs, the recent work of An (2011b) demonstrates for generic density and anisotropy that a separable system \\citep{Ciotti} solves the Jeans degeneracy problem completely by specifying moments at all orders with the potential and anisotropy parameter alone. This is however by no means a general solution and without a strong physical motivation its practical use is difficult to evaluate. An alternative hierarchy of \\textit{pseudomoments} \\citep{King} in spherical systems \\citep{Amendt}, tailored to minimise the increasing dependence of standard moments to the tails of the distribution, breaks the degeneracy with physical arguments for weak nonisothermality. A key issue with the pseudomoment method is accessibility of the observable standard moments which has not yet been shown to be universal. With practical intent we thus persist with the standard moments for direct comparison with the data. Recently there have been a number of alternative original methodologies presented for breaking the degeneracy. \\cite{penarrubia} and \\cite{amoriscomulti} utilise the existence of chemodynamically distinct stellar populations in the Fornax and Sculptor galaxies to exploit the robustness of the mass profile to degenerate anisotropy at the stellar half radius and are able to derive an estimate of the mass slope $\\Gamma = d\\log M / d\\log r$ that places stringent constraints on cusped density profiles. When taken at face value this observation, together with the additional apparent problem of missing satellites \\citep{pen12}, is in tension with the $\\Lambda$CDM model (although the role of baryons in the shaping the inner core of dark matter halos is very complicated \\citep{governato}). Though powerful this method relies on multiple populations that may not exist in other dwarf spheroidals such as Carina \\citep{penarrubia} and assumes Gaussianity in the line-of-sight velocity profiles that is inconsistent with the Jeans equations at orders higher than two. Whilst an application of the Jeans analysis to multiple populations is straightforward, at higher orders the number of input parameters quickly becomes impractical and more importantly still, splitting the population increases the errors associated with limited sampling. One way to mitigate this problem that is employed in the analysis of elliptical galaxies \\citep[see e.g][]{Bender} is to use Gauss-Hermite moments \\citep{Gerhard, Franx} that efficiently measure the shape of the distribution with less reference to the tails of the distribution than the traditional kurtosis. An extension of this method to discrete data sets \\citep{Amorisco} enables an efficient extraction of non-Gaussian shape parameters suitable for an analysis of dwarf spheroidals. Though such an analysis is statistically preferable to conventional moment analysis it is difficult to ensure that the parametrised prior distributions are both physical and exhaustive. As hinted by \\cite{Gerhard} it would be interesting to see whether a Jeans-like analysis to the Gauss-Hermite moments is viable which would ensure that the shape parameters are fitted to equilibrated systems. If information from non-Gaussianities can break the degeneracy then we choose as a simple first step to investigate what the well-established Jeans formalism can tell us for spherically symmetric systems such as dwarf spheroidal galaxies. Numerical methods such as the orbit-superposition algorithm \\citep{Schwarzschild} have also recently been applied to dwarf spheroidal galaxies \\citep[see for e.g][]{Breddels,Jardel} that make no assumption on the form of the anisotropy and guarantee physical distribution functions thus providing an interesting complement to the weaknesses of the traditional analytic methods described above. The layout of our paper is as follows. In section \\ref{sec2} we review the mathematics behind the Jeans equation and line of sight calculations. To utilise the constraining power of the fourth order statistics we are thus motivated to provide the full analytical set of fourth order solutions which is achieved with the introduction of an analog to the Binney anisotropy parameter at fourth order. This is outlined in section \\ref{sec3} wherein we show how to construct a generic model for the projected fourth order moment to be incorporated into a joint likelihood analysis, extending to full generality the over-constrained method employed by \\cite{Lokas05}. In section \\ref{sec4} we outline the joint likelihood analysis of dispersion and kurtosis that through the Jeans equations allows a fit of the density and anisotropy parameters to moments extracted from LOS velocity data. Section 5 sees the method tested for a set of simulated dwarf spheroidal data sets and as a proof of concept we contrast directly the performance of the traditional and joint analysis in constraining the anisotropy and crucially the density parameters. Finally we will make some concluding remarks and outline our future research program. \\begin{section}{Preliminary}\\label{sec2} \\subsection{Moments of the Distribution Function} In the study of stellar systems, a 6-dimensional function $f$ is used to specify the distribution \\citep{jean} of stars in position and velocity space. For a spherically symmetric system this is related to the underlying gravitational potential $\\Phi(r)$ by the time-independent and collisionless Boltzmann equation \\citep{merry}, \\begin{center} \\begin{eqnarray}\\label{boltz} \\frac{\\partial f}{\\partial t} &=& v_r \\frac{\\partial f}{\\partial r} + \\left( \\frac{v^{2}_{\\theta}+v^{2}_{\\phi}}{r}-\\frac{d\\Phi}{dr}\\right) \\frac{\\partial f}{\\partial v_r} \\nonumber\\\\ &+& \\frac{1}{r}(v^{2}_{\\phi}\\cot \\theta - v_rv_{\\theta})\\frac{\\partial f}{\\partial v_{\\theta}} \\\\ &-& \\frac{1}{r}(v_{\\phi}v_r+v_{\\phi}v_{\\theta}\\cot \\theta )\\frac{\\partial f}{\\partial v_{\\phi}} \\nonumber\\\\ &=& 0. \\nonumber \\end{eqnarray} \\end{center} Multiplying \\eqref{boltz} by ${v^{l}_{r}v^{m}_{\\theta}v^{n}_{\\phi}}$ and then integrating over all velocities restates it in terms of its \\textit{true} velocity moments \\begin{equation}\\label{moms} \\nu \\overline{v^{2i}_r v^{2j}_{\\theta} v^{2k}_{\\phi}} = \\int v^{2i}_r v^{2j}_{\\theta} v^{2k}_{\\phi} f(r,\\textbf{v}) d^3v. \\end{equation} where $\\nu(r)$, as an effective zeroth moment that marginalises the distribution function in velocity space, is the local density of stars. Due to the spherical symmetry of the system it is trivial to show by averaging over the azimuthal angles that the odd moments vanish and that many of the true even moments are related by constant prefactors. To make the notation more compact we again follow the example of \\cite{merry} and introduce the \\textit{intrinsic} moments of the tangential velocity $v_t=(v^{2}_{\\theta}+v^{2}_{\\phi})^{1/2}$, \\begin{eqnarray} \\overline{v^{2i}_r v^{2j}_{\\theta} v^{2k}_{\\phi}} = \\frac{1}{\\pi} B(j+\\frac{1}{2},k+\\frac{1}{2})\\overline{v^{2i}_r\\; v^{2(j+k)}_t} \\end{eqnarray} where $B(x,y)$ is the Beta function. In the subsequent analysis we find that it is not necessary to explicitly refer to the true moments at any stage and to simplify the mathematics the tangential moments will be used exclusively from here on in. \\subsection{Jeans Equations} To isolate the second order moments i.e the radial and tangential dispersions which in practice have the smallest statistical errors, the Boltzmann equation is traditionally multiplied by $v_r$ and integrated over all velocities \\citep{binney} to give, \\begin{equation}\\label{jeans} \\frac{d(\\nu \\sigma^{2}_{r})}{dr} + \\frac{2\\beta}{r}\\nu \\sigma^{2}_{r} +\\nu \\frac{d \\Phi}{dr} = 0. \\end{equation} where $\\nu(r)$ is the local stellar density, $\\Phi(r)$ is the gravitational potential that depends on the total density of the system $\\rho(r) = \\nu(r)+\\rho_{\\text{dm}}(r)$ via, \\begin{equation}\\label{phi} \\Phi(r) = \\frac{4\\pi G}{r} \\int^{r}_{0} r^2 \\rho(r)dr \\end{equation} and the Binney anisotropy parameter \\citep{binney} $\\beta(r)$, \\begin{equation} \\label{anis} \\beta(r) \\equiv 1-\\frac{\\sigma^{2}_{t}(r)}{2\\sigma^{2}_{r}(r)}, \\end{equation} measures the deviation of the dispersions from the isotropic system $(\\sigma^{2}_{r}=\\sigma^{2}_{\\theta}=\\frac{1}{2}\\sigma^{2}_{t})$ wherein all directions in velocity space are equally probable\\footnote{We choose for mathematical convenience to adopt the 2D tangential dispersion rather than the 1D employed by e.g \\cite{Lokas02} which accounts for the additional factor of 2.}. For a dSph, where the mass-luminosity ratios are often greater than 10 \\citep{mateo} the dark matter component is very significant. To illustrate the higher order analysis we consider first the fourth order where multiplying equation \\eqref{boltz} by $v^{3}_{r}$ and $v_r v^{2}_{\\theta}$ respectively relates the three intrinsic moments at fourth order $\\overline{v^{4}_{r}},\\overline{v^{4}_{t}}$ and $\\overline{v^{2}_{r}v^{2}_{t}}$ by the two fourth order Jeans equations \\citep{merry}, \\begin{equation} \\label{hojeans1} \\frac{d(\\nu \\overline{v^{4}_{r}})}{dr} - \\frac{3}{r}\\nu \\overline{v^{2}_{r}v^{2}_{t}} + \\frac{2}{r}\\nu \\overline{v^{4}_{r}} + 3 \\nu \\sigma^{2}_{r} \\frac{d \\Phi}{dr} = 0 \\end{equation} \\begin{equation} \\label{hojeans2} \\frac{d(\\nu \\overline{v^{2}_{r}v^{2}_{t}})}{dr} - \\frac{1}{r}\\nu \\overline{v^{4}_{t}} + \\frac{4}{r}\\nu \\overline{v^{2}_{r}v^{2}_{t}} + \\nu \\sigma^{2}_{t} \\frac{d \\Phi}{dr} = 0. \\end{equation} The advent of the augmented density system by \\cite{dejonghe86} has greatly enhanced the mathematical description of the Jeans analysis and it is within this framework that the complete set of Jeans equations has been presented \\citep{an11b}, \\begin{eqnarray}\\label{genjean} \\frac{d(\\nu \\overline{v^{2p}_{r}v^{2q}_{t}})}{dr} &=& -\\frac{2}{r}\\left[(q+1)\\nu \\overline{v^{2p}_{r}v^{2q}_{t}}-(p-\\frac{1}{2})\\nu \\overline{v^{2p-2}_{r}v^{2q+2}_{t}}\\right] \\nonumber\\\\ &&-(2p-1)\\nu \\overline{v^{2p-2}_{r}v^{2q}_{t}}\\frac{d\\Phi}{dr}. \\end{eqnarray} The number of equations at 2$n$th order is therefore $n$, the number of permutations of $(p,q)$ for which $p+q=n$ and $1\\leq p \\leq n,\\; 0 \\leq q \\leq n$. Each of the $n$ moments at 2$n$th order enter the derivative of a corresponding equation apart from $v^{2n}_{t}$. We also note that as the distribution function is linear upon disassembling into $M$ stellar sub-components, it follows from \\eqref{moms} that the composite moments are related to the constituent moments via \\begin{equation} \\nu \\overline{v^{2i}_{r}\\; v^{2j}_{t}}=\\sum^{M}_{c=1}\\nu_c (\\overline{v^{2i}_r\\; v^{2j}_t})_c \\end{equation} where $\\nu_c$ is the constituent stellar density that intuitively weights the relative contribution from each sub-population. One can then show that the Jeans equations are also linear and thus that solving a Jeans equation for a composite distribution function is equivalent to simultaneously solving the sum of equations for each individual population. Indeed this must be true for stars orbiting in a common potential (like the dark matter dominated potential of a dwarf spheroidal galaxy) else one could not arbitrarily select sub-samples from the data that were equilibrated. For chemically distinct populations with different histories however we reemphasise this point and retain the larger composite population to reduce the sampling errors that are critical in a higher moment analysis. \\subsection{Projected Moments} Unfortunately due to the distant nature of astronomical objects the stellar positions and velocities may only be observed along the line-of sight and rather than the true moments and the local density one must instead use the \\textit{projected} moments and the surface density profile $ \\Sigma $ as functions of the projected radius $R$. The surface density profile is the projection along the line of sight of $\\nu(r)$, \\begin{equation} \\Sigma(R) = 2\\int^{\\infty}_{R} \\frac{\\nu(r)r}{\\sqrt{r^2-R^2}}dr. \\end{equation} which may be inverted directly via the Abel Inversion for the local density. The component of the velocity along the line of sight, which for convenience is chosen as the z-direction, may be expressed as \\begin{equation}\\label{losz} v_{\\text{los}} = v_r \\cos(a) - v_{\\theta} \\sin(a) = v_r \\sqrt{1-\\frac{R^2}{r^2}} - v_{\\theta}\\frac{R}{r}. \\end{equation} where $\\sin a = R/r$ is the unobservable depth angle that determines the extent to which each stellar velocity is radial or tangential. Finding the projected moment at 2$n$th order \\citep[see also][]{dejonghe92} is akin to averaging over the 2$n$th power of equation \\eqref{losz}, \\begin{eqnarray}\\label{nlosgen} \\Sigma\\overline{v^{2n}_{\\text{los}}}(R) &=& 2\\int^{\\infty}_{R} \\overline{(v_r \\cos a-v_{\\theta}\\sin a)^{2n}} \\frac{\\nu(r) r}{\\sqrt{r^2-R^2}}dr \\nonumber\\\\ &=& 2\\int^{\\infty}_{R} \\sum^{n}_{k=0} C_{n,k} \\overline{v^{2(n-k)}_{r}v^{2k}_{t}} \\frac{\\nu(r) r}{\\sqrt{r^2-R^2}}dr \\end{eqnarray} and one thus requires all of the intrinsic moments to calculate the projection with the coefficients, \\begin{equation} C_{n,k} = \\binom{2n}{2k}\\frac{B(k+\\frac{1}{2},\\frac{1}{2})}{\\pi} \\left(1-\\frac{R^2}{r^2}\\right)^{n-k}\\left(\\frac{R^2}{r^2}\\right)^{k}. \\end{equation} To simplify the projected dispersion the anisotropy parameter is again introduced in place of the tangential dispersion, \\begin{equation}\\label{LOSsecond} \\Sigma \\sigma^{2}_{\\text{los}}(R) = 2\\int^{\\infty}_{R} (1-\\beta\\frac{R^2}{r^2}) \\frac{\\nu \\sigma^{2}_{r} r}{\\sqrt{r^2-R^2}}dr \\end{equation} such that specifying the potential and anisotropy parameter is sufficient to first calculate the radial dispersion via equation \\eqref{jeans} and then the projected moment for comparison with observation. An explicit calculation of the projected fourth moment \\begin{equation} \\label{LOSfour} \\Sigma\\overline{v^{4}_{\\text{los}}}(R) = 2\\int^{\\infty}_{R} \\left(C_{2,0} \\overline{v^{4}_{r}} + C_{2,1}\\overline{v^{2}_{r}v^{2}_{t}} +C_{2,2}\\overline{v^4_t}\\right) \\frac{\\nu(r)r}{\\sqrt{r^2-R^2}}dr \\end{equation} \\begin{equation} C_{2,k} = \\left\\{ \\begin{array}{cc} (1-\\frac{R^2}{r^2})^2 & k=0\\\\ 3 \\frac{R^2}{r^4}(r^2-R^2) & k=1\\\\ \\frac{3}{8} \\frac{R^4}{r^4} & k=2 \\end{array}\\right. \\end{equation} recovers the result in \\cite{merry}. \\subsection{The Degeneracy Problem} The normal Jeans analysis, which approximates the distribution function by its second order moments, has traditionally been the only viable means of modeling the limited sample sizes afforded by kinematic surveys of dwarf spheroidal galaxies. Unfortunately it has been demonstrated \\citep[e.g][]{merritt87} that the integral equation \\eqref{LOSsecond} can be highly degenerate with no way of distinguishing the entangled intrinsic dispersions. In a typical statistical treatment, where a model for the dark matter density is fitted with a hand-picked form of the anisotropy, there are numerate parameter sets $p=\\{\\beta(r),\\Phi(r)\\}$ that yield identical line of sight dispersions within statistical errors. As the anisotropy parameter is a completely unknown degree of freedom it is impossible to uniquely specify the potential with the line-of-sight dispersion alone and such a treatment is liable not only to imprecision but also to inaccuracy. With a recent improvement in the data, there has been a renewed interest in the higher order moments that may be used to distinguish those parameter sets degenerate at second order. Unfortunately whilst the fourth moment is practically within reach, a theoretical problem persists. As suggested by \\cite{an11b} we note that at each successive order the Jeans analysis introduces $n+1$ moments and only $n$ constraining Jeans equations such that the intrinsic moments are not specified by the second order parameters $p$. A minimal requirement to define the system at fourth order is to specify one of the fourth order moments or, in analogy with Binney's anisotropy parameter, to specify the ratio of two. To lift the degeneracy one additionally desires that the projected fourth moment depend only upon the second order parameters such that a new net constraint is added to the system without expanding the parameter space. The anisotropy parameter in particular, whilst inherent to the dispersions, has no direct bearing on the higher order moments without artificial insertion. It can be concluded therefore that to utilise the higher order moments one must present a new constraint to the system via a simplifying assumption or empirical observation that optimally ties the higher order moments to the anisotropy parameter and thus may be used to constrain it. To lift the degeneracy completely one requires all of the projected moments which as proved by \\cite{dejonghe92} is equivalent to knowledge of the distribution function. \\subsection{Incompatibility of Equilibrium and the Assumption of Gaussian Velocity Distributions} Whilst an additional constraint is required for unique solutions to the fourth order Jeans equations, imposing more than one over constrains the equations such that there is no consistent equilibrium solution. If one assumes that the joint distribution of radial and the tangential velocities is normal and separable, \\begin{equation} f(v_r,v_\\theta,v_\\phi) = \\frac{1}{(2\\pi)^{3/2} \\sigma^{2}_{\\theta} \\sigma_r }\\exp\\left(-\\frac{v^{2}_{r}}{2\\sigma^{2}_{r}}-\\frac{v^{2}_{\\theta}+v^{2}_{\\phi}}{2\\sigma^{2}_{\\theta}}\\right) \\end{equation} then all higher orders are fixed. At fourth order the moments follow from \\eqref{boltz} \\begin{equation}\\label{gmoms} \\overline{v^{4}_{r}} = 3 \\sigma^{4}_{r} \\end{equation} \\begin{equation} \\overline{v^{2}_{r}v^{2}_{t}} = 2\\overline{v^{2}_{r}v^{2}_{\\theta}} = 2\\sigma^{2}_{r}\\sigma^{2}_{\\theta} = (1-\\beta)\\sigma^{4}_{r} \\end{equation} \\begin{equation} \\overline{v^{4}_{t}} = \\frac{8}{3}\\overline{v^{4}_{\\theta}} = 8 \\sigma^{4}_{\\theta}= \\frac{1}{2}(1-\\beta)^2\\sigma^{4}_{r}. \\end{equation} imposing three constraints to the system. Let's now take a step back and say that we enforce only the first, that the radial velocity distribution is Gaussian. From \\eqref{hojeans1} there is a unique solution for the comoment $\\overline{v^{2}_{r}v^{2}_{t}}$ where the RHS is zero. This in turn yields $\\overline{v^{4}_{t}}$ from \\eqref{hojeans2} and thus the one additional constraint specifies a unique equilibrated system. Introducing the other two constraints however completely specifies the LHS of the Jeans equations which is not guaranteed to be (in fact it is exceptionally rarely) zero. Recalling that to generate \\eqref{hojeans1}, the CBE \\eqref{boltz} is multiplied through by $v^{3}_{r}$ and integrated over all velocities we perform the same to the time derivative which yields, \\begin{equation} \\int v^{3}_{r} \\frac{\\partial f}{\\partial t} d^3 v = \\frac{d \\nu \\overline{v^{3}_{r}}}{dt} \\end{equation} i.e the rate of change (when normalised by the dispersions) of skewness in the radial velocity distribution. With the Gaussian assumption then plugging \\eqref{gmoms} into the fourth order Jeans equations enables a calculation of this effect for a given set of parameters, \\begin{equation} \\frac{d\\nu \\overline{v^{3}_{r}}}{dt} = 3\\nu\\sigma^{2}_{r}\\left[\\left(\\frac{1-3\\beta}{r}-\\frac{1}{\\nu}\\frac{d\\nu}{dr}\\right)\\sigma^{2}_{r}-\\frac{d\\Phi}{dr}\\right]. \\end{equation} \\end{section} ", "conclusions": "The Jeans analysis is extremely useful in identifying the gravitational potential from the line of sight velocities of stars moving in that potential and is used to learn more about many different kinds of astrophysical objects. The information obtained in this way is limited by degeneracies which exist as a result of our ignorance of the velocity anisotropy within the stellar tracer populations. Following \\cite{Lokas02} we looked not only at the width (2nd moment of velocity) of the stellar velocities but also their kurtosis (ratio between 2nd and 4th moment) in a bid to break some of these degeneracies. With increasingly large samples we are also motivated to consider the higher order moments because the assumption of Gaussianity is inconsistent with an equilibrium solution. A major problem with utilising the two fourth order Jeans equations however is that without additional information about the fourth moments then the system is under constrained, i.e there is no unique prediction for the fourth moment of the LOS velocity distribution given the anisotropy parameter $\\beta(r)$ and density $\\rho(r)$ alone. In the literature this has been addressed by assuming a particular form for the distribution function such as the separable augmented density wherein the ratio of fourth order moments is correlated with $\\beta$ the ratio of second moments. The nature of this assumption is, from a physical standpoint, arbitrary and restricts the range of solutions such that systems with tangentially biased variances ($\\beta<0$) must necessarily have tangential velocity distributions with flatter tops than their radial counterparts. In this paper we have presented a new mathematical framework for calculating the higher order moments of the Jeans equation based upon introducing an analogue of the Binney Anisotropy parameter at each higher order and we have demonstrated that this determines the complete set of solutions to the Jeans equation at each order. At fourth order it is shown that the introduction of the analog $\\beta^{\\prime}$ allows for a free variation of the shape parameters with $\\beta^{\\prime}>0$ naively implying a more flat topped tangential distribution and $\\beta^{\\prime}<0$ implying a more flat topped radial distribution. With the fourth order anisotropy $\\beta^{\\prime}$ as an independent parameter one can scan the entire range of distribution functions in a likelihood analysis and we show that not only is the necessary extension to the Jeans equations and projected moments straightforward but by design the separable augmented density system is the limiting case $\\beta^{\\prime} = \\beta$. With no \\textit{a priori} intuition for a correlation between the anisotropy parameters $\\beta^{\\prime} = f(\\beta)$ however the degeneracy problem is not guaranteed to be improved as the additional free parameter simply introduces a new degeneracy in the kurtosis measurement. To test whether this new degeneracy was as affecting as its notorious second order counterpart we developed a method to compare the constraints on density parameters for eight simulated dwarf spheroidal data sets with sample size, experimental errors and stellar surface density comparable to Fornax the largest existing data set. Einasto profiles were assumed for the density of dark matter with four A-D exhibiting an extended inner core and the other four E-H having a vanishing core that mimics an NFW at resolvable distances from the galactic centre. Under the assumption of constant anisotropy $\\beta$ and for the first time introducing an independent constant $\\beta^{\\prime}$ to minimally close the fourth order Jeans equation we performed a traditional dispersion-only and dispersion-kurtosis analysis of the simulated data to monitor the relative performance in recovering the input density parameters. The mass slopes corresponding to the input parameters were recovered inside the $95\\%$ confidence intervals for the MCMC output in seven out of eight dispersion-only and dispersion-kurtosis analyses for different data sets which demonstrates that the likelihood developed in Section 4 provides an accurate account of the uncertainties. Whilst the outlier for the dispersion-kurtosis is marginally excluded the dispersion-only outlier completely excludes the input core parameters. This anomaly is however rectified with the joint analysis which emphasises that an additional reference to the data can help to reduce spurious sampling effects. As expected the anisotropy parameter $\\beta$ was more tightly constrained for all data sets. In six out of eight data sets the constraints on the mass slope are tighter for the joint analysis with one of these exceptions being the dispersion-only outlier discussed above. For the most realistic cases where the limited sampling caused a scatter that completely obscured the anisotropy and mass slope with a dispersion-only analysis (B, D and E) the inclusion of the kurtosis was particularly effective and the degeneracy between the extended and vanishing core parameters was broken. Whilst a detailed investigation with other parameterisations is required to confirm this finding it provides strong motivation for further study into the higher order Jeans analysis and demonstrates that even a freely varying fourth order model could prove more constraining than a second order analysis that makes an assumption without reference to the fourth order data. Ideally we would like to test how constraining the additional fourth order information is in the case where we allow the relationship between $\\beta$ and $\\beta'$ to vary to a greater or lesser extent. It would also be interesting to try and use the results of N-body simulations to motivate physical choices for the relationship between the two anisotropy parameters. We are working on all these issues and hope to present new results for real dwarf spheroidal data sets in the near future." }, "1207/1207.3226_arXiv.txt": { "abstract": "{} {Recent studies of the optical/UV and X-ray ephemerides of X1822-371 have found some discrepancies in the value of the orbital period derivative. Because of the importance of this value in constraining the system evolution, we comprehensively analyse all the available optical/UV/X eclipse times of this source to investigate the origin of these discrepancies.} {We collected all previously published X-ray eclipse times from 1977 to 2008, to which we added the eclipse time observed by Suzaku in 2006. This point is very important to cover the time gap between the last RXTE eclipse time (taken in 2003) and the most recent Chandra eclipse time (taken in 2008). Similarly we collected the optical/UV eclipse arrival times covering the period from 1979 to 2006, adding a further eclipse time taken on 1978 and updating previous optical/UV ephemeris. We compared the X-ray and the optical/UV ephemeris, and finally derived a new ephemeris of the source by combining the eclipse arrival times in the X-ray and optical/UV bands. } {The X-ray eclipse time delays calculated with respect to a constant orbital period model display a clear parabolic trend, confirming that the orbital period of this source constantly increases at a rate of $\\dot{P}_{\\rm{orb}} =1.51(7) \\times 10^{-10}$ s/s. Combining the X-ray and the optical/UV data sets, we find that $\\dot{P}_{\\rm{orb}} =1.59(9) \\times 10^{-10}$ s/s, which is compatible with the X-ray orbital solution. We also investigate the possible presence of a delay of the optical/UV eclipse with respect to the X-ray eclipse, finding that this delay may not be constant in time. In particular, this variation is compatible with a sinusoidal modulation of the optical/UV eclipse arrival times with respect to the long-term parabolic trend. In this case, the optical/UV eclipse should lag the X-ray eclipse and the time-lag oscillate about an average value.} { We confirm that the orbital period derivative is three orders of magnitude larger than expected from conservative mass transfer driven by magnetic braking and gravitational radiation. } ", "introduction": "X1822-371 is an eclipsing compact binary system with a period of 5.57 hr hosting a 0.59 s X-ray pulsar. Several authors have reported new orbital ephemeris of the source using observations performed in different energy bands. \\citet[hereafter BU10]{Burderi_2010} analysed X-ray data of X1822-371 covering the period from 1996 to 2008 to determine the eclipse times of the source and improved the previous X-ray ephemeris of X1822-371 reported by \\citet[hereafter PA00]{parmar2000} that covered the period from 1977 to 1996. BU10 added their data to those used by PA00 finding a positive derivative of the orbital period of $(1.499 \\pm 0.071) \\times 10^{-10}$ s/s that is compatible with the previous one given by PA00 but with a smaller associated error. \\citet[hereafter BA10]{Bayless2010} obtained the optical/UV ephemeris of X1822-371 using data covering the period from 1979 to 2006. They obtained a value of the orbital period derivative of $(2.12 \\pm 0.18) \\times 10^{-10}$ s/s, which is compatible with that reported by PA00 but slightly larger than the value proposed by BU10. \\citet[hereafter JI11]{Ji2011}, using the X-ray eclipse arrival times reported by PA00 and the eclipse arrival times inferred by the two Chandra/HETG observations of X1822-371 performed in 2000 (Obs ID: 671) and in 2008 (Obs ID: 9076 and 9858), already included in the work of BU10, estimated a value of the orbital period derivative of $(0.83 \\pm 0.16) \\times 10^{-10}$ s/s, with the error at the 90\\% confidence level, almost a factor of two smaller than the value reported by BU10. We summarise the values of the eclipse reference time $T_0^e$, the orbital period $P_{\\rm{orb\\; 0}}$, and the orbital period derivative $\\dot{P}_{\\rm{orb}}$ obtained by PA00, BU10, BA10, and JI11 in Table \\ref{Tab1}. In this work, we comprehensively examine both X-ray and optical/UV eclipse arrival times to give the most updated ephemeris of X1822-371, adding to the eclipse arrival times reported by BU10 the one obtained from a Suzaku observation performed in 2006. We also include a data point from a Ginga observation performed in 1989, and a data point from a ROSAT observation performed in 1992. We critically examine the discrepancies that have emerged in calculating the orbital ephemeris in previous papers and, finally, show the ephemeris of the X1822-371 by combining the optical/UV and X-ray data-sets. \\begin{table*}[ht] \\caption{Journal of the ephemerides of X1822-371 discussed in this work.} \\label{Tab1} % \\begin{center} \\begin{tabular}{l r r r r} % \\hline\\hline % Parameters & \\cite{parmar2000}&\\cite{Burderi_2010}& \\cite{Bayless2010}&\\cite{Ji2011}\\\\ \\hline $T_0^e$ (MJD$_{\\odot}$) & 45614.80964(15)& 45614.80948(14)& 45614.81166(74)& 45614.80927(25)\\\\ $P_{\\rm{orb\\; 0}}$ (s)& 20054.1990(43)& 20054.2056(22)& 20054.1866(69)& 20054.2181(41)\\\\ $\\dot{P}_{\\rm{orb}}$ ($\\times 10^{-10}$ s/s) & 1.78(20) & 1.499(71)& 2.12(18)& 0.827(95)\\\\ $\\chi^2/(d.o.f.)$& 21.4/16& 38.69/25& 70.04/32& 35.99/19 \\\\ \\hline % \\end{tabular} \\end{center} {\\small \\sc Note} {\\footnotesize---Uncertainties are at the 68\\% c. l. for a single parameter. We show the reference time $T_0^e$ of the eclipse arrival times in units of MJD, the orbital period $P_{\\rm{orb\\; 0}}$ in units of seconds calculated at $T_0^e$, the derivative of the orbital period $\\dot{P}_{\\rm{orb}}$ in units of s/s, and finally the $\\chi^2/(d.o.f.)$ obtained fitting the eclipse arrival times with a quadratic function.} \\end{table*} ", "conclusions": "We have revisited and discussed the X-ray and optical/UV ephemerides of X1822-371. Fitting simultaneously the optical/UV and X-ray time delays, we have found that the optical/UV eclipses of X1822-371 lag the X-ray eclipses by $ 127 \\pm 52$ s with a significance level of 2.4 $\\sigma$. However, this time-lag may not be constant in time. Fitting the optical/UV time-lags, we have found a statistically significant variation, which is compatible with a sinusoidal modulation at two different periods, $\\sim 18$ yr and 239 d. In the first case, the optical/UV eclipses lag the X-ray eclipse by an average time of $161 $ s (significance 6.7$\\sigma$) and this delay oscillates in time around this value with an amplitude of $194$ s (significance 6.7$\\sigma$). In the second case, the optical/UV eclipses lag the X-ray eclipse by an average time of 105 s (significance 4.4$\\sigma$), and this delay oscillates in time around this value with an amplitude of $267$ s (significance 6.2$\\sigma$). Owing to the relatively small number of points over a long-time span of 30 yr, we cannot be sure of the period of this modulation, because we cannot exclude much shorter periods. Long and relatively continuous optical/UV observations are necessary to prove or disprove the presence of this periodicity in the optical/UV eclipse time-lags. Our results confirm the value of the orbital solution derived by the X-ray eclipse times given by BU10 and that the orbital period derivative is three orders of magnitude larger than expected on the basis of the conservative mass transfer driven by magnetic braking and gravitational radiation. We have also confirmed this result by combining the X-ray data and the optical/UV data of X1822-371." }, "1207/1207.6403_arXiv.txt": { "abstract": "The parallel code NMAGIC is an implementation of a particle-based method to create made-to-measure models in agreement with observations of galaxies. It works by slowly correcting the particle weights of an evolving N-body system, until a satisfactory compromise is achieved between the goodness of the fit to a given set of observational data, and some degree of smoothness (regularization) of the underlying particle model. We briefly describe the method together with a new regularization scheme in phase-space, which improves recovering the correct orbit structure in the models. We also mention some practical applications showing the power of the technique in investigating the dynamics of galaxies. ", "introduction": "In the last decades, an increasing amount of high quality photometric and kinematic data for galaxies have become available. The dynamical state of the observed galaxy, however, cannot be directly inferred from observations, due to projection effects, and modelling is essential to learn about the distribution of stellar orbits and the total (i.e. due to luminous and dark matter) gravitational potential in the observed galaxy. Therefore, several methods to model the observational data and create \"made-to-measure\" models have been devised. For instance, assuming that all the integrals of motion are known, one can fit observations with parametrized distribution functions \\citep{dejonghe84,dejonghe86,bishop87,gerhard91,hunterdezeeuw92,carollo95,kuijken95, mago95, merritt96,dehger93, matger99}, or solve the Jeans equations subject to the observational constraints \\citep{binmam82,binney90,magobin94,lokas02,cappellari08}. Another technique which is widely used is the Schwarzschild orbit superposition method \\citep{sch79, sch93}: a large library of orbits is computed in a fixed potential, and then the weights of individual orbits are adjusted until the model matches the photometry and kinematics of the target galaxy \\citep{richstone85, cappellari02, gebhardt03, Valluri04, thomas05, vandenbosch10}. \\citet{st96} proposed an alternative, particle-based method. A modified version suitable for modelling observational data with errors was designed by \\citet{dl07} and implemented in the parallel code NMAGIC. More recent implementations of this method can be found in \\citet{dehnen09} and \\citet{longmao10}. The basic idea is to evolve a system of particles while slowly correcting the individual weights of particles until the $N$-body system reproduces the observational data. The method is very powerful, since no orbit library needs to be computed or stored, no symmetry restrictions need to be imposed, and the potential can be evolved self-consistently from the particles. Moreover, the algorithm properly accounts for observational errors, and a great variety of observational data can be used simultaneously in the weight adaptation, including photometry, long-slit spectroscopic data, integral field absorption-line kinematics, and PNe velocities. So far, the particle made-to-measure method has been used to investigate the dynamics of the Milky Way's bulge and disk \\citep{bissantz04}, and the dark matter fraction and orbital structure in the outer halos of elliptical galaxies \\citep{dl08,dl09,das11}. ", "conclusions": "The parallel code NMAGIC is a powerful tool to build made-to-measure galaxy models, which reproduce a wide variety of observational data. NMAGIC works by adapting the weights of an $N$-body particle system until the target observables are well matched, subject to additional regularization constraints. We have briefly described a new, improved regularization scheme in phase-space, and shown how well the intrinsic properties of the target galaxy can be recovered independently of initial conditions, and how NMAGIC models can help us in learning about the orbital structure and gravitational potentials in galaxies." }, "1207/1207.6129_arXiv.txt": { "abstract": "{ We study the directional effect of the expected axion dark matter signal in a resonant cavity of an axion haloscope detector, for cavity geometries not satisfying the condition that the axion de Broglie wavelength $\\lambda_a$ is sufficiently larger than the cavity dimensions $L$ for a fully coherent conversion, i.e. $\\lambda_a \\gtrsim 2\\pi L$. We focus on long thin cavities immersed in dipole magnets and find, for appropriately chosen cavity lengths, an O(1) modulation of the signal with the cavity orientation with respect the momentum distribution of the relic axion background predicted by the isothermal sphere model for the galactic dark matter halo. This effect can be exploited to design directional axion dark matter detectors, providing an unmistakable signature of the extraterrestrial origin of a possible positive detection. Moreover, the precise shape of the modulation may give information of the galactic halo distribution and, for specific halo models, give extra sensitivity for higher axion masses. } ", "introduction": "Axions are light pseudoscalar particles that arise in theories in which the Peccei-Quinn U(1) symmetry solves the strong CP problem \\cite{Peccei:1977hh}. They are produced in the early universe as coherent field oscillations, the so-called misalignment mechanism~\\cite{Sikivie:2006ni,Wantz:2009it}. If the PQ symmetry is restored by reheating after inflation, axion strings and domain walls form and decay, providing an additional source of non relativistic axions. By these means, axions could provide all or part of the cold dark matter of the Universe. Together with WIMPs, axions are the most attractive solution to the dark matter problem. For these relic axions to account for the right amount of cold dark matter needed by current cosmological models, the axion mass need to be in the range $10^{-6} - 10^{-3}$ eV, the ``classic axion window''~\\cite{Wantz:2009it}. Much lower $\\sim$neV masses are still possible in fine-tuned models, the so-called ``anthropic axion window'' \\cite{Linde:1987bx,Hertzberg:2008wr}. QCD axions with masses above the classic window are still possible, but not as a solution to the dark matter problem (unless non-standard cosmological scenarios are invoked~\\cite{Visinelli2010}). Masses above 20 meV or so become in tension with current supernova modeling. Masses above $\\sim$1 eV are ruled out by cosmology (too high thermal production of axions~\\cite{Hannestad:2010yi}) and by astrophysical observations~\\cite{Raffelt:1999tx}. Masses of a few meV could also explain the apparent anomalous energy loss of white dwarfs~\\cite{Isern:2010wz,Isern1992,Isern:2008nt,Isern:2008fs}. Axion-like particles (ALPs) or more generic weakly interacting sub-eV particles (WISPs), appearing in a number of extensions of the standard model (SM)~\\cite{Jaeckel2010}, like for example string theory~\\cite{Cicoli:2012sz}, may share some of the properties of the axions. Most notably, under some circumstances, they could also be a dark matter candidate~\\cite{Arias:2012mb}. Some of the detection mechanisms devised for axions are applicable to other ALPs. Most axion detection strategies invoke the a-$\\gamma$-$\\gamma$ interaction, present in every axion model, and that gives rise to axion-photon oscillation inside magnetic fields. Axion helioscopes~\\cite{Sikivie:1983ip} look for the large axion flux emitted by our Sun with keV energies, using a powerful magnet pointing to the Sun. The most powerful helioscope built, CAST~\\cite{Zioutas:2004hi,Andriamonje:2007ew,Arik:2008mq,Aune:2011rx}, is currently sensitive to QCD axion models at the $\\sim$ eV scale. The future IAXO~\\cite{Irastorza:2011gs} will have a sensitivity to scan a large fraction of the models in the high mass range $10^{-3}-1$ eV. ALPs can also be searched for purely in the laboratory~\\cite{Ehret:2010mh}, but with insufficient sensitivity to reach axion models. In the classic window, and under the assumption that axions are the dominant component of dark matter, relic axions can be directly detected using Sikivie's haloscope technique~\\cite{Sikivie:1983ip}. Due to the relic axions being non-relativistic, axion conversion gives photons with energies equal to the axion mass, in the microwave regime. If the conversion happens in a microwave cavity resonant to the axion mass, due to the low velocity dispersion of the axions, the conversion is substantially enhanced, and a sufficiently high sensitivity can be obtained to explore realistic QCD axion models. This technique has been used already by a number of experiments, being ADMX the most powerful haloscope to-date \\cite{Asztalos:2001tf,Asztalos:2003px}, with sensitivity to QCD axions of $\\mu$eV mass. According to the standard formalism of the haloscope technique~\\cite{Sikivie:1983ip}, the power delivered into a cavity immersed in a magnetic field due to conversion from relic axions is: \\begin{equation} P_0 = \\gagamma^2 V B^2 C \\frac{\\rho_a}{m_a} Q \\label{pout} \\end{equation} \\noindent where $V$ is the volume of the cavity, $B$ its magnetic field strength, $Q$ its the loaded quality factor of the cavity (that we have assumed that it is lower than the relative energy spread of the axion energy $Q_a \\sim 10^6$), and $C$ is a geometry factor involving the precise electric field of relevant resonant mode in the cavity $\\textbf{E}_{\\textrm{cav}}(\\textbf{x})$ and the magnetic field $\\textbf{B}(\\textbf{x})$: \\begin{equation} C= \\frac{\\left(\\int dV \\textbf{E}_{\\textrm{cav}}(\\textbf{x}) \\textbf{B}(\\textbf{x})\\right)^2}{V|\\textbf{B}|^2 \\int dV \\epsilon(\\textbf{x}) \\textbf{E}^2_{\\textrm{cav}}(\\textbf{x})} \\label{cfactor} \\end{equation} The previous equations are valid under the basic assumption that the typical de Broglie wavelength of the relic axions $\\lambda_a$ is longer that the characteristic size of the cavity $d$. \\begin{equation}\\label{condition} \\lambda_a \\gtrsim d \\end{equation} Only in this case the resonant conversion giving rise to (\\ref{pout}) takes place. The axions composing the dark matter halo have a velocity distribution given by the specific halo model but in general with typical values of $\\sim$ 300 km/s. So approximately the de Broglie wavelength of the relic axions depends on the axion mass in the following way: \\begin{equation}\\label{debroglie} \\lambda_a = \\frac{2\\pi}{p_a} \\sim 12.4 \\textrm{ m} \\left(\\frac{10^{-4}\\textrm{ eV}}{m_a}\\right) \\left(\\frac{300 \\textrm{ km/s}}{v_a}\\right) \\end{equation} As can be seen from (\\ref{debroglie}), for axion masses well below 10$^{-4}$ eV, any magnet geometry in the $\\sim$m scale is well within the condition (\\ref{condition}). ADMX is indeed using a cylindrical cavity of 1 m length and 0.6 m diameter inside a solenoidal magnet, tunable to axion masses at the few $\\mu$eV scale, and enjoys sensitivity sufficient to exclude one of the two main axion benchmark models (KSVZ axions)~\\cite{Kim:1979if, Shifman:1979if}. Improvements to enhance the sensitivity to lower $\\gagamma$ values are ongoing as well as R\\&D to build setups resonant at higher axion masses (above $10^{-5}$ eV)~\\cite{Asztalos2010,2012APS..APRD14001M}. For higher axion masses, the main challenge relies on the fact that the needed resonant cavity geometries are smaller and so is their corresponding sensitivity ($P$ being proportional to the volume of the cavity $V$). Recently the use of long thin cavities (waveguides) inside strong dipole magnets have been proposed as a possibility to achieve competitive sensitivity in the 10$^{-5}$-10$^{-4}$ eV range~\\cite{Baker:2011na,Caspers:2010zz}. The use of small cross-section, but long and powerful, dipole magnets like the ones used in particle accelerators (and already recycled for axion physics in, e.g., CAST~\\cite{Aune:2011rx} or ALPS~\\cite{Ehret:2010mh}) can accommodate cavities resonant at these higher frequencies (driven by the small dimension of the waveguide) while at the same time keeping large enough $V$ and $B$. Moreover, this option may give additional technical advantages with respect the conventional approach, e.g., regarding mode crossing or mode localization~\\cite{Baker:2011na}. The possibility of using few-meter long cavities for detection of few $\\times 10^{-5}$ eV relic axions bring us closer to the limitation expressed by~(\\ref{condition}). In this paper we explore in some detail the effect of this limitation on the predicted relic axion signal for realistic dark matter momentum distributions. In particular, we note that for thin and long geometries like the ones mentioned before, the orientation of the cavity with respect the main incoming axion direction affects the signal intensity. For carefully chosen lengths this effect can be maximized and used as a powerful identification signature of the extraterrestrial origin of an eventual positive detection. Moreover, like in the case of WIMP directional experiments~\\cite{Ahlen:2009ev}, the resulting signal modulation will show specific features of the galactic dark matter halo, enabling us to discriminate between different options currently considered. Finally, we want to stress that directionality in ALP detection (and emission) antennas were first discussed in~\\cite{Caspers:2010zz}, where concepts like phased arrays, and excitation of higher order or traveling wave modes were put forward. They were qualitatively discussed in the context of microwave-shinning-through-wall experiments sensitive to hidden photons and ALPs, for which directionality would give extra sensitivity, although of course also usable in relic axion searches. Here we focus ourselves on the relic axion case, and on the simple setup consisting on a long rectangular cavity in a constant magnetic field and resonant in a fundamental mode. As mentioned, already in such simple case directional effects appear when the condition~(\\ref{condition}) is not satisfied. We will compute the signal strength in that particular situation, study the directional effect and derive prescriptions for the length of the cavity needed to maximally exploit it. We will convolute the result with the momentum distribution of relic axions at the Earth, according to two different halo models, in order to assess its potential in realistic conditions and perform some preliminary sensitivity computation. Of course these concepts could be generalized to arbitrary cavity geometries, that could evolve into true relic axion antennas, however, the restricted scenario studied here is specially appealing because of the feasibility of constructing such setups in the near future. The article is structured as follows. In section \\ref{sec:signal} we discuss the basic concept for an idealized situation with a single axion incoming direction. In section \\ref{sec:halo} we present our numerical calculations for a realistic axion momentum distribution following the standard isothermal isotropic model. In section \\ref{sec:caustics} we present the calculations for a particular halo model, that of Sikivie's caustic rings. In section~\\ref{sec:sensitivity} we perform some sensitivity calculations. We finish with our conclusions in section \\ref{sec:conclusions}. ", "conclusions": "We have studied the expected axion dark matter signal in a resonant cavity of an axion haloscope detector, not satisfying the condition that the axion de Broglie wavelength is sufficiently larger than the cavity dimensions for a fully coherent conversion, i.e. $\\lambda_a \\gtrsim 2\\pi L$. In this case, and for cavity geometries largely non-spherical, the expected signal develops a dependency on the direction of the incoming axion or, equivalently, the orientation of the cavity with respect the distribution of relic axion momenta. This is the first time, to our knowledge, that a directional effect in axion dark matter detection is studied, effect that can be exploited to design directional axion dark matter detectors, conceptually equivalent (although technologically rather different) than the much sought WIMP dark matter directional detectors\\footnote{It is interesting to note that WIMP directionality is achieved by measuring the WIMP-induced nuclear recoil direction in a suitable particle detector. It is a completely different mechanism, whose directionality is not dependent on the size of the detector~\\cite{Ahlen:2009ev}.}. The directional signal of dark matter, proposed 25 years ago for the first time~\\cite{Spergel:1987kx} in the context of WIMP searches, is acknowledged to provide an unmistakable signature of the extraterrestrial origin of a possible positive detection. We have focused our study to simple long and thin cavity geometries immersed in powerful dipole magnets, like the ones recently proposed to achieve competitive sensitivity in the $10^5-10^4$ eV axion mass range. We found that indeed a O(1) modulation of the signal with respect to the cavity orientation is present if adequate cavity lengths are employed, even when a realistic axion momentum distribution is assumed using the isothermal sphere model for the galactic dark matter distribution. Other halo models will yield a different modulation signals, although we do not expect a difference in the qualitative conclusions here stated for halo models that suppose a smooth departure from the isothermal sphere. For more radically different models, like for example the caustic rings model, this effect is even more identificative and interesting. The presence of low dispersion flows in the dark matter distribution make the effect studied here valid even for high axion masses, pushing the range of validity well above the conventional limitation provided the cavity is oriented perpendicular to the incoming flow direction. In any case, this effect no only provides a clear signature of dark matter but would give extra information on the particular galactic distribution of dark matter, an invaluable information on the galaxy halo formation and evolution. The case studied involving long cavities in dipole magnets is particularly appealing because it could be realized in the near future given that this type of magnets, developed for accelerator technology, are already in use by the axion community in experiments looking for solar axions, like CAST~\\cite{Aune:2011rx}, or axions (or ALPs) produced in the laboratory, like ALPS~\\cite{Ehret:2010mh}. These results add motivation to the use of these setups also as haloscopes in search for relic axions. Although the movement of rotation of the Earth could allow a static magnet to scan the sky in search for the modulation studied here, we note that magnets used or foreseen for helioscopes are equipped with movable platforms (to orient the magnet to the Sun) that would also be of additional value in a directional relic axion campaign, providing flexibility to scan the sky more efficiently." }, "1207/1207.2781.txt": { "abstract": "We present observations of 36 late-M~dwarfs obtained with the KeckII/NIRSPEC in the J-band at a resolution of $\\sim$ 20,000. We have measured projected rotational velocities, absolute radial velocities, and pseudo-equivalent widths of atomic lines. 12 of our targets did not have previous measurements in the literature. For the other 24 targets, we confirm previously reported measurements. We find that 13 stars from our sample have $v$~sin~$i$ below our measurement threshold (12 km s$^{-1}$) whereas four of our targets are fast rotators ($v$~sin~$i$ $>$ 30 km s$^{-1}$). As fast rotation causes spectral features to be washed out, stars with low projected rotational velocities are sought for radial velocity surveys. At our intermediate spectral resolution we have confirmed theidentification of neutral atomic lines reported in \\citet{mclean07}. We also calculated pseudo-equivalent widths (p-EW) of 12 atomic lines. Our results confirm that the p-EW of K I lines are strongly dependent on spectral types. We observe that the p-EW of Fe I and Mn I lines remain fairly constant with later spectral type. We suggest that those lines are particularly suitable for deriving metallicities for late-M dwarfs. ", "introduction": "Radial velocity studies of early bright M dwarfs have yielded planets, including a planetary system \\citep{marcy99} and rocky planets \\citep{rivera05, udry07, mayor09}. Though M dwarfs present themselves as promising candidates for rocky planet searches in the habitable zone, the effort required to measure precise radial velocities are often thwarted by their higher projected rotational velocities ($v$~sin~$i$ $>$ 30 km s$^{-1}$) and stellar activities. Radial velocity measurement precision is limited by stellar rotation. An increase in stellar rotation causes narrow deep lines to become broad, shallow, and blended. Such lines reduce available radial velocity information, thereby reduce the precision. \\citet{reiners10b} find that approximately 50\\% of their sample of 63 M dwarfs with spectral types M7-M9.5 show projected rotational velocities greater than 10 km s$^{-1}$. A comprehensive study of projected rotational velocities of early- to mid-M dwarfs suggests an increasing trend in projected rotational velocity with later spectral type \\citep{jenkins09}. The relation between stellar activity and projected rotational velocity is well established in literature \\citep{noyes84, delfosse98a, pizzolato03}. This relation also holds true for stars at the end of the main sequence \\citep{mohanty03}. Stellar activity produces stellar spots. These spots distort the line profiles of the stellar absorption lines that are critical for radial velocity measurement. This distortion leads to a change in the bisectors of the absorption lines. Such lines can impersonate an unseen companion thus resulting in a false detection \\citep{henry02}. An extensive simulation work on stellar spots by \\cite{reiners10a} suggests that spots can cause radial velocity shifts of 100 m s$^{-1}$. A comprehensive study of stellar activity of M dwarfs (West et al. 2004) indicates that the fraction of active stars increases from early-M dwarfs (10$\\%$) to late-M dwarfs (75$\\%$). Hence, late-M dwarfs are likely to be fast rotating and active. With the upcoming infrared radial velocity surveys dedicated to search for planets around M dwarfs, such as HPF \\citep{mahadevan10} and CARMENES \\citep{quirrenbach10}, precise measurements of stellar rotation among late-M~dwarfs becomes crucial. Measurement of radial and projected rotational velocities requires a good understanding of lines in the stellar atmosphere. Over the last seven years, identification and characterization of neutral atomic lines have been done at low (R $\\sim$ 2000) \\citep{cushing05} and intermediate resolutions (R $\\sim$ 20,000) \\citep{mclean07}. However, both these studies used a small sample of M dwarfs and calculated pseudo-equivalent width (p-EW) of a few atomic lines. With 36 late M~dwarfs (M5.0 - M9.5), we have more than doubled the number of objects for which p-EW of 12 neutral atomic lines have been measured. In this paper, we verify the identification of neutral atomic lines using the Vienna Atomic Line Database \\citep{kupka02}. We calculate rotational and absolute radial velocities (in a companion paper, we reported relative radial velocities for stars with multiple observations \\citet{rodler12}), and measure p-EWs of 12 neutral atomic lines. The paper is organized as follows: In $\\S$2, we describe our sample, compare it with existing data, provide our instrumental setup for observations, and data reduction procedure. In $\\S$3, we list the results of our observations that we analyze and discuss in $\\S$4. In $\\S$5, we summarize our work and provide conclusions. ", "conclusions": "\\subsection{Projected Rotational and Absolute Radial Velocities} We have identified 13 targets in our sample with $v$~sin~$i$ below our detection threshold ($\\sim$ 12 km s$^{-1}$), while four stars (LP44-1627, LP349-25, 2MJ0004-1709, and 2MASSJ1835+3259) show velocities greater than 30 km s$^{-1}$. Figure 4 (right panel) compares $v$~sin~$i$ values of our sample with those in literature (\\citet{stauffer86, marcy92, delfosse98b, gizis02, mohanty03, bailer04, fuhrmeister04, jones05, reiners07, west08, reiners10b, jenkins09}). Our results indicate low residuals for stars with projected rotational velocities below 30 km~ s$^{-1}$. Though, the three fast rotating stars show larger velocity dispersion their measurements are within 3-$\\sigma$ of the literature values. The measurement of rotational velocity of the two components of LP349-25 using NIRPSEC and LGS AO \\citep{konopacky12} in conjunction suggests that LP349-25A (M8.0) has a velocity of 55 $\\pm$ 2 km s$^{-1}$ while the fainter LP349-25B has a velocity of 83 $\\pm$ 3 km s$^{-1}$. Our measurement for the combined spectra is closer to the primary than the secondary. As the secondary is faint, the projected rotational velocity is dominated by the primary. This is particularly evident in works of \\citet{reiners10a} where, using R $\\sim$ 31,000, their measured projected rotational velocity is also closer to that of the primary. The same scenario is observed for the case of 2M2206-2047, where the projected rotational velocities of the two components, A and B are 19 $\\pm$ 2.0 km~ s$^{-1}$ and 21 $\\pm$ 2.0 km~ s$^{-1}$ , respectively. We find the velocity of the combined spectrum to be 22.2 $\\pm$ 2.0 km~ s$^{-1}$. However, as both components are M8.0, their contribution is similar. \\subsection{Neutral Atomic Lines in M dwarfs} Figure 2 plots a sample of our reduced spectra by NIRSPEC echelle orders with order 65 in the upper left panel and order 58 in the lower right panel. Atomic absorption lines are identified by dashed lines. Spectral types are listed on the right side of each plot. These spectra are similar to those shown in \\citet{mclean07}, however with larger sample size each panel shows a continuous transition from M5.0 to M9.5 in steps of half spectral type. Tables 4 - 6 list the p-EW of 12 neutral lines. The K~I doublet at 11690 \\AA$ \\ $and the K~I triplet 11771 \\AA$ \\ $in order 65 (Figure 2, top left panel) show increasing line width with later spectral type. By M9.5 the width of the K~I triplet increases to the extent that it blends into a single line. The widening of lines is induced by collisional broadening with H$_2$ molecules in the atmospheres of M~dwarfs \\citep{mclean07}. We also find that the K~I lines in order 61 have a similar trend like those in order 65. These results shown in Figure 5 agree with the trend found by \\citet{mclean07}, and \\citet{cushing05}. A weak Fe~I line in order 65 is observed at longer wavelengths. This line weakens with later spectral type and by M9.5 it is almost indiscernible. The modeled spectra of late M~dwarfs by \\citet{lyubchik04} and \\citet{rice10}, indicate a presence of weaker Fe~I and Ti~I lines in order 65. We identified their positions, but found them to be heavily blended and hence are not clearly distinguishable in our spectra. Better spectral resolution is required in order to distinguish these lines. The Fe~I lines (doublet 11886 \\AA$ \\ $, 11887 \\AA$ \\ $ and a singlet 11976 \\AA$ \\ $) dominate order 64 (Figure 2, top right panel). At M5.0, Fe~I lines are seen as narrow sharp absorption lines. However, with later spectral type, the Fe~I line at shorter wavelength broadens and blends with surrounding lines. Our results agree with those in \\citet{cushing05} within the errors. A quantitative comparison is discussed in the next section. Order 64 also contains the Mg I line at 11828 \\AA$ \\ $and Ti~I line at 11893 \\AA. Both lines are weak and weaken with later spectral type. The Ti~I line disappears around M9 while the Mg I line becomes indistinguishable around M8.5. The Mn~I line (12899 \\AA) in order 59 (Figure 2, bottom left panel) maintains its strength with later spectral type while the two Ti~I lines are observed to weaken with later spectral types. As seen in the figure, the Mn~I line maintains its sharp and narrow feature with later spectral type. The two Al~I lines (13127 \\AA$ \\ $, 13150 \\AA) in order 58 (Figure 2, bottom right panel) weaken with later spectral types. The Na I line at 12682 \\AA$ \\ $ (Figure 2, middle right panel) also shows a decline in strength with later spectral type. By M8.5 the Na I is completely indistinguishable. \\subsection{Pseudo-Equivalent Widths} As a trend of increasing width with later spectral type is qualitatively observed in the neutral atomic lines in Figure 2, a quantitative comparison is desired. We calculate the p-EWs of neutral atomic lines shown in Figure 2. Figures 5-7 plot the computed p-EWs for our targets with respect to the spectral types. Comparing our data (open circles) with those from \\citet{cushing05} in figures 5 and 6, we see that Cushing and colleagues' values (filled diamonds) are within our expected errors. For clarity, in figures 5 and 6, \\citet{cushing05} values have been shifted to the right by 0.1 spectral types. However, p-EW values by \\citet{mclean07} for two late M dwarfs do not match with ours or with those of \\citet{cushing05}. These values are not plotted as they are well beyond the range of the plot. In each panel of Figures~5-7, a linear regression fit to the data is shown as a short-dash line. The linear relation including its y-offset and slope are listed in Table 7 for the 12 neutral atomic lines. These relations quantitatively illustrate the variation of each atomic line with respect to spectral type. The top two panels of Figure 5 show an increasing trend of p-EWs with later spectral type for K~I lines. Figure~5 compares p-EW of our sample with those from \\citep{cushing05}. We find that our p-EW measurements agree with theirs. The p-EW calculations of Fe~I lines in figure 6 top panel suggests small variation with spectral type. However, the Al~I line shows a decreasing trend. Mn~I line along with two Ti~I lines, are observed in Order 59. As seen in Figure 2, the Mn~I line maintains its narrow sharp feature with later spectral type. This line has been studied in greater detail by \\citet{lyubchik07}. They model the Mn~I line in five ultracool dwarfs ranging from M6 to L0. Using the WITA6 program \\citep{pavlenko00} for the NextGen model structure they have been able to fit the Mn~I line, and their p-EW calculations carried out using the DECH20 package \\citep{galazutdinov92} on three dwarfs between M6 and M9 indicate a variation of 2 m\\AA. Our results confirm that the Mn~I line is not sensitive to variation in spectral type." }, "1207/1207.0539_arXiv.txt": { "abstract": "Hydrogen fluoride has been established to be an excellent tracer of molecular hydrogen in diffuse clouds. In denser environments, however, the HF abundance has been shown to be approximately two orders of magnitude lower. We present Herschel/HIFI observations of HF~$J=1-0$ toward two high-mass star formation sites, NGC6334~I and AFGL~2591. In NGC6334~I the HF line is seen in absorption in foreground clouds and the source itself, while in AFGL~2591 HF is partially in emission. We find an HF abundance with respect to H$_2$ of $1.5\\cdot 10^{-8}$ in the diffuse foreground clouds, whereas in the denser parts of NGC6334~I, we derive a lower limit on the HF abundance of $5\\cdot 10^{-10}$. Lower HF abundances in dense clouds are most likely caused by freeze out of HF molecules onto dust grains in high-density gas. In AFGL 2591, the view of the hot core is obstructed by absorption in the massive outflow, in which HF is also very abundant ($3.6\\cdot 10^{-8}$) due to the desorption by sputtering. These observations provide further evidence that the chemistry of interstellar fluorine is controlled by freeze out onto gas grains. ", "introduction": "Hydrogen fluoride (HF) is an exceptional molecule, because of its peculiar chemistry. Fluorine is one of the few atoms which has a greater affinity to hydrogen than hydrogen itself, and thus the reaction $$ H_2 + F \\rightarrow HF+H $$ is highly exothermic ($\\rm \\Delta E\\approx 16000$~K, although the activation of this reaction energy is 400~K and quantum mechanical tunneling through this barrier has to be considered). The destruction of HF occurs by the relatively slow photo-dissociation ($\\rm1.17\\cdot10^{-10}~s^{-1}$) and reactions with ions, such as He$^+$, H$_3^+$, and C$^+$, which are rare compared to H$_2$. Thus HF formation is much more efficient than its destruction under most interstellar conditions. As a consequence HF is considered to be the main reservoir of fluorine in molecular regions, i.e, where $\\rm H_2/H>1$, and therefore it is thought to be an ideal tracer of H$_2$. Astrochemical models predict an abundance of HF relative to H$_2$ (all fractional abundances given in this paper are relative to H$_2$) of $3.6\\cdot10^{-8}$ in such clouds (\\citealt{NWS05}). The first detection of HF was reported by \\cite{NZS97}, who observed the HF~$J=2-1$ transition ($\\nu=2463.43$~GHz) with relatively low resolution (R=9600) in absorption towards Sgr~B2, using the {\\it{Infrared Space Observatory (ISO)}}. However, because high densities or strong radiation fields are needed to populate the HF~$J=1$ level, the HF column density derived from this observation is subject to large uncertainties. With the advent of the {\\it{HIFI}} instrument \\citep{dHP10} aboard the \\emph{Herschel Space Observatory} \\citep{PRP10}, high resolution spectroscopy ($\\rm R>10^6$) of the ground state rotational transition of HF has become possible for the first time. Several HF observations in diffuse cloud along the lines of sight toward strong sub-millimeter dust continuum sources, such as W31C (\\citealt{NSP10}), W49N and W51 (\\citealt{SNP10}), and Sgr~B2(M) (\\citealt{MEP11}) have been reported. HF~$J=1-0$ is detected in absorption toward all of these continuum sources, which was expected considering the extreme conditions necessary to populate the excited rotational states of this molecule. These observations prove that HF is an excellent tracer of H$_2$ in diffuse clouds, although the abundances of HF relative to H$_2$, derived by comparison with CH, are $\\sim 1.5\\cdot10^{-8}$, with an uncertainty of 46\\%, and thus a factor of two below the chemical model predictions. Furthermore, HF turned out to be a much more sensitive tracer of H$_2$ than the commonly used CO, which is in agreement with the model prediction that HF is more abundant than CO in diffuse clouds with low extinction (\\citealt{NWS05}). \\cite{SNP10}, for example, detected a diffuse cloud on the line of sight toward W51, not seen in CO. The only HF abundance calculated in a high-mass star-forming region, and not in a diffuse foreground cloud, has been reported by \\cite{PBL10}, who detected HF~$J=1-0$ in absorption towards Orion~KL. Orion~KL is a peculiar source, heated from the front, which shows hardly any lines in absorption at submillimeter wavelengths. The HF abundance derived by Phillips et al. is a lower limit of $1.6\\cdot10^{-10}$, approximately two orders of magnitude lower than the HF abundances reported in diffuse clouds. This result can be either explained by geometrical effects, due to the incomplete coverage of the continuum source by absorbing material, or, under the assumption that most fluorine is bound in HF, is caused by HF molecules sticking onto the surfaces of dust grains under high-density conditions. The only three detections of HF in emission so far are in the immediate vicinity of the AGB-star IRC+10216 \\citep{ACW11}, the Orion Bar \\citep{vON12} and towards the Seyfert 1 Galaxy Mrk~231 \\citep{vIM10}. To investigate the abundances of HF in dense regions in more detail, we analyze here the HF~$J=1-0$ spectra towards two high-mass star-forming regions, NGC6334~I and AFGL~2591. NGC6334 is a relatively nearby (1.7 kpc; \\citealt{N78}) high-mass star-forming region, which harbors sites of many stages of protostellar evolution (\\citealt{SH89}). Single-dish continuum observations at sub-millimeter wavelength show that a total mass of 200~$M_\\odot$ is associated with NGC6334 I (\\citealt{S00}). The molecular hot core NGC6334~I, studied extensively over the last decades (e.g. \\citealt{BWT08, BWT07}; \\citealt{HBM06}), is associated with an ultra compact HII region (\\citealt{dRD95}) and shows a very line-rich spectrum (\\citealt{SCT06}, \\citealt{TWM03}). {\\it{HIFI}} observations in NGC6334~I have led to the detection of CH \\citep{vvL10}, H$_2$O \\citep{ELB10}, H$_2$O$^+$ (\\citealt{OML10}), and H$_2$Cl$^+$ (\\citealt{LPN10}), showing that hydrides are common toward this source. Furthermore, a bipolar outflow (\\citealt{LSP06}; \\citealt{BWT08}) and H$_2$O, OH, CH$_3$OH class~II, and NH$_3$ masers have been detected (e.g., \\citealt{KJ95}; \\citealt{EvM96}; \\citealt{WLT07}). \\cite{RSW11} studied the radial structure of several high-mass star-forming cores, including NGC6334~I. This investigation, mainly based on HCN observations with the APEX telescope, revealed a density law of $n\\propto r^{-1.5}$. The temperature also follows a power law, with an index of 0.83 and 0.4 in the inner and the outer part, respectively. Interferometric data (SMA), however, revealed that the internal structure of NGC6334~I is much more complex. The hot core consists of four compact condensations within a 10$''$ region, emitting about 50\\% of the continuum flux, whereas the other 50\\% stem from the extended envelope \\citep{HBM06}. AFGL~2591 is an isolated high-mass star-forming region located within the Cygnus-X region. The distance toward AFGL~2591 was reported to be between 0.5~kpc and 2.0~kpc, and for modeling purposes a distance of 1~kpc was assumed (e.g., \\citealt{vvE99}). However, new measurement base on VLBI parallax measurements of H$_2$ masers suggest that the distance might be as large as 3.3~kpc \\citep{RBS12}. Based on a distance of 1~kpc, the luminosity of AFGL~2591 is $\\rm\\sim2\\cdot10^4~L_{\\odot}$, however, adopting the newly measured distances the luminosity would be approximately a factor ten higher. The embedded central star, which is obscured in the visible by the massive envelope, has an estimated mass of 16~M$_\\odot$ and an effective temperature of 33,000~K \\citep{vM05}. Besides the circumstellar envelope a massive outflow with an outflow velocity of $\\sim 20$~km\\,s$^{-1}$ with respect to the systemic velocity has been detected \\citep{LTS84}. In addition, \\cite{HM95} found a second, faster ($\\rm v_{out}\\approx 40 km\\,s^{-1}$), but weaker outflow. Many attempts to model the structure of AFGL~2591 have been carried out assuming spherical symmetry and a power-law density and temperature structure (e.g., \\citealt{vvE99}; \\citealt{vvE00}; \\citealt{Dvv02}; \\citealt{dHF09}). A cavity, caused by the outflow might by present (\\citealt{vvE99}). \\cite{vvS11} mapped AFGL~2591 in the frequency range from 330--373~GHz with the JCMT. In this frequency range they found $\\sim 160$ lines, of which 35 show extended emission ($>15''$). These observations reveal a small scale structure ($\\leq 10^4$~AU), and a line of sight velocity gradient is apparent in most molecules. Their modeling suggests that a non isotropic structure or a velocity gradient must be present on $\\sim10^4$~AU scales in order to explain the observations. ", "conclusions": "\\label{DIS} We determine the abundance of HF in very different physical environments, and obtain very different results. In agreement with previous studies, we find the HF abundance of $\\sim 1.5\\cdot10^{-8}$ in diffuse clouds, implying that about half of the available fluorine is in the form of HF in the gas phase. An exception is the +8.0~km\\,s$^{-1}$ component toward NGC6334~I, in which HF is less than half as abundant as in other diffuse clouds. The chemical composition of this foreground component, however, seems to be different from the other components (e.g., strongly enhanced CH$^+$ abundance; Lis et al.~in prep.), which requires further studies. In the high-mass star-forming regions, at much higher densities compared with diffuse clouds, the HF abundance seems to drop by about two orders of magnitude. The correlation with the physical conditions, especially density and temperature, would be crucial to determine the nature of this HF depletion, but it still remains to be found. \\cite{PBL10} argues that the low abundance of HF at higher densities is due to freeze out of HF onto dust grains. This argument is drawn from the fact that almost all fluorine should be bound in HF due to its high proton affinity, and that the desorption energy of HF is assumed to be quite large due to its polar nature. In dense ($\\rm 10^{5}-10^6~cm^{-3}$) and warm (100--150~K), but dynamically active regions such as the outflow in AFGL~2591, we find HF abundances close to the values found in diffuse clouds. In the picture of HF freeze out at high densities this can be explained by HF desorption due to sputtering by high velocity particles, similar to the desorption of water in outflows (\\citealt{KVv10}; \\citealt{ELB10}). In both sources HF is the only fluorine bearing species we find in the HIFI line survey, and we can confirm non-detections for CF, CF$^+$, DF, and H$_2$F$^+$. In NGC6334~I the upper limit for CF$^+~J=5-4$, a molecule which has been detected in the Orion Bar with an abundance of a few times $10^{-10}$ \\citep{NSM06}, is 0.012~K. Assuming cloud properties, as determined by C$^{18}$O fits (Plume priv. communication), yields an upper limit of the CF$^+$ abundance of $6\\cdot 10^{-12}$. This low upper limit, 50 times lower than the abundance found in the Orion Bar, can be explained by the different nature of the two objects. C$^+$, a precursor of CF$^+$, and thus CF$^+$ itself, are expected to be much more abundant in photo-dominated regions, such as the Orion Bar, than in massive cores. The upper limit of the abundance of these four fluorine bearing species together is $\\sim 7\\cdot 10^{-11}$, and thus these molecules hold less than 0.4\\% of the total fluorine. The non-detection of other fluorine bearing species is another indication that freeze out onto dust grains causes the low HF abundances observed in dense, quiescent gas. This analysis demonstrates that the utility of HF as a tracer of H$_2$, as predicted by chemical models, is only applicable to diffuse clouds. The fluorine chemistry in dense clouds is not yet fully understood and further theoretical and observational studies are required." }, "1207/1207.0822_arXiv.txt": { "abstract": "We study the formation of photospheric emission lines in O stars and show that the rectangular profiles, sometimes double peaked, that are observed for some stars are a direct consequence of rotation, and it is unnecessary to invoke an enhanced density structure in the equatorial regions. Emission lines, such as \\niv\\ \\lb 4058 and the \\niii\\ \\lb\\lb4634-4640-4642 multiplet, exhibit non-standard ``limb darkening'' laws. The lines can be in absorption for rays striking the center of the star and in emission for rays near the limb. Weak features in the flux spectrum do not necessarily indicate an intrinsically weak feature -- instead the feature can be weak because of cancellation between absorption in ``core'' rays and emission from rays near the limb. Rotation also modifies line profiles of wind diagnostics such as \\heii\\ \\lb4686 and H$\\alpha$ and should not be neglected when inferring the actual stratification, level and nature of wind structures. ", "introduction": "Many O stars exhibit rapid rotation that broadens their line profiles. In the absence of a wind, the influence of rotation can be accurately computed by integrating the intensities from the stellar disk with allowance for the projected velocity in the direction of the observer. For convenience, and speed, it is customary to allow for the effects of rotation by convolving the flux spectrum with a rotational broadening function that accounts for limb darkening. This procedure, which assumes that the line profile does not change across the disk (i.e., the continuum and line have the same limb darkening laws) \\citep[e.g.,][]{Gray92_book}, generally works very well. As our own tests have shown, the simple convolution procedure is sufficiently accurate for the majority of photospheric absorption profiles in O stars, although for the most precise work the disk integration procedure should be used. In general, O stars exhibit a wind which modifies the photospheric H$\\alpha$ profile (possibly driving H$\\alpha$ into emission) and generates numerous UV P~Cygni profiles. Because the wind is extended, and because the rotation rate declines with radius, the simple convolution technique is not valid. In such a case it is necessary to take the 2D structure of the wind into account. In the simplest case one can assume that rotation simply affects the velocity structure of the wind while maintaining an almost spherical structure, while for rapid rotation it is expected that the density structure will also be altered. How the 2D density structure is altered is very complex, and depends on the rotation rate relative to the critical rotation rate, the rotation rate relative to the terminal wind velocity, and the closeness of the star to the Eddington limit. Depending on their values, it is possible to get either density enhancements in the equatorial or polar flows \\citep{MM00_Edd_lim}. A prolate wind structure is expected unless the number of lines (or more correctly the flux [line] mean opacity) increases sufficiently rapidly towards the equator to overcome the decreasing equatorial flux arising from gravity darkening \\cite[e.g.,][]{MM00_Edd_lim}. In many O stars intense \\niii\\ \\lb4634--4642 emission is seen and is used as a criterion to define the so-called ``f\" class \\citep{Wal71_Of, SMW11_class}. When resolved, these lines can show a rectangular (or trapezoidal-like) profile\\footnote{For simplicity we refer to these profiles as non-Gaussian, while other profiles will be referred to as Gaussian, even if they do not exhibit extended wings.}, with the stronger components exhibiting a double peaked structure. In the spectra of O supergiants presented by \\cite{BHL12_Osg}, \\zpup, HD~16691, and HD~210839 show broad, non-gaussian profiles for the \\niii\\ multiplet, although HD~15570, HD~163758 and HD~192639 exhibit more gaussian-like profiles. The mechanisms driving the \\niii\\ lines into emission have been discussed previously by \\cite{RPN11_NIII}. The dominant mechanism depends on the precise stellar parameters -- both continuum fluorescence (e.g., the Swings mechanism, \\citealt{BM71_NIII}) and dielectronic recombination (previously discussed for these lines by \\citealt{MHC72_NIII}) can be important. \\cite{RPN11_NIII} also indicate that the strength of the lines can be influenced by interactions between the \\niii\\ and \\oiii\\ resonance lines. In this paper we discuss the profiles of photospheric emission lines in O stars, and show that rotation can produce rectangular-like profiles that often exhibit double-peaked profiles. The phenomena is not restricted to the \\niii\\ \\lb4634--4642 complex --- in \\zpup \\niv\\ \\lb 4058 and Si\\,{\\sc iv} \\lb\\lb 4089,4116 also exhibit non-gaussian profiles. We show that the observed line profiles are a direct consequence of the lines being photospheric and being in emission. The lines do not show ``classic'' limb darkening -- rather the intensity distribution across the star can be flat, and may even show limb brightening. In addition, we further examine the influence of rotation on the H$\\alpha$ and He\\,{\\sc ii} 4686 line profiles. Rotation produces observable effects on the line profiles, and these effects are seen. Because of the influence of rotation on the wind velocity law (which is now 2D) wind line profiles potentially depend on both $v$ and $\\sin i$ rather than $v \\sin i$. To model these emission lines on O stars, it is absolutely essential to properly account for the effects of rotation in O stars that exhibit rapid rotation. It is especially true for rapid rotators such as $\\zeta$ Pup which has $v \\sin i \\sim 210 \\,\\kms$. ", "conclusions": "We have demonstrated that rotation alone can explain the rectangular profiles, sometimes double-peaked, that are observed for some optical emission lines in O stars. It is not necessary to invoke departures of the density structure from spherical symmetry to explain these profiles. To properly compute the emission line profiles, it is necessary to integrate the intensities across the rotating stellar disk -- the classic convolution of the flux does not work because the line profiles vary substantially across the stellar disc (i.e., with impact parameter). In some cases the line is in absorption for rays striking the center of the star, while it is strongly in emission for rays near the limb. For \\niv\\ \\lb 4058 in \\zpup, the observed weak double peaked profile is actually a result of near perfect cancellation between absorption at low impact parameters, and emission at high impact parameters. This illustrates that the variation in intensity with impact parameter can provide additional insights into line formation. From their study of the \\niii\\ lines in HD 16691, \\cite{DRL09_prof} concluded that the lines arise from close to the star in a large scale corotating structure. While we confirm that rotation is crucial in explaining the observed line shapes, we find that the \\niii\\ lines are primely of photospheric origin in these supergiants and that rotation alone can explain their observed shapes. For most analyses it is adequate to use the normal convolution technique when modeling absorption lines. However, there are subtle differences between disk profiles, and the profiles obtained with convolution, and these could be important when using absorption lines to study line asymmetries and the possible influence of the velocity field at the base of the wind. Following a suggestion by Francisco Najarro, and due to the intense interest with the mass of the most massive stars, we also ran a model for R136a3, which was found by \\cite{CSH10_R136} to have an initial mass of $\\sim 165\\,$\\Msun\\ and a $v \\sin i \\sim 200\\,$\\kms. Profiles computed using a 2D model to take into account rotation showed only small differences from profiles computed using the convolution technique, and are too small to significantly affect conclusions drawn by \\cite{CSH10_R136}. Although already well known, we also highlighted the influence rotation has on main optical wind diagnostics -- H$\\alpha$ and \\heii\\ \\lb 4686. While the correct allowance for rotation does not fundamentally alter the line strength (and hence, for example, inferred mass-loss rates) it does significantly alter the line profile. This is of crucial importance -- it means that H$\\alpha$ and \\heii\\ \\lb 4686 emission line profiles in rapidly rotating O stars cannot be used to infer variations in clumping with distance and the importance of velocity-porosity unless rotation is accurately taken into account. Further, rotation will influence the accuracy of parameters, such as $\\beta$ (in the classic wind-law $v(r)=(1-r/\\Rref)^\\beta$). The correct treatment of rotation is also important for UV resonance lines formed in the wind. As photospheric lines can influence the resonance line profiles, it is important to take rotation into account. However, a simple convolution overestimates the effect of rotation on the wind lines and can adversely affect estimates of the terminal velocity, and microturbulence within the wind. Since it is relatively easy to treat the gross influence of rotation accurately, this should become the norm in future studies of O stars. Issues related to the influence of rotation on the azimuthal density structure, and the dependence of clumping parameters on rotation, will require further studies. \\label{Sec_future}" }, "1207/1207.5894_arXiv.txt": { "abstract": "The solar corona and heliosphere are visible via sunlight that is Thomson-scattered off of free electrons, yielding a radiance against the celestial sphere. In this second part of a three-article series, we discuss linear polarization of this scattered light parallel and perpendicular to the plane of scatter in the context of heliopheric imaging far from the Sun. The difference between these two radiances, (\\emph{pB}), varies quite differently with scattering angle, compared to the sum that would be detected in unpolarized light (\\emph{B}). The difference between these two quantities has long been used in a coronagraphic context for background subtraction and to extract some three-dimensional information about the corona; we explore how these effects differ in the wider-field heliospheric imaging case where small-angle approximations do not apply. We develop an appropriately-simplified theory of polarized Thomson scattering in the heliosphere, discuss signal-to-noise considerations, invert the scattering equations analytically to solve the three dimensional object location problem for small objects, discuss exploiting polarization for background subtraction, and generate simple forward models of several classes of heliospheric feature. We conclude that \\emph{pB} measurements of heliospheric material are much more localized to the Thomson surface than are \\emph{B} measurements, that the ratio \\emph{pB/B} can be used to track solar wind features in three dimensions for scientific and space weather applications better in the heliosphere than corona; and that, by providing an independent measurement of background signal, \\emph{pB} measurements may be used to reduce the effect of background radiances including the stably polarized zodiacal light. ", "introduction": "Introduction} Coronagraphs and heliospheric imagers observe sunlight that has been Thomson scattered off free electrons in the corona and solar wind. This potential was realized with the invention of the coronagraph \\citep{Lyot1939}, and solar wind transients such as coronal mass ejections (CMEs) have been observed with ground-based coronagraphs since the 1950s \\citep[e.g.][]{DeMastus1973}. These were accompanied by spacecraft coronagraphs in the 1970s \\citep[e.g.][]{Koomen1975,MacQueen1974} and the coronagraph legacy continues to this day. The physics by which this light is scattered are well established, with the original theory predating the discovery of the electron \\citep{Schuster1879}. Other important developments include the work of \\citet{Minnaert1930} and \\citet{Billings1966}. The latter is the publication most commonly referred to when discussing Thomson scattering theory with regard to white light observations. The utility of this theory to identify physical properties (such as mass) in solar wind transient phenomena (such as CMEs) observed by coronagraphs is well known. Early works include \\citet{Gosling1975,Hildner1975,Rust1979}; and \\citet{Webb1980}. In the last decade, coronagraphs have been accompanied by another type of white light observer. Heliospheric imagers, first SMEI \\citep{Eyles2003} and then the \\emph{STEREO/}HIs \\citep{Eyles2009}, observe Thomson scattered light at much larger angles ($>20^{\\circ}$) from the Sun, and their ability to track transients such as CMEs has been demonstrated \\citep[e.g.][]{Tappin2004,Howard2006,Webb2006,Harrison2008,Davis2009}. All current heliospheric imagers observe unpolarized Thomson scattered sunlight, in part because only recently \\citep{DeForest2011} has it become clear that the bright stellar background could be subtracted with sufficient precision for quantitative analysis of the faint Thomson-scattered signal from an imaging instrument. Although heliospheric imagers use the same scattering physics as coronagraphs, the wide viewing angle leads to significantly different geometry and requires different treatment. For example, Thomson scattering becomes much simpler in the heliospheric case because the Sun can be treated as a near-point source, eliminating the need to carry van de Hulst coefficients \\citep{Minnaert1930,vandehulst1950} when performing scattering/radiance calculations. More immediately, coronagraphs are often assumed to operate near the sky plane, but the relevant figure for a wide-field imager is the ``Thomson surface'' defined by the locus of the point of closest approach to the Sun of each line of sight from the observer. That locus is the sphere with diameter passing between the Sun and the observer \\citep{Vourlidas2006}. Paper I of this series \\citep{Howard2012} covered the applied theory of Thomson scattering to describe the relationship between this broad-field geometry, illumination, and scattering efficiency in unpolarized heliospheric imaging, and demonstrates that a fortuitous cancellation yields a broad plateau (the ``Thomson Plateau'') of nearly uniform radiance sensitivity to electron density. In the present paper, II of a planned series of three, we explore the consequences of Thomson scattering theory for polarized light in the heliospheric context, and discuss scientific applications of the theory. In particular, we invert the scattering equations analytically to show how polarized Thomson scattering imagery can be used to determine the three-dimensional location of individual small heliospheric features, without the front/back ambiguity present in similar efforts with coronagraphs \\citep[e.g.][]{Dere2005}, and explore analytically the limits of the technique. Further, we demonstrate via a simple forward model that the polarization signal remains present even for large features that whose position cannot be solved for analytically. We also discuss the stability of the polarization signal from the Zodiacal light and its implications for measuring the absolute radiance, rather than merely feature-excess radiance, of Thomson scattered light from heliospheric electrons. Polarization measurement of Thomson scattered light observation has existed since the dawn of the coronagraph \\citep{Lyot1933} and has been used for three dimensional analysis of CMEs since shortly after their discovery \\citep[e.g.][]{Poland1976,Wagner1982,Crifo1983}. The \\emph{Skylab} coronagraph, \\emph{Solwind}, C/P on board \\emph{SMM,} and LASCO on board \\emph{SOHO} all had polarizing capabilities. Perhaps because polarized coronagraph imagery, requiring photometry, is harder to work with than is unpolarized imagery \\citep[e.g.][]{MacQueen1993}, it has not been fully exploited in the spaceflight context although recent work \\citep{Dere2005,Moran2010,deKoning2011} may indicate a renaissance of polarized image exploitation in the corona. Polarized detection of Thomson scattered light from CMEs at wide angles from the Sun dates back much farther than direct heliospheric imaging cameras such as SMEI and \\emph{STEREO}/HI. The Helios spacecraft photometers \\citep[e.g.][]{Leinert1975} were used both to characterize the zodiacal light \\citep[e.g.][]{Leinert1981,Leinert1989} and also to detect heliospheric structures via time-domain analysis of the polarized intensity signal measured by three photometers on the spinning spacecraft \\citep{Jackson1986,Webb1987}. Hardware on board \\emph{Helios} sorted detected photon events into accumulator bins based on their temporal phase relative to the 1 Hz spacecraft spin. These bins were accumulated to yield sky brightnesses, including polarization and color signals, on time scales of several hours. These angularly separated photometric signals were used to generate synoptic maps of the Thomson scattering surface brightness of the solar wind \\citep{Hick1991}, to estimate CME mass \\citep{JacksonWebb1995,Webb1995}, and even to constrain coarse tomographic reconstructions of the three dimensional structure of the heliosphere \\citep{Jackson1995}. In Section \\ref{sec:Elementary}, below, we discuss heliospheric imaging with polarized light in the same context as we have previously with unpolarized light \\citep[Paper I:][]{Howard2012}, including: basic theory; importance of the Thomson surface; a summary description of instrument sensitivity, detectability, and signal-to-noise; and discussion of two key scientific applications of polarized imaging. Section \\ref{sec:Forward-Modeling-of} moves to simulations of these effects applied to non-infinitesimal features including CMEs and corotating interaction regions (CIRs). In Section \\ref{sec:Discussion} we discuss the applied theoretical results and their implications for future missions. ", "conclusions": "Discussion} We have developed an appropriately-simplified theory of polarized Thomson-scattered imaging in the heliosphere, and shown that sensitivity of $pB$ measurements to electron density is well localized along the line of sight at the TS, in contrast to the Thomson plateau observed via unpolarized light ($B$). For unpolarized detection, the illumination function and scattering efficiency have equal and opposite second derivatives at the TS, leading to the Thomson plateau; in the case of excess polarized radiance detection they have equal second derivatives at the TS, leading to a sharper maximum in intensity than would be observed via an $s-$scattering process with no angular dependence. The difference in spatial kernel between $pB$ and $B$ enables location of individual solar wind features in three dimensions, and we have developed a theory of small feature location including the importance of SNR and kinematic effects in determining the precision to which location may be measured. The curvature of the TS breaks the front/back asymmetry that hinders attempts to accomplish the same thing in coronagraph images near the Sun. Features far from the TS are more readily located in three dimensions than are features near the TS, because the $pB/B$ ratio varies more with the sky angle $\\xi$ far from the TS. An instrument that could measure the two relevant polarizations with comparable SNR to that of the \\emph{STEREO}/SECCHI imagers could locate the exit angle of small features within well under 10\\degr{}, greatly enhancing interpretation of the remote solar wind signal and potentially improving space weather prediction through direct location of Earth-directed features that are hard to measure geometrically \\citep[e.g.][]{Lugaz2010}. The $pB$ signal is a differential signal that, by construction, eliminates unpolarized bright features from the field of view. This effect has been used historically in coronagraphs to remove stray light and sky light from the Thomson scattering signal. Direct application in the heliosphere is more complex because the zodiacal light is polarized far from the Sun and hence is not removed by the \\emph{pB} calculation in the wide-field heliospheric imaging case. However, the zodiacal light is extremely stable both in degree of polarization and in overall intensity, leading to the possibility of direct absolute measurements of the Thomson scattering signal using long baselines and \\emph{in situ} measurement of the average wind density; the additional stability of the polarization ratio further reduces the residual background variation below what may be achieved with a single measurement. Other diffuse light sources from near Earth include orbital ram airglow and high altitude aurora, both of which we anticipate to be unpolarized and hence invisible in \\emph{pB.} We have performed simple forward modeling of geometric structures similar to known solar wind structures to develop intuition about how a polarized heliospheric imager, if developed, will respond to those features. We find that one may expect significant, unambiguous $pB$ signatures of the 3-D location of large features, although analysis of such features is more complex than compact features whose geometry may be neglected. Techniques for, and the limits of, large scale feature location in 3-D from $pB$ measurements require more detailed treatment and will be covered in the third paper of this series." }, "1207/1207.5770_arXiv.txt": { "abstract": "We argue that the multiple shells of circumstellar material (CSM) and the supernovae (SN) ejecta interaction with the CSM starting 59 days after the explosion of the Type Ia SN (SN Ia) PTF~11kx, are best described by a violent prompt merger. In this prompt merger scenario the common envelope (CE) phase is terminated by a merger of a WD companion with the hot core of a massive asymptotic giant (AGB) star. In most cases the WD is {{ disrupted}} and accreted onto the more massive core. However, in the rare cases where the merger takes place when the WD is denser than the core, the core will be {{ disrupted}} and accreted onto the cooler WD. In such cases the explosion might occur with no appreciable delay, i.e., months to years after the termination of the CE phase. This, we propose, might be the evolutionary route that could lead to the explosion of PTF~11kx. This scenario can account for the very massive CSM within $\\sim 1000 \\AU$ of the exploding PTF~11kx star, for the presence of hydrogen, and for the presence of shells in the CSM. ", "introduction": "\\label{sec:intro} Observations and theoretical studies cannot teach us yet whether all three scenarios for the formation of Type Ia supernova (SN Ia), or only one or two of them can work --- e.g., \\cite{Livio2001}, \\cite{Maoz2010}, \\cite{Howell2011}. These three basic theoretical scenarios can be described as follows. ($i$) In the single-degenerate (SD) scenario (e.g., \\citealt{Whelan1973}; \\citealt{Nomoto1982}; \\citealt{Han2004}) a WD grows in mass through accretion from a non-degenerate stellar companion. However, the mass increase of the WD seems to be very limited (e.g., \\citealt{Idan2012}). ($ii$) In the double-degenerate (DD) scenario (\\citealt{Webbink1984, Iben1984}; see \\citealt{vanKerkwijk2010} for a paper on sub-Chandrasekhar mass remnants) two WDs merge after losing energy and angular momentum through the radiation of gravitational waves \\citep{Tutukov1979}. Some list the ``double-detonation'' mechanism \\citep{Woosley1994, Livne1995} as a separate channel, although it involves two WDs. ($iii$) In the core-degenerate (CD) scenario for the formation of SN Ia the Chandrasekhar ($M_{\\rm Ch}$) or super-Chandrasekhar mass WD is formed at the termination of the CE phase, or during the planetary nebula phase, from a merger of a WD companion with the hot core of a massive AGB star \\citep{KashiSoker2011, IlkovSoker2012a, Soker2011}. The merger of a WD with the core of an AGB star was studied in the past (\\citealt{Sparks1974, Livio2003, Tout2008}). \\citet{Livio2003} suggested that the merger of the WD with the AGB core leads to a SN Ia that occurs at the end of the CE phase or shortly after, and can explain the presence of hydrogen lines. {{ However,}} due to its rapid rotation (e.g., \\citealt{Anand1965}; \\citealt{Ostriker1968}; \\citealt{Uenishi2003}; \\citealt{Yoon2005}; \\citealt{Lorenetal2009}), and possibly very strong central magnetic fields (e.g., \\citealt{Kundu2012}), the explosion of a WD with $M\\ge M_{\\rm Ch}$ might be substantially delayed. The recently observed SN Ia PTF~11kx \\citep{Dilday2012} has narrow lines, including hydrogen lines, and indications of interaction with a massive circumstellar medium (CSM) which starts 59 days after the explosion. \\cite{Dilday2012} argued that PTF~11kx can be explained by the SD scenario for SN Ia. They dismiss the merger {{ scenario}} as suggested by \\citet{Livio2003} on several grounds. {{ In particular,} they claim it cannot account for several CSM shells. We find this unjustified in section \\ref{sec:CSM}. We further discuss the {{ properties}} of the CSM in section \\ref{sec:radiated}, where we calculate the total radiated energy {{ arising}} from the collision of the exploding gas with the CSM. The destruction of the WD and its accretion onto the core while the core is still hot might prevent an early ignition of carbon \\citep{Yoon2007}, which is one of the theoretical problems of the DD scenario (e.g., \\citealt{SaioNomoto2004}). However, in the case of PTF~11kx ignition of the merger product is required quite early, within $\\sim 30 \\yr$ from the CE ejection. We therefore consider in section \\ref{sec:massive} the possibility that in the case of PTF~11kx the core was destructed onto the cooler WD. In {{ those}} cases where {{ an} explosion occurs shortly after merger in the CE, there is no delay by spin-down or emission of gravitational waves. {{{{ { This is called a violent-prompt merger, which might be considered as a sub-version of the DD and CD scenarios.} }}}} In section \\ref{sec:ignition} we show that {{ in}} these cases a massive envelope {{ can be ejected}}, and {{ that this scenario}} can account for the frequency of such SN Ia. Our summary is in section \\ref{sec:summary}. ", "conclusions": "\\label{sec:summary} The conclusion that the presence of any CSM around a SN Ia implies its association with the single-degenerate (SD) scenario --- e.g., \\cite{Sternberg2011} --- is problematic, in particular in cases where the CSM mass within $\\sim 1000 \\AU$ is $\\ga 0.01~{\\rm M_{\\sun}}$. An explosion set by the {{{{ {violent-prompt} }}}} merger of a WD companion with the core of the giant star naturally occurs within a massive CSM --- the ejected common envelope \\citep{Livio2003}. We find the association of the SN Ia PTF~11kx \\citep{Dilday2012} with the SD scenario to be unlikely due to the massive CSM. Instead, we found that the {{{{ {violent-prompt} }}}} merger of a massive WD with a massive core, as marked by the thick line on the upper right corner of Figures~\\ref{fig:massenv} and \\ref{fig:massratio}, to be able to account for the properties of PTF~11kx. We predict that interaction of the ejecta with the CSM will take place over the coming decades and that the amount of mass in the CSM will be accumulated to few times of ${\\rm M_{\\sun}}$. We also found from our population synthesis calculations that the number of systems with large total mass of the WD and core match the number of SN Ia with massive CSM as deduced from observations." }, "1207/1207.3295_arXiv.txt": { "abstract": "We test analytic predictions from different models of magnetospheric accretion, which invoke disk--locking, using stellar and accretion parameters derived from models of low resolution optical spectra of 36 T Tauri stars (TTSs) in NGC 2264 (age$\\sim$3 Myrs). Little evidence is found for models that assume purely dipolar field geometries; however, strong support is found in the data for a modified version of the X--wind model \\citep{shu94} which allows for non--dipolar field geometries. The trapped flux concept in the X--wind model is key to making the analytic predictions which appear supported in the data. By extension, our analysis provides support for the outflows predicted by the X--wind as these also originate in the trapped flux region. In addition, we find no support in the data for accretion powered stellar winds from young stars. By comparing the analysis presented here of NGC 2264 with a similar analysis of stars in Taurus (age$\\sim$1--2 Myr), we find evidence that the equilibrium interaction between the magnetic field and accretion disk in TTS systems evolves as the stars grow older, perhaps as the result of evolution of the stellar magnetic field geometry. We compare the accretion rates we derive with accretion rates based on U--band excess, finding good agreement. In addition, we use our accretion parameters to determine the relationship between accretion and H$\\beta$ luminosity, again finding good agreement with previously published results; however, we also find that care must be used when applying this relationship due to strong chromospheric emission in young stars which can lead to erroneous results in some cases. ", "introduction": "T Tauri stars (TTSs), first classified and named by \\citet{Joy45}, are low mass, pre--main sequence (PMS) objects that are roughly grouped into two classifications: 1. classical T Tauri stars (CTTSs), which show evidence of a circumstellar disk and mass accretion onto the central star in the form of excess emission in the X--ray, UV, optical, and infrared \\citep{bert89,argiroffi07}, and 2. weak--line T Tauri stars (WTTSs), or naked T Tauri stars (NTTSs), which are also PMS stars but do not show evidence for significant mass accretion and do not have inner disks \\citep{walt87}. WTTSs are thought to be the evolutionary product of CTTSs that have ceased accreting material from their disks and are continuing their gravitational contraction to the main sequence, though there is significant overlap in HR diagrams. Understanding how CTTSs interact with and ultimately disperse their disks is vital to our knowledge of how these systems evolve, supposedly into the WTTS phase, and eventually become sun--like stars, some of which are likely surrounded by planetary systems similar to our own. The goal of this study is to test predictions of magnetospheric accretion theories (discussed below) in the $\\sim$3 Myr old open cluster NGC 2264. The slightly greater age of NGC 2264 compared to well--studied regions like Orion and the Taurus--Auriga star forming region, both with an age $\\sim$1--2 Myr, allows an exploration of magnetospheric accretion in an older population of TTSs which permits an investigation of how magnetoshperic accretion may evolve over time. It is now well established that the excess emission observed in the spectra of CTTSs is due to the accretion of gas from a circumstellar disk \\citep[e.g.][]{bert88,phart91}. The original proposition for how the excess emission forms was put forth by \\citet{lbp74} who imagined the interaction between star and disk occuring in a boundary layer: the fast--rotating disk dissipates its energy in a narrow region which is in contact with the surface of the more slowly rotating central star, resulting in the excess emission observed at blue and UV wavelengths. While the boundary layer paradigm is able to explain the hot continuum emission \\citep[e.g.][]{bert88,bb89}, it fails to account for the large velocity--shifts in the line profiles of CTTSs and the large equivalent widths of H$\\alpha$ emission that are commonly observed \\citep{hart98}. Further evidence against the boundary layer model includes observed inner--disk holes around some CTTSs \\citep{meyer97}, which nevertheless shows signs of accretion \\citep{hart98}. Chromospheric models explaining the observed excesses were proposed by \\cite{cram79} and \\cite{calvet84}, and these were successful in reproducing some of the spectral chracteristics of CTTSs. However, like the boundary layer paradigm, these models fail to account for velocity shifts and widths in spectral lines and could not reproduce near--IR (disk) excesses, as well as becoming increasingly unfeasible when used to describe highly veiled spectra \\citep{calvet84}. Magnetospheric accretion, in which gas from the disk is loaded onto stellar magnetic field lines and impacts the surface of the star at near free--fall velocities, provides the best explanation for the observed properties of CTTSs \\citep{bouvier07}. Support for magnetospheric accretion as the dominant accretion process on CTTSs is strong. Magnetospheric accretion onto astrophysical objects was first investigated by \\cite{gl77} who modeled pulsating neutron stars accreting mass from a binary companion. \\cite{us84} first suggested a similar accretion mechanism for TTSs. The observational link between young stellar objects (YSOs) and the role magnetic fields play in their evolution was first established based on images of strong, collimated bipolar outflows and more highly collimated jets, and the energy required to power them \\citep[see][]{apenmundt89}. The first investigations of mass loss from TTSs, based on observations of energetic winds, were performed by \\cite{Varsavsky}, \\cite{herbig61}, and \\cite{kuhi64} by assuming that the flows were driven from the surface of the star. Some of the first evidence for the \\textit{circumstellar} origin of outflows in CTTSs, and as a result the conclusion that these outflows might be intrinsically linked to mass accretion onto the star, was discovered by \\cite{edwards89}. They suggested that the winds and outflows were direct results of accretion and thus could not be purely stellar in origin. We now know that bipolar outflows are ubiquitous phenemona in YSOs \\citep{bally07}. These outflows require the existence of strong, hourglass--shaped magnetic field lines to transport and collimate material from the disk or star into the observed bipolar flows \\citep{pudritz83,pudritz86,shu88,cam90}; winds driven purely by thermal or radiation pressure cannot account for the large outflow velocities that are observed \\citep{lada85}. Magnetic field lines can also act as a collimating agent, bounding the outflows at large distances from the central star \\citep[e.g.][]{shu88}. Magnetospheric accretion theories can also account for the observed variability of the UV excess \\citep{bert88,alencar01}: accretion columns terminating at the stellar surface co--rotate with the star, thus producing variability on the same timescale as the stellar rotation period. Velocity shifts of emission and absorption lines are naturally explained by the $\\sim$250 km/s velocities of the infalling material \\citep[e.g.][]{calvet98}. One of the most important questions concerning the evolution of TTSs is the problem of how these objects are able to shed angular momentum and rotate with velocities well below break--up \\citep[e.g.][]{vogel81,herbst02,lamm04,mak04}. The magnetospheric accretion model provides a potential answer to this question due to the role of the stellar magnetic field in transferring angular momentum from the star to the disk: if the magnetic field couples with a sufficiently ionized disk and acts as a braking torque on the star, it will rotate more slowly than if no braking mechanism were present \\citep{kon91,shu94}. This interaction essentially locks the star to the disk (i.e. \\textit{disk--locking}) and prevents the star from rotating at break--up velocity. This scenario can account for many of the observed rotation rates of PMS stars \\citep[][hereafter L04]{edwards93,kearns98,lamm04}. Strong evidence for disk--locked stellar rotation has been found in the Orion Nebula Cluster (ONC) \\citep{choi96,herbst02} and NGC 2264 \\citep[][hereafter L05]{lamm05}, though contradictory results have been found by the studies of \\cite{stass99} and \\cite{rebull01} for the ONC, and \\cite{mak04} for NGC 2264. Evidence supporting magnetic disk--locking manifests itself in the form of a bimodal period distribution \\citep[e.g.][]{attridge92} and the detection of disk signatures \\citep{edwards93,rebull06}, and hence circumstellar disks, around long--period stars. \\cite{herbst02}, based on the work of \\cite{choi96}, point out that the shorter period peak in the bimodal distribution is simply a binning artifact (i.e. the rotational evolution of stars that have ceased to be regulated by their disks should span a range of shorter periods) while the longer period peak ($\\sim$8 days) is a direct result of magnetic disk--locking \\citep{attridge92,herbst01,herbst02}. To date, studies by \\cite{rebull06}, \\cite{dahm2011}, and \\cite{cieza07} provide the strongest evidence for correlations between slowly rotating stars and the presence of a circumstellar disk, though the \\cite{dahm2011} study finds high confidence levels ($>$99\\%) for only the M--dwarfs in their sample. The work of \\cite{herbst02} also shows strong evidence that longer period stars have a higher incidence of circumstellar disks. These correlations support the hypothesis of star--disk interactions being the main culprit in removing angular momentum from TTS systems.\\defcitealias{lamm05}{L05}\\defcitealias{lamm04}{L04} \\citetalias{lamm05} also report a bimodal period distribution for a large sample of stars in NGC 2264, similar to what has been reported for the ONC. However, \\cite{mak04} find no evidence for a bimodal period distribution in NGC 2264. In fact, \\cite{mak04} find no significant difference between the period distributions of the Orion region and NGC 2264, which differ in age by a factor of $\\sim$2 \\citepalias{lamm05}. This result directly conflicts with the results of \\citetalias{lamm05} who find that the period distribution of angular momentum conservation for NGC 2264 is consistent with stellar contraction from the period distribution of the ONC based on fully convective PMS models. \\citetalias{lamm05} explain that the conclusions of \\cite{mak04} are a result of the latter study comparing the period distribution of NGC 2264 with a larger, inhomogeneous region in Orion. \\citetalias{lamm05}, on the other hand, only compare the period distribution of NGC 2264 with the younger homogeneous region of the ONC. Studies of both clusters find that the rotation period distributions for higher mass stars ($M_*$ $>$ 0.25 M$_\\odot$) and lower mass stars ($M_*$ $<$ 0.25 M$_\\odot$) peak at different locations, with the lower mass stars peaking at a shorter period than those with greater mass \\citep{herbst01,herbst02,mak04,lamm05}. Despite some conflicting results, observational evidence seems to point to disk--star interactions as being the primary candidate for explaining TTS rotation rates. Further support for magnetic disk--locking can be found in the measurements of magnetic fields on TTSs. Though the sample size of measured fields is relatively small, \\cite{JK99}, \\cite{guenther99}, \\cite{JK07}, and \\cite{yangJK11} measure relatively uniform values of $\\sim$1-3 kG for surface magnetic fields on TTSs. The relative constancy of their magnetic fields, combined with the ubiquitous presence of circumstellar disks, points to interactions between the stellar magnetic field and disks as being a prime candidate for angular momentum regulation in CTTSs. However, \\cite{JK07} points out that the magnetic fields of TTSs may not be strong enough to enforce disk--locking if a dipolar field geometry is not assumed at the stellar surface. On the other hand, \\cite{JK02}, hereafter JG02, investigated correlations predicted by several different theories of magnetospheric accretion (see below) and found support for models that assume a \\textit{non}--dipolar surface field geometry and only weak support for models assuming purely dipolar fields.\\defcitealias{JK02}{JG02}\\defcitealias{OS95}{OS95} In this study, we extend the analysis performed by \\citetalias{JK02} to a larger sample of PMS stars in NGC 2264. \\citetalias{JK02} examined several analytic relationships predicted by four different magnetospheric accretion theories, specifically those of \\cite{kon91}, \\cite{camcamp93}, \\cite{shu94}, and a modified version of the \\cite{OS95} model (hereafter OS95), using multiple sets of observations of CTTSs in the Taurus--Auriga molecular cloud complex. The theories and their predictions will be discussed in Section 2. Modeling the accretion columns producing the excess emission in CTTSs as hot slabs of hydrogen \\citep[e.g.][]{valenti93} enables us to derive reliable estimates of stellar and accretion parameters for our sample in NGC 2264. We then test the same correlations predicted by the aforementioned theories for this set of stars. NGC 2264 is an ideal candidate for this study due to the availability of rotation periods and spectral type determinations for a large number of stars. As mentioned above, NGC 2264 is slightly older than the Taurus--Auriga region and the ONC (3 Myrs vs. 1-2 Myrs; \\citetalias{lamm05}) and so should provide a good testing ground for the models using a more evolved PMS stellar population. Exploring young clusters at different ages allows us to see when the equilibrium magnetspheric accretion relationships and disk--locking break down. Section 3 describes our observations and reductions and the classification of the stellar sample into accreting and negligibly accreting stars. In Section 4 we discuss our models and how they are used to derive stellar and accretion parameters (see Appendix A for a detailed description of our fitting procedure and model assumptions). Section 5 includes the application of the accretion theory predictions to our sample and tests of the resulting correlations. A discussion of the results is given in Section 6 and our conclusions are presented in Section 7. ", "conclusions": "We have derived stellar and accretion parameters for 36 TTSs in NGC 2264 using spectrophotometric measurements taken in 2004 and 2005. Our estimate for the age of the sample ($\\sim$6.4 Myrs), calculated using the pre--main sequence evolutionary tracks of \\citet{siess00}, is older than several more statistically significant age determinations which average to $\\sim$3 Myrs. This is due in large part to: 1. our small sample size, and 2. our use of the \\citet{siess00} PMS models which tend to produce older ages compared with the commonly used tracks of \\citet{baraffe98} and \\citet{dantona94} \\citep{rebull02,dahm05}. The mass accretion rates for our stars are similar to estimates by \\citetalias{rebull02}, though the agreement is worse at lower $\\dot{M}$. Using the derived parameters we test analytic predictions from the purely dipolar disk--locking models of \\citet{kon91} and \\citet{shu94} (eq. 1) and the non--dipolar field prediction of \\citetalias{JK02}'s modified version of the \\citet{OS95} model (eq. 3). We find good support for the modified \\citet{OS95} model of magnetospheric accretion and disk--locking, although the correlation is influenced to some degree by a strong relationship between $\\dot{M}$ and $R_*^2$. A lack of support for the dipolar theories, however, highlights the need for an extra constraint in the theory or the abandonment of disk--locking. This constraint is provided by the inclusion of $f_{acc}$ in eq. (3). Although the support we find for the \\citet{OS95} theory is not without uncertainty, our results confirm the findings of \\citetalias{JK02} and provide more evidence for disk--locking than against it. In addition, our results find no support for theories that assume a dipolar magnetic field geometry at the stellar surface. Recent evidence has shown that this assumption is probably not valid for TTSs \\citep[e.g.][]{daou06,mohanty08,donati08,donati11b}. This scenario does not exclude the possibility of accretion powered stellar winds as the agent which removes angular momentum from TTSs \\citep[e.g.][]{matt05,cranmer}. In fact, both stellar winds and disk winds launched from near the truncation point probably play a role in the removal of angular momentum, as suggested by \\citet{edwards06}. The primary assumption made in our initial analysis is that the magnetic field strength does not vary significantly from star to star. We have attempted to roughly account for differing dipolar magnetic field strengths from star to star by assigning values based on the size of the star's radiative core as suggested by the recent results of Donati et al.; however, doing so actually weakened the correlations present in the data. In order to make this analysis more robust, better statistics on the strength of dipolar magnetic field components in TTSs, especially those at the $\\sim$3 Myr age, are needed. Whether the constant field strength assumption is justified or not for stars at similar evolutionary stages may help explain the slope evolution observed in our data. In addition, we do not find any relationship between the theoretical dipolar magnetic field component and radiative core mass in our sample. Future studies of emission and absorption lines from TTSs environments will help place constraints on the launching region of the outflows, thus helping to confirm or reject the hypothesis of disk--locking, and subsequent removal of angular momentum by a disk wind from the truncation radius, supported in this work. Until more observational evidence can be gathered, it appears that the non--dipolar magnetspheric accretion model of \\citet{OS95} remains a strong candidate for explaining the observed relationship between stellar and accretion parameters of CTTSs at the $\\sim$3 Myr evolutionary stage and, by extension, that the disk--locking scenario is taking place in young stars. \\appendix" }, "1207/1207.4683_arXiv.txt": { "abstract": "The aim of this work is to propose a joint exploitation of heterogeneous datasets from high-resolution/few-channel experiments and low-resolution/many-channel experiments by using a multiscale needlet Internal Linear Combination (ILC), in order to optimize the thermal Sunyaev-Zeldovich (SZ) effect reconstruction at high resolution. We highlight that needlet ILC is a powerful and tunable component separation method which can easily deal with multiple experiments with various specifications. Such a multiscale analysis renders possible the joint exploitation of high-resolution and low-resolution data, by performing for each needlet scale a combination of some specific channels, either from one dataset or both datasets, selected for their relevance to the angular scale considered, thus allowing to simultaneously extract high resolution SZ signal from compact clusters and remove Galactic foreground contamination at large scales. ", "introduction": "Multifrequency observations of the microwave sky are a mixture of various diffuse and compact components \\citep{1995SSRv...74...37B,1999NewA....4..443B}: the Cosmic Microwave Background (CMB) emission, the SZ effect from galaxy clusters, the foreground emission from the Galactic interstellar medium (ISM), the Cosmic Infrared Background (CIB) emission, the emission from compact extra-galactic radio sources, the instrumental noise. Each of these components has a distinctive frequency signature (which can be known or not), a distinctive spatial distribution on the celestial sphere, and a distinctive spectral distribution on the angular scales (power spectrum). The separation of the sky components in such multifrequency observations of the CMB is an important part of the processing and {analysis} of such observational data. The thermal Sunyaev-Zeldovich (TSZ) effect from galaxy clusters is a spectral distortion of the CMB black body radiation due to inverse Compton scattering of the CMB photons off hot electrons contained in the intra-cluster gas \\citep{1972CoASP...4..173S}. It is responsible for secondary temperature anisotropies, which introduce an excess of power in the primordial CMB temperature anisotropies at small angular scales. This spectral distortion is independent of the cosmological redshift so that the measure of the TSZ effect is a powerful and unique tool to detect new galaxy clusters at any redshift \\citep{1999PhR...310...97B, 2002ARA&A..40..643C}. The separation and the extraction of the TSZ effect from the other sky emissions is made possible by a multifrequency coherent analysis because of the prior knowledge of the frequency signature of the SZ effect. The main limitation in detecting SZ clusters comes from the achievable resolution of the instrument and the level of contamination by the other sky emissions and the instrumental noise. Various component separation methods have been developed to extract the thermal SZ signal from the observed frequency maps of a single experiment. Some of them require a prior assumption on the template SZ profile such as the matched filtering methods in \\citet{1996MNRAS.279..545H}, \\citet{2002ApJ...580..610H,2002MNRAS.336.1057H}, and \\citet{2006A&A...459..341M}. Other methods are blind such as the methods based on statistical independence (ICA in \\citet{1999ISPL....6..145H,2002MNRAS.334...53M}, Spectral Matching ICA in \\citet{2003MNRAS.346.1089D,2008arXiv0803.1814C}, GMCA in \\citet{2008StMet...5..307B} which also benefits from the sparsity of the components). Some methods are parametric and require physical modeling of some components (MEM in \\citet{1998MNRAS.300....1H}, Commander in \\citet{2008ApJ...676...10E}). Non-parametric methods include needlet Internal Linear Combination (NILC) in \\citet{2008A&A...491..597L} and in \\citet{2009A&A...493..835D}, MILCA in \\citet{2010arXiv1007.1149H}, multidimensional NILC in \\citet{2011MNRAS.410.2481R} and \\citet{2011MNRAS.418..467R}). The ILC, which has been used first on the data of the WMAP mission to reconstruct the CMB emission \\citep{2003ApJS..148...97B}, is a particularly simple component separation method which does not assume any particular parametrization for foreground emission, and for the other emissions considered as contaminants. The frequency scaling vector of the component of interest to extract is the only parameter needed for implementing ILC. For all the component separation methods, the problem is to find the best compromise between the most accurate unbiased estimation of the component of interest and the best minimization of the residuals from the other sky emissions. In this work we address the problem of jointly exploiting the datasets of multiple experiments with various specifications for optimizing the reconstruction of the thermal SZ effect at high resolution. No single instrument provides the best measurement of the SZ effect. Space-borne instruments, such as Planck for instance \\citep{planck2011-5.1b, 2011arXiv1112.5595P, 2012arXiv1205.3376P}, can observe the whole sky in the millimetre through many frequencies (nine channels from $30$ GHz to $857$ GHz). The large number of frequency channels covered by such space-borne instruments enables component separation methods, such as ILC, to improve the minimization of various contaminations (Galactic foregrounds, CMB, instrumental noise) in the reconstructed SZ map. However, the instrumental resolution of Planck for instance is limited to $5$ arcmin, therefore the extraction of SZ signal from clusters of arcminute angular size, by using such dataset only, remains unachievable. Conversely, higher resolution ground-based telescopes, such as the Atacama Cosmology Telescope (ACT) \\citep{2011ApJ...731..100M} and the South Pole Telescope (SPT) \\citep{2011ApJ...738..139W, 2012arXiv1203.5775R} for instance, are designed to measure the SZ temperature anisotropies up to arcminute angular scales but the presence of the atmosphere limits the frequency coverage of those experiments. The few number of frequency channels observed by such ground-based telescopes ($148$ GHz, $218$ GHz, and $277$ GHz in ACT) limits the ability of these instruments to remove the Galactic foreground contamination in the reconstruction of the SZ signal of nearby clusters. Here we propose to perform SZ component separation from the joint exploitation of heterogeneous sets of maps (both maps from many-channel/lower-resolution instrument and maps from few-channel/higher-resolution instrument) by using needlet ILC as a \\emph{multiscale}, or \\emph{multiresolution}, approach for combining multiple instrument datasets. The paper is organised as follows. In Sect. \\ref{sec:method} we describe the needlet ILC method on patches of the sky and show how the needlets can be exploited to combine multiple instrument datasets with various resolutions for optimizing the SZ component separation. The results are presented on simulations of the sky in Sect. \\ref{sec:results} where we apply the method jointly on two heterogeneous instrument datasets. We discuss the robustness of the approach in Sect. \\ref{sec:discussion} and conclude in Sect. \\ref{sec:conclusion}. ", "conclusions": "\\label{sec:conclusion} The reconstruction of galaxy cluster profiles through thermal SZ effect is an exciting challenge in data analysis of current CMB experiments because the SZ profile is a powerful probe of the baryon physics in the cluster outskirts, beyond the radius of virialisation in the intra-cluster medium. In this region, X-ray measurements fail to reconstruct the thermal pressure profile. For the first time, the SZ profiles of some galaxy clusters have been reconstructed from the surveys of the high resolution telescopes ACT \\citep{2011ApJ...732...44S} and SPT \\citep{2010ApJ...716.1118P}. Results on the SZ profile from the Planck survey have also been published very recently in \\citet{2012arXiv1207.4061P}. Both in ACT and SPT results, the SZ signal has been reconstructed by using a matched spatial filtering, which is an effective way of reconstructing the SZ effect when few frequency maps are available from the observations. However, matched filtering relies on priors on the expected template profile, which may appear controversial when one is interested in reconstructing the cluster profile. The needlet ILC on patches developed in this work is a blind approach which does not suffer from any prior to reconstruct the SZ profiles. Apart from the prior issue, the accuracy of the SZ reconstruction also relies on the achievable resolution of the instrument and on the level of contamination either by sky emissions (Galactic foregrounds, CMB, etc) or by the instrumental noise. We have highlighted that both problems can be tackled simultaneously by combining extra frequency maps from multiple instrument datasets with different resolutions. We have shown in this work how needlet ILC has the ability to adapt the needlet windows in Fourier space to the beams of the channel-maps exploited for component separation, thus offering the possibility of combining heterogeneous datasets coming from multiple instruments with various resolutions and different frequency coverages. On the one hand, the properties of localization of the needlets both in scale and in space enable component separation to adapt to the local conditions of contamination (physical at large scales, instrumental at small scales), on the other hand they enable component separation to aggregate heterogeneous instrument datasets with different achievable resolutions. We have applied needlet ILC on three different simulated datasets (single Planck-like, single ACT-like, and joint Planck-like/ACT-like) for different patches of the sky centered on various extended and compact clusters. The performance of the reconstruction has been validated both on particular clusters and on samples of selected clusters. The aggregation of multiple datasets with needlet ILC allows us to reconstruct the SZ profile over a large radius with both high resolution and robust foreground and noise cleaning. We have also explored the limitations of the method. Needlet ILC applied on a small-size patch of the sky encounters a bias in the reconstructed SZ signal. However, combining a subset of channels in that case enables ILC to tackle this problem of bias. An optimized selection of the subset of channels would be required to improve the reconstruction on very small patches. It would also be very instructive to optimize in a future work the wavelet localization (type of wavelets, width of spectral windows) for multi-instrument SZ component separation. Needlet ILC performed on the joint dataset including both Planck-like and ACT-like datasets successfully reconstructs the TSZ effect from compact clusters (${\\theta_{200} < 5~ \\mathrm{arcmin}}$) beyond the beam of the single Planck-like instrument while cleaning foreground residuals better than a component separation applied on the single ACT-like instrument. The multiresolution approach presented in this work appears to be promising for a future multi-instrument analysis of the SZ effect." }, "1207/1207.3892_arXiv.txt": { "abstract": "Gaussian random fields pervade all areas of science. However, it is often the departures from Gaussianity that carry the crucial signature of the nonlinear mechanisms at the heart of diverse phenomena, ranging from structure formation in condensed matter and cosmology to biomedical imaging. The standard test of non-Gaussianity is to measure higher order correlation functions. In the present work, we take a different route. We show how geometric and topological properties of Gaussian fields, such as the statistics of extrema, are modified by the presence of a non-Gaussian perturbation. The resulting discrepancies give an independent way to detect and quantify non-Gaussianities. In our treatment, we consider both local and nonlocal mechanisms that generate non-Gaussian fields, both statically and dynamically through nonlinear diffusion. ", "introduction": "\\label{sec_crit} To gain some insights into the physical mechanisms that generate non-Gaussianities, consider first how an isotropic Gaussian field $H_G(\\vec{r})$ arises from the random superposition of waves (or equivalently, Fourier modes) \\begin{equation} H_G(\\vec{r}) = \\sum_{\\vec{k}} A(k) \\cos( \\vec{k} \\cdot \\vec{r} + \\phi_{\\vec{k}} ), \\label{eq1} \\end{equation} \\noindent with an amplitude spectrum $A(k)$\\footnote{The power spectrum $A(k)^2$ is the Fourier transform of the two-point correlation function.} that depends only on the magnitude of the wave vectors, $k = |\\vec{k}|$. The phases $\\phi_{\\vec{k}}$ are uncorrelated random variables uniformly distributed in the range $[0, 2\\pi]$. The statistical properties of the Gaussian field $H_G(\\vec{r})$ are entirely encoded by the function $A(k)$, or equivalently, its moments $K_{2n}=\\sum_{\\vec{k}} k^{2n} \\frac12 A(k)^2$. The most basic difference between Gaussian and non-Gaussian variables is that Gaussian ones are always symmetric about their mean. As a consequence, irrespective of its power spectrum, a Gaussian field has equal densities of maxima and minima. Hence, a nonvanishing imbalance $\\Delta n$ between these two types of extrema serves as a probe to detect and quantify the non-Gaussian component of a signal, provided that it can be measured directly. For example, consider the primordial curvature perturbation field, $\\Phi$, a nearly Gaussian field of central interest to modern cosmological studies \\cite{cite_Dodelson}. Within a \\emph{local} approximation, the primordial field is obtained from a Gaussian field $\\Phi_G$ via a nonlinear relation $\\Phi = \\Phi_G + f_{nl} \\Phi_G ^2 + g_{nl} \\Phi_G ^3$. Determining the parameters $f_{nl}$ and $g_{nl}$ is one of the central tasks in the study of cosmological non-Gaussianities. As we shall see, the quadratic coefficient can be determined from the imbalance $\\Delta n$ between maxima and minima of $\\Phi$. The imbalance can be derived in the more general context of a non-Gaussian field $h$ that is obtained from a Gaussian field $H_G$ via any nonlinear deformation $h=F_{NL}(H_G)$. If $F_{NL}$ is a monotonic function, the maxima and minima do not change -- only a nonmonotonic behavior of $F_{NL}$ can alter this balance. The critical points of $h$ are given by $ \\vec{\\nabla} h = F_{NL}'(H_G) \\vec{\\nabla} H_G = 0$, where the dash indicates the derivative of $F_{NL}(H_G)$ with respect to $H_G$. This condition shows that $h$ and $H_G$ have the same critical points. Note however that, if $F_{NL}'(H_G(\\vec{r_0})) < 0$ at a critical point $\\vec{r}_0$, then a maximum (minimum) at $H_G(\\vec{r}_0)$ will be turned into a minimum (maximum) at $h(\\vec{r}_0)$. The number of saddle points does not change because of topological constraints (it is equal to the invariant sum of maxima and minima). If the transformation has a bias towards converting minima into maxima, then $h$ will have more maxima than minima; for example, $h=H_G+\\varepsilon H_G^2$ reverses its slope at sufficiently negative values of $H_G$, which are most likely to be minima. Following this logic, the first step toward calculating the imbalance between maxima and minima is to determine the probability $g(z)$ that $H_G(\\vec{r_0}) = z$ for a minimum $\\vec{r_0}$ of $H_G$. The symmetry properties of $H_G$ imply that the analogous probability distribution for maxima is $g(-z)$. The fraction of minima of $H_G$ that become maxima of $h$ is obtained by integrating $g(z)$ over the range of $z$ for which $F_{NL}'(z) < 0$. Likewise, the fraction of maxima of $H_G$ that are turned into minima is given by the integral of $g(-z)$ over the same range. The overall imbalance in the densities of the maxima and minima of $h$ can be readily obtained by adding these opposite contributions. The result reads \\begin{equation} \\begin{split} \\Delta n\t& \\equiv \\frac{n_{\\mathrm{max}} - n_{\\mathrm{min}}}{n_{\\mathrm{max}} + n_{\\mathrm{min}}} \\\\ & = \\int_{z:F_{NL}'(z)<0} \\! \\rd z \\, \\big( g(z) - g(-z) \\big). \\end{split} \\label{eq_maxmin} \\end{equation} For two dimensions, the exact analytical expression for $g(z)$ is explicitly derived in Appendix A -- it depends only on the moments $K_0$, $K_2$ and $K_4$. Rescaling $h(\\vec{r}) \\rightarrow h(a\\vec{r})$ does not affect the function $g(z)$, since it increases the density of maxima with any value of $H_G$ by the same proportion. Hence, only $K_0=\\langle H_G^2\\rangle$ (which sets the scale of the distribution) and the dimensionless parameter $\\lambda \\equiv \\frac{K_2^2}{K_0K_4}$ can enter in the expression for $g(z)$ (see plot in the inset of Fig.~\\ref{fig_maxmin}). \\begin{figure} \\centering \\includegraphics{fig2a} \\caption{The relative difference $\\Delta n$ between the densities of maxima and minima of $h = H_G + \\varepsilon H_G^2$, where $H_G$ is a Gaussian field with $\\lambda = 3/4$, as a function of $\\varepsilon$. The data points are results from computer-generated fields, the solid curve is the theoretical result (Eq.~\\eqref{eq_maxmin}). The inset shows the corresponding distribution of minima $g(z)$ (which forms the basis of our theoretical result), both the theoretical curve and a histogram of data gathered from computer-generated fields.} \\label{fig_maxmin} \\end{figure} As an illustration, apply Eq.~\\eqref{eq_maxmin} to the perturbed Gaussian field $h = H_G + \\varepsilon H_G^2$. Figure~\\ref{fig_maxmin} shows our theoretical formula as a continuous line, validated by numerical data (dots) obtained from computer-generated random surfaces with amplitude spectrum $A(k) \\sim \\theta(k_D-k)$, having $\\lambda = \\tfrac34$. The imbalance between maxima and minima is particularly useful to track large deviations from Gaussianity. This can be seen explicitly for the quadratic perturbation considered above. Note that, since $H_G$ itself has an equal number of maxima and minima, the imbalance $\\Delta n$ in $h$ is only created when one of these critical points is inverted. Whether $H_G$ has a high likelihood of having a negative $F_{NL}'(H_G)=1+2\\varepsilon H_G$ is controlled by $\\varepsilon \\sqrt{\\langle H_G^2\\rangle}$. The probability of $|2 \\varepsilon H_G|$ exceeding $1$ is exponentially small. Thus, $\\Delta n$ rises roughly as $e^{-c/({\\varepsilon^2}\\langle H_G^2\\rangle)}$, where $c \\approx \\frac{1}{8}$. Note that our approach in deriving Eq.~\\eqref{eq_maxmin} and $g(z)$ is nonperturbative and hence capable of describing large deviations from Gaussianity (for the specific type of deviation considered here), as demonstrated by the agreement between our formula and numerics over the entire $\\varepsilon$-range probed in Fig.~\\ref{fig_maxmin}. ", "conclusions": "" }, "1207/1207.6153_arXiv.txt": { "abstract": "Most Gamma-ray bursts (GRBs) have erratic light curves, which demand that the GRB central engine launches an episodic outflow. Recent Fermi observations of some GRBs indicate a lack of the thermal photosphere component as predicted by the baryonic fireball model, which suggests a magnetic origin of GRBs. In view that powerful episodic jets have been observed along with continuous jets in other astrophysical black hole systems, here we propose an intrinsically episodic, magnetically-dominated jet model for GRB central engine. Accumulation and eruption of free magnetic energy in the corona of a differentially-rotating, turbulent accretion flow around a hyperaccreting black hole lead to ejections of episodic, magnetically dominated plasma blobs. These blobs are accelerated magnetically, collide with each other at large radii, trigger rapid magnetic reconnection and turbulence, efficient particle acceleration and radiation, and power the observed episodic prompt gamma-ray emission from GRBs. ", "introduction": "Observations of gamma-ray bursts (GRBs) suggest that the GRB central engine is able to launch an ultra-luminous, highly relativistic jet. Most GRBs have erratic, rapidly varying lightcurves \\citep{fishman95}, typically lasting 10s to 100s of seconds for long-duration GRBs, and less than 2 seconds for short-duration GRBs. Recent observations of GRBs pose some important constraints on the models of GRB central engine. First, recent Fermi Large Area Telescope (LAT) observations suggest that most GRBs have featureless smoothly joint broken power-law spectra (i.e. Band-function spectra \\citep{band93}) in a wide energy band \\citep{abdo09,zhang11}. The standard fireball model predicts a strong thermal emission component from the fireball photosphere \\citep{pac86,meszaros00,peer06}. The non-detection likely suggests that the GRB outflow is magnetically dominated \\citep{zhang09,fan10}. Second, data analysis suggests that the GRB lightcurves not only can be decomposed into many pulses \\citep{norris96,hakkila03}, but most of them can be also decomposed into the superposition of a fast, spiky component and a slow, smooth component \\citep{gao12,vetere06}. A radiation model that invokes magnetic turbulent reconnection triggered by collisions of magnetically-dominated blobs has been proposed \\citep{zhang10}, which can interpret the new observational data. This radiation model requires a central engine that can eject an episodic, magnetically dominated jet. The leading model of GRB central engine invokes a hyper-accreting black hole \\citep{narayan92,meszaros99,narayan01}. In most previous GRB central engine models, the rapid variability in the erratic light curves is attributed either to the intermittency of the accretion flow, i.e. a time-dependent accretion rate $\\dot M$ \\citep{macfadyen99}, or to the instability during the propagation of the jet in the stellar envelope \\citep{zhang03,morsony10}. For a magnetically dominated central engine, the leading model is the Blandford-Znajek (BZ) mechanism \\citep{blandford77}, which requires a large-scale open magnetic field connecting the black hole and an external astrophysical load, whose origin is an open question. This model also tends to generate a continuous jet, unless the accretion rate is highly variable. Screw or kink instabilities \\citep{li00,mizuno09} are invoked to disrupt a continuous jet and produce discrete blobs. Since the BZ mechanism is powered by accretion, the immediate advantage of involving BH spin rather than accretion disk as the source of jet power is not evident. On the other hand, in addition to the continuous jets, episodic jets, which is intermittent and in the form of discrete moving blobs, have been observed in other black hole systems. A magnetohydrodynamical model has been proposed to explain the formation of these episodic jets \\citep{yuan09}. In this paper, we propose a central engine model for GRBs, based on this idea. A review on observations and the model of episodic jets are presented in Section 2. Our model is delineated in detail in Section 3. The salient features of the model are summarized in Section 4. ", "conclusions": "Observations to relatively well-studied black hole systems such as AGNs and black hole X-ray binaries show the existence of two types of jets, namely continuous and episodic ones. GRB observational data require that the central engine launches a magnetically dominated episodic jet. The traditional BZ jet model requires that the accretion rate is highly variable or that a continuous jet is disrupted by instabilities during propagation. In this paper we propose an intrinsically episodic central engine model for GRBs by invoking ejections of episodic magnetic blobs from a hyper-accretion flow around a black hole. Our basic calculations suggest that the predicted energetics and timescales are consistent with the observations. More detailed numerical simulations are called for to validate the scenario. This model has several appealing features. Firstly, the episodic jets invoked in our model have obtained strong observational supports in other black hole systems \\citep*[e.g.,][and references therein]{yuan09}. Observations show that they are more powerful than continuous jets. Secondly, this model naturally satisfies the observational requirement for the lack of a bright photosphere component in the GRB spectra \\citep{zhang09}, since the engine natually launches a magnetized outflow. Thirdly, episodic jets intrinsically consist of individual blobs, whose collisions are natually expected. These collisions are an important ingredient to interpret GRB variability within the magnetically dominated model \\citep{zhang10}. Fourthly, the directions of the magnetic fields surrounding the two adjacent blobs in general have some angles. This greatly eases triggering fast reconnection and the subsequent turbulence, which can efficiently convert magnetic energy into radiation \\citep{zhang10,lazarian99,mckinney11}. We add two notes here. First, we did not discuss collimation of episodic jets in the progenitor stellar envelope. We point out that the same physics invoked in the magnetic tower model \\citep{uzdensky2006} should equally apply to our model. Second, the scenario we propose should also work for a neutron star, not necessarily a black hole, as has been discussed by some previous authors \\citep[e.g.][]{blackman96,kluzniak98,dai06,metzger11}. It may be difficult to differentiate this model from other GRB central engine models from the GRB observational data. Due to their large distances and small scales, it is essentially impossible to witness ejection of magnetic blobs from GRBs. Nonetheless, observing episodic blobs may be possible for nearby AGNs and X-ray binaries with high spatial and temporal resolution observations in the future. These observations may be used to verify this generic episodic central engine model." }, "1207/1207.3630_arXiv.txt": { "abstract": "We use a simple analytical model to derive a closed form expression for the bolometric light-curve of super-luminus supernovae (SLSNe) powered by a plastic collision between the fast ejecta from core collapse supernovae (SNe) of types Ib/c and IIn and slower massive circum-stellar shells, ejected during the late stage of the life of their progenitor stars preceding the SN explosion. We demonstrate that this expression reproduces well the bolometric luminosity of SLSNe with and without an observed gamma ray burst (GRB), and requires only a modest amount ($M < 0.1\\,M_\\odot$) of radioactive $^{56}$Ni synthesized in the SN explosion in order to explain their late-time luminosity. Long duration GRBs can be produced by ordinary SNe of type Ic rather than by 'hypernovae' - a subclass of superenergetic SNeIb/c. ", "introduction": "The progenitor stars of core-collapse supernovae of type Ib/c, are stripped of their envelope through strong winds and/or major eruptions during the final stages of their life before their explosion (e.g., Pastorello et al.~2010; Quimby et al.~2011, and references therein). The ejected massive shells in these eruptions sweep up slower stellar winds, which were blown before these eruptions, and create a matter-clean space surrounding the progenitor star. The interaction of the radiation from the SN explosion with such slowly moving circumstellar (CS) shells often produces delayed emission of narrow lines, which stops when the SN ejecta collide with the CS shell (Chugai~1990,1992). Light from the progenitor star back-scattered by CS shell(s) into the matter-free space around the progenitor star produces a glory - a halo of scattered light surrounding the progenitor star. In the cannonball (CB) model of gamma ray bursts (GRBs), long duration GRBs are produced by inverse Compton scattering of glory photons by the electrons in the highly relativistic bipolar jets of plasmoids (cannonballs) of ordinary matter, which presumably are ejected in mass accretion episodes of fall-back material on the newly formed central object (neutron star or black hole) in stripped-envelope supernova explosions (e.g., Dar \\& R\\'ujula~2000,2004; Dado et al.~ 2009). In this scenario, SN explosions that produce long GRBs (SNe-GRB) are ordinary core-collapse SNe of type Ic where the kinetic energy of the ejecta is typically a few $10^{51}$ ergs, rather than 'hypernovae' - hypothetical super energetic core collapse SNIc explosions, where the kinetic energy of the ejecta exceeds a few $10^{52}$ ergs and their bolometric light-curve is powered by the radioactive decay of $M\\gg 0.1\\, M_\\odot$ of $^{56}$Ni synthesized in the explosion (e.g., Iwamoto et al.~1998; Nakamura et al.~2001). So far, the CB model has been very successful in predicting/reproducing the main observed properties of long duration GRBs (e.g., their rate, location in star formation regions, association with core collapse SNe of type Ib/c, typical photon energy, multi-pulse structure, pulse shape and duration, spectral evolution of the individual pulses and large photon polarization) and the observed correlations between them (see, e.g., Dado et al.~2009; Dado \\& Dar 2012, and references therein). The CB model scenario implies that in SNeIc the fast SN ejecta may collide with a slowly expanding massive CS shell ejeted some time before the explosion. Such a collision may produce a very luminous SNIc and even super-luminous (SLSN), which are powered mainly by the collision rather than by a large mass of $^{56}$Ni synthesized in the explosion. Because of relativistic beaming, most of the GRBs that are produced in SNeIc and SLSNe are beamed away from Earth and are not observed. Indeed, most SNeIc and SLSNe are not accompanied by an observed GRB. The first discovered SLSN without an associated GRB was SN1999as (Knop et al.~1999) at a redshift z=0.127, which was much more luminous than the very bright SNe of type Ib/c that produced observed GRBs such as SN1998bw (Galama et al.~1998), that produced GRB 980425 (Soffitta et al.~1998, Pian et al.~2000), SN2003dh (Stanek et al.~2003; Hjorth et al.~2003) that produced GRB 030329 (Vanderspek et al.~2003), and SN2006aj (Campana et al.~2006; Pian et al.~2006; Sollerman et al.~2006) that produced GRB 060218 (Cusumano et al.~2006). More recently, transient surveys that were monitoring many square degrees of the sky every few nights have discovered several additional SLSNe in the nearby Universe without an observed GRB. The first one was SN2005ap (Quimby et al.~2007). The absence of hydrogen in its spectrum, its very broad lines, and its energetics led Quimby et al.~ (2007) to propose that SN2005ap could have been produced by the same mechanism that produces SNe with an observable GRB. Other discoveries of SLSNe without an observed GRB include SN2003ma behind the Large Magellanic Cloud at z =0.289 by the SuperMACHO microlensing survey (Rest et al.~2011), SN2006gy (Smith et al.~2008,2010), SN2007bi (Gal-Yam et al.~ 2009), SN2008am (Chatzopoulos et al.~2011), SN2010hy (Kodros et al.~ 2010; Vinko et al.~2010), SN2010gx (Pastorello et al.~2010), SN 2010jl (Stoll et al.~2011), and SCP 06F6 (Quimby et al.~2011). Alternative mechanisms which were invoked in order to explain the observed luminosity of SLSNe include:\\\\ I. Radioactive decay of large amounts (several $M_\\odot$) of radioactive $^{56}$Ni produced in pair-instability explosions of extremely massive stars (Rakavy \\& Shaviv~1967; Barkat et al.~1967; Heger \\& Woosley~2002; Waldman~2008; Gal-Yam et al.~2009; Yoshida \\& Umeda~2011), which are efficiently mixed in the SN ejecta. \\\\ II. Efficient conversion of kinetic energy of the SN ejecta into thermal energy in SN explosions inside optically thick winds (Falk \\& Arnett~1973,1977; Ofek et al.~2007; Smith \\& McCray~2007; Smith et al.~2010; Balberg \\& Loeb~2011; Chevalier \\& Irwin~ 2011,2012; Moriya et al.~2012; Chatzopoulos et al.~2012; Ginzburg \\& Balberg~2012; Ofek et al.~2012).\\\\ III. Collision(s) of the fast SN ejecta with slowly expanding dense circum-stellar shell(s) ejected by the progenitor star sometime before the SN explosion (Grassberg et al.~1971; Moriya et al.~2012), supplemented by energy release in the radioactive decay chain\\\\ ${\\rm^{56}Ni\\rightarrow^{56}Co\\rightarrow^{56}Fe}$ of $M(^{56}{\\rm Ni}) \\ll M_\\odot$ synthesized in the explosion. The rapid decay of the bolometric light-curve of SLSNe such as SN2010gx, and the very large mass, $\\sim 10\\, M_\\odot$, of $^{56}$Ni needed to explain its peak luminosity, however, indicate that the pair instability mechanism where a large mass of $^{56}$Ni is produced (scenario I) is unlikely to be its power source (Pastorello et al.~2010). In scenario II, the progenitor star explodes into a dense wind, and the strong shock that presumably explodes it breaks out into the wind. This strong shock is assumed to convert the kinetic energy which it imparts to the ejecta in SN explosions into internal thermal energy of the wind. Numerical simulations of core-collapse SNe, however, so far have not produced consistently strong enough shocks that can reproduce the observed SNe where the typical kinetic energy of the debris is a few $10^{51}$ ergs. But, by adjusting the wind parameters and the energy deposited in it, and by introducing many simplifying assumptions, Ginzburg and Balberg (2012) were able to calculate bolometric light curves for some SLSNe, which look like those observed. Scenario III has not been studied yet in detail with numerical hydrodynamical codes (Ginzburg \\& Balberg 2012). However, scenario III is strongly suggested by observations of SNn of type Ib/c and by the success of the CB model of GRBs in predicting the main observed properties of long GRBs produced in SNe of Type Ic. In this letter we use a simple analytical model based on only a few general assumptions to derive a closed form expression for the bolometric luminosity of SLSNe in scenario III. It involves only few adjustable parameters. We use it to demonstrate that collisions between the fast ejecta from core collapse SNe of types Ib/c and their massive circum-stellar shells, which were ejected in eruptions of their massive progenitors in the years preceeding their SN explosion, together with a modest amount of radioactive isotopes, which were synthesized in the SN explosion and deposited in the ejecta, can reproduce quite well the bolometric light-curves of both the supernovae that were observed in association with GRBs and those of the recently discovered SLSNe without an observed GRB. We conclude with a short discussion of the implications for SN explosions and GRBs. ", "conclusions": "Very large photospheric velocities were inferred from the early-time broad line spectra of SNe associated with an observed long GRB (SNe-GRB) such as SN1998bw, SN2003dh, and SN2003lw. Also unusual large quanities of $^{56}$Ni synthesized in these explosions were inferred from both their bolometric lightcurves and spectra (e.g., Nakamura et al. 2001; Mazzali et al. 2001, 2006). The unusual large value of the kinetic energy of the SN explosion that was inferred from the very large early-time photospheric velocities led Iwamoto et al. (1998) and the above authors to conclude that SNe-GRB belong to a class of hyper-energetic SNe (\"hypernovae\"), where the kinetic energy of the ejecta is typically $5\\times 10^{52}$ ergs, and the synthesized mass of $^{56}$Ni is $\\sim 0.5\\, M_\\odot$. However, the early time-photospheric velocity that was inferred, e.g., from the broad lines in the spectrum of SN1998bw (Patat et al. 2000) decreased by a factor 4 from $\\sim$ 40,000 km s$^{-1}$ to $\\sim 10,000$ km s$^{-1}$ within the first 30 days after the explosion. If these velocities were the bulk motion c.m. velocities of the whole SN ejecta, such a deceleration would have required collision with a mass $M\\sim 3\\,M_{ej}\\sim 30\\, M_\\odot$ enclosed within $R\\sim 5\\times 10^{15}$ cm. But, a typical ISM baryon density of 1 cm$^{-3}$, yields an ISM mass $M\\sim ~ 4 \\pi\\, m_p\\, R^3/3 \\sim 1\\times 10^{-10}\\, M_\\odot$ within such a radius, while a wind environment will have typically only $\\dot{M}\\, R/V \\lsim 10^{-3}\\, M_\\odot$ within $R\\sim 5\\times 10^{15}$ cm. Hence, either the observed velocity was only of thin photospheric layer of the SN shell which decelerated rapidly by collision with a massive wind/shell (while the mean velocity of the SN shell was its $\\sim 5000$ km/s) or the broad absorption lines were due to line broadenig e.g., by Compton scattering inside an optically thick the SN shell). In both cases the kinetic energy of the explosion should have been estimated from the late-time nebular velocity of $M_{ej}+ M_{css}$ rather than from modeling the early time photospheric velocity (Nakamura et al. 2001; Mazzali et al. 2001, 2006). The typical observed expansion velocities, ~5000 km/s, during the nebular phase of SNe associated with observed long duration GRB imply kinetic energy release of only $\\sim 2.5\\times 10^{51}\\,(M_{ej}+M_{css})/10 M_\\odot$ erg typical to ordinary SN explosions and do not support an \"hypernovae\" origin of SN-GRBs. Also the true values of the mass of $^{56}$Ni which were synthesized in SN1998bw and other SN-GRBs may be much smaller than inferred, e.g., by Nakamura et al.(2001) and Mazzali et al.(2001,2006). These large masses were inferred mainly from the peak-luminosity in the photospheric phase, while in our model, and in reality, a large part of it could be supplied by the collision between the SN shell and the CS shell. The inferred mass from the nebular phase is highly model dependent since it depends on the fraction of the radioactive energy release that is absorbed in the SN shell that depends on the unknown mass of the shell and its density distribution as function of time, and on the density distribution of $^{56}$Ni within it: The fraction of the total $\\gamma$-ray energy that is absorbed in the SN shell is $1-1/\\tau_\\gamma + e^{-\\tau_\\gamma}/\\tau_\\gamma$. For nearly-transparent SN shells, $\\tau_\\gamma\\ll 1$, and only a fraction $\\tau_\\gamma/2\\ll 1$ of the total $\\gamma$-ray energy is deposited in the SN shell. Hence, in the models of Mazzali and collaborators, a much larger mass of $^{56}$Ni was required in order to power the observed luminosity of SNe-GRB in both the photospheric and nebular phases. However, because the radius of of the SN shell expands like $R\\approx V_{ej}\\, t$, $\\tau_\\gamma$ decreases like $t^{-2}$. Neglecting other losses, the luminosity powered by the decay of $^{56}$Co must decline then like $t^{-2}\\, e^{-t/111.4\\,{\\rm d}}$. The observed bolometric light-curve of SN1998bw during the time interval 300-778 day, however, displayed an exponential decay consistent with that of $^{56}$Co without the $t^{-2}$ modulation. This is possible if either the CS shell is opaque to both the $\\gamma$-rays and the positrons from the decay of $^{56}$Co, or opaque only to the positrons. In our model the SN shell deposits its radio-isotopes at the bottom of a CS shell. If the lightcurve of SN1998bw is powered also by the SN shell collision with a CS shell, which is nearly opaque to both the $\\gamma$-rays and the positrons from the decay of $^{56}$Co it implies a rather small, $\\sim\\, 0.06\\, M_odot$ (i.e., normal). SLSNe are probably stripped-envelope SNe, most of which are SNeIc and SNeIIn (e.g., Pastorello et al.~2010; Quimby et al.~2011 and references therein). Their luminosity is powered mainly by plastic collision between their fast ejecta and slowly expanding massive circum-stellar shells formed in eruptions during the final stage of their life before their explosion. SLSNe may also produce GRBs, but like SNeIc-GRBs, most of them are not observed because they are beamed away from our line of sight to the SN explosion. The CS environments of core-collapse SNe provide evidence of massive winds and ejection episodes of massive shells, probably in thermonuclear eruptions preceding their SN explosion. Although it defies current paradigms of stellar evolution theory, perhaps the expulsion of a large fraction of the stellar mass by winds and thermonuclear eruptions preceding the SN explosion of massive stars make it possible for the energy deposition by the shock and the neutrinos from their collapsing core to unbind the left-over external mass and impart to it a kinetic energy of the order of $E_k\\sim $ several $10^{51}$ ergs." }, "1207/1207.6223_arXiv.txt": { "abstract": "We review the evidence behind recent claims of spatial variation in the fine structure constant deriving from observations of ionic absorption lines in the light from distant quasars. To this end we expand upon previous non-Bayesian analyses limited by the assumptions of an unbiased and strictly Normal distribution for the ``unexplained errors'' of the benchmark quasar dataset. Through the technique of reverse logistic regression we estimate and compare marginal likelihoods for three competing hypotheses---\\textsc{(i)} the null hypothesis (no cosmic variation), \\textsc{(ii)} the monopole hypothesis (a constant Earth-to-quasar offset), and \\textsc{(iii)} the monopole+dipole hypothesis (a cosmic variation manifest to the Earth-bound observer as a North--South divergence)---under a variety of candidate parametric forms for the unexplained error term. Our analysis reveals weak support for a skeptical interpretation in which the apparent dipole effect is driven solely by systematic errors of opposing sign inherent in measurements from the two telescopes employed to obtain these observations. Throughout we seek to exemplify a `best practice' approach to Bayesian model selection with prior-sensitivity analysis; in a companion paper we extend this methodology to a semi-parametric framework using the infinite-dimensional Dirichlet process. ", "introduction": "Recent claims by \\citet{web11} and \\citet{kin12} of a cosmic dipole signal in the fine structure constant deriving from their extensive compilation of Keck and VLT quasar absorption line measurements have been greeted with a healthy mix of excitement and skepticism by cosmologists at large. While undoubtably controversial with respect to the prevailing picture of a Copernican Universe, homogeneous in its composition and physical laws on the largest scales, the possibility of \\textit{some} late-time variation in the ``fundamental constants'' has been expressly identified within a number of well-studied cosmological theories. The touchstone for skeptical reaction to these claims was in fact the close alignment between the equator of the alleged dipole and the North--South divide between the sightlines of the two telescopes used to collect these data; the skeptical interpretation being that systematic errors of opposing sign are to blame for the apparent signal. Crucially, while the presence of an unexplained error term in the quasar dataset has been acknowledged by the Webb et al.\\ team, to-date all estimates of statistical significance for the dipole (i.e., those by the Webb et al.\\ team and their colleagues at UNSW; \\citealt{ber12}) have been computed under the assumption of strictly unbiased, Normal errors. Hence, the motivation for our Bayesian re-analysis of the dataset under a variety of biased and unbiased, Normal and non-Normal error models. The Bayesian model selection (BSM) framework used herein aims to identify the most plausible model to explain the observed data (with the hope of minimizing the posterior predictive error) from amongst a pre-defined set of hypotheses \\citep{ber96,kad04}. The quantitative basis for the BSM procedure is, characteristically, a ratio of marginal likelihoods \\citep{jef61,jay03}; the resulting ``Bayes factor'' operating much like an automatic Ockham's Razor favouring simplicity over complexity (cf.\\ \\citealt{jef92}). As a well-motivated, ``automatic'' procedure to distinguish between rival theories, given even limited or heterogeneous data, BSM has become highly popular in cosmological (and astronomical) research with novel applications \\citep{tro07,van09} now abounding in the literature and user-friendly software packages for marginal likelihood estimation \\citep{wei12x,fer13} readily available. However, with this powerful machinary at hand it can be all too easy to fall into the trap of na\\\"ively/lazily applying the BSM technique without due regard for its limitations, most notably the sensitivity of Bayes factors to the chosen priors on the internal parameters of each candidate model. Hence, in the present study we give particular attention to demonstrating the key elements of principled BSM with prior-sensitivity analysis \\citep{kas95,gel03}; in this endeavour we hope to provide a minimal template for future cosmological BSM studies. Though we restrict the present investigation to the usual case of parametric model selection, in a companion work (Paper II) we extend our methodology to a potentially more challenging semi-parametric formulation. The structure of this paper is as follows. In the remainder of the Introduction we give a detailed, historical overview of the observational evidence for and against a cosmic variation in the fine structure constant. In Section \\ref{quasardataset} we describe the Webb et al.\\ team's publically-available ``quasar dataset'' and give a preliminary investigation into the nature of the unexplained error term. In Section \\ref{generativemodels} we propose a number of candidate mathematical forms for the latter, explain our choice of prior densities on their governing parameters, and examine the resulting posterior distribution of each. In Section \\ref{rlr} we briefly review the reverse logistic regression procedure for marginal likelihood estimation, before proceeding to report and compare the latter for each hypothesis plus error model pairing (and to conduct our prior-sensitivity analysis) in Section \\ref{marginallikelihoods}. Finally, we summarize our conclusions and discuss the merits of possible follow-up observational strategies for resolving this debate in Section \\ref{conclusions}. \\subsection{The Fine Structure Constant in a Cosmological Context} Introduced in 1916 by Arnold Sommerfeld to abbreviate a recurring, dimensionless factor in his relativistic extension of the Bohr model for the hydrogen atom, the fine structure constant, $\\alpha$, is now recognized as one of the fundamental coupling terms of quantum electrodynamics (QED). In this context $\\alpha$ serves to characterize the strength of the electromagnetic interaction and has been humorously dubbed, ``the Peter Sellers of QED'' \\citep{bor86}, owing to the many different roles it plays within this theory: setting the scale for all electromagnetic cross-sections, binding energies, and decay rates; and, of course, the scale of fine structure splitting in atomic and ionic spectral lines (cf.\\ \\citealt{gri05}). Defined by the ratio of the squared elementary charge, $e^2$, to the permittivity of free space, $\\varepsilon_0$, the reduced Planck constant, $\\hbar$, and the vacuum speed of light, $c$, \\begin{equation} \\alpha = \\frac{e^2}{(4 \\pi \\varepsilon_0)\\hbar c}, \\end{equation} the fine structure constant has an on-Earth laboratory value, known to remarkable precision through exacting experimentation on quantum scale systems, of $\\alpha^{-1} \\approx 137.035999037(91)$ \\citep{bou11,han08}. Though nominally designated a ``constant'' there are various theoretical foundations to support a time-/space-varying $\\alpha$, if empirical evidence of such can be definitively established; most notably, within electrodynamic scalar field models \\citep{bek82,car98} and grand unification theories \\citep{mar84,bra03}. An historical evolution (with respect to cosmological time) of the fine structure constant driven by a decreasing speed of light could even offer a compelling solution to the infamous ``Horizon Problem''\\footnote{Namely, the remarkable homogeneity of the Universe beyond even the scale of casual connection under the standard model; that is, over physical separations exceeding the maximum distance traversable since the beginning of time at the speed of light.} of Big Bang cosmology without inflation \\citep{mof93,bar99,alb99}. At present there exist just a handful of different approaches, both terrestrial and astronomical, through which one may search for this experimentalist's ``Holy Grail''. These include: (\\textsc{i}) the in-laboratory comparison of optical atomic clocks \\citep{for07,ros08}; (\\textsc{ii}) the nuclear modelling of samarium isotopes extracted from the Oklo natural fission reactor \\citep{shl76,dam96,gou06} or rhenium extracted from Earth-fallen meteorite samples \\citep{oli04}; (\\textsc{iii}) the cosmological modelling of angular fluctuations in the temperature and polarization of the Cosmic Microwave Background \\citep{roc04,nak08,men09,gal10,cal11}; and/or (\\textsc{iv}) the identification of telltale frequency shifts in $\\alpha$-sensitive features of astronomical spectra, most notably the fine structure emission lines of ionized carbon observed in radio waves from distant lensed galaxies \\citep{lev12,wei12} and the ionic absorption lines of various species observed in the visible/near-visible light from distant quasars \\citep{web99,web01,mur01,mur03,web11,kin12}. After a number of initial disagreements between rival teams---see \\citet{lam04}, for instance---now resolved, only the post-Millennial quasar-based studies at present claim to have recovered any significant evidence for an evolving fine structure constant. First identified in 1966 \\citep{bur66,sto66} the distinctive absorption lines apparent to Earth-bound observers at UV-to-optical wavelengths in the light from distant quasars arise from the interaction of this light with ions of various species encountered during its passage through intervening gas clouds; these residing in the extensive halos of both bright, star-forming galaxies and dark proto-galaxies. The potential of these absorption lines (especially the singly and triply ionized silicon [Si\\textsc{ii} and Si\\textsc{iv}] doublets) as probes of the fundamental constants at extra-galactic distances, complementary to the already-used emission line technique \\citep{sav56,bah65}, was quickly realized by \\citet{bah67} and used to constrain the cosmic variation of $\\alpha$ to within $\\sim$10\\% of its on-Earth value ($\\Delta \\alpha / \\alpha = 0.98$[0.05])\\footnote{Throughout both the astronomical literature and our analysis herein the particular notation, $\\Delta \\alpha / \\alpha$, is used to represent the fractional offset of the fine structure constant in the extra-galactic system under study from its on-Earth laboratory value, i.e., $\\Delta \\alpha / \\alpha = [\\alpha - \\alpha_\\mathrm{Earth}]/\\alpha_\\mathrm{Earth}$.} in the direction of one particular quasar (3C 191). Improvements in astronomical instrumentation have since allowed far stronger constraints on $\\alpha$ variation to be established through this approach. \\citet{iva99}, for instance, have demonstrated $|\\Delta \\alpha/\\alpha| < 2.3 \\times 10^{-4}$ (95\\% CI) in the early Universe (at $\\sim$4 Gyr after the Big Bang; $\\sim$9.4 Gyr ago) through Si\\textsc{iv} doublet observations along nine quasar sightlines from a ground-based Russian telescope (the BTA-6 at the Special Astrophysical Observatory). To search for cosmic $\\alpha$ variation below this limit (of roughly one part in ten thousand) with quasar absorption line spectra requires application of the ``many multiplet'' (MM) technique \\citep{dzu99} in which the relative frequency shifts of \\textit{multiple} ionic species are simultaneously compared; the transitions of those species relatively insensitive to $\\alpha$ variation (e.g.\\ singly ionized magnesium [Mg\\textsc{ii}]) effectively serving as calibration benchmarks for the stronger shifters (e.g.\\ singly ionized iron [Fe\\textsc{ii}]). Implementation of the MM technique necessarily involves a multi-parameter fit to constrain not only $\\Delta \\alpha / \\alpha$ but also a suite of nuisance parameters accounting for the physical structure of the intervening gas cloud (i.e., column density, kinetic temperature, and the dominant line broadening mechanism). Pioneering this technique in 1999 the Webb et al.\\ team \\citep{web99} recovered tentative evidence that $\\alpha$ may have been lower in the past from a sample of 30 extra-galactic absorbers spanning $0.5 < z_\\mathrm{abs} < 1.6$ (here $z_\\mathrm{abs}$ denotes the cosmological redshift of the absorber(s) under study) observed with the Keck telescope ($\\Delta \\alpha /\\alpha = -1.1\\ [0.4] \\times 10^{-5}$). This exciting result was quickly heralded as support for a certain class of cosmological model admitting a decelerating speed of light \\citep{bar00,dav02}, though the claimed theoretical basis for the well-publicized Davies et al.\\ interpretation---derived from a (mistaken) consideration of black hole thermodynamics---was soon disproven \\citep{car03}. More interestingly from a statistical point of view was the Webb et al.\\ team's concern for a possible underestimation of the true errors in their quoted uncertainties, with a remarkably strong dip in their $\\alpha$ estimates over a narrow redshift interval ($0.9 \\lesssim z_\\mathrm{abs} \\lesssim 1.1$) suggesting the presence of an unexplained source of observational error. Further post-Millennial studies by the same team \\citep{web01,mur01,mur03} with the Keck telescope---expanding their original sample to a total of 141 absorbers and their redshift baseline to $0.2 < z_\\mathrm{abs} < 3.7$---ultimately strengthened the apparent weight of evidence for a time-varying fine structure constant beyond the ``4$\\sigma$ level''; given the particular assumptions of the statistical analysis employed. Chief amongst these the absence of any \\textit{bias} in the afore-mentioned unexplained error term---its presence still inferred from a marked excess of the sample variance over that expected from known sources of observational noise, seemingly greatest within the high redshift population. At this time \\citet{mur03} proposed a number of potential explanatory factors for an additional error term unique to high redshift systems---including the entry of damped Lyman-alpha absorbers (dense clouds of neutral hydrogen featuring complex velocity structures more challenging to model via the MM technique) into the sample at $z_\\mathrm{abs} \\gtrsim 2$ (where the rest-frame ultra-violet of their characteristic spectral lines is redshifted within the optical window accessible to on-Earth observers). An alternative hypothesis, the imprint of a spatial variation in $\\alpha$ (as we describe below), was discounted at this stage as only weakly supported by the available data (according to a bootstrap significance test); see \\citet{mur03}. It was thus a great surprise in 2010 when new $\\Delta \\alpha / \\alpha$ estimates for 154 (primarily Southern hemisphere) absorption systems along 60 quasar sightlines (52 new and 8 in common with the original Keck sample) derived by the same team \\citep{kin12,web11}---but this time using archival spectra from the Very Large Telescope (VLT) in Chile---appeared to indicate the opposite evolutionary trend with redshift. Namely, that the fine structure constant was in fact higher in the past for these absorbers ($\\Delta \\alpha /\\alpha = 0.154\\ [0.132] \\times 10^{-5}$). To resolve this contradiction the team were forced to resurrect the spatial variation hypothesis, proposing a smooth transition in $\\alpha$ across the Universe manifest to the on-Earth observer as a ($z$-invariant) monopole plus a ($z$-dependent\\footnote{\\citet{kin12} in fact consider a variety of candidate functional forms for their dipole model, including a (redshift) $z$-invariant dipole, a $z^\\beta$-dependent dipole, and an $r(z)$-dependent dipole (with $r(z)$ the cosmological lookback distance). All exhibit (bootstrap randomization-based) statistical significances of $\\sim$4$\\sigma$ over a monopole-only null. However, we note that: \\textsc{(i)} a $z$-invariant dipole implies a strong breaking of the Copernican principle---namely, that the Earth-bound observer does not occupy a privileged position within the cosmos; and \\textsc{(ii)} the $z^\\beta$- alternative (which encompasses a close approximation to the $r(z)$- model at $\\beta\\approx0.3$) already appears from the \\citet{kin12} study to be an over-fitting of the available data. Hence, we focus exclusively on the (monopole+)$r(z)$-dipole scenario in the present analysis. For reference, this was also the approach taken by \\citet{ber12}.}) dipole field. We illustrate graphically in Figure \\ref{kingmodel} the nature of the spatial variation in $\\Delta \\alpha / \\alpha$ under the best-fit model of this form (cf.\\ Section \\ref{generativemodels}) from \\citet{kin12} on a color-/symbol-coded map of the celestial sphere. Note the scale of the inferred variation, which is at the $\\Delta\\alpha/\\alpha \\lesssim 10^{-5}$ level only accessible (as noted earlier) via the MM technique. \\begin{figure*} \\vspace{0pc} \\begin{center} \\includegraphics[width=0.4\\textwidth]{fig1a} \\includegraphics[width=0.4\\textwidth]{fig1b}\\end{center} \\vspace{-0.275cm}\\caption{Visualization of the monopole+$r(z)$-dipole model (cf.\\ Section \\ref{generativemodels}) for explaining the apparent spatial variation of the fine structure constant proposed by the Webb et al.\\ team. The fractional difference of the fine structure constant from its on-Earth value under this model, $\\Delta \\alpha / \\alpha_\\mathrm{mod}$, as a function of the observational sightline for the best-fit solution of \\citet{kin12} is shown at $z=1$ in the lefthand panel and that at $z=3$ in the righthand panel (with the color-/symbol-coding explained in the top-right legend of each). The maximum, minimum, and equator of the key North--South dipole component of this model are overlaid as well for reference.} \\label{kingmodel} \\end{figure*} In Figure \\ref{kingdipole} we present a comparable visualization of the ``raw'' $\\Delta\\alpha/\\alpha$ variation in the Webb et al.\\ team's quasar dataset from which the apparent dipole effect was inferred. To this end we compute in each bin of right ascension and declination containing at least one quasar sightline (but typically two or more; noting as well that there are on average 2-3 absorbers per sightline) the weighted mean, \\begin{equation} \\overline{\\Delta \\alpha / \\alpha} = \\frac{\\sum (\\Delta \\alpha / \\alpha_i) / (\\sigma_{\\mathrm{obs},i}^2+\\sigma_{\\mathrm{sys},i}^2)}{ \\sum 1/(\\sigma_{\\mathrm{obs},i}^2+\\sigma_{\\mathrm{sys},i}^2)}.\\end{equation} Here $\\sigma_\\mathrm{obs}$ and $\\sigma_\\mathrm{sys}$ represent the standard deviations of the explained and unexplained error terms, respectively, in the Webb et al.\\ team's proposed generative model (which we detail in Section \\ref{quasardataset}). Despite the substantial degree of noise evident in these measurements (note the increase in scale with respect to that of Figure \\ref{kingmodel}) one may yet discern an excess of negative $\\Delta \\alpha / \\alpha$ estimates in the far North and an excess of positive $\\Delta \\alpha / \\alpha$ estimates in the far South as per the dipole hypothesis. Though, in anticipation of the skeptical interpretation of these results (i.e., that biases of opposite sign in the measurements from each telescope are to blame for the apparent North--South divergence), we note also that the equator of the alleged dipole provides a near perfect subdivision of the sample into its VLT and Keck constituents. \\begin{figure*} \\vspace{0pc} \\includegraphics[width=0.825\\textwidth]{fig2} \\vspace{+0.1cm}\\caption{Visualization of the apparent dipole signature in the \\citet{kin12} / \\citet{web11} quasar dataset. Each quasar sightline in the sample is marked here on a projection of the sky based on the J2000 equatorial coordinate system; with solid-symbols denoting VLT observations and open symbols Keck observations. Note that for most of these quasars there are multiple intervening absorbers providing independent measurements of $\\Delta \\alpha / \\alpha$ along the sightline. The apparent spatial variation of this observable is illustrated via a color-/symbol-coding (detailed in the top-right legend) of its weighted mean in each of the indicated subdivisions. The maximum, minimum, and equator corresponding to the best-fit monopole+$r(z)$-dipole solution of the \\citet{kin12} paper are overlaid as well for reference.} \\label{kingdipole} \\end{figure*} A ``4$\\sigma$ level'' significance for the spatial variation hypothesis was calculated by \\citet{kin12} through a bootstrap randomization of $\\Delta \\alpha / \\alpha$ estimates across sightlines and surveys against the simple monopole (i.e., time-/space-invariant with constant [non-zero] Earth-to-quasar offset) alternative---a result echoed by \\citet{ber12} in their review of the dataset using the Akaike Information Criterion, the $F$ statistic, and the ``error ellipsoid method''. With the new VLT sample apparently just as subject to unexplained measurement noise as the original Keck dataset, however, a number of strong assumptions were required to justify the significance testing procedures employed. In particular, it was considered necessary to treat the aforesaid as strictly Normal and strictly unbiased (viz.\\ zero mean and mode). Assessing the impact of these assumptions, which are easily relaxed within a Bayesian framework as we demonstrate herein, thus represents an essential ``next step'' in the analysis of this observational benchmark given the profound implications for both fundamental physics and cosmology if the dipole interpretation is to become accepted. Following the extensive world-wide media coverage of their work\\footnote{See, for example, the 23 October 2010 issue of New Scientist magazine (\\#2783), or the 2 September 2010 issue of The Economist magazine.} the Webb et al.\\ team's spatial variation hypothesis was strongly criticized in a number of public forums, including the popular science blogs, ``Uncertain Principles'' by \\href{}{Chad Orzel} and ``Cosmic Variance'' by \\href{}{Sean Carroll}. The former citing the evident alignment of the alleged dipole's equator with the coverage overlap of the Keck and VLT telescopes (cf.\\ Figure \\ref{kingdipole}) as an indicator that biases of opposite sign in the observations from each might well be to blame; and the latter citing the extraordinarily low mass-energy budget ($\\sim$$10^{-42}$ GeV) for the underlying scalar field implied by the cosmic expanse of the fitted dipole. From a Bayesian perspective these objections may be viewed as disagreements over the relative prior probabilities of competing proposals. Namely, the degree to which the particular model for the unexplained error term adopted by the experimenters should be favoured over some alternative model (or family of models), and the degree to which the null hypothesis (of a time-/space-invariant $\\alpha$) should be favoured ``theoretically'' over the proposed dipole hypothesis. ", "conclusions": "Although there exist a number of well-developed cosmological theories within which a time- and/or space-varying fine structure constant can be readily admitted \\citep{bek82,car98,mar84,bra03}, for many physicists the prior probability of such must be considered small given both our faith in certain long-standing physical principles and the null results of previous experiments designed to test these. In particular, past analyses of the Oklo natural fission reactor \\citep{shl76,dam96,gou06}, Earth-fallen meteorite samples \\citep{oli04}, and optical atomic clocks \\citep{for07,ros08} have strongly favoured a scenario of negligible \\textit{temporal} variation in $\\alpha$ locally; with \\citet{ros08}, for instance, determining $|\\Delta \\alpha/\\alpha|<1.6 \\times 10^{-17}$ per year at present on Earth. These results may ``daringly'' (with respect to the strong assumption of zero higher derivatives) be extrapolated to $|\\Delta \\alpha/\\alpha| < 2\\times10^{-7}$ since the Big Bang ($\\sim$13.4 Gyr ago); with more prosaic but reliable constraints of $|\\Delta \\alpha / \\alpha| < 0.1$ over this time provided by contemporary analyses of fluctuations in the Cosmic Microwave Background \\citep{roc04,nak08,men09,gal10,cal11}. \\textit{Spatial} variation in the fine structure constant, on the other hand, has been previously constrained to $|\\Delta \\alpha/\\alpha| < 10^{-4}$ on cosmic scales via both emission and absorption line based quasar studies \\citep{bah65,bah67,iva99}. Although the Webb et al.\\ team's claimed dipole strength contributes an effect well below this level the earlier absorption line results may nevertheless be supposed to have reinforced the prior beliefs of many astronomers (by way of the cosmological principle) that the physical laws of the Universe are spatially invariant. A related expectation that the Universe should appear essentially homogeneous when viewed on sufficiently large scales has also been hitherto well-supported by measurements of the fractal dimension (converging toward three) at large scales in galaxy redshift surveys \\citep{mar98,scr12} and by the isotropy of the Cosmic Microwave Background (\\citealt{man86}; though see \\citealt{eri07} and \\citealt{cla99}). The proposal of a large-scale spatial variation of the fine structure constant across the observable Universe runs into direct conflict with this established paradigm. To quantify our inherent preference for the strict null hypothesis over the dipole we might therefore, for argument's sake, suppose a prior log odds ratio of 5, such that $P_\\mathrm{prior}(\\mathrm{strict\\ null})/P_\\mathrm{prior}(\\mathrm{dipole}) \\approx 150$. In this conservative scenario, even the Bayes factor of $\\sim$$300$ in favour of the dipole under a strictly unbiased, Normal error model is reduced to a posterior odds ratio of just $2:1$, providing insufficient evidence to convince ourselves beyond doubt that the fine structure constant does indeed vary across the Universe in the manner described. Such a prior hypothesis weighting against the dipole also appears reflected in the relative caution with which both independent cosmologists \\citep{oli11} and the Webb et al.\\ team themselves have at times discussed their results. Moreover, as we have demonstrated herein (cf.\\ Section \\ref{marginallikelihoods}) the skeptical interpretation of null variation under a \\textit{biased} error model (in which the current $\\Delta \\alpha / \\alpha$ estimates from the Keck and VLT carry systematic biases of opposing sign) may in fact be preferred from a Bayesian model selection perspective. Thus new observations are undoubtably required if the Webb et al.\\ team or others are to convince a majority of cosmologists to accept the dipole hypothesis. To this end a dedicated campaign targetting quasars \\textit{at the poles of the inferred dipole} has been proposed \\citep{web11,kin12}, and is, in effect, already in progress with a first series of high resolution spectra suitable for $\\Delta \\alpha/\\alpha$ estimation recently obtained on the VLT under Large Program 185.A-0745 \\citep{mol13}. Assuming these measurements are subject to both explained and unexplained error terms consistent with those of the original quasar dataset---and supposing further the availability of an equivalent new Keck sample---one can easily forecast the power of such an experimental strategy (and potential alternatives) for settling the null vs.\\ dipole debate. One way to do this is via Monte Carlo simulation of future Bayes factors, assuming the dipole plus (unbiased) Normal error model pairing as truth. We perform such computational simulations here by repetition over the following procedure. First, we draw a $\\{\\bm{\\theta}_m,\\bm{\\theta}_e\\}$ pair of dipole hypothesis plus $\\textsc{\\bf{A}}_0$$\\textsc{\\bf{A}}_0$$\\textsc{\\bf{A}}_0$ error model parameter vectors from the current posterior. For each of $n_\\mathrm{new}$ proposed targets we draw a $\\Delta \\alpha/\\alpha_\\mathrm{obs}$ estimate accordingly with reference to its nominated telescope (i.e., $\\epsilon_\\mathrm{sys}$ group), sightline, and redshift, plus a $\\sigma_\\mathrm{obs}$ sampled (with replacement) from the current quasar dataset. We then compute the marginal log likelihoods of these mock observations under both the strict null and dipole hypotheses, proposing a biased error model for the former and an unbiased error model for the latter, and treating our current posteriors for each as priors. In the case of limited new data this computation can be efficiently and accurately performed via importance sampling from our earlier reverse logistic regression draws; that is, for $n_\\mathrm{new}$ small enough the updated posterior will not differ too markedly from the current posterior. For reference we also perform these simulations with unbiased error models assigned to both hypotheses. The distribution of future Bayes factors under a proposed observing strategy approximated in this manner can be termed the ``predicted posterior odds distribution'' (or PPOD; \\citealt{tro07}). \\begin{figure*} \\vspace{0pc} \\includegraphics[width=0.825\\textwidth]{fig10} \\vspace{-0.05cm}\\caption{The predicted power of future quasar observations for increasing the log Bayes factor in favour of the dipole hypothesis (supposed here as truth) under the nominal (unbiased, Normal) error model in the lefthand panel and under our default skeptical (biased, Normal) error model in the righthand panel. In each case we suppose the availability of 25 new $\\Delta\\alpha/\\alpha$ estimates from each of the VLT and Keck telescopes subject to both explained and unexplained error terms consistent with those of the original quasar dataset. Three alternative targetting strategies have been simulated: \\textsc{(i)} all targets placed at the current maximum a posteriori sightline of the pole (for the VLT) or anti-pole (for the Keck); \\textsc{(ii)} all targets placed on the dipole's equator; and \\textsc{(iii)} a compromise with ten measurements on the equator and the remainder on the pole or anti-pole (as appropriate). Note that the future log posterior odds ratio for each case should be considered the sum of the future log Bayes factor considered above with the current log Bayes factor from the original quasar dataset (lefthand panel: $+$5.7; righthand panel: $-$2.5) minus our log prior odds ratio of 5 against the dipole.} \\label{future} \\end{figure*} In Figure \\ref{future} we present the (log) PPOD for three basic targetting strategies, all supposing 25 detected absorbers with $\\Delta \\alpha / \\alpha$ measurements from observations on each of the Keck and VLT telescopes.\\footnote{The choice of 25 new absorbers from each telescope (i.e., $n_\\mathrm{new} = 50$) is consistent with expectations for the ongoing VLT Large Program \\citep{mol13}.} In the first, all targets are supposed placed at the current maximum a posteriori sightline of the pole (for the VLT) or anti-pole (for the Keck); in the second, all targets are supposed placed on the dipole's equator (accessible to both telescopes); and in the third a compromise is made with ten measurements on the equator and the remainder on the pole or anti-pole (as appropriate). As expected it is indeed the first strategy that gives the greatest chance of recovering the desired additional support of a further $+$5 on the log Bayes factor of the null vs.\\ dipole comparison with the unbiased error model applied to both. The mixed targetting strategy performs a little worse than the dipole-only case, while the equator-only strategy performs worst of all. With regards to choosing between the null and dipole when allowing for a biased error model on the former the performance of our three strategies is in fact reversed, such that the equatorial design offers the greatest chance of a conclusive outcome. An intuitive interpretation of this result is simply that for distinguishing between a biased and unbiased error model under two competing hypotheses the most effective design is to observe at locations where those underlying hypotheses are already in close agreement. Finally it is worth noting the wide range of predicted Bayes factors under all these possible observating strategies, which are so broad as to include a non-trivial possibility of rejecting the true hypothesis despite an optimal experimental design. This phenomenon arises principally from the broad range allowed for the strength parameter of the dipole, $B$, under the current posterior (see Figure \\ref{mcmcMB}), which gives non-negligible weight to the possibility that the dipole signal is far too small (with respect to the observational errors) to confidently recover from an $n_\\mathrm{new}=50$ sample. Thus, we reach the final conclusions of our Bayesian reanalysis of the quasar dataset. Namely, that: \\textsc{(i)} given both our incomplete understanding of the observational errors and our limited theoretical (prior) expectations regarding the properties of any spatial variation in the fine structure constant, the present observational coverage (featuring limited overlap of the two telescopes for which opposing biases might be suspected) must be deemed inadequate to properly distinguish the Webb et al.\\ team's proposed dipole field from the (strict or monopole) null; and \\textsc{(ii)} one cannot afford to overlook the importance of observations along the equator of the alleged dipole in additon to those proposed along its poles when planning future campaigns with the Keck and VLT telescopes for the purpose of settling this debate. We have also demonstrated in this study the utility of the reverse logistic regression technique for marginal likelihood computation with efficient prior-sensitivity analysis." }, "1207/1207.1129_arXiv.txt": { "abstract": "We present \\Chandra\\ ACIS-I X-ray observations of \\aJ1311 and \\bJ1653, the two brightest high Galactic latitude ($|b|>$10\\deg) $\\gamma$-ray sources from the 3 month \\Fermi-LAT bright source list that are still unidentified. Both were also detected previously by EGRET, and despite dedicated multi-wavelength follow-up, they are still not associated with established classes of $\\gamma$-ray emitters like pulsars or radio-loud active galactic nuclei. X-ray sources found in the ACIS-I fields of view are catalogued, and their basic properties are determined. These are discussed as candidate counterparts to \\aJ1311\\ and \\bJ1653, with particular emphasis on the brightest of the 9 and 13 \\Chandra\\ sources detected within the respective \\Fermi-LAT 95$\\%$ confidence regions. Further follow-up studies, including optical photometric and spectroscopic observations, are necessary to identify these X-ray candidate counterparts in order to ultimately reveal the nature of these enigmatic $\\gamma$-ray objects. ", "introduction": "} Since its launch in 2008, the \\Fermi\\ Large Area Telescope \\citep[LAT;][]{atw09} has enabled substantial progress in our understanding of the high-energy (HE; $>$100 MeV) $\\gamma$-ray Universe. A long-standing problem in the field has been in the secure identification of discrete HE $\\gamma$-ray sources, with most objects detected previously by COS B \\citep{swa81} and EGRET \\citep{har99,cas08}, remaining unidentified prior to the present \\Fermi-era. These sources have eluded identification mainly due to high source confusion in the poorly localized $\\gamma$-ray regions with typical 95$\\%$ confidence radii, \\r95 $\\sim 0.4\\deg-0.7\\deg$ in the 3rd EGRET catalog \\citep[3EG;][]{har99}. Also, for variable HE emitters \\citep[e.g.,][]{tav97,nol03}, there was a lack of prompt response to the $\\gamma$-ray events, with much of the multi-wavelength follow-up work pursued many years later \\citep[e.g.,][]{par08}. Multi-wavelength follow-up observations of unidentified 3EG sources attracted much effort, but met with mixed success \\citep[see][for a summary]{muk04}. The LAT's all-sky monitoring capability (20$\\%$ of the sky at all times), its increased sensitivity ($>$10$\\times$ better than EGRET), and dramatically improved localizations over EGRET, has enabled many of the unidentified EGRET sources to be successfully associated with lower-energy counterparts. The dominant $\\gamma$-ray emitting population consists of radio-loud active galactic nuclei (AGN), including blazars and radio galaxies. Perhaps unexpectedly, a substantial population of new $\\gamma$-ray pulsars have been identified ($\\gamma$-ray pulsations detected via a blind search or using known radio ephemerides), along with a handful of pulsar wind nebula, supernova remnants, and $\\gamma$-ray binaries. While the population of EGRET unidentified sources has quickly diminished, the number of fainter unidentified \\Fermi-LAT $\\gamma$-ray objects has been increasing \\citep[see][for a summary]{2fgl}. \\begin{table*} \\footnotesize \\begin{center} \\caption{} \\begin{tabular}{lccccccc} \\hline\\hline \\multicolumn{1}{c}{Name} & \\multicolumn{2}{c}{Pointing Center} & \\multicolumn{2}{c}{LAT Centroid} & \\multicolumn{1}{c}{ObsID} & \\multicolumn{1}{c}{Start Time} & \\multicolumn{1}{c}{Net Exp.} \\\\ \\cline{2-3}\\cline{4-5} \\multicolumn{1}{c}{}& \\multicolumn{1}{c}{R.A.} & \\multicolumn{1}{c}{Decl.} & \\multicolumn{1}{c}{$l$ (deg)} & \\multicolumn{1}{c}{$b$ (deg)} & \\multicolumn{1}{c}{} & \\multicolumn{1}{c}{in 2010 (UT)} & \\multicolumn{1}{c}{(ks)}\\\\ \\hline 0FGL~J1311.9$-$3419 & 13 11 49.68 &--34 29 31.2 & 307.686 & +28.195 & 11790 & Mar 21 15:38:54 & 19.87 \\\\ 0FGL~J1653.4$-$0200 & 16 53 43.68 &--01 58 30.0 & 16.593 & +24.931 & 11787 & Jan 24 06:21:27 & 20.77 \\\\ \\hline \\hline \\end{tabular} \\end{center} \\Chandra\\ observational summary for the two targets. The positions are the pointing centers (J2000.0 equinox) set at the time of the observations and the LAT centroids are the 2FGL catalog values in Galactic coordinates. The \\Chandra\\ observation ID (ObsID), start time, and net exposure (Net Exp.) are also provided. \\label{table-1} \\end{table*} As part of a systematic study of \\Fermi-LAT unidentified sources in X-rays, including \\Suzaku\\ \\citep[e.g.,][]{mae11,tak12a}, \\Swift\\ \\citep{fal11}, and \\XMM-Newton \\citep[e.g.,][]{wol10}, we obtained new \\Chandra\\ observations in cycle-11 covering the fields of five unidentified high Galactic latitude sources from the initial 3 month \\Fermi-LAT bright source list \\citep[0FGL;][]{bsl}. With its large field of view ($17' \\times 17'$, sufficient to cover the LAT 95$\\%$ confidence regions completely) and excellent sensitivity (5$\\sigma$ flux sensitivity of $\\sim 1.5 \\times 10^{-14}$ \\cgsflux, $0.5-8$ keV), \\Chandra\\ observations allow for a sensitive study of all X-ray sources within the LAT localization regions. For the three targets subsequently identified as pulsars, PSR~J2214+3002/\\cJ2214 \\citep{ran11} and PSR~J2241$-$5236/\\dJ2241 \\citep{kei11}, and a possible radio quiet millisecond pulsar (MSP) in \\eJ2339\\ \\citep{kon12}, we reported the results of our \\Chandra\\ observations in those papers. Here, we present the results of the \\Chandra\\ study of the remaining two objects (\\aJ1311\\ and \\bJ1653)\\footnote{Throughout, we retain the 0FGL names although newer information, particularly from the 2FGL release, are used. For reference, the various \\Fermi-LAT catalog names for the sources studied in this paper are: 0FGL~J1311.9$-$3419 = 1FGL~J1311.7$-$3429 = 2FGL~J1311.7$-$3429, and 0FGL~J1653.4$-$0200 = 1FGL~J1653.6$-$0158 = 2FGL~ J1653.6$-$0159.}. These \\Fermi-LAT sources were detected previously by EGRET, as 3EG~J1314$-$3431/EGR~J1314$-$3417 and 3EG~J1652$-$0223/EGR~J1653$-$0249 \\citep{har99,cas08}, and are two of the brightest remaining unidentified sources from that era. In the following, we describe the analysis of the \\Chandra\\ observations in Sec.~2, including the X-ray detection and characterization methods (Sec.~\\ref{sec-fields}) and a more detailed analysis and discussion of individual X-ray sources found within and near the LAT localization regions (Sec.~\\ref{sec-bright}). Results from a search for positional matches with archival optical, near-infrared, mid-infrared, and radio catalogs are also summarized in Sec.~\\ref{sec-bright}. We then discuss the general population of X-ray sources as potential counterparts to the $\\gamma$-ray objects with particular emphasis on the aforementioned brightest \\Chandra\\ sources (Sec.~3), and conclude with a summary of the results (Sec.~4). ", "conclusions": "} The $\\gamma$-ray sources studied, \\aJ1311\\ and \\bJ1653, are the two brightest unidentified \\Fermi-LAT sources found at high Galactic latitudes. In fact, both have been detected by EGRET (Sec.~1), and the improved localizations provided now by the \\Fermi-LAT with \\r95 $\\simeq~2.0' - 3.6'$ \\citep[compared to the corresponding EGRET values $\\simeq~34' - 44'$;][]{har99} allow us to address the counterparts with more certainty. From an X-ray perspective, our \\Chandra\\ observations detected sources down to a $0.5-8$ keV flux threshold of $\\sim (0.2-0.3) \\times 10^{-14}$ \\cgsflux\\ (Table~\\ref{table-2}). This is much improved over typical flux limits of $\\sim 10^{-13}$ \\cgsflux\\ achieved in all-sky surveys like the RASS \\citep[0.1$-$2.4 keV;][]{vog99,vog00}. Also, existing pointed \\Swift\\ and \\Suzaku\\ observations (Sec.~\\ref{sec-bright}) detected only the single brightest source within the LAT error ellipses, compared to the 9 and 13 \\Chandra\\ detected X-ray sources that can now be considered as potential counterparts to the LAT $\\gamma$-ray source. Due to their high Galactic latitudes, the most obvious candidate counterparts to these unidentified $\\gamma$-ray sources would be extragalactic objects, specifically, blazars fainter than currently catalogued. The LAT error circles of the high latitude objects have in fact been searched for blazars down to a radio flux limit of $\\sim$30 mJy \\citep{lbas}, but even fainter blazars are now being found in large numbers, e.g., from cross-correlations of the SDSS/RASS/FIRST databases \\citep[e.g.,][]{plo08}. In this context however, we found that none of the detected X-ray sources in either field had radio counterparts in the NVSS catalog (Sec.~\\ref{sec-bright}). In fact, only one (extended) radio source was found within the 95$\\%$ LAT error ellipse of \\bJ1653\\ (below), with none detected within the \\aJ1311\\ localization. The NVSS flux limit of $\\sim$2.5 mJy at 1.4 GHz is $\\sim 10 \\times$ fainter than the faintest typical radio sources currently associated with \\Fermi-LAT $\\gamma$-ray blazars \\citep{2lac,radgamma}. This absence of radio sources in general, and the lack of point source counterparts to the X-ray detected objects specifically, allow us to rule out faint blazars as possible counterparts of these sources. Radio-quiet AGN are not currently known to be $\\gamma$-ray emitters, except for a few examples of nearby galaxies with substantial starburst contributions \\citep{2lac,len10,ten11,latseyferts}, so even if the \\Chandra\\ detected X-ray sources are confirmed to be AGN, they are likely unrelated to the $\\gamma$-ray source. From a $\\gamma$-ray perspective, the faintest radio blazars tend to have harder HE $\\gamma$-ray spectra ($\\Gamma <2$), sometimes extending into TeV energies. Unlike these blazars, the two $\\gamma$-ray sources discussed are characterized by soft $0.1-100$ GeV spectra, being parameterized with single power-laws in the 1FGL catalog analysis \\citep{1fgl} with slopes, $\\Gamma = 2.25 \\pm 0.05$ (\\aJ1311) and $2.29 \\pm 0.06$ (\\bJ1653). In fact, a more detailed analysis of the longer 2 year LAT dataset presented in the 2FGL catalog indicated the HE spectra were best characterized as log-parabolas \\citep{2fgl}\\footnote{http://heasarc.gsfc.nasa.gov/FTP/fermi/data/lat/catalogs/ source/lightcurves/2FGL$\\_$J1311d7m3429$\\_$spec.png http://heasarc.gsfc.nasa.gov/FTP/fermi/data/lat/catalogs/ source/lightcurves/2FGL$\\_$J1653d6m0159$\\_$spec.png \\label{footnote-spec}}, not typically observed in $\\gamma$-ray blazars \\citep{lbasspectra}. Moreover, unlike typical bright $\\gamma$-ray emitting blazars, their variability indices of $17-19$ are below the 41.6 threshold in the 2FGL catalog analysis \\citep{2fgl}, indicating no significant $\\gamma$-ray variability within the 24 month dataset\\footnote{http://heasarc.gsfc.nasa.gov/FTP/fermi/data/lat/catalogs/ source/lightcurves/2FGL$\\_$J1311d7m3429$\\_$lc.png http://heasarc.gsfc.nasa.gov/FTP/fermi/data/lat/catalogs/ source/lightcurves/2FGL$\\_$J1653d6m0159$\\_$lc.png \\label{footnote-lc}}. Nearby radio galaxies have been found to be likely $\\gamma$-ray emitters \\citep[e.g.,][and references therein]{2lac} and it is possible that some fraction of the unidentified high Galactic latitude LAT sources could be unknown radio galaxies (i.e., `misaligned blazars') that could be faint X-ray emitters \\citep[e.g.,][]{can99}. In particular, steady $\\gamma$-ray emission from radio lobes is possible \\citep{che07}, as was observed in the nearby radio galaxy Centaurus~A \\citep{cena} and possibly in NGC~6251 \\citep{tak12b}. In this context, the single radio source (NVSS~J165348.44$-$015958.7; 11.5 $\\pm$ 1.7 mJy at 1.4 GHz) found within the \\bJ1653\\ error ellipse is extended, with measured elliptical dimensions, major axis = 125.4\\arcsec, minor axis $<46.9\\arcsec$ (i.e., unresolved in this direction), at position angle = 36\\deg. Inspecting Figure~\\ref{image2}, we see that the \\Chandra\\ sources N58 and N57 (within the 2FGL ellipse) and N69 (outside the 2FGL ellipse) are located near the extended tips of this radio source. If any of these \\Chandra\\ sources are related to NVSS~J165348.44$-$015958.7, it could plausibly mark a radio outflow from the X-ray source. However, the lack of an optically bright extended giant elliptical galaxy counterpart to any of these X-ray sources eliminates the possibility that these are nearby radio galaxies. Relatedly, young radio sources are expected to be steady $\\gamma$-ray emitters \\citep[e.g.,][]{mcc11}, but such objects are typically bright compact cm-wavelength sources but no such sources were found in the NVSS image within the LAT error regions in the present cases. As AGN likely can be ruled out as potential counterparts, the remaining possibilities are open to debate. As both $\\gamma$-ray sources are located at mid Galactic latitudes ($|b| =10\\deg - 30\\deg$; Table~\\ref{table-1}), there is the interesting possibility of isolated neutron stars associated with the $\\gamma$-ray sources, as was posited for 3EG~J1835+5918 \\citep[][and references therein]{hal02}, and subsequently confirmed with \\Fermi-LAT observations of this \\citep{j1836} and other pulsars \\citep{psrcat}. Indeed, \\aJ1311\\ and \\bJ1653\\ are two of nine total high Galactic latitude ($|b|>10\\deg$) $\\gamma$-ray sources from the LAT bright source list \\citep[][]{bsl} that were unidentified at the time. The other seven sources have all since been identified as pulsar powered sources. Specifically, they were predominantly identified as radio/$\\gamma$-ray emitting MSPs -- PSR~J0614$-$3329, PSR~J1231$-$1411, PSR~J2214+3000 \\citep{ran11}, PSR~J2241$-$5236 \\citep{kei11}, and PSR~J2302+4442 \\citep{cog11} -- with one normal young pulsar \\citep[PSR~J2055+25;][]{saz10}. The two subjects of the present study have been similarly searched for pulsating radio emission with null results in the past \\citep[based on their EGRET localizations;][]{cra06} and in new searches of the \\Fermi\\ error circles \\citep{ran11}. In the remaining case (\\eJ2339), the brightest \\Chandra\\ X-ray source within the \\Fermi-LAT error ellipse was found to be a black widow-type MSP and is the likely counterpart of the $\\gamma$-ray source \\citep[][see also \\citet{rom11}]{kon12}. As in these other seven 0FGL cases, the LAT spectra of our remaining two unidentified objects display significant curvature$^{\\ref{footnote-spec}}$ and are steady $\\gamma$-ray emitters$^{\\ref{footnote-lc}}$, so may point to a pulsar origin for them as well. In our \\Chandra\\ observations of the MSPs identified with the LAT sources, PSR~J2214+3000 \\citep{ran11} and PSR J2241$-$5236 \\citep{kei11}, the MSPs were spatially coincident with the brightest X-ray sources within the LAT error ellipses. The X-ray spectral analysis indicated a thermal origin with best fit blackbody temperatures, $kT \\sim 0.2-0.3$ keV, with essentially no photons above 2 keV. Similar X-ray spectral results for the other identified pulsars were derived from \\XMM\\ observations of PSR~J2302+4442 \\citep{cog11} and \\Swift\\ observations of PSR~J0614$-$3329 and PSR~J1231$-$1411 \\citep{ran11}; see also past \\XMM\\ observations of the 0FGL~J0614.3$-$3330 case \\citep{lap06}. In fact, we found that several X-ray sources (Sec.~\\ref{sec-bright}) have soft spectra that appear thermal in origin. Taking the 2FGL \\citep{2fgl} observed $0.1-100$ GeV $\\gamma$-ray fluxes for \\aJ1311\\ ($F_{\\gamma} = 6.17 \\times 10^{-11}$ \\cgsflux) and \\bJ1653\\ ($F_{\\gamma} = 3.43 \\times 10^{-11}$ \\cgsflux), the \\Chandra\\ observations probe a range of X-ray\\footnote{To ease comparison with other works \\citep[see, e.g.,][]{mar11}, we extrapolated our $0.5-8$ keV X-ray fluxes ($F_{\\rm 0.5-8~keV}$) into $0.3-10$ keV X-ray ones ($F_{\\rm X}$) assuming, $F_{\\rm X} = 1.25 \\times F_{\\rm 0.5-8~keV}$.} to $\\gamma$-ray flux (or luminosity) ratios, $F_{\\rm X}/F_{\\gamma} \\simlt (2-8) \\times 10^{-3}$ corresponding to the brightest X-ray sources, and down to $\\sim (0.04 - 0.1) \\times 10^{-3}$ for the faintest ones. For a pulsar interpretation, this probes the conversion ranges of spin-down powers to (non-thermal or thermal) X-rays for possible neutron star candidates \\citep[e.g.,][]{mir00}. In contrast to the thermal X-ray sources discussed above, the brightest \\Chandra\\ sources in our two cases show X-ray spectra that extend to $\\sim 7-8$ keV and were best fit with single power-law spectra (Sec.~\\ref{sec-bright}). If these are the counterparts to the LAT sources, the objects are likely not normal pulsars or MSPs. The non-thermal X-ray spectra of these sources are similar to the case of CXOU~J233938.7$-$053305, the putative counterpart of the (formerly) unidentified high Galactic latitude source \\eJ2339\\ \\citep{kon12}. In the latter case, the spectrum is best characterized by an absorbed power-law model with $\\Gamma = 1.1$ and a $0.3-10$ keV X-ray flux of $3 \\times 10^{-13}$ \\cgsflux. Its black widow-like MSP nature was discovered after optical follow-up of the detected X-ray variable \\Chandra\\ source revealed a 4.63 hr period of the binary derived via optical photometric \\citep{kon12} and spectroscopic \\citep{rom11} monitoring. In our \\Chandra\\ observations, we found only one faint X-ray source (CXOU~J165337.2$-$020020, N37) within the \\bJ1653\\ error ellipse to be possibly variable, albeit with limited statistics. In the case of \\wJ1311a, the bright source within the 2FGL ellipse of \\aJ1311\\ that had a \\Suzaku\\ discovered short term flare (factor of 10 increase in the the first 20 ks of the $\\sim$100 ks observation span), we found no significant variability within our 20 ks \\Chandra\\ observation. However, these observations together with a \\Swift\\ snapshot, provided evidence for variability on months timescales also (Sec.~\\ref{sec-bright}). In this context, it may be fruitful to photometrically monitor the optical counterparts to these X-ray sources (Sec.~\\ref{sec-bright}) for similar variability as in the case of CXOU~J233938.7$-$053305. \\begin{figure*} \\begin{center} \\includegraphics[width=6.25cm,angle=-90]{figure4a.eps}\\includegraphics[width=6.25cm,angle=-90]{figure4b.eps} \\end{center} \\caption{\\Chandra\\ ACIS-I spectrum of the brightest X-ray source within the 2FGL error ellipse of \\aJ1311 (\\wJ1311a, N35; left), and of the bright X-ray source just outside (to the south of) the error ellipse (\\xJ1311b, N39; right). The top panels show the data points with the line indicating the best fit absorbed single power-law models and the bottom panels show the ratio between the data and the models. } \\label{spectrum1} \\end{figure*} \\begin{figure*} \\begin{center} \\includegraphics[width=6.25cm,angle=-90]{figure5a.eps}\\includegraphics[width=6.25cm,angle=-90]{figure5b.eps} \\end{center} \\caption{\\Chandra\\ ACIS-I spectrum of the brightest (\\yJ1653a, N38; left) and the second brightest (\\zJ1653b, N44; right) X-ray sources within the 2FGL error ellipse of \\bJ1653. The top panels show the data points with the line indicating the best fit absorbed single power-law and blackbody models, respectively, and the bottom panels show the ratio between the data and the models. } \\label{spectrum2} \\end{figure*} } \\Chandra\\ observations of the two brightest unidentified high Galactic latitude $\\gamma$-ray sources from the 3 month \\Fermi-LAT bright source list were obtained. Both sources were previously detected by EGRET, and remain two of the most enigmatic $\\gamma$-ray sources to date. The basic X-ray properties of all sources observed in the ACIS-I FOV were determined, with 9 to 13 X-ray sources detected within the respective \\Fermi-LAT error ellipses. Using existing catalogs (USNO~B1.0, 2MASS, WISE), the sub-arcsecond \\Chandra\\ positions enabled us to locate optical, near-infrared and mid-infrared counterparts to several of the X-ray sources. The ensemble of X-ray sources, focused primarily on the ones within the \\Fermi-LAT localizations, were further discussed as potential counterparts of the $\\gamma$-ray sources. Although existing all-sky optical/infrared catalogs returned only a handful of X-ray counterpart matches, our work enables future optical identifications if deeper images can be obtained. In the case of \\aJ1311, the brightest X-ray source (\\wJ1311a) within the \\Fermi-LAT error ellipse is the most credible counterpart. This source was detected in a previous \\Suzaku\\ observation with X-ray variability on sub-day timescales, and our newly presented \\Chandra\\ (and \\Swift) observations suggest variability on longer timescales. Together with the power-law nature of the X-ray spectrum, it appears similar to the case of CXOU~J233938.7$-$053305, a candidate black widow MSP that is the likely counterpart of a similar \\Fermi-LAT source, \\eJ2339\\ \\citep{kon12}. Another bright X-ray source, \\xJ1311b, found just outside the LAT error ellipse and also previously detected by \\Suzaku, did not show any significant X-ray variability, and is a less probable counterpart to the $\\gamma$-ray source. In the case of \\bJ1653, the brightest X-ray source (\\yJ1653a) within the \\Fermi-LAT localization also displays a non-thermal X-ray spectrum from our \\Chandra\\ observation. However, there was no significant X-ray variability found within our 20 ks observation, and on the longer timescale probed by comparing to a \\Swift\\ observation obtained $\\sim$10 months earlier. A faint optical counterpart to this X-ray source is found, and optical photometric and spectroscopic monitoring may be fruitful to test if this source is similar to the case of CXOU~J233938.7$-$053305. The second brightest \\Chandra\\ source (CXOU~J165341.4$-$015927) found within the 2FGL localization also has an optical/infrared counterpart, and has a likely thermal origin for the X-rays. This, and the other fainter X-ray sources revealed in our \\Chandra\\ observation, can be similarly investigated to ultimately probe the nature of this unidentified $\\gamma$-ray source." }, "1207/1207.3662.txt": { "abstract": "{}{}{}{}{} % 5 {} token are mandatory \\abstract % context heading (optional) % {} leave it empty if necessary {We present a theoretical calibration of a new metallicity diagnostic based on the \\strom index $m_1$ and on visual -- near-infrared (NIR) colors to estimate the global metal abundance of cluster and field dwarf stars. % aims heading (mandatory) % {XX} % methods heading (mandatory) To perform the metallicity calibration we adopt $\\alpha$-enhanced evolutionary models transformed into the observational plane by using atmosphere models computed adopting the same chemical mixture. % results heading (mandatory) We apply the new visual - NIR Metallicity--Index--Color (MIC) relations to two different samples of field dwarfs and we find that the difference between photometric estimates and spectroscopic measurements is on average smaller than 0.1 dex, with a dispersion smaller than $\\sigma$ = 0.3 dex. We apply the same MIC relations to a metal-poor (M~92) and a metal-rich (47~Tuc) globular cluster. We find a peak of -2.01$\\pm$0.08 ($\\sigma$ = 0.30 dex) and -0.47$\\pm$0.01 ($\\sigma$ = 0.42 dex), respectively.} % conclusions heading (optional), leave it empty if necessary % {XX} ", "introduction": "The intermediate-band \\strom photometric system \\citep{strom66} has, for stars with spectral types from A to G, several indisputable advantages when compared with broad-band photometric systems. {\\em i\\/}) The ability to provide robust estimates of intrinsic stellar parameters such as the metal abundance (the $m_1=(v-b)-(\\bmy)$ index, \\citealt[hereafter CA07]{twa00, hilker00,io07}), the surface gravity (the $c_1=(u-v)-(v-b)$ index), and the effective temperature (the $H_{\\beta}$ index, \\citealt{nissen88, ols88, twa00}). The $H_{\\beta}$ index is marginally affected by reddening, and therefore can also be compared to a simple color such as \\bmy\\ to provide individual estimates of reddening corrections \\citep{nissen91}. The same outcome applies to the reddening free $[c_1]$ index, and indeed theoretical and empirical evidence \\citep{ste91, nissen94, io05} suggests that a color such as \\umy\\ compared to $[c_1]$ --which is a temperature index for stars hotter than 8,500 K-- provides a robust reddening index for blue horizontal branch stars. {\\em ii\\/}) The use of the $m_1/c_1$ versus color plane can also be safely adopted to distinguish cluster and field stars \\citep{twa00, rey04, faria, aden09, arnadottir}. {\\em iii\\/}) Accurate \\strom photometry can also be adopted to constrain the ensemble properties of stellar populations in complex stellar systems like the Galactic bulge \\citep{feltz} and the disk \\citep{hay01}. {\\em iv)} The $v$ filter is strongly affected by two $CN$ molecular absorption bands ($\\lambda=4142$, $\\lambda=4215$ \\AA). Stars with an over-abundance of carbon ($C$) and/or nitrogen ($N$), i.e. $CH$- and/or $CN$-strong stars, will have, at fixed color, a larger $m_1$ value, a fundamental property for identifying stars with different $CNO$ abundances in Globular Clusters (GCs, CA07, \\citealt{io09, io11}). On the other hand, the \\strom system presents two relevant drawbacks. {\\em i\\/}) the $u$ and $v$ bands have short effective wavelengths, namely $\\lambda_{eff}\\,=\\,3450$ and $\\lambda_{eff}\\,=\\,4110\\,$\\AA. As a consequence the ability to perform accurate photometry with current CCD detectors is hampered by their reduced sensitivity in this wavelength region. {\\em ii\\/}) The intrinsic accuracy of the stellar parameters, estimated using \\strom indices, strongly depends on the accuracy of the absolute zero-point calibrations. This typically means a precision better than 0.03 mag. This precision could be easily accomplished in the era of photoelectric photometry, but it is not trivial effort in the modern age of CCDs. The calibration of \\strom photometric indices to obtain stellar metal abundances is not a new technique. Empirical calibrations based on such a method have been given by \\citet{strom64, bond70, craw75, nissen81}. In these works, most of the stars adopted to perform the calibration are nearby ($d \\lesssim$ 100 pc) F and early-type G dwarfs, with $\\feh >$ -0.8. The adopted samples include a significant fraction of young and intermediate-age disk stars, and a minority of low-mass, old stars. Moreover, these calibrations are based on differential indices $\\delta_{m1}$ and $\\delta_{c1}$, i.e. $\\delta_{m1} = m_{1,Hyades} (\\beta) - m_{1,star} (\\beta) $, where $\\beta$ stands for $H_{\\beta}$ and $m_{1,Hyades} (\\beta)$ is the standard relation between $m_1$ and $\\beta$ for the Hyades given by \\citet{craw75} and by \\citet{ols84}. The $\\delta_{m1}$ can also be defined with \\bmy as indipendent parameter, but as pointed out by Crawford, $\\delta_{m1} (b-y)$ is less sensitive to metallicity than $\\delta_{m1} \\beta$, because the \\bmy color is also affected by blanketing. \\begin{figure*} % \\centering \\includegraphics[width=16truecm, height=10truecm]{fig1_calamida.ps} \\caption{$F110W,\\ (y-F110W)$ CMD for a metal-poor --M92 (left)-- and a metal-rich --47~Tuc (right)-- globular. The green and red solid lines display cluster isochrones. The adopted chemical compositions, age, true distance modulus and reddening are labeled. Tracks were computed by assuming $\\alpha-$enhanced chemical mixtures (PI06) and transformed into the observational plane by adopting atmosphere models with the same $\\alpha-$enhancement.} \\label{fig1} \\end{figure*} The differential indices $\\delta_{m1}$ and $\\delta_{c1}$ have the advantage that they measure mostly metallicity and surface gravity, respectively, and are free of temperature effects. However, they are affected by the uncertainty on the photometric zero-point of the Hyades standard relations and require accurate $\\beta$ photometry. \\citet{ols84} derived a metal abundance calibration using high dispersion spectroscopic measurements of \\feh from \\citet{cayrel83} and new \\strom photometry for a sample of F, G and K dwarfs \\citep{ols83}. Olsen provided a linear solution for $\\delta_{m1} (b-y)$ in the range -0.8 $ < \\feh <$ 0.4 for F-G0 dwarfs, and a parabolic solution in the range -2.6 $ < \\feh <$ 0.4 for G0-K1 dwarfs. However, only one calibrating star has $\\feh <$ -1.9 and only three have $\\feh <$ -1.5. \\citet[hereafter SN89]{schu89} performed new intrinsic color and metallicity calibrations based on a sample of 711 high-velocity and metal-poor stars with \\strom photometry from \\citet{schu88}. The stars have been selected to have spectral types in the range F0-K5, surface gravity 3.4 $ < log(g) < $ 5.4 and \\feh abundances are based on high-resolution spectra of \\citet{cayrel83, cayrel85, francois86}. In addition, these stars have $uvby$ photometry in the system of \\citet{ols83, ols84}. The SN89 metallicity calibrations are based on the indices $m_1$ and $c_1$ and they are valid in the color range 0.22 $< (b-y)_0 <$ 0.59 mag and are reddening dependent. \\citet{hay02} claimed that systematic discrepancies of $\\sim$ -0.1/0.3 dex affected the SN89 photometric metallicity determinations of metal-rich stars in the quoted color range. They showed that this was a consequence of a mismatch between the standard sequence $m_1,\\ (b - y)$ of the Hyades used by SN89 to calibrate their metallicity scale, and the Olsen system. A new calibration was proposed by \\citet{hay02}, on the basis of an enlarged spectroscopic data set, that makes the SN89 calibration applicable to the Olsen's photometric catalogs. \\citet{ramirez} proposed a new metallicity calibration for dwarf stars based on an updated spectroscopic catalog by \\citet{cayrel01} and the $m_1$ and $c_1$ indices, with different equations for three $b-y$ color ranges (0.19 $< \\bmy <$ 0.35, 0.35 $< \\bmy <$ 0.50, 0.50 $< \\bmy <$ 0.80 mag). These relations are valid over a broad metallicity range (-2.5 $ < \\feh < $ 0.4) and for $log(g) > $ 3.4. \\citet[hereafter AR10]{arnadottir}, based on their new compilation of dwarf stars, tested some of the most recent and/or more popular metallicity calibrations \\citep[SN89]{ols84,hay02,ramirez}. They found that the calibrations by SN89 and by \\citet{ramirez} perform equally well, but the latter covers a larger parameter space. They suggested to adopt the calibration by \\citet{ramirez} and the calibration by \\citet{ols84} for dwarfs redder than $\\bmy =$ 0.8 mag. More recently, \\citet[hereafter CAS11]{casagrande} presented a new empirical metallicity calibration for dwarf stars based on the $m_1$ and $c_1$ indices, covering the metallicity range -2.0 $ \\lesssim \\feh \\lesssim$ 0.5 and 0.23 $< \\bmy <$ 0.63 mag. They validated the new calibrations by estimating the metallicity of field dwarfs and open clusters, and showed that they give reliable abundance estimates with a dispersion smaller than 0.1 dex. All these relations to estimate the metal abundance of dwarf stars are hampered by the presence of molecular $CN$,$CH$ and $NH$-bands that affect the \\strom $uvb$ filters, and in turn the global metallicity estimates. Moreover, the quoted relations include the $c_1$ index, that is based on the $u$ filter. As already mentioned, observations in the $u$ filter are very demanding concerning the telescope time and the photometry in this band is less accurate due to the reduced CCD sensitivity in this wavelength region. We derive, for the first time, a theoretical calibration of a metallicity diagnostic based on the $m_1$ index and on visual--near-infrared (NIR) colors for dwarf stars. The visual--NIR ($y$, $J,H,K$) colors adopted in our new Metallicity-Index-Color (MIC) relations have two clear advantages when compared with optical colors: {\\em i) \\/} they are not hampered by the presence of $CN$,$CH$ and $NH$-bands; {\\em ii) \\/} strong sensitivity to effective temperature. This means that the quoted molecular bands still affect the MIC relations, but only through the \\strom $m_1$ index. The new MIC relations are based on $\\alpha$-enhanced evolutionary models and $\\alpha$--enhanced bolometric corrections and color-temperature transformations and are valid in the metallicity range -2.5 $ < \\feh <$ 0.5, and for dwarf stars in the mass range 0.5 $< M < $ 0.85 $M_\\odot$ (4.5 $< log(g) < $ 5). The structure of the current paper is as follows. In \\S 2 we discuss in detail the photometric catalogs of Galactic Globular clusters (GGCs) adopted to validate evolutionary models in the visual--NIR colors. Section 3 deals with the approach adopted to calibrate the visual--NIR MIC relations, while in \\S 4 we present the different tests we performed to validate the current theoretical calibrations together with the comparison between photometric estimates and spectroscopic measurements of metal abundances. We summarize the results and briefly discuss further improvements and applications of the new MIC relations in \\S 5. %__________________________________________________________________ ", "conclusions": "We have presented a new theoretical metallicity calibration based on the $m_1$ index and on visual--NIR colors to estimate the global metal abundance of cluster and field dwarf stars. We adopted $\\alpha$-enhanced evolutionary models transformed into the observational plane by using atmosphere models with the same chemical mixture to derive the new MIC relations. This is the first time that visual--NIR colors are adopted to estimate photometric metallicities of dwarf stars. The main advantages of the new MIC relations are the following: {\\em i\\/}) the molecular bands $CH$, $CN$ and $NH$ affect the $m_1$ index, but the color indices are unaffected by their presence; {\\em ii\\/}) the metallicity sensitivity is larger in the metal-rich regime, i.e. for $\\mh \\gtrsim$ -1.0, where isochrones in the $m_1$ versus $y$ - NIR color planes are well separated compared to the same isochrones in the $m_1$ versus \\strom color planes. The sensitivity is larger also in the metal-poor regime, where isochrones overlap in the $m_1$ versus \\strom color planes at bluer colors, i.e. for $\\bmy \\le$ 0.4 mag; {\\em iii\\/}) the slopes of the MIC relations based on visual--NIR colors are on average shallower than the MIC relations based on \\strom colors. This means that the former indices have, at fixed $m_1$ value, a stronger temperature sensitivity; {\\em iv\\/}) they do not include the $u$ filter, which is very time consuming from the observational point of view. Moreover, \\strom photometry in the $u$-band is usually less accurate given the reduced sensitivity of CCDs in this wavelength region. In order to validate the new theoretical metallicity calibration we adopted two sample of field dwarfs with \\strom and NIR photometry and high-resolution spectroscopy available. The first sample includes 96 dwarfs selected from the study by AR10. The mean difference between photometric and spectroscopic abundance is $-0.02\\pm0.10$ dex, with a mean intrinsic dispersion of $\\sigma$ = 0.31 dex ($m_1,\\ y-J$, $m_1,\\ y-H$, $m_1,\\ y-K$ relations), while is $0.07\\pm0.06$ dex, with $\\sigma$ = 0.31 dex ($[m],\\ y-J$, $[m],\\ y-H$, $[m],\\ y-K$ relations). A further check has been performed by adopting 185 dwarfs selected from the study by CAS11. In this case the mean difference between photometric and spectroscopic abundance is $-0.06\\pm0.07$ dex, with a mean intrinsic dispersion of $\\sigma$ = 0.22 dex ($m_1,\\ y-J$, $m_1,\\ y-H$, $m_1,\\ y-K$ relations), while is $-0.01\\pm0.06$ dex, with $\\sigma$ = 0.25 dex($[m],\\ y-J$, $[m],\\ y-H$, $[m],\\ y-K$ relations). The quoted independent comparisons indicate that the new theoretical MIC relations provide accurate metal abundances for field dwarf stars with a dispersion smaller than 0.3 dex. We tested the calibration by adopting also MS stars of two GGCs covering a broad range in metal abundance, i.e. M~92 ($\\feh$ = - 2.31) and 47~Tuc ($\\feh$ = -0.72), for which both \\strom and WFC3 NIR photometry were available. We found that the metallicity distributions of 47~Tuc based on the new visual--NIR MIC relations are larger than suggested by spectroscopic measurements of both iron and $\\alpha$-element abundances. Most of the spread is given by photometric errors but a fraction of it, and in particular the asymmetry in the metallicity distribution might be caused by the occurrence of multiple populations in this cluster \\citep{milone12}. Unfortunately, the stars for which we estimated the metallicity do not have, to our knowledge, spectroscopic measurements of iron, $\\alpha$ and light elements. Therefore, we cannot constrain, on a quantitative basis, whether the large spread and the asymmetry in the metallicity distribution of 47~Tuc MS stars is caused by peculiar abundance patterns. According to current estimates photometric errors either in optical or in NIR magnitudes increase the spread in metallicity by at most for 0.30 dex. New medium and high-resolution spectra for the selected cluster MS stars can help us to shed new light on the culprit(s) causing the large spread in metal abundance. In the case of M~92 the spread of the metallicity distributions is given by the large photometric errors of the catalog at this magnitude level. We also plan to provide independent calibrations of the new visual-NIR diagnostic by using either semi-empirical CTRs for both \\strom \\citep{clem} and NIR colors and/or colors predicted by different sets of atmosphere models (PHOENIX, Hautschild et al. 1999a,b; MARCS, Gustafsson et al. 2008) or different sets of evolutionary models \\citep{dotter07,dotter08, girardi, vandenberg}." }, "1207/1207.0633_arXiv.txt": { "abstract": "Pulsar glitches are traditionally viewed as a manifestation of vortex dynamics associated with a neutron superfluid reservoir confined to the inner crust of the star. In this Letter we show that the non-dissipative entrainment coupling between the neutron superfluid and the nuclear lattice leads to a less mobile crust superfluid, effectively reducing the moment of inertia associated with the angular momentum reservoir. Combining the latest observational data for prolific glitching pulsars with theoretical results for the crust entrainment we find that the required superfluid reservoir exceeds that available in the crust. This challenges our understanding of the glitch phenomenon, and we discuss possible resolutions to the problem. ", "introduction": " ", "conclusions": "" }, "1207/1207.0296_arXiv.txt": { "abstract": "We calculate the observable signature of a black hole accretion disk with a gap or hole created by a secondary black hole embedded in the disk. We find that for an interesting range of parameters of black hole masses ($\\sim10^6$--$10^9\\ \\msun$), orbital separation ($\\sim1\\ \\units{AU}$ to $\\sim 0.1\\ \\units{pc}$), and gap width (10--190 disk scale heights), the missing thermal emission from a gap manifests itself in an observable decrement in the spectral energy distribution. We present observational diagnostics in terms of power-law forms that can be fit to line-free regions in AGN spectra or in fluxes from sequences of broad filters. Most interestingly, the change in slope in the broken power-law is almost entirely dependent on the width of gap in the accretion disk, which in turn is uniquely determined by mass ratio of the black holes, such that it scales roughly as $q^{5/12}$. Thus one can use spectral observations of the continuum of bright active galactic nuclei to infer not only the presence of a closely separated black hole binary but also the mass ratio. When the black hole merger opens an entire hole (or cavity) in the inner disk, the broad band SED of the AGN or quasar may serve as a diagnostic. Such sources should be especially luminous in optical bands but intrinsically faint in X-rays (i.e., not merely obscured). We briefly note that viable candidates may have already been identified, though extant detailed modeling of those with high quality data have not yet revealed an inner cavity. ", "introduction": "\\label{intro} The prevalence of supermassive black holes (SMBHs) at the centers of galaxies is well established \\citep{richstoneetal98}. In the context of a hierarchical merging universe, merging galaxies can lead to the merging of the black holes \\citep[BHs;][]{vhm03}. Mergers of BHs are important as one potential pathway for the growth of BHs and have even been suggested as the principal cause of the scaling relations between BH mass and host galaxy properties \\citep{2011ApJ...734...92J}. Mergers are also strong gravitational wave emitters, and the asymmetric emission of gravitational waves, especially from spinning BHs, can lead to a recoil of the merged BH large enough to kick it out of the host galaxy potentially influencing the BH occupation fraction of galaxies \\citep[e.g.,][]{2008MNRAS.384.1387V, 2008ApJ...682L..29B, merrittetal04}. Even if two galaxies with BHs merge and the BHs sink to the center of the merged galaxy, however, it is not a given that the BHs will merge within a Hubble time \\citep[e.g.,][]{2003ApJ...596..860M, 1980Natur.287..307B}. The search for precursors to BH mergers has naturally focused on active galactic nuclei (AGNs) that are either spatially resolvable at separations of $\\sim1\\,\\units{kpc}$ \\citep[e.g.,][]{2012ApJ...753...42C, 2012arXiv1201.1904B} or have spectroscopically distinct broad line regions at separations of $\\sim0.1\\,\\units{pc}$ \\citep[e.g.,][]{1996ApJ...464L.107G, 2009Natur.458...53B, 2010ApJ...716..866S, 2011arXiv1106.2952E}. One avenue for finding BH pairs at very small separations comes from an analogy with protoplanetary disks in which the presence of a planet may be inferred from its influence on the protoplanetary disk, particularly through gaps and holes in the disk carved out by the planet \\citep{1980ApJ...241..425G, 1996ApJ...460..832T, 2005astro.ph..7492A, 2008ApJ...682L.125E}. Here, we consider the observable consequences of an accretion disk gap caused by a secondary BH as has been considered before \\citep{1991MNRAS.250..505S, 2009ApJ...700.1952H, 2010MNRAS.407.2007C}. Owing to the complexity of AGN and quasar spectra, we have developed observational diagnostics in terms of power-law forms that can be fit to line-free regions in spectra, or potentially even to fluxes derived from sequences of broad filters. We have also given consideration to signatures that may be evident in broad-band SEDs of quasars and AGN that span the optical and X-ray regimes. In section \\ref{setup} we describe our assumptions regarding the accretion disk and secondary BH in which we are interested. Then we outline the conditions for a gap to be opened and its subsequent evolution in section \\ref{gapopening} and \\ref{gapevolution}, respectively. In section \\ref{gapsed} we present our calculations of the spectral energy distribution from a gapped accretion disk and the physical inferences that can be made from observations. In section \\ref{holesed} we describe the basic electromagnetic appearance of an accretion disk with a large inner hole and suggest some candidate sources. Finally, we discuss our results and future work in section \\ref{discuss}. ", "conclusions": "\\label{discuss} We have described the basic observational appearance of a BH with a gapped or holed accretion disk, both of which are observable through their broad-band SEDs. We have focused our consideration on parameter space that results in observable signatures between 2000\\ \\AA\\ and 2\\ ${\\mu}m$. At shorter wavelengths, absorption of ultraviolet photons will prevent unambiguous determination of details of the AGN continuum. At longer wavelengths, reprocessed emission from dust grains will similarly contaminate thermal emission from an accretion disk. In this range, however, it is possible to measure the continuum emission to the precision needed to infer the presence of a gap. Current and future large surveys such as SDSS, Pan-Starrs, and LSST may be able to exploit the observational signatures we have developed. The primary outcome of this paper is to delineate clear diagnostics of accretion disks with gaps and holes. To do so, we have made some simplifying assumptions that allow for the clearest picture of what will happen to the accretion disk SED, but there are, of course, some potential complications that will need to be considered in the future. One such complication is that we have treated the emission as only coming from the top of the accretion disk. Emission coming from the walls of the gap will complicate the signal and, depending on the vertical temperature structure, either increase or decrease the observability of the gap, potentially by a large amount given the possibility of dramatically increased $h/r$ at the outer boundary from gas pile up \\citep{2012arXiv1205.4714K}. Tidal features produced by the secondary as seen in simulations of protoplanetary disks will also need to be considered and may lead to additional corroborating observational signals. The periods of the secondaries that are observable are months to several years, allowing for a potential detection of a variable or modulating signal \\citep{2010AJ....140..642T, 2010ApJ...714..404T, 2012MNRAS.420..860S, 2012ApJ...755...51N}. One such periodic signature from an accreting secondary could be \\ion{Fe}{25}\\ and \\ion{Fe}{26} narrow emission lines that would move with the secondary. Such a measurement should be possible with the \\emph{Astro-H} X-ray observatory \\citep{2010SPIE.7732E..27T}. Another potential complication is that the population of observed quasar SEDs have a large variance \\citep[e.g.,][]{2006ApJS..166..470R, 2011ApJS..196....2S} that arises from a number of unknown causes not thought to be connected to the occurrence of any secondary black hole. For example, typical optical-to-X-ray spectral indices ($F_\\nu \\propto \\nu^{\\alpha_\\mathrm{ox}}$) can range from $\\alpha = -1.8$ to $-1.2$. Unambiguously discerning between low X-ray activity arising from quasars coming from the low-end of a natural range and low X-ray activity arising from quasars with a missing inner disk requires sufficiently deep X-ray observations. Regarding gapped accretion disks, our focus on local deviations from an overall power-law continuum emission mitigates the potential confusion from variance in SEDs arising from other causes. That is, it is possible to discern a dip in a power-law SED portion without \\emph{a priori} knowing what the power-law slope is. A full theoretical understanding of normal quasar continuum emission would, of course, allow one to model gaps in accretion disks with utmost fidelity. We thank Cole Miller and Mike Eracleous for extremely useful discussions as well as Zoltan Haiman, Bence Kocsis, Mateusz Ruszkowski, Alberto Sesana, and Taka Tanaka for thoughtful comments. We thank Jim Saborio and staff for their hospitality during which a large portion of this work was completed. KG acknowledges support provided by NASA through Chandra Awards GO0-11151X and G02-13111X and through Hubble Award HST-GO-12557.01-A awarded by the STScI. \\def\\href#1#2{#2}" }, "1207/1207.2770_arXiv.txt": { "abstract": "When the Sun ascends the red giant branch (RGB), its luminosity will increase and all the planets will receive much greater irradiation than they do now. Jupiter, in particular, might end up more highly irradiated than the hot Neptune GJ~436b and, hence, could appropriately be termed a ``hot Jupiter.'' When their stars go through the RGB or asymptotic giant branch (AGB) stages, many of the currently known Jupiter-mass planets in several-AU orbits will receive levels of irradiation comparable to the hot Jupiters, which will transiently increase their atmospheric temperatures to $\\sim$1000~K or more. Furthermore, massive planets around post-main-sequence stars could accrete a non-negligible amount of material from the enhanced stellar winds, thereby significantly altering their atmospheric chemistry as well as causing a significant accretion luminosity during the epochs of most intense stellar mass loss. Future generations of infrared observatories might be able to probe the thermal and chemical structure of such hot Jupiters' atmospheres. Finally, we argue that, unlike their main-sequence analogs (whose zonal winds are thought to be organized in only a few broad, planetary-scale jets), red-giant hot Jupiters should have multiple, narrow jets of zonal winds and efficient day-night redistribution. ", "introduction": "\\label{sec:intro} The ``hot Jupiter'' class of exoplanets was not generally anticipated prior to the discoveries of the first planets around main sequence stars \\citep{mayor+queloz1995, marcy+butler1996}. Their existence, however, could have been predicted long before then, since the Sun's luminosity will increase by a factor of several thousand as it ascends the red giant branch (RGB), thereby turning Jupiter into a hot Jupiter in several billion years. \\begin{figure*}[t!] \\plotonesc {f1.eps} \\caption{Orbital separations where red-giant hot Jupiters can be found around a solar-type star. The evolution of a 1-$M_\\sun$ star's radius (magenta curve) is shown as a function of age in gigayears (Gyr), and the consequent levels of irradiation as a function of stellar age and orbital distance (color scale). Green curves show regions of (age, orbital radius) space where an object would receive the same incident irradiation as known various known highly irradiated objects: GJ~436b, HD~189733b, HD~209458b, TrES-4, HAT-P-7b, and WASP-33b. The red dashed curve shows how Jupiter's orbit will evolve purely under the influence of stellar mass loss, with a red dot at the end showing the position achieved at the end of our stellar model. Note that, in order to more clearly show the final, very brief part of the star's evolution, the time axis is stretched by a factor of 40 during the ascent of the AGB phase (from 12.092~Gyr until the end of the model). The location where the time axis changes scale is marked with a thin blue line and hatch marks along the horizontal axis.} \\label{fig:irradiation1} \\end{figure*} Roughly 20\\% of the more than 700 currently known exoplanets\\footnote{See http://exoplanet.eu \\citep{schneider_et_al2011}, or see http://exoplanets.org \\citep{wright_et_al2011} for a differently vetted list.} have masses greater than half of Jupiter's, orbital radii greater than 1~AU, and will become hot Jupiters (i.e., for the present purposes, this means they will receive at least as much irradiation as the hot Neptune GJ~436b) before the end of stars' lives. Depending on the efficiency of tidal dissipation in RGB and AGB stars, many of these planets might eventually be tidally engulfed by their stars \\citep{carlberg_et_al2009, villaver+livio2009, nordhaus_et_al2010}, where they could play a role in shaping planetary nebulae (PNe) \\citep{soker_et_al1984, nordhaus+blackman2006, nordhaus_et_al2007} or creating highly magnetized white dwarfs \\citep{nordhaus_et_al2011}. Regardless of whether these planets are eventually swallowed by their stars, at some point in their futures they will be highly irradiated. Some searches for evidence of planets and other companions to post-main-sequence stars have already been undertaken. White dwarf atmospheres polluted by tidally shredded asteroids indicate the presence of distant planetary companions around $\\sim$30\\% of white dwarfs \\citep{zuckerman_et_al2010}. A post-common-envelope 50-$M_J$ (Jupiter-mass) object was found in a tight orbit around a white dwarf \\citep{maxted_et_al2006}. Tentative evidence has also suggested the existence of several other systems, including two small planets in tight orbits around a subdwarf star \\citep{charpinet_et_al2011}, and a jovian body around the pulsating white dwarf GD-66 \\citep{mullally_et_al2007, mullally_et_al2008, mullally_et_al2009} --- although recent data have complicated the planetary hypothesis for the latter system (\\citealt{farihi_et_al2012}; Hermes, private communication). Direct-imaging searches for warm companions to white dwarfs have yet to find any \\citep{hogan_et_al2009}, but have not yet surveyed large numbers of stars. Finally, \\citet{johnson_et_al2011} has found a number of giant planets around slightly evolved, subgiant stars (and some of these planets will soon become hot Jupiters). Here, we consider the properties of ``red-giant hot Jupiters'' (RGHJs), a term that we use somewhat loosely to refer generally to hot Jupiters around post-main-sequence stars that are on longer period orbits than their cousins around main-sequence-star. In \\S\\ref{sec:hot}, we calculate the range of orbital distances at which a gas giant planet around a post-main-sequence star might be considered a hot Jupiter, and we consider possible heating of a planet's bulk interior. In \\S\\ref{sec:acc}, we examine how the accretion of stellar wind (and perhaps rocky material) onto a RGHJ might pollute its atmosphere. In \\S\\ref{sec:chemspec}, we describe changes in Jupiter's chemistry and spectrum that will occur as the Sun evolves beyond the main sequence. In \\S\\ref{sec:motions}, we describe some qualitative differences between the wind patterns expected on RGHJs and on main-sequence hot Jupiters. In \\S\\ref{sec:conc}, we conclude and speculate on the observability of RGHJs. ", "conclusions": "\\label{sec:conc} It has long been appreciated that as a star expands and grows more luminous after its main-sequence lifetime, the orbital radii corresponding to a given equilibrium temperature move farther out from the star. \\citet{lopez_et_al2005} and \\citet{schroder+connonsmith2008}, for instance, considered the outward expansion of the habitable zone after a star leaves the main sequence, although whether life would be likely to have enough time to develop from abiotic conditions in the limited time available remains an open question \\citep{spiegel+turner2012}. A large number of known exoplanets are far enough from their stars that they are not currently strongly irradiated but they will be in the future. To investigate the properties that such planetary companions will have in the future, when their stars are much more luminous, we computed a suite of post-main-sequence stellar models to see how far out an object could be and still plausibly be called a hot Jupiter around a sufficiently evolved star. We found that, as far out as $\\sim$35~AU from a 3-$M_\\sun$ star, a companion might transiently be highly enough irradiated to merit the moniker ``hot Jupiter.'' We argued that massive planets might accrete an amount of stellar wind that is large compared with their own atmospheres, resulting in significant atmospheric pollution if the stellar wind's composition is different from the planet's, and perhaps creating the atmospheric signature of a CRP (even if the planet's bulk composition is not carbon-rich). We showed that an angular-Rhines-scale argument suggests that that RGHJs might be expected to have narrow-banded atmospheric jets, in contrast to main-sequence hot Jupiters. Detection of RGHJs might prove to be observationally challenging. The probability that a planet transits its star increases dramatically as the star's radius increases by a factor of 100 or more, but recognizing transits (or occultations) of RGHJs could prove nearly impossible, since the transit depth becomes tiny ($\\sim$10$^{-6}$) and the duration could last weeks to months \\citep{assef_et_al2009}. High-contrast imaging is a potential avenue for detecting photons from RGHJs, but the solar neighborhood is not rich in likely suitable targets and the required angular resolution to resolve companions to distant giant stars might be daunting, given present technology. It might actually be easier to find evidence of former RGHJs than of current RGHJs, perhaps by surveying white dwarfs for companions at several AU \\citep{gould+kilic2008}. Purely because of the timescale of photon diffusion, energetic arguments suggest that a companion should remain hot, after an AGB star fades into oblivion, for roughly as long as it had been highly irradiated (a few hundred million years). Such a reheated companion, perhaps identified by its infrared excess relative to the white dwarf's spectral energy distribution, might be mistaken for a more massive companion. However, this potential source of confusion is quantitatively significant only for very young white dwarfs, since the thermal relaxation timescale of a reheated planet is also of order a few hundred million years, short in comparison to a Hubble time. The searches for exoplanets in the last two decades have uncovered many surprises. Here, we investigated the properties of a category of giant planets that might be found soon and that have not been studied in detail previously. Hot Jupiters around post-main-sequence stars must exist in the Galaxy and represent a future evolutionary stage of many of the known jovian-mass planets, including our own Jupiter. \\vspace{0.5in} {\\sc Acknowledgments} We thank many people for useful discussions, in particular Rodrigo Fernandez, JJ Hermes, Ruobing Dong, Scott Tremaine, Ed Turner, Jason Nordhaus, Adam Burrows, Jeremy Goodman, Doug Lin, Subo Dong, Jay Farihi, Jonathan Mitchell, Adam Showman, Kristen Menou, Scott Gaudi, and John Johnson. We also thank an anonymous referee for helpful comments that materially improved the manuscript. DSS gratefully acknowledges support from NSF grant AST-0807444 and the Keck Fellowship. NM acknowledges support from the Yale Center for Astronomy and Astrophysics through the YCAA postdoctoral Fellowship." }, "1207/1207.2764_arXiv.txt": { "abstract": "We present results from the first combined study of variable stars and star formation history (SFH) of the Milky Way (MW) ``ultra-faint\" dwarf (UFD) galaxy Leo\\,T, based on F606W and F814W multi-epoch archive observations obtained with the Wide Field Planetary Camera 2 on board the Hubble Space Telescope. We have detected 14 variable stars in the galaxy. They include one fundamental-mode RR Lyrae star and 10 Anomalous Cepheids with periods shorter than 1 day, thus suggesting the occurrence of multiple star formation episodes in this UFD, of which one about 10 Gyr ago produced the RR Lyrae star. A new estimate of the distance to Leo~T of 409 $^{+29}_{-27}$ kpc (distance modulus of 23.06 $\\pm$ 0.15 mag) was derived from the galaxy's RR Lyrae star. Our $V, V-I$ color-magnitude diagram of Leo\\,T reaches $V \\sim$ 29 mag and shows features typical of a galaxy in transition between dwarf irregular and dwarf spheroidal types. A quantitative analysis of the star formation history, based on the comparison of the observed $V, V-I$ CMD with the expected distribution of stars for different evolutionary scenarios, confirms that Leo\\,T has a complex star formation history dominated by two enhanced periods about 1.5 and 9 Gyr ago, respectively. The distribution of stars and gas shows that the galaxy has a fairly asymmetric structure. ", "introduction": "Numerical simulations and semi-analytical models operating in the $\\Lambda$-cold-dark-matter ($\\Lambda$-CDM) scenario of galaxy formation (e.g. Bullock \\& Johnston 2005) suggest that the halos of large spirals like the Milky Way (MW) and the Andromeda (M31) galaxies could have entirely been built up by merging of disrupted satellites. Despite an intensive search, the surviving ``building blocks\" of this formation process so far have remained elusive. However, new hope to the quest has been triggered in the last few years, by the discovery of many faint dwarf galaxies in the outskirts and halos of the MW and M31 spirals. The census of the new MW companions counts so far 17 newly discovered dwarf galaxies, that were detected mainly from the analysis of the Sloan Digital Sky Survey (SDSS) data (see e.g. Belokurov et al. 2007, 2010; Watkins et al. 2009; Koposov et al. 2009, and references therein). Similarly, 12 dwarf galaxies were known to be M31 companions until 2004, of which only 6 are dSphs, but a wealth of 19 previously undetected new satellites were discovered in the last 5 - 6 years, by the panoramic surveys of the M31 halo carried out with the Isaac Newton and the Canada-France-Hawaii Telescopes (Ferguson et al. 2002; Ibata et al. 2007; McConnachie et al. 2009; Richardson et al. 2011) and, more recently, also by the SDSS (Slater et al. 2011; Bell et al. 2011). The new dSphs are fainter than those previously known, with typical surface brightness generally around 28 mag/arcsec$^2$ or less, hence they were named ``ultra-faint\" dwarfs (UFDs). With luminosities that reach as low as 10$^3$ L$_{\\odot}$, the UFDs provide the ultimate opportunity to test models of the formation and chemical enrichment of the first bound structures, and their implications for the formation of larger galaxies (e.g. Bovill \\& Ricotti 2011a,b and references therein; Tumlinson 2010). Fig. 1 shows the location of classical (bright) dSphs and UFDs in the absolute magnitude versus half-light radius (${{\\rm M_V} - \\log r_h}$) plane. The plot is an adapted and updated version of Belokurov et al.'s (2007) Figure 8. The MW and some of the M31 GCs are also shown in the plot, for comparison, as well as the GCs of NGC~5128 and the ultra compact dwarf ellipticals in the Virgo cluster. Likely due to an observational selection effect, the Andromeda's new satellites are generally brighter than the MW UFDs. Nevertheless, with faint luminosities typical of the bulk of globular clusters (GCs) and large spatial dimensions typical of dSphs, both the MW and the M31 new satellites sample a totally unexplored region of the ${{\\rm M_V} - \\log r_h}$ plane (see Fig.~\\ref{figbelokurov}). \\begin{figure*} \\begin{center} \\includegraphics[width=12cm,bb=45 185 535 665,clip]{f1.ps} \\caption{Absolute magnitude (${\\rm M_V}$) {\\it versus} half-light radius ($r_h$) for different objects. Lines of constant surface brightness are marked. Open circles and squares are the classical dSphs surrounding the MW and M31, respectively. Filled circles are the MW UFDs, and 3 extremely low luminosity GCs, discovered mainly from the analysis of the SDSS data (see e.g. Belokurov et al. 2007, 2010; Watkins et al. 2009; Koposov et al. 2009; and references therein). The Leo~T UFD is marked by a large filled star (red, in the electronic edition of the journal). Filled squares (grey) are the M31 dSph satellites discovered after 2004 (see, e.g., Richardson et al. 2011, and references therein; Slater et al. 2011; Bell et al. 2011). Open triangles are the Galactic GCs from Harris (1996), and Mackey et al. (2006). Small filled stars are the GCs of the Andromeda galaxy, from Federici et al. (2007). Small open stars are the GCs in the outer halo of M31, from Mackey et al. (2007) and Martin et al. (2006). Asterisks are the extended M31 GCs, from Huxor et al. (2005) with parameters measured by Mackey et al (2006). Filled circles (cyan) are the GCs of NGC~5128, from McLaughlin et al. (2008). Filled inverted triangles (magenta) are the nuclei of dwarf elliptical galaxies in the Virgo cluster, from Cot\\'e et al. (2006). Diamonds are the ultra-compact dwarfs in the Fornax cluster, from De Propris et al. (2005).} \\label{figbelokurov} \\end{center} \\end{figure*} The UFDs appear to be clustered in groups on the sky (see, e.g., Figure 1 of Richardson et al. 2011). Their velocity dispersions are small, 3-4 km/s, and they are quite extended (see Fig.~\\ref{figbelokurov}). The MW UFDs seem to be particularly dark matter dominated (see Table~1 of Wolf et al. 2010, and references therein) and their chemical abundances are extreme, with large dispersions, and stars as metal poor as [Fe/H]=$-4$ (see Tolstoy, Hill \\& Tosi, 2009 and references therein). Often they have irregular shape or are elongated, as Canes Venatici II, Segue I, Ursa Major II, Hercules (see Fig.5 of Belokurov et al. 2007; Fig. 11 of Mu{\\~n}oz et al. 2010; and Figs. 6 and 7 of Musella et al. 2012), as likely distorted by the tidal interaction with the MW. Some of them (Bootes II and Segue I) seem to be entangled in complex kinematic streams as the Sagittarius tails (see, map on V. Belokurov's web page: \\url{http://www.ast.cam.ac.uk/~vasily/sdss/field_of_streams/dr6/fos_dr6_marked_names.tif}), see Law \\& Majewski (2010) for an in depth discussion of this issue. All UFDs host an ancient population as old as about 10 Gyr, and generally have GC-like CMDs, resembling, although more dispersed, the CMDs of metal poor Galactic clusters like M92, M15 and M68 (Belokurov et al. 2007; Moretti et al. 2009; Musella et al. 2012; Brown et al. 2012). The only remarkable exception to this general behavior is represented by the Leo~T UFD (Irwin et al. 2007). Discovered as a stellar overdensity in the SDSS Data Release 5, the galaxy's CMD, based on follow-up observations with the 2.5 m Isaac Newton Telescope, % revealed that Leo~T is characterized by an intermediate-age stellar population with a metallicity of [Fe/H] $\\sim -1.6$ dex, together with a young population of blue stars of age $\\sim$ 200 Myr (Irwin et al. 2007). These authors estimated for the galaxy a distance modulus of $(m-M)_0 = 23.1$ mag (corresponding to a distance D= 417 kpc), an half-light radius $ r_h$ (Plummer\\footnote{Irwin et al.'s (2007) half-light radius is derived by fitting the radial profile of Leo~T with a standard Plummer law (Plummer 1911).})=1.4$^\\prime$, a surface brightness $\\mu_{0, V}$ (Plummer)= 26.9 mag/arcsec$^2$, and an integrated magnitude $M_{tot, V} = -$7.1 mag. According to these parameters the galaxy locates at the bright end of the MW UFD distribution in the ${{\\rm M_V} - \\log r_h}$ plane (see Fig.~\\ref{figbelokurov}), ranking second only to Canes Venatici~I (CVn~I), and in the transition region between the MW and the M31 UFDs. The deeper ($g \\lesssim$ 26.5 mag) CMD published by de Jong et al. (2008) confirms the presence in Leo~T of very young stars with ages between $\\sim$ 200 Myr and 1 Gyr, as well as an older stellar population ($>$5 Gyr and [Fe/H]$\\sim -1.7$ dex). The galaxy is embedded into an HI cloud with an heliocentric velocity of 38.6 km s$^{-1}$, a velocity dispersion of 6.9 km s$^{-1}$ and an estimated mass of 2.8 $\\times 10^5 M_{\\odot}$ (Ryan-Weber et al. 2008), or 4.3$\\times 10^5 M_{\\odot}$ (Grcevich \\& Putman 2009). Leo~T remains, so far, the only UFD found to contain a significant amount of neutral gas. The presence of HI above a certain column density is often associated with star formation, and indeed Leo T is the only UFD, and the lowest luminosity galaxy, with ongoing star formation known to date. The presence of gas also suggests that Leo~T may be affected by internal differential reddening. Simon \\& Geha (2007) obtained Keck DEIMOS spectra of 19 red giants in Leo~T for which they estimated a mean metal abundance of [Fe/H]=$-2.29 \\pm 0.10$ (on the Rutledge et al. 1997 metallicity scale) with a dispersion of $\\sigma_{\\rm [Fe/H]}$=0.35 dex, from the equivalent width of the Ca triplet absorption lines. Kirby et al. (2008) have re-analyzed Simon \\& Geha (2007) spectroscopic material for Leo~T by applying an automated spectral synthesis technique. They derived metallicities in the range of [Fe/H]=$-$0.12 to $-3.22$ dex, with an average value of $\\langle {\\rm [Fe/H]} \\rangle $ = $-2.02 \\pm 0.05$ dex and dispersion $\\sigma_{\\rm [Fe/H]}$=0.54 dex (where individual stellar [Fe/H] values were first weighted by the inverse square of the errors and then averaged to obtain the mean [Fe/H] value). This value was later revised to $\\langle {\\rm [Fe/H]} \\rangle $ = $-1.99 \\pm 0.05$ dex, $\\sigma_{\\rm [Fe/H]}$=0.52 dex, by Kirby et al. (2011). The spectroscopic mean metallicities are lower than the values adopted by Irwin et al. (2007) and de Jong et al. (2008). Simon \\& Geha (2007) also measured a mean stellar velocity of 38.1 $\\pm$ 2.0 km s$^{-1}$ with a dispersion of 7.5 $\\pm$ 1.6 km s$^{-1}$, in excellent agreement with the HI velocity and gas velocity dispersion of 6.9 km s$^{-1}$ measured by Ryan-Weber et al. (2008). From the velocity dispersion Simon \\& Geha (2007) infer a dark halo mass of $(8.2 \\pm 3.6) \\times 10^6 M_{\\odot}$ and a mass-to-light ratio of $138 \\pm 71 M_{\\odot}/L_{\\odot,V}$ for Leo T, while from the HI observations, Ryan-Weber et al. (2008) infer a lower limit for the total dynamical mass of $\\sim 3 \\times 10^6 M_{\\odot}$ and a mass-to-light ratio of $\\gtrsim 56 M_{\\odot}/L_{\\odot,V}$. Rocha et al. (2012) have estimated infall times ($t_{infall}$) for the MW satellites based on the Via Lactea II cosmological simulation (Diemand, Kuhlen \\& Madau 2007; Diemand et al. 2008; Kuhlen 2010). They find that Leo T was accreted much more recently ($t_{infall} \\leq$ 1 Gyr) than the other MW dSph and UFD satellites (see Figure 3 of Rocha et al.'s paper) and would be just falling into the Milky Way for the first time. This could explain why Leo~T managed to retain its gas and kept forming stars until a few hundreds Myr ago. As part of a project aimed at understanding the evolution of the UFDs, their connection with the MW and M31, and with the classical dwarfs and the GCs, we have already studied 7 MW UFDs (Bootes~I, Dall'Ora et al. 2006; CVn~I, Kuehn et al. 2008; Canes Venatici~II, Greco et al. 2008; Coma, Musella et al. 2009; Leo~IV, Moretti et al. 2009; Ursa Major~II, Dall'Ora et al. 2012; and Hercules, Musella et al. 2012), and in this paper we present results from the combined study of variable stars, star formation history (SFH), and spatial distribution of the resolved stellar populations in Leo~T. \\begin{figure*} \\includegraphics[width=12cm,clip]{f2.ps} \\caption{Map of the Leo~T UFD showing the gas distribution, according to Ryan-Weber et al. (2008) and the position of the field observed with the HST. Adapted from Ryan-Weber et al. (2008) Fig. 1.} \\label{figHImap} \\end{figure*} The properties of the Leo~T variable stars provide us information on the conditions at the epochs of formation of these variables (young ages: 50 -- 200 Myr, from short and intermediate-period Classical Cepheids $-$ CCs; intermediate-age: $\\sim$ 1 -- 2 Gyr, from the Anomalous Cepheids $-$ ACs; and old ages: t $>$ 10 Gyr, from the RR Lyrae stars) and, combined with deep CMDs of the stars resolved on the stacked images, have allowed us to perform a quantitative analysis of the SFH history occurred in Leo~T, using the synthetic CMD method (Cignoni \\& Tosi 2010, and references therein). Observations, data reduction and calibration of the Leo~T photometry are presented in Section 2. Results on the variable stars, the catalogue of light curves, and the distance to the galaxy they lead to are discussed in Section 3. The CMD of Leo\\,T and the galaxy SFH are presented in Sections 4 and 5, respectively. The comparison of the spatial distribution of the Leo~T gaseous and stellar components is discussed in Section 6. Finally, a summary of the main results is presented in Section 7. ", "conclusions": "We have studied the variable stars of the Leo~T UFD and performed a quantitative analysis of the galaxy's SFH. We have detected 14 variables in Leo~T, they include 1 fundamental-mode RR Lyrae star, 10 ACs, and two putative binaries. The average period of the RR Lyrae star ($P$=0.6027 d) suggests an Oosterhoff-Intermediate classification for Leo~T, similarly only to CVn~I, among the UFD galaxies. The magnitude difference between ACs and the RR Lyrae star suggests a metallicity lower than Z=0.0004, and more likely around 0.0002 ([Fe/H]=$-2.0$ dex) for the intermediate-age and the oldest stellar components in Leo~T. Adopting this metal abundance, a reddening value $E(B-V)$=0.03 mag, and an absolute visual magnitude for the RR Lyrae star of $M_V(RR)$ =0.44 mag (at [Fe/H]=$-$2.0 dex) the distance modulus inferred for Leo~T is $(m-M)_0=23.06 \\pm 0.15$ mag, in excellent agreement with the modulus obtained by fitting the First Overtone Blue Edge (FOBE) of the ACs' sample, $(m-M)_0=23.05 \\pm 0.10$ mag, and only slightly shorter than the value provided by the SFH analysis, $(m-M)_0=23.16$ mag. In spite of the low mass, Leo T underwent a complex SFH, with two major star forming episodes, about 7-9 Gyr ago and 1-2 Gyr ago, overimposed on a continuous star formation activity. The average star formation rate per pc$^2$ was around $9\\times 10^{-11}\\,M_{\\odot}\\,\\mathrm{yr}^{-1}\\,\\mathrm{pc}^{-2}$ allover the galaxy lifetime, with a range of variation between $\\approx 10^{-11}\\,M_{\\odot}\\,\\mathrm{yr}^{-1}\\,\\mathrm{pc}^{-2}$ and $\\approx 4\\times 10^{-10}\\,M_{\\odot}\\,\\mathrm{yr}^{-1}\\,\\mathrm{pc}^{-2}$. Overall, these results are in good agreement with previous determinations (e.g., de Jong et al. 2008, and Weisz et al. 2012), with the exception of our lower star formation rate earlier than 9 Gyr ago. The rather isolated location, and the suggestion that the galaxy would be just falling into the Milky Way for the first time (Rocha et al. 2012), might explain why Leo~T managed to retain its gas and kept forming stars until 1-2 Gyr ago, for which observational evidence is provided by our identification of a conspicuous population of AC variables, and even later (a few hundreds Myr ago), as the presence in the CMD of upper-MS and BL stars, and our reconstruction of the galaxy SFH clearly demonstrate. Leo~T's youngest stars (upper-MS and BL stars) are spatially confined in the North-East part of the galaxy (WF2 camera), whereas intermediate-age (ACs) to old populations (RR Lyrae and RGB stars) are evenly distributed across cameras WF3 and WF4, and almost absent in the WF2 camera. Similarly, the distribution of neutral hydrogen shows an asymmetric spatial distribution, with a low density protrusion extending to the North-East. The proximity of youngest stellar component and gas protrusion suggests they are coupled, leaving us with the open question of what triggered this common anomaly in such a rather isolated UFD. Observations extending over a much larger field of view will be mandatory to address this question. \\bigskip" }, "1207/1207.2552_arXiv.txt": { "abstract": "f(R) gravity is thought to be an alternative to dark energy which can explain the acceleration of the universe. It has been tested by different observations including type Ia supernovae (SNIa), the cosmic microwave background (CMB), the baryon acoustic oscillations (BAO) and so on. In this Letter, we use the Hubble constant independent ratio between two angular diameter distances $D=D_{ls}/D_s$ to constrain f(R) model in Palatini approach $f(R)=R-\\alpha H^2_0(-\\frac{R}{H^2_0})^\\beta$. These data are from various large systematic lensing surveys and lensing by galaxy clusters combined with X-ray observations. We also combine the lensing data with CMB and BAO, which gives a stringent constraint. The best-fit results are $(\\alpha,\\beta)=(-1.50,0.696)$ or $(\\Omega_m,\\beta)=(0.0734,0.696)$ using lensing data only. When combined with CMB and BAO, the best-fit results are $(\\alpha,\\beta)=(-3.75,0.0651)$ or $(\\Omega_m,\\beta)=(0.286,0.0651)$. If we further fix $\\beta=0$ (corresponding to $\\Lambda$CDM), the best-fit value for $\\alpha$ is $\\alpha$=$-4.84_{-0.68}^{+0.91}(1\\sigma)_{-0.98}^{+1.63}(2\\sigma)$ for the lensing analysis and $\\alpha$=$-4.35_{-0.16}^{+0.18}(1\\sigma)_{-0.25}^{+0.3}(2\\sigma)$ for the combined data, respectively. Our results show that $\\Lambda$CDM model is within 1$\\sigma$ range. ", "introduction": "$} One of the most striking things in modern cosmology is the universe undergoing an accelerated state \\cite{accelertion}. In order to explain this phenomenon, people have introduced new component which is known as dark energy. The simplest model is cosmological constant ($\\Lambda$CDM). It is consist with all kinds of observations while it indeed encounters the coincidence problem and the \"fine-tuning\" problem. Besides, there are many other dark energy models including holographic dark energy \\cite{holographic}, quintessence \\cite{quintessence}, quintom \\cite{quintom}, phantom \\cite{phantom}, generalized Chaplygin gas \\cite{GCG} and so on. Besides dark energy, the acceleration can be explained in other ways. If the new component with negative pressure does not exist, General Relativity (GR) should be modified. Until now, at least two effective theories have been proposed. One is considering the extra dimensions which is related to the brane-world cosmology \\cite{DGP}. The other is the so-called f(R) gravity \\cite{f(R)}. It changes the form of Einstein-Hilbert Lagrangian by f(R) expression. These theories can give an acceleration solution naturally without introducing dark energy. There are two kinds of forms about the f(R), the metric and the Palatini formalisms \\cite{formalisms}. They give different dynamical equations. They can be unified only in the case of linear action (GR). For the Palatini approach, the form $f(R)=R-\\alpha H^2_0(-\\frac{R}{H^2_0})^\\beta$ is chosen so that it can result in the radiation-dominated, matter-dominated and recent accelerating state. Furthermore, it can pass the solar system and has the correct Newtonian limit \\cite{Newtonian}. In this Letter, we consider the Palatini formalisms. Under this assumption, the f(R) cosmology has two parameters. What we want to emphasize is, among the parameters $(\\alpha,\\beta,\\Omega_m)$, only two of them are independent. Therefore, we can exhibit the constraint results on either $(\\alpha,\\beta)$ space or $(\\Omega_m,\\beta)$ space. Various observations have already been used to constrain f(R) gravity including SNIa, CMB, BAO, Hubble parameter (H(z)) and so on. Among these works, parameter $\\beta$ has been constrained to very small value. In these papers \\cite{constraint1}, they get $\\beta\\sim 10^{-1}$; in \\cite{constraint2}, the matter power spectrum from the SDSS gives $\\beta\\sim 10^{-5}$; in \\cite{constraint3}, the $\\beta$ was constrained to $\\sim 10^{-6}$. From these results, the f(R) gravity seems hard to be distinguished from the standard theory, where $\\beta=0$. One effective way to solve this problem in astronomy is combining different cosmological probes. Strong lensing has been used to study both cosmology \\cite{lensing1} and galaxies including their structure, formation and evolution \\cite{lensing2}. The observations of the images combined with lens models can give us the information about the ratio between two angular diameter distances, $D_{ls}$ and $D_s$. The former one is the distance between lens and source, the latter one is the distance from observer to the source. Because the angular diameter distance depends on cosmology, the $D_{ls}/D_s$ data can be used to constrain the parameters in f(R) gravity. In this Letter, we select 63 strong lensing systems from SLACS and LSD surveys assuming the singular isothermal sphere (SIS) model or the singular isothermal ellipsoid (SIE) model is right. Moreover, a sample of 10 giant arcs is also contained. Using these 73 data, we try to give a new approach to constraining f(R) gravity. This Letter is organized as follows. In Section 2, we briefly describe the basic theory about f(R) gravity and the corresponding cosmology. In Section 3, we introduce the lensing data we use, the CMB data and the BAO data. The constraint results are performed in Section 4. At last, we give a summary in Section 5. Throughout this work, the unit with light velocity $c=1$ is used. ", "conclusions": "$} In this Letter, we use $D_{ls}/D_s$ data from lensing systems to constrain f(R) gravity in Palatini approach $f(R)=R-\\alpha H^2_0(-\\frac{R}{H^2_0})^\\beta$. Compared with references, we can see the constraint effects that $D_{ls}/D_s$ data give can be compatible with other data (SNe Ia, H(z), BAO, CMB and so on). Moreover, we find although the best-fit values of the parameters are different from various observations, the directions of the contours in $(\\alpha,\\beta)$ space are very similar, thus needing different observations to break the degeneracy. The $D_{ls}/D_s$ data propose a new way to probe the cosmology \\cite{dlds}. As we expect, the lensing data alone cannot give a stringent constraint. There are at least three aspects that contribute to the error. First, the assumption that the lens galaxies satisfy SIS or SIE model may have some issues especially for four images. Second, the measurements of velocity dispersions have some uncertainties. Finally, the error exists due to the influence of line of sight mass contamination \\cite{sight}. Combining with CMB and BAO, it gives $\\beta\\sim 10^{-1}$, which contains the $\\Lambda$CDM model. Until now, we cannot distinguish it from the standard cosmology, where $\\beta=0$. For future lensing study, in order to improve the constraint, we hope large survey projects can find more strong lensing systems. At the same time, a better understand about the lens model and more precise measurements can give us more stringent results and more information about f(R) gravity. \\textbf{\\ Acknowledgments } This work was supported by the National Natural Science Foundation of China under the Distinguished Young Scholar Grant 10825313, the Ministry of Science and Technology national basic science Program (Project 973) under Grant No.2012CB821804, the Fundamental Research Funds for the Central Universities and Scientific Research Foundation of Beijing Normal University." }, "1207/1207.5975_arXiv.txt": { "abstract": "\\noindent A set of hydrodynamical models based on stellar evolutionary progenitors is used to study the nature of SN~2011dh. Our modeling suggests that a large progenitor star ---with $R \\sim 200$ $R_\\odot$---, is needed to reproduce the early light curve of SN~2011dh. This is consistent with the suggestion that the yellow super-giant star detected at the location of the SN in deep pre-explosion images is the progenitor star. From the main peak of the bolometric light curve and expansion velocities we constrain the mass of the ejecta to be $\\approx$ $2$ $M_\\odot$, the explosion energy to be $E= 6-10\\times 10^{50}$ erg, and the \\Ni\\, mass to be approximately $0.06$ $M_\\odot$. The progenitor star was composed of a helium core of 3 to 4 $M_\\odot$ and a thin hydrogen-rich envelope of $\\approx$ $0.1$ $M_\\odot$ with a main sequence mass estimated to be in the range of 12--15 $M_\\odot$. Our models rule out progenitors with helium-core masses larger than 8 $M_\\odot$, which correspond to $M_{\\mathrm{ZAMS}} \\gtrsim 25$ $M_\\odot$. This suggests that a single star evolutionary scenario for SN~2011dh is unlikely. ", "introduction": "\\label{sec:intro} Type IIb supernovae (SNe~IIb) are transitional objects within the family of core-collapse SNe (CCSNe), as their spectroscopic classification evolves from Type II (i.e. with H lines), to type Ib (i.e. dominated by helium lines). SNe~IIb were recognized as a new class with the discoveries of SN~1987K \\citep{1988AJ.....96.1941F} and SN~1993J \\citep[e.g.][]{1993Natur.364..507N,1997ARA&A..35..309F}. As the rest of CCSNe---which includes H-rich types II-P and II-L, H-less type Ib, and He-less type Ic---, they are believed to arise from the violent death of stars with initial masses greater than 8 $M_\\odot$. Massive stars may suffer considerable mass loss during their evolution, due to strong stellar winds or mass transfer to a binary companion. Therefore, the vast spectroscopic and photometric diversity observed among CCSNe is related to the ability of the progenitor star to retain its outermost layers before the explosion. Despite the efforts to improve our understanding of the progenitor of each subclass of CCSNe, many questions remain open. In this sense, the study of a very well-observed object can shed light on the nature of CCSNe and their massive progenitor systems. SN~2011dh was discovered on May 31.893 by amateur astronomers and immediately confirmed by the Palomar Transient Factory (PTF) in the nearby spiral galaxy M51 \\citep{2011CBET.2736....1G,2011ATel.3398....1S,2011ApJ...742L..18A} which had also hosted three other CCSNe in the past 17 years. Strong constraints on the date of explosion, to better than $0.6$ days were established using pre- and post-SN imaging \\citep{2011ApJ...742L..18A}. SN~2011dh was extensively monitored in a wide wavelength range, including early detections in radio and X-rays \\citep{2012ApJ...752...78S}. It was classified as type IIb \\citep{2011ApJ...742L..18A}, a relatively rare subclass of CCSNe, based on the optical spectrum. Soon after discovery, a source was identified as a possible progenitor of SN~2011dh in archival, multi-band HST images \\citep{2011ApJ...739L..37M,2011ApJ...741L..28V}. Photometry of the source is compatible with a yellow super-giant (YSG) star. The detection of pre-supernova (pre-SN) objects in high-resolution imaging has provided important information on the possible progenitors of several SNe. For SNe~II-P, these observations confirm the red supergiant nature of the progenitor, as previously suggested by theory. An upper limit of $M_\\mathrm{ZAMS}= 16.5 \\pm 1.5 \\, M_\\odot$ was suggested for this type of SNe \\citep{2009MNRAS.395.1409S}. However, such upper limit has been recently revised by \\citet{2012MNRAS.419.2054W}. They found a higher value for the maximum mass of SN~II-P progenitors of $M_\\mathrm{ZAMS}= 21^{+2}_{-1}\\, M_\\odot$ when additional extinction due to dust produced in the red supergiant wind is taken into account. For other subtypes of CCSNe, detections have not been as common and only in a few cases have the progenitors been conclusively identified, e.g the type II-pec SN~1987A \\citep{1987Natur.328..318G}, the type IIb SN~1993J \\citep{1994AJ....107..662A,2004Natur.427..129M}, and the type IIn SN~2005gl \\citep{2007ApJ...656..372G}. To determine the main-sequence mass ($M_\\mathrm{ZAMS}$) from pre-SN imaging, it is necessary to derive a luminosity and intrinsic colour from the photometry, and compare with some evolutionary track. This was done for SN~2011dh by \\citet{2011ApJ...739L..37M} and \\citet{2011ApJ...741L..28V}. But although both studies obtained consistent effective temperature and luminosity, and they employed the same evolutionary model, the derived $M_\\mathrm{ZAMS}$ was different. Because color uncertainties are expected to arise from unknown mass-loss history of the progenitor, \\citet{2011ApJ...739L..37M} assumed only the luminosity was reliable, and derived $M_\\mathrm{ZAMS}= 13\\pm 3 \\; M_\\odot$ from the end point of the evolutionary track that matched the luminosity. Meanwhile \\citet{2011ApJ...741L..28V} derived $M_\\mathrm{ZAMS}= 17-19 \\; M_\\odot$ by choosing the closest track that matches the luminosity and color of the source in the Hertzsprung-Russell (H-R) diagram although this point does not correspond with the final position of the star at the end of its evolution. Alternatively, \\citet{2011ApJ...742L...4M} compared stellar population synthesis with the stellar association surrounding SN~2011dh to derive $M_\\mathrm{ZAMS}$. Assuming that the stars in the vicinity of the SN are coeval, they concluded that the value of $M_\\mathrm{ZAMS}$ is most likely close to the estimate of \\citet{2011ApJ...739L..37M}. Some authors have suggested that the YSG star detected in the pre-SN images is not the actual progenitor of SN~2011dh but its binary companion or even an unrelated object. These authors suggest, instead, that the exploding object was a compact star. The arguments for this are based, first, on the shock velocity derived from radio and sub-millimeter observations \\citep{2012ApJ...752...78S,2012arXiv1201.0771B,2012arXiv1201.0770K}. The relatively high shock velocity found, $v_{\\mathrm{sw}} \\approx 0.1\\, c$, would indicate a compact progenitor, or a type cIIb SN, in the scheme proposed by \\citet{2010ApJ...711L..40C}. The second argument for a compact progenitor was introduced by \\citet{2011ApJ...742L..18A} and is based on the quick decline of the optical light curve (LC) soon after the explosion, as compared with SN~1993J, along with a low temperature derived from an early-time spectrum, as compared with analytic expressions by \\citet{2011ApJ...728...63R}. Such a compact progenitor would be inconsistent with the YSG star identified by \\citet{2011ApJ...739L..37M} and \\citet{2011ApJ...741L..28V}, which is expected to have a radius of $\\approx 270 \\,R_\\odot$ (from $\\mathrm{log}\\, L= 4.92$ $L_\\odot$, $T_{\\mathrm{eff}}= 6000 \\,K$, given $L= R^2\\, T_{\\mathrm{eff}}^4$ in solar units). Early observations such as those available for SN~2011dh provide a unique opportunity to analyze the physical properties of the progenitor. In particular, one of the most direct ways to estimate the size of the progenitor is by modeling the LC during the cooling phase that occurs after the shock break-out and before the re-heating by radioactive decay. Observations during this phase are very scarce due to its short duration and so they are very valuable. The studies described above do not present a specific modeling of the early light curve. With the aim of assessing the nature of the progenitor we set out to perform a more detailed modeling of the available observations. A well-known method to estimate the properties of the progenitor object, as well as the explosion energy, and the amount and distribution of the radioactive material, is to compare the observations with predictions from hydrodynamical models. In this paper, we present a set of hydrodynamical models applied to stellar evolutionary progenitors that aim to elucidate the compact or extended nature of the progenitor of SN~2011dh, as well as the main physical parameters of the explosion. Our initial and LC models are presented in \\S~\\ref{sec:models}. A comparison between models and observations is done in \\S~\\ref{sec:LC}. The global physical parameters are studied in \\S~\\ref{sec:SLC}, and the sensitivity of the LCs on the initial radius is investigated in \\S~\\ref{sec:ELC}. In \\S~\\ref{sec:discussion} we discuss the results and summarize our main conclusions \\S~\\ref{sec:conclusion}. ", "conclusions": "\\label{sec:conclusion} We have calculated a set of hydrodynamical models applied to stellar-evolution progenitors in order to study the nature of SN~2011dh. Comparing our models with the observed bolometric LC during the second peak and with line expansion velocities we found that a progenitor with He core mass of $3.3$--4 $M_\\odot$, an explosion energy of $6-10\\times 10^{50}$ erg, and a \\Ni\\, mass of $\\approx$ $0.06$ $M_\\odot$ reproduce very well the observations assuming a distance of $7.1$ Mpc to M51. This type of model is consistent with a main sequence mass between 12 and 15 $M_\\odot$. Remarkably, the range of mass found with our hydrodynamical modeling is in very good agreement with the estimates from two other independent methods, i.e. pre-SN imaging and stellar population analysis. This is different from the situation generally encountered for SNe~IIP where LC modeling estimates main sequence masses that are higher up to a factor of two than those estimated from pre-SN imaging \\citep{2008A&A...491..507U,2009ARA&A..47...63S}. We have studied the effect of the progenitor radius on the early LC and temperature evolution. We found that a progenitor with radius similar to that of the YSG star detected in the pre-SN images is compatible with the early observations of SN~2011dh without contradiction with the temperature that is derived from the spectrum, as opposed to what \\citet{2011ApJ...742L..18A} found using analytic models. Furthermore, progenitors with radii $< 200 \\, R_\\odot$ fail to reproduce the early $g'$-band LC. Although our hydrodynamical models show differences in the temperature evolution between extended and compact progenitors, these differences are less marked than those predicted by analytic expressions and they are almost unnoticeable for $t\\gtrsim$ 2 days. Therefore, the spectrum temperature at $t\\approx 2.8$ days is not useful in this case to discriminate between compact and extended progenitors. We have tested {\\em and ruled out} progenitors with He core masses $\\gtrsim 8 M_\\odot$ which correspond to $M_{\\mathrm{ZAMS}} \\gtrsim 25 \\, M_\\odot$. Considering the limitations at such stellar masses of single-star winds to expel the H-rich envelope almost entirely, as required for SNe~IIb, this result is in favor of a binary origin for SN~2011dh. We have also performed binary evolution calculations with mass transfer to test the possibility of systems that are compatible with the pre-SN observations of SN~2011dh. We have shown that a system with 16 $M_\\odot$ + 10 $M_\\odot$ and an initial period of 150 days predicts that the primary star ends its evolution in the H-R diagram at the right position as compared with the YSG star detected in the pre-SN images \\citep{2011ApJ...739L..37M,2011ApJ...741L..28V}. Furthermore, the He mass of the primary at the end of the evolution was $\\approx 4 M_\\odot$, which is consistent with our hydrodynamical modeling. The binary evolution calculations further predict that some hydrogen mass of $\\approx 4 \\times 10^{-3} M_\\odot$ is left in the envelope of the primary star, which is required to produce a SN~IIb. To test the binary scenario we have studied the effect of the putative companion star on the pre-explosion photometry in comparison with the observations. Because the secondary star is predicted to be much hotter than the primary star, we found that the largest effect appears in the blue and UV filters. The contribution of the secondary to the flux in the F336W band, however, is marginal, at the $1.5$ $\\sigma$ level. The contribution is further decreased to the $0.6$ $\\sigma$ level when non-conservative mass accretion is considered. However, our predictions can be tested in a few years time by a search for a blue object at the location of the SN." }, "1207/1207.7141_arXiv.txt": { "abstract": "We present the results of the fifth Interferometric Imaging Beauty Contest. The contest consists in blind imaging of test data sets derived from model sources and distributed in the OIFITS format. Two scenarios of imaging with CHARA/MIRC-6T were offered for reconstruction: imaging a T~Tauri disc and imaging a spotted red supergiant. There were eight different teams competing this time: Monnier with the software package MACIM; Hofmann, Schertl and Weigelt with IRS; Thi\\'ebaut and Soulez with MiRA ; Young with BSMEM; Mary and Vannier with MIROIRS; Millour and Vannier with independent BSMEM and MiRA entries; Rengaswamy with an original method; and Elias with the radio-astronomy package CASA. The contest model images, the data delivered to the contestants and the rules are described as well as the results of the image reconstruction obtained by each method. These results are discussed as well as the strengths and limitations of each algorithm. ", "introduction": "\\label{sec:intro} % The IAU Interferometry Beauty Contest is a competition aimed at encouraging the development of new algorithms in the field of interferometric imaging, by showcasing the current performance of image reconstruction packages. The contest is being conducted by the Working Group on Image Reconstruction of IAU Commission 54. The principle of the contest is the following. One or several science cases are first selected by the organizers, then realistic models of the science targets are used to generate synthetic images. These ``truth'' images are then turned into data sets, by simulating the acquisition of interferometric observables by a typical interferometer. Finally, the contestants attempt to reconstruct images from the data sets without knowledge of the original truth images beyond the nature of the target. The reconstruction closest to the truth image is then declared the winner. The previous contests took place in 2004\\cite{BC2004}, 2006\\cite{BC2006}, 2008\\cite{BC2008} and 2010\\cite{BC2010}, thus the 2012 Interferometry Beauty Contest described here is the fifth contest. The contest results were announced on July 5th during the 2012 SPIE Astronomical Telescopes and Instrumentation conference in Amsterdam. ", "conclusions": "The general agreement amongst contestants is that both targets were hard to reconstruct. This gave rise to more variance in reconstruction quality than what was witnessed during previous contests; but also demonstrated that decent imaging quality can be obtained on very resolved objects in realistic conditions. Alp~Fak was a difficult target, being probably too resolved for any current software to reconstruct it very well. MiRA and IRS obtained the best results, managing to reconstruct a smooth central star. The regularization used by MACIM favored uniform patches of fluxes, and thus was most probably not adapted to recover the original distribution. Bet~Fak was overall reconstructed well in terms of size, but the actual location of the spots was very dependent on the algorithm. Perhaps because stellar surface imaging is a more familiar application, image priors such as limb-darkened disks were used by most reconstruction. As both the $(u,v)$ coverage and signal-to-noise of Bet~Fak were derived from real CHARA/MIRC-6T data, these results demonstrate caution will be needed when reconstructing spots from real data. MiRA achieved the best scores on Alp~Fak, but its lower performance on Bet~Fak (as well as the current specificities of the Beauty Contest metrics that put more weight on this target) prevented it from getting the best overall scores. With record participation and overall convincing reconstructions, most contestants felt that the fifth Beauty Contest was successful at showcasing the diversity and strengths of the current imaging packages in monochromatic mode. As several packages are planed to add multi-wavelength imaging capabilities in 2012-2013, this contest may be indeed the last one to figure only monochromatic data. As multi-wavelength image reconstruction is both a difficult algorithmic problem and a necessity for new science, the next Beauty Contests should definitively prove exciting..." }, "1207/1207.7007_arXiv.txt": { "abstract": "\\begin{singlespace} \\noindent \\textcolor{black}{Galaxies and galaxy clusters have rotational velocities (v) apparently too fast to allow them to be gravitationally bound by their visible matter (M). This has been attributed to the presence of invisible (dark) matter, but so far this has not been directly detected. Here, it is shown that a new model that modifies inertial mass by assuming it is caused by Unruh radiation, which is subject to a Hubble-scale ($\\Theta$) Casimir effect predicts the rotational velocity to be: $v^{4}=2GMc^{2}/\\Theta$ (the Tully-Fisher relation) where G is the gravitational constant, M is the baryonic mass and c is the speed of light. The model predicts the outer rotational velocity of dwarf and disk galaxies, and galaxy clusters, within error bars, without dark matter or adjustable parameters, and makes a prediction that local accelerations should remain above $2c^{2}/\\Theta$ at a galaxy's edge.}\\end{singlespace} ", "introduction": "{Introduction}} \\begin{singlespace} \\noindent \\textcolor{black}{Zwicky (1933) first noticed that galaxies in galaxy clusters were moving too fast to be held together gravitationally by their visible matter, and proposed the existence of an invisible (dark) matter that provides the extra required gravitational pull. A similar problem in disc galaxy rotation was proven by the accurate rotation curves of Rubin }\\textit{\\textcolor{black}{et al}}\\textcolor{black}{. (1980). Dark matter is still the most popular explanation for these problems, but, after decades of searching, it has not been directly detected, though many efforts are ongoing, such as CDMS-II (2009) and XENON10 (2009).} \\noindent \\textcolor{black}{Milgrom (1983) proposed an alternative explanation for galaxy rotation. He speculated that either 1) the force of gravity may increase or 2) the inertial mass ($m_{i}$) may decrease for the low accelerations at a galaxy's edge. His empirical scheme, called Modified Newtonian Dynamics (MoND), can fit disc galaxy rotation curves, and has the advantage of being less tunable than dark matter. However, it does require one arbitrary parameter, the acceleration $a_{0}$, and it does not predict the dynamics of galaxy clusters (Aguirre et al., 2001, Sanders, 2002). The model proposed in this paper has no adjustable parameters, and fits these clusters more closely (it fits spirals less well, but is still within error bars).} \\end{singlespace} \\noindent \\textcolor{black}{Haisch et al. (1994) proposed that inertia might be due to a reaction to the magnetic component of Unruh radiation, which is seen only by accelerating bodies (the work of Haisch et al. was an initial inspiration in the development of the model of inertia presented in this paper, although the actual mechanism that produces inertia is not specified here, and need not be that of Haisch et al.).} \\noindent \\textcolor{black}{The wavelength of Unruh radiation lengthens as acceleration reduces, and Milgrom (1994) noted that as galactic radius increases, the rotational acceleration reduces, the Unruh waves lengthen, and, at the radius where galactic dynamics start to become non-Newtonian, the Unruh waves reach the Hubble scale. He speculated, without assigning a specific cause, that this event may abruptly reduce inertia, affecting dynamics and perhaps explaining MoND. However, an abrupt loss of inertia at a particular galactic radius is not what is seen. The observations show a more gradual deviation from Newtonian behaviour.} \\begin{singlespace} \\noindent \\textcolor{black}{McCulloch (2007) proposed a model for inertia that could be called a Modification of inertia resulting from a Hubble-scale Casimir effect (MiHsC) or Quantised Inertia. MiHsC assumes that the inertial mass of an object is caused by Unruh radiation resulting from its acceleration with respect to surrounding matter, and that this radiation is subject to a Hubble-scale Casimir effect. This means that only Unruh waves that fit exactly into twice the Hubble diameter are allowed, so that an increasingly greater proportion of the Unruh waves are disallowed as accelerations decrease and these waves get longer, leading to a new gradual loss of inertia as acceleration reduces. This loss of inertia is far more gradual than Milgrom's proposal, discussed above. In MiHsC the inertial mass becomes} \\noindent \\textcolor{black}{ \\begin{equation} m_{I}=m_{g}\\left(1-\\frac{\\beta\\pi^{2}c^{2}}{|a|\\Theta}\\right)\\sim m_{g}\\left(1-\\frac{2c^{2}}{|a|\\Theta}\\right) \\end{equation} } \\noindent \\textcolor{black}{where $m_{g}$ is the gravitational mass, $\\beta=0.2$ (part of Wien's displacement law), c is the speed of light, and $\\Theta$ is the Hubble diameter ($2.7\\times10^{26}m$, from Freedman, 2001). For the derivation of Eq. 1 see McCulloch (2007) and for a justification for the use of the modulus of the acceleration see McCulloch (2008) and McCulloch (2011). MiHsC has now been tested quite successfully on several anomalies that have been observed in environments where accelerations are small (see McCulloch, 2007, 2008, 2010, 2011). MiHsC violates the equivalance principle, but not in a way that could have been detected in a torsion balance experiment (McCulloch, 2011).} \\noindent \\textcolor{black}{McGaugh et al. (2009) studied the baryonic mass of disc galaxies and showed that there were none with a baryonic mass of less than $10^{9}M_{\\odot}$. This minimum mass is also predicted by MiHsC (McCulloch, 2010) since in MiHsC mutual accelerations must always be above $2c^{2}/\\Theta$ (close to the acceleration attributed to dark energy). Using this prediction of a minimum acceleration, in this paper MiHsC is applied to the rotation of a wider range of cosmic structures.} \\end{singlespace} ", "conclusions": "" }, "1207/1207.1621_arXiv.txt": { "abstract": "Recently, several groups identified a tentative $\\gamma$-ray line signal with energy $\\sim 130$ GeV in the central Galaxy from the Fermi-LAT data. Such a $\\gamma$-ray line can be interpreted as the signal of dark matter annihilation. However, the offset $\\sim 220$ pc ($1.5^{\\circ}$) of the center of the most prominent signal region from the Galactic center Sgr A$^{\\star}$ has been thought to challenge the dark matter annihilation interpretation. Considering the fact that such a 130 GeV $\\gamma$-ray line signal consists of only $\\sim14$ photons, we suggest that the ``imperfect'' consistency of these photons with the expected dark matter distribution is due to the limited statistics. The offset will be smaller as more signal photons have been collected in the near future. Our Monte Carlo simulation supports the above speculation. ", "introduction": " ", "conclusions": "" }, "1207/1207.4374_arXiv.txt": { "abstract": "We develop a Principal Component Analysis aimed at classifying a sub-set of 27,350 spectra of galaxies in the range $0.4 10^9 M_{\\odot}, M_{\\star} > 10^7 M_{\\odot}$) undergo rapid and frequent bursts of star formation. The associated SNe feedback from these SF episodes results in structural changes to the central mass distribution of these dwarf galaxies, resulting in reduced DM densities and shallower inner DM density profiles than DM-only galaxies. } {\\item The progenitors of lower luminosity satellites ($M_{\\star} \\le 10^7 M_{\\odot}$ at infall) reduce their gas content, and hence become less efficient at forming stars, earlier than more massive dwarfs. This gas loss is partly due to heating by the uniform UV background, as well as subsequent gas loss in early star forming/feedback events, preventing them from having multiple strong bursts of star formation that lead to DM core creation. Low luminosity satellites, therefore, tend to retain steep DM density profiles that are comparable to DM-only runs. } {\\item For SPH satellites across all masses, the overall reduction prior to infall in total $v_c$ is, on average, less than 5 km/s. Although SNe-driven outflows have reduced the central DM mass in halos with $M_{vir} > 10^9 M_{\\odot}$, the presence of gas keeps the overall central mass comparable to the DM-only case. However, this gas is stripped after infall. Hence, the major reduction in central mass is set in place within the DM component prior to infall, but the removal of gas is necessary to reduce the total mass between the SPH and DM-only runs by z=0.} {\\item Once accreted, SPH satellites experience more mass loss due to tidal stripping than DM-only satellites, the amount of which is dependent on infall time and orbit. While both SPH and DM-only satellites are affected by tidal stripping, the presence of a baryonic disk in the SPH runs results in a greater reduction in the central $v_c$ in SPH satellites. The influence of a baryonic disk is especially strong for satellites with shallow DM density profiles. We find that SPH satellites with $z_{infall} > 1$ experience a reduction in their central circular velocities that is $11 - 62\\%$ more than the reduction in $v_c$ at 1 kpc experienced by their DM-only counterparts. } {\\item By $z=0$, the combined effects of DM core creation and enhanced tidal stripping for luminous satellites results in a significant discrepancy between the circular velocity profiles of SPH and DM-only satellites. We find that the $v_c$ at 1 kpc predicted for satellites by DM-only simulations should be reduced by $\\Delta(v_c,1kpc) \\sim 0.2 v_{max, DM-only}-0.26$ km/s, for satellites with $20 < v_{max} < 50$ km/s at infall. } \\end{itemize} High resolution that allows simulators to limit star formation to high density peaks is an essential requirement to reproduce the baryonic effects in this paper. However, \\citet{Governato2010} and \\citet{Guedes2011} showed that, even at high resolution, if star formation is allowed to occur diffusely across the disk, no large scale outflows are generated. Restricting star formation to high density peaks instead leads to overpressurized regions of hot gas when stars go SNe, leading to outflows. These overpressurized regions expand faster than the local dynamical time. When this occurs in the central $\\sim$1 kpc, the potential flattens as this hot gas expands, leading to an irreversible expansion of the DM orbits \\citep{Pontzen2012}. Restricting star formation to high density peaks is comparable to allowing stars to only form in giant molecular clouds, rather than across the entire disk at any given time. The simulations used in this work are the first at these high resolutions to include metal line cooling \\citep{Shen2010} and a presciption for self-shielding of cold gas, allowing star formation to be tied to the shielded regions where H$_2$ can form \\citep{Christensen2012}. It is important to note that simulations that do not resolve the effect of feedback at high densities will be unable to reproduce the results of this paper. However, the feedback model employed in this paper has been shown to match the observed mass -- metallicity relation for galaxies as a function of redshift \\citep{Brooks2007, Maiolino2008}, the baryonic Tully-Fisher relationship (Christensen et al. in prep.), the size -- luminosity relation of galaxy disks \\citep{Brooks2011}, the $z=0$ stellar mass to halo mass relation \\citep{Munshi2012}, and the central mass as a function of stellar mass for galaxies in the luminosity range in this paper \\citep{Governato2012}. This large number of successes in matching the observed scaling relations of galaxies lends credence to the particular feedback model employed in this paper to study satellite galaxies. In addition to reproducing the above scaling relations, the star formation and feedback model used in this work has been used to simulate bulgeless dwarf disk galaxies \\citep{Governato2010}. Importantly, the processes that lead to bulgeless galaxies also transform cuspy DM density profiles into cored profiles, leading these simulations to match the central dark matter densities derived by the {\\sc THINGS} and {\\sc Little THINGS} surveys \\citep{Oh2011, Governato2012}. In other words, the simulations used in this work have been shown to reconcile the cusp/core problem in CDM. In this paper, we have shown that this same model can alleviate the tension between the dense, massive satellites predicted by CDM with the observations of lower density, luminous dSph satellites \\citep{Boylan-kolchin2012, Wolf2012, Hayashi2012}. We note that the processes described in this paper also act to reduce the overall number of massive subhalos that exist at $z=0$, potentially solving the missing satellites problem in CDM. Hence, it remains possible to resolve the small scale problems of CDM with a proper model for baryonic physics, and without invoking exotic forms of DM. The results presented here show that using DM-only CDM simulations to study the internal dynamics of luminous satellites will lead to erroneous results. While it is safe to assign the most luminous satellites to the originally most massive halos, those massive halos will experience evolution that CDM DM-only runs don't account for. In \\citet{Brooks2012}, we address how the model presented here affects the interpretation of kinematic observations of the Milky Way's dSph population." }, "1207/1207.2833_arXiv.txt": { "abstract": "We present a systematic study of the evolution of intermediate- and low-mass X-ray binaries consisting of an accreting neutron star of mass $1.0-1.8 M_{\\odot}$ and a donor star of mass $1.0-6.0 M_{\\odot}$. In our calculations we take into account physical processes such as unstable disk accretion, radio ejection, bump-induced detachment, and outflow from the $L_{2}$ point. Comparing the calculated results with the observations of binary radio pulsars, we report the following results. (1) The allowed parameter space for forming binary pulsars in the initial orbital period - donor mass plane increases with increasing neutron star mass. This may help explain why some MSPs with orbital periods longer than $\\sim 60$ days seem to have less massive white dwarfs than expected. Alternatively, some of these wide binary pulsars may be formed through mass transfer driven by planet/brown dwarf-involved common envelope evolution. (2) Some of the pulsars in compact binaries might have evolved from intermediate-mass X-ray binaries with anomalous magnetic braking. (3) The equilibrium spin periods of neutron stars in low-mass X-ray binaries are in general shorter than the observed spin periods of binary pulsars by more than one order of magnitude, suggesting that either the simple equilibrium spin model does not apply, or there are other mechanisms/processes spinning down the neutron stars. ", "introduction": "Millisecond pulsars (MSPs) are neutron stars (NSs) characterized by short spin periods ($P_{\\rm spin} < 20 $ ms) and weak surface magnetic fields ($B < 10^{9}$ G), which are often found in binaries with a white dwarf (WD) companion. It is generally agreed that MSPs are old NSs, recycled by accretion of mass and angular momentum from the donor stars through Roche-lobe overflow (RLOF) during the previous low-mass X-ray binary (LMXB) evolution. The NS was spun up, due to mass accretion, into a MSP, while the donor evolved to be a He or CO WD \\citep[see e.g., ][for reviews]{bv91,tv06}. The stability of mass transfer in an X-ray binary depends on the ratio of the masses of the donor star and the NS, and the initial orbital period of the system. Traditionally, it was thought that, if the mass ratio is large enough ($\\gtrsim 1.5$), the mass transfer is likely to be unstable, resulting in a common envelope (CE) evolution \\citep{p76,w84,il93}, where the NS spirals into the envelope of the donor on a very short timescale ($<10^{3}$ yr). More recent investigations \\citep[e.g.,][]{tvs00,p00,kdkr00,prp02,prp03} show that X-ray binary systems with intermediate-mass (up to $\\sim 5 M_{\\odot}$) donor stars (i.e., IMXBs) can avoid the spiral-in phase and experience rapid mass transfer on a thermal timescale, successfully evolve to become LMXBs. For LMXBs, a CE evolution is even unavoidable if the separation of binary components is large enough, so that the donor reaches the asymptotic giant branch (AGB) phase with a deep convective envelope before RLOF (i.e. Case C RLOF). Stable mass transfer in LMXBs usually occurs on a timescale $\\sim 10^{8}-10^{10}$ yr when the donor star is on the main sequence or (sub)giant branch at the onset of RLOF (i.e., Case A or B RLOF), and forming MSPs seems to be feasible in this process. It was found by \\citet{ps88,ps89} that there exists a critical bifurcation orbital period ($P_{\\rm bif}$), the initial orbital period that separates the formation of converging LMXBs (which evolve with decreasing orbital periods until the donor star becomes degenerate or an ultra-compact binary is formed) from diverging LMXBs. The value of $P_{\\rm bif}$ is found to be $\\sim 1$ d, but depends heavily on the processes of tidal interactions and the mechanisms and efficiency of orbital angular momentum loss \\citep{e98,prp02,vs05,ml09}. The final products of LMXB evolution are binary pulsars. The distributions of the spin periods of the pulsars, the orbital periods and the WD masses can be used to testify the models of I/LMXB evolution. \\citet{d08} compared the theoretical expectations of I/LMXB evolution to the populations of Galactic binary pulsars. He showed that a significant population of binary pulsars with 1 d $\\lesssim P_{\\rm orb} \\lesssim 100$ d are generally consistent with being the descendants of long-period LMXBs or IMXBs. However, there remain quite a few unresolved puzzles. For example, binary pulsars with $P_{\\rm orb}\\gtrsim$ 60 d seem to have WD companions less massive than predicted by theory, as pointed out previously by \\citet{ts99}, and those with 0.1 d $\\lesssim P_{\\rm orb}\\lesssim$ 1 d are inconsistent with any I/LMXB evolution. In the previous studies on I/LMXB evolution, a canonical NS (of mass $\\sim 1.3-1.4\\,M_{\\odot}$) was usually adopted. This seems to be supported by the finding that the NS mass distribution is consistent with a narrow Gaussian at $1.35\\pm 0.04\\,M_{\\odot}$ \\citep{tc99}. However, both observations \\citep[see][and references therein]{z11,k10,spr10,o12} and theories \\citep{n84,tww96,h03,wj05} suggest that the initial masses of NSs may occupy a large range, probably originating from two different mechanisms of forming NS: iron-core collapse supernovae and electron-capture supernovae. NSs in high-mass X-ray binaries (HMXBs) have experienced very little accretion because of their young ages, so their masses should be very close to those at birth. The measured masses of NSs in HMXBs range from $1.06^{+0.11}_{-0.10} \\,M_{\\odot}$ for SMC X$-$1 \\citep{v07} to $1.86\\pm 0.16\\,M_{\\odot}$ for Vela X$-$1 \\citep{b01,q03}. Recently \\citet{r11} present an improved method for determining the mass of NSs in eclipsing X-ray pulsar binaries and apply it to six systems. They find that the NS masses range from $0.87\\pm 0.07\\,M_{\\odot}$ (eccentric orbit) or $1.00\\pm0.10\\,M_{\\odot}$ (circular orbit) for 4U1538$-$52 to $1.77\\pm 0.08\\,M_{\\odot}$ for Vela X$-$1. So in a proper investigation, the influence of the NS masses should be included. Recent evolutionary calculations by \\citet{db10} have provided evidence that the evolution of I/LMXBs depends upon the NS mass. In this work we perform systematic calculations of I/LMXB evolution and discuss the properties of the produced binary and millisecond pulsars (BMSPs), taking into account different initial NS masses. We adopt the initial donor masses to be $1.0 - 6.0 M_{\\odot}$, and two different initial masses (1.0 and 1.8 $M_{\\odot}$) for the NS. During the mass transfer processes, part of the transferred mass from the donor star may escape from the binary system, carrying away the orbital angular momentum. The formation of BMSPs is closely related to the mechanisms of mass and angular momentum loss, and the stability of mass transfer is also dependent on the angular mementum loss rate \\citep{spv97}. Generally, it is assumed that all the mass transferred from the donor will accrete onto the NS unless in a super-Eddington mass transfer phase, during which the NS will accrete at the Eddington rate ($\\dot{M}_{\\rm Edd}\\sim 1.5\\times 10^{-8}\\,M_{\\odot}$\\,yr$^{-1}$ for a $1.4\\,M_{\\odot}$ NS), and the residual mass escapes from the binary system, carrying the NS's specific orbital angular momentum. This is the so-called ``isotropic reemission model''. Actually, there may also be mass loss even during the sub-Eddington mass transfer phase. For example, the accretion disk can become thermally and viscously unstable when the orbital period is larger than a critical value \\citep{v96,dlhc99}, leading to limit cycle behavior of the mass transfer rate. The NS can accrete mass only during outbursts, while most matter may be ejected out of the binary systems during quiescence by the radiation and magnetic pressure of the rapidly rotating NS \\citep{rst89,bdb02}. In our calculations, mass loss due to super-Eddington mass transfer, radio ejection caused by an unstable disk \\citep{bdb02} or bump-induced detachment \\citep{dvl06}, and in some cases outflow from the $L_{2}$ Lagrangian point are included. This paper is organized as follows. In Section 2 we describe the stellar evolution code and the binary model used in this paper. We present the calculated results in Section 3 and discuss their possible applications in the formation of BMSPs in Section 4, and summarize in Section 5. ", "conclusions": "This work is motivated by the fact the evolution of I/LMXBs with canonical NSs seems to meet difficulties in explaining some of the observational characteristics of BMSPs, and the measurements of the NS masses indicate a wide distribution $\\sim 1-1.8\\,M_{\\odot}$. We have preformed numerical calculations of the evolution of I/LMXBs consisting of a 1.0 or $1.8 M_{\\odot}$ NS and a $1.0-6.0 M_{\\odot}$ donor star, to investigate its dependence on the initial NS mass, and the properties of the descendent BMSPs. The main results can be summarized as follows. 1. The allowed parameter space in the initial $P_{\\rm orb} -M_2$ diagram for forming recycled pulsars increases with increasing NS mass. This may help explain the formation of BMSPs with $P_{\\rm orb}\\gtrsim 60$ d and their distribution in the $P_{\\rm orb}-M_{\\rm WD}$ diagram. Alternatively, some of these wide binary pulsars may be formed through mass transfer driven by planet/brown dwarf-involved CE evolution. 2. The equilibrium spin periods $P_{\\rm eq}$ of accreting NSs in LMXBs derived from the standard magnetosphere-accretion disk interaction model are in general shorter than the observed spin periods of BMSPs by more than one order of magnitude. This implies that either the simple equilibrium spin model does not apply for the spin evolution in accreting NSs, or there are other mechanisms/processes to spin down the NSs when forming BMSPs. 3. Some of the compact IMBPs might have evolved from IMXBs in which the companion star were strongly magnetized Ap/Bp stars with enhanced MB. 4. Our calculations doubt the suggestion that the orbital period gap ($\\sim20 -60$ d) of BMSPs is related to the occurrence of the bump-related detached phase in the LMXB evolution, since the accretion disk would become unstable at earlier time." }, "1207/1207.6417_arXiv.txt": { "abstract": "Using high resolution off-band \\ha\\ data from the New Solar Telescope and Morlet wavelet analysis technique, we analyzed transverse motions of type II spicules observed near the North Pole of the Sun. Our new findings are that i) some of the observed type II spicules display kink or an inverse ``Y'' features, suggesting that their origin may be due to magnetic reconnection, and ii) type II spicules tend to display coherent transverse motions/oscillations. Also, the wavelet analysis detected significant presence of high frequency oscillations in type II spicules, ranging from 30 to 180~s with the the average period of 90~s. We conclude that at least some of type II spicules and their coherent transverse motions may be caused by reconnection between large scale fields rooted in the intergranular lanes and and small-scale emerging dipoles, a process that is know to generate high frequency kink mode MHD waves propagating along the magnetic field lines. ", "introduction": "\\noindent Spicules of type II \\citep{2007PASJ...59S.655D} are ubiquitous chromospheric plasma upflows that are though to transfer of mass into the corona and be an important part of coronal heating and solar wind acceleration process \\citep[e.g.,][]{2011Sci...331...55D,bart_roots,scot_upflows}. Type II spicules have a disk counterpart, the so called rapid blue-shifted excursions associated with the network fields \\citep{langangen_2008, counterparts}. Type II spicules (hereafter spicules II) possess physical properties different from classical spicules: a shorter lifetime (10~s -- 100~s), smaller width (150 -- 700~km) as well as much higher line-of-sight (50 -- 150~km s$^{-1}$) and transverse (10~km s$^{-1}$) velocities. Spectroscopic studies indicate that the plasma in these events is heated throughout their lifetime and the related upflows exhibit jet like properties \\citep{2009ApJ...705..272R}. Spicules II are also known to exhibit transverse/swaying, volume filling motions \\citep[e.g.,][]{Tomczyk_2007, Scott_2011Natur} with amplitudes of 10-30~km s$^{-1}$ and periods of 100-500~s (see review by \\cite{2009SSRv..149..355Z}), which are interpreted as upward or downward propagating ``Alfvenic'' waves \\citep{Okamoto_Bart}, or MHD kink mode waves \\citep[e.g.,][]{2009ApJ...705L.217H, 0004-637X-749-1-30, Kuridze_2012}. There is little consensus among researchers as to how spicules II originate and what is the source of their transverse oscillations. While some works suggest that reconnection process \\citep{2008ApJ...679L..57I,2007PASJ...59S.655D,archontis_jet_model} and the oscillatory reconnection \\citep{0004-637X-749-1-30} in particular may account for their origin, others propose that strong Lorentz force \\citep{2011ApJ...736....9M} or propagation of the p-modes \\citep{2009ApJ...702L.168D} may be at the origin. Moreover, \\cite{2011ApJ...730L...4J} argues that spicules II could be warps in 2D sheet like structures (as opposed to tube-like structures), while \\cite{zhang12revision} questions the existence of spicules II as a distinct class altogether. The related difficulties in interpretation of solar data mainly arise from limited spatial resolution and the complexity of the chromosphere \\citep[e.g.,][]{2011ApJ...730L...4J,0004-637X-749-2-136}. Although a substantial amount of work has been done in the field, the majority of studies dealt with off-limb data when data interpretation may be influenced by effects of line of sight integration and event selectivity, since only tallest and the most prominent spicules can be reliably measured. In particular, \\cite{Moortel_2012} notes that the superposition of coronal loops may impair identification of both the oscillating structure and the mode of the observed oscillation. Here we focus on oscillations of type II spicules observed against the solar disk near the North Pole. In Section 3 we present results of wavelet analysis of off-band H$\\alpha$ data obtained with the New Solar Telescope \\citep[NST,][]{goode_nst_2010, Cao_IRIM} installed at the Big Bear Solar Observatory (BBSO), as well as discuss evolution of two spicule II events that exhibited signatures of a kink and inverted ``Y'' configuration. ", "conclusions": "Using high resolution images of spicules II and Morlet wavelet analysis technique, we analysed transverse motions of type II spicules observed by the NST at the North Pole regions. The data suggests that the spicule II activity quite often exhibit coherent transverse motions. The width of the band of coherently moving spicules is about 1.5-2.0~Mm. We report periods of the transverse motions as detected by the wavelet technique in the range of 30-180~s with the most probable period of about 90~s. MHD simulations by \\cite{2008ApJ...679L..57I} show that such high frequency waves with periods of 90~s are likely to be driven by small-scale reconnection processes. They also could be the result of photospheric oscillations leaking in to the chromosphere, since the detected periods are below the cut-off periods of chromospheric oscillations and kink-waves \\citep{Kuridze_2012}. Similar short period transverse motions were recently detected by \\cite{Okamoto_Bart} using of limb observations in the Ca II H line \\citep[see also ][]{2009SSRv..149..355Z, 2007Sci...318.1574D, Scott_2011Natur}. Present wavelet results emphasise significant presence of short period oscillations associated with type II spicules. We also found that some spicules of type II display kink and inverted ``Y'' shape. These features were detected at the very root of spicules II, where they probably originate. We thus argue that at least some spicules II events and their transverse motions may be due to the interchange type of reconnection between the ``open'' (or large scale closed) field lines and the emerging small-scale dipoles accompanied by generation of high frequency kink mode MHD waves propagating along the magnetic field lines. At a first glance, the interchange reconnection approach may not be agreeing with the coherent oscillations of spicules. One possibility is that a reconnection event removes flux from a bundle of flux tubes thus inducing rapid shuffling of footpoints and displacement of flux tubes within a cluster as this suddenly disturbed magnetic system undergoes equilibrium reconfiguration. The disturbance may generate both small scale (component) reconnection and high frequency MHD waves as proposed by \\cite{0004-637X-736-1-3}, which may account for the appearance of smaller and thinner spicules II as well as their group oscillations. The majority of existing off-limb studies were focused on tallest and most prominent features, possibly associated with ``open'' and/or closed large loops that tend to oscillate with lower frequencies due to their larger extend and lower intensity of the magnetic field \\citep[see e.g.,][]{2001A&A...372L..53N}. However, it appears that the majority of the spicules II in our data set appears to be associated with closed loops that probably span one or two super-granular cells (30-60~Mm). Their fields on, average, may be stronger and the observed transverse motions may be dominated by shorter period oscillations. Moreover, the number density of spicules II in our data set is much lower that that in the off-limb data, so we had an opportunity to resolve and more accurately detect individual events. For instance, we reliably measure transverse motions of these well resolved \\ha\\ structures at heights below 3-5~Mm. While we detect in many cases at least one period of oscillations, the off-limb data at these heights tend to produce linear tracks \\citep{2007Sci...318.1574D} that may be less reliable in determining periods. Combination of these facts may be the reason that the short period oscillations are prominently present in our study \\citep[also see relevant conclusions by][]{Moortel_2012}. This work was conducted as part of the NASA Living with a Star ``Jets'' Focused Science Team. Authors thank T.~J. Okamoto, B. de Pontieu, T.~J. Wang and H. Tian for valuable discussions. Authors acknowledge the usage of the Heliophysics Events Knowledgebase and data from NASA SDO mission. We thank BBSO observing and engineering staff for support and observations. This research was supported by NASA grants GI NNX08AJ20G, LWS NNX08AQ89G, NNX11AO73G, as well as NSF AGS-1146896 grant." }, "1207/1207.6598_arXiv.txt": { "abstract": "We show how an appropriate stationary crystalline structure of the magnetic field can induce a partial fragmentation of the accretion disk, generating an axial jet seed composed of hot plasma twisted in a funnel-like structure due to the rotation of the system. The most important feature we outline is the high degree of collimation, naturally following from the basic assumptions underlying the crystalline structure. The presence of non-zero dissipative effects allows the plasma ejection throughout the axial jet seed and the predicted values of the accretion rate are in agreement with observations. ", "introduction": "A relevant and puzzling feature of compact astrophysical objects (\\emph{e.g.}, Gamma Ray Burst \\cite{Pi99} and Active Galactic Nuclei \\cite{Kr99}) is their capability to generate highly energetic and collimated axial jets, which are well observed in the various electromagnetic bands. The way such jets remain collimated over a very long astrophysical path is a non trivial question, but indications exist that a pressure balance can take place as an effect of the medium in which they are propagating \\cite{Pi99, BKL01}. However, the mechanism which is responsible for the generation of such a peculiar emission of matter and radiation remains almost unidentified. A measurement of the difficulty to find a satisfactory explanation of the jet phenomenon is provided by the fact that one of the most promising proposals postulates magnetic monopole effects \\cite{BKL01, Be10, LWS87, LBC91, Lo02}. Indeed, the possibility to formulate a reliable model for the generation of matter jet (without involving exotic physics) must be regarded as a significant achievement in understanding the behavior of the accreting material near a compact astrophysical source. An alternative framework for the investigation of the stellar wind and jet origin was offered by the studies pursued in \\cite{Co05, CR06}, where it was shown how a local equilibrium configuration allows the existence of a periodic structure for the magnetic flux surfaces, as long as a proper account for the plasma magnetic backreaction is provided. This approach was extensively studied in \\cite{LM10, MB11GRG, BMP11} and generalized to the global profile of the disk equilibrium in \\cite{MB11}. In this scheme, the equilibrium of a rotating stellar disk could favor the emergence of wind and jet seeds already in the steady state, as discussed in \\cite{MC10}. Such a work upgrades the original idea of an ideal purely rotating disk in which a crystalline magnetic field arises, by including non zero poloidal velocities and matter fluxes. This feature implies that the original local equilibrium must now deal with a non zero azimuthal and electron-momentum-balance equations. The reason that this scheme is able to account for high peaks of the vertical velocity (\\emph{i.e.}, the seeds of winds and jets) consists just in the structure that the azimuthal component of the electron force balance takes in the presence of poloidal velocities. In fact, since the azimuthal electric field must identically vanish by virtue of the axial symmetry, one obtains a vertical radial velocity that contains (in its denominator) the radial component of the magnetic field (almost vanishing in the background dipole field). Since the backreaction induces an oscillating radial profile in the radial component of this field (and, in the non linear case, also in the disk mass density), it can be immediately and qualitatively recognized that, in the $O$-point of the magnetic configuration, the vertical velocity takes extremely large values. This specific picture of the disk equilibrium configuration is certainly very intriguing and promising, in view of the generation of winds or jets along the axial direction. Nonetheless, it contains three significant limitations: \\emph{(i)} the analysis is performed in a local radial scenario only, which prevents a localization of the peaks within the disk; \\emph{(ii)} the emerging peak array appears to have a periodic nature, in place of what is commonly observed in real sources; \\emph{(iii)} the matter-flux lines are closed, since they follow the magnetic profile and therefore no real ejection of matter is possible. The present analysis is aimed at overcoming these difficulties and, by a global study of the plasma profile in the disk, we are able to construct an essentially single peak picture (the remaining ones being strongly suppressed outward). Furthermore, by including small dissipative effects in the plasma (according to the possibility to preserve the corotation theorem \\cite{BMP11}), we show how the matter-flux lines become open, allowing a real ejection of material out of the source. In this respect, our work constitutes a significant step toward a settled model for a collimated jet in the scenario introduced by Coppi \\cite{Co05}. Moreover, it offers a valuable theoretical framework for facing a more phenomenological characterization of the wind or jet generation in real astrophysical sources (in stellar as well as in galactic contexts). Regarding the phenomenological impact of the model, three main merits of this study must be put in order. First of all, the collimated nature of the jet seed is guaranteed by the very short scale of the magnetic field crystalline structure in the disk (see \\cite{MB11}) and by the self-consistence of the equilibrium configuration equations. Since the region in which the radial component of the magnetic field (and also the mass density) is approximately zero is very tiny, the outcoming jet is obliged to have a very small cross depth. In the scenario of periodic jet seeds discussed above \\cite{MC10}, this fact is a shortcoming too, implying a small scale sequence of peaks, and therefore it is not clear how they appear when a macroscopical average is taken. Here the primary jet seed is only one and its small scale origin is directly related with its collimation. Second, it is possible to infer the location of the jet seed from the temperature profile of the disk. In fact, the latter quantity is related to a parametric polynomial whose zeros fix the position of the seed. Finally, the present model is able to account for a non-zero accretion rate, which is compatible with the typical observed values in a real system. The paper is structured as follows. In the first Section, the basic equations and assumptions of the model for a thin accretion disk are introduced. In the second Section, the equilibrium equations are integrated and the behaviors of the mass density and vertical velocity of the disk are outlined. As a result, vertical velocity peaks, corresponding to the seed of winds and jets, are obtained. In the third Section, the consistency of the model is outlined and the temperature profile of the system is derived. In the fourth Section, we show that by introducing dissipative effects (in particular, the resistivity), the matter-flux lines become open, implying an effective ejection of material. Brief concluding remarks follow. ", "conclusions": "As in \\cite{MB11} the crystalline structure of the magnetic field in a thin accretion disk (reacting to the dipole field of the central object) has been extended from a local to a global profile, in this paper we have upgraded the analysis of wind and jet seeds performed in \\cite{MC10} toward a global picture. This allows one to reproduce important observational features such as the isolated nature of the jet and a non zero accretion rate of the disk as a whole. Within such a theoretical paradigm, we have developed a stationary equilibrium configuration of the disk which is able to incorporate the presence of a highly collimated jet profile, corresponding to those regions of the disk where the matter density acquires an absolute minimum in its crystalline radial dependence. The main assumptions of our model are the relative small values of the advective forces and of the requested dissipative coefficients (with an associated Prandtl number greater than unity), according to the preservation of the corotation theorem. A non zero resistivity coefficient is required to account for real ejection of material from the disk, since the matter flux lines are no longer isomorphic to the magnetic field profile. The model contains also a free radial polynomial function $g(r)$, which fixes the position of the jet ($g=1$ would imply a jet too close to the central object, \\emph{i.e.}, almost at $r=0$). However, such a function can be directly connected to the thermodynamical properties of the region where the jet arises, in particular with its temperature. The radial domain where the jet takes place corresponds to a peak of the temperature profile, which, in turns, increases with the distance from the equatorial plane. The most appealing feature of the proposed model is that the jet collimation comes out from the compatibility of the equilibrium configuration system and, at the same time, the jet radial extension is very tiny due to the small scale of the crystalline profile. This property, associated to a non zero accretion rate of the disk (the values are in the observed range as soon as the typical orders of magnitude of a stellar system are considered), makes the jet model derived here of very promising phenomenological impact and surely an intriguing perspective in view of a non stationary extension of the equilibrium configuration. The main merit of this work is the significant upgrading provided with respect to the local analysis \\cite{MC10}, which enforces the idea that the crystalline profile of the backreaction is a suitable scenario to implement the jet formation from an axisymmetric accretion structure. {\\small ** This work was partially developed within the framework of the \\emph{CGW Collaboration} (www.cgwcollaboration.it). **} \\newpage" }, "1207/1207.4281_arXiv.txt": { "abstract": "{\\noindent $N$-body simulations predict that dark matter haloes are described by specific density profiles on both galactic- and cluster-sized scales. Weak gravitational lensing through the measurements of their first and second order properties, shear and flexion, is a powerful observational tool for investigating the true shape of these profiles. One of the three-parameter density profiles recently favoured in the description of dark matter haloes is the Einasto profile. We present exact expressions for the shear and the first and second flexions of Einasto dark matter haloes derived using a Mellin-transform formalism in terms of the Fox $H$ and Meijer $G$ functions, that are valid for general values of the Einasto index. The resulting expressions can be written as series expansions that permit us to investigate the asymptotic behaviour of these quantities. Moreover, we compare the shear and flexion of the Einasto profile with those of different mass profiles including the singular isothermal sphere, the Navarro-Frenk-White profile, and the S\\'ersic profile. We investigate the concentration and index dependences of the Einasto profile, finding that the shear and second flexion could be used to determine the halo concentration, whilst for the Einasto index the shear and first and second flexions may be employed. We also provide simplified expressions for the weak lensing properties and other lensing quantities in terms of the generalized hypergeometric function.} ", "introduction": "\\label{sec:Intro} A more accurate description of the elements that constitute our universe, such as the dark matter haloes that are believed to exist around galaxies and clusters, is of crucial importance for our understanding of cosmological structural formation. Recent results from $N$-body simulations of cold dark matter (CDM) \\citep{2004MNRAS.349.1039N,2006AJ....132.2685M,2008MNRAS.387..536G,2008MNRAS.388....2H,2009MNRAS.398L..21S,2010MNRAS.402...21N,2011MNRAS.415.3177R,2012arXiv1202.6061V} indicate that nonsingular three-parameter models such as the \\citet{1965TIAAA.17..01} profile, fit a wide range of dark matter haloes better than singular two-parameter models, e.g. the Navarro, Frenk, and White (NFW) profile \\citep*{1996ApJ...462..563N,1997ApJ...490..493N}. \\noindent The Einasto profile is given by \\noindent \\begin{equation} \\rho(r)=\\rho_{\\text{s}}\\exp\\left\\{ -d_{n}\\left[\\left(\\frac{r}{{r_{{\\text{s}}}}}\\right)^{1/n}-1\\right]\\,\\right\\} ,\\label{eq:einasto_generic} \\end{equation} \\noindent where $r$ is the spatial radius, the shape parameter $n$ is called the Einasto index, ${r_{{\\text{s}}}}$ represents the radius of the sphere that contains half of the total mass, $\\rho_{\\text{s}}$ is the mass density at $r={r_{{\\text{s}}}}$, and $d_{n}$ is a function that ensures that ${r_{{\\text{s}}}}$ is indeed the half-mass radius. An analytical expansion for the function $d_{n}\\approx3n-1/3+8/1215n+\\mathcal{O}\\left(n^{2}\\right)$ is provided by Retana-Montenegro et al. \\citep[][hereafter \\citetalias{2012arXiv1202.5242R}]{2012arXiv1202.5242R}. One important characteristic of this profile is that its power-law logarithmic slope, $\\gamma\\left(r\\right)=-{\\text{d}}\\text{ln}\\rho/{\\text{d}}\\text{ln}r\\sim r^{1/n}$, depends on the Einasto index, which provides a profile that more accurately fits in the inner regions of simulated dark matter haloes than other profiles such as the NFW profile. In the study of real galaxies, several authors have used multi-component Einasto models, consisting generally of two or more Einasto components for each galaxy, where each component represents a homogeneous stellar population with its own set of parameters. For example, some of the first galaxies to be modelled using multi-component Einasto models were M31 by \\citet{1969Afz.....5..137E} with values of $0.25\\leq n\\leq1$, and other nearby galaxies such as Milky Way, M87, M32, Fornax, and Sculptor, and M31 by \\citet{1974smws.conf..291E} with $0.5\\leq n\\leq4$. Later, in a series of papers multi-component Einasto models were employed to model the luminous components of several galaxies such as the Milky Way \\citep{1989A&A...223...89E}, M87 \\citep{1991A&A...248..395T}, M31 \\citep{1994A&A...286..753T}, and M81 \\citep{1998A&A...335..449T}; in these papers, the Einasto index is characterised by values of $0.36\\leq n\\leq7.1$. The seven distant spiral galaxies GSS 074-2237, GSS 064-4412, GSS 094-2210, GSS 104-4024, GSS 064-4442, MDS uem0-043, and HDFS J223247.66-603335.9 were studied by \\citet{2003A&A...403..529T} and \\citet{2005A&A...433...31T}, respectively. As in the earlier works mentioned, they modelled each visual component with a Einasto profile, the authors found values of $0.25\\leq n\\leq0.91$ and noted that the Einasto indices for the disk component of the galaxies at high redshift follow a trend of having smaller values than the ones at lower redshift. \\citet{2006MNRAS.371.1269T} fitted a multi-component Einasto model to the Sombrero galaxy, with $0.78\\leq n\\leq3$ for the visual components. \\citet{2007arXiv0707.4375T} and \\citet{2007arXiv0707.4374T} presented a multi-component Einasto law study of M31: using photometric data and metallicity measurements, they obtained the matter distribution of luminous components with $0.70\\leq n\\leq4.20$, then tried to fit several models for the dark matter halo using kinematical data from the literature to construct a dynamical model and derive the dark matter density of the galaxy, they concluded that Einasto and NFW profiles give the best fits. \\citet{2011arXiv1112.3120D} fitted the surface brightness density of a sample of elliptical galaxies using a multi-component Einasto profile, finding values of $1\\leq n\\leq3$ for the central components, and $5\\leq n\\leq8$ for the outer components. \\citet{2011AJ....142..109C}, who studied the rotation curves of low mass spiral galaxies, modelled the dark matter halo with a Einasto profile, and obtained smaller values of $n$ than predicted by computational simulations. On the other hand, according to $N$-body numerical calculations the Einasto index dependens on both the halo mass and redshift \\citep{2008MNRAS.388....2H,2008MNRAS.387..536G}. Typical values of the Einasto index are in the range $5\\leq n\\leq8$ according to the results of $N$-body simulations \\citep{2004MNRAS.349.1039N,2008MNRAS.387..536G,2008MNRAS.388....2H,2010MNRAS.402...21N}. \\citet{2012arXiv1202.6061V} analysed dark matter haloes of Milky Way-like systems and concluded that the Einasto model with values of $2\\leq n\\leq5$ is preferred over the NFW profile. An alternative form of the density often used in dark matter halo studies is \\noindent \\begin{equation} \\rho\\left(r\\right)=\\rho_{-2}\\,\\exp\\left\\{ -2n\\left[\\left(\\frac{r}{r_{-2}}\\right)^{1/n}-1\\right]\\,\\right\\} ,\\label{eq:einasto_halo_vers} \\end{equation} \\noindent where $r_{-2}$ is the radius at which the logarithmic slope of the density distribution has a value of $-2$ and $\\rho_{-2}=\\rho\\left(r_{-2}\\right)$. A useful quantity to define is the concentration $c_{E}=r_{200}/r_{-2}$, where $r_{200}$ is the virial radius of a halo of mass $M_{200}$, whose density is 200 times the critical density of the Universe at the halo redshift. One of the advantages of the Einasto profile over other profiles is that it has excellent agreement with the conditions outlined by \\citet{1969AN....291...97E} for constructing real galactic models,\\textbf{ }specifically, some moments must be finite. In particular, for this profile some moments, such as the total mass, central gravitational potential, and effective radius, are finite. In contrast, other profiles have logarithmic moments that must be truncated at some radius to ensure that the profile remains finite. Gravitational lensing provides a direct way to study the mass distribution of large structures in the universe, such as galaxies and clusters, without making any assumptions about their dynamical state or composition. Lensing studies taking advantage of high\\textendash{}quality imaging have proven to be successful in mapping the distribution of dark matter in clusters and galaxies \\citep{2003ApJ...598..804K,2006ApJ...648L.109C,2008ApJ...687..959B,2009ApJ...702..603A,2010PASJ...62..811O,2010MNRAS.405.2215O,2011ApJ...741..116O,2011A&A...529A..93H,2012ApJ...744...94R,2012ApJ...747...96J,MNR:MNR20248}. There are two lensing regimes: the strong regime, where multiple images or strong distortions of a galaxy can be produced by an intervening distribution of matter, and the weak regime, where the lensed galaxy image is only slightly distorted, causing the intrinsic elliptical galaxy to appear as a distorted elliptical image. Weak lensing is a valuable and accurate tool for determining the shapes of dark-matter-halo density profiles, such as ellipticity \\citep{2004ApJ...606...67H,2006MNRAS.370.1008M,2007ApJ...669...21P,2009ApJ...695.1446E,2010ApJ...721..124D,2010MNRAS.405.2215O} and triaxiality \\citep{2005ApJ...632..841O,2005A&A...443..793G,2007MNRAS.380..149C,2009MNRAS.393.1235C,2012MNRAS.420..596F}. Until now, most weak lensing studies have considered only linear-order effects such as the weak shear, the quantity responsible for the induced ellipticity in the galaxy (see e.g. \\citet{1995ApJ...449..460K,2001PhR...340..291B,9783540303091,2008ARNPS..58...99H} for reviews). In the past few years, the study of high-order lensing properties has grown in importance (\\citealp{2002ApJ...564...65G}, \\citealp{2005ApJ...619..741G}, \\citet[][hereafter \\citetalias{2006MNRAS.365..414B}]{2006MNRAS.365..414B}). These properties written as high-order derivatives of the deflection potential can be recognized as convergence and shear gradients. The convergence gradient, called the first flexion $\\mathcal{F}$, induces a centroid shift in the lensed image with respect to the source or {}``skewness''. The shear gradient, called second flexion $\\mathcal{G}$, generates an arc-like shape in the lensed image or {}``arcness''. Weak flexion provides useful information about dark matter haloes on galactic- and cluster-sized scales, particularly when probing substructure on smaller-scales where flexion is more sensitive to shear-only studies \\citep{2009MNRAS.395.1438L,2010MNRAS.409..389B,2010arXiv1008.3088E}. Several methods have been developed to measure the flexion of a lensed image, e.g. shapelets (\\citetalias{2006MNRAS.365..414B}; \\citealt{2007MNRAS.380..229M}; \\citealt{2012MNRAS.tmp.2376F}) and surface brightness moments \\citep{2006ApJ...645...17I,Irwin200583,Irwin2007,2007ApJ...660.1003G,2007ApJ...660..995O,2008ApJ...680....1O,2009ApJ...699..143O,2008A&A...485..363S}. \\citet{2011ApJ...736...43C} introduced a new method, called the analytic image model (AIM), to study flexion in astronomical images. Observational measurements of flexion include the detection of mass substructure in the Abell 1689 cluster by \\citet{2007ApJ...666...51L} and \\citet{2011ApJ...736...43C} using observations of the Hubble Space Telescope (HST), as well as \\citet{2008ApJ...680....1O} employing Subaru images; galaxy-galaxy flexion detection in the ground-based survey Deep Lens Survey \\citep{2005ApJ...619..741G} and the space-based HST COSMOS survey \\citep{2011MNRAS.412.2665V}. In addition, flexion has been proposed as a powerful cosmological tool: \\citet{MNR:MNR17838} suggested the use of convergence shear and flexion maps to decrease errors in the measuring standard candles distances, \\citet{2011arXiv1104.3955C} studied how the flexion signal-to-noise ratio could be used to discern between cosmological models, \\citet{MNR:MNR17838} and \\citet{MNR:MNR20051} proposed the use of cosmic flexion to probe large-scale structure. \\citet{2009MNRAS.400.1132H} studied the halo ellipticity on galactic scales, and found that the inclusion of flexion yields tighter constraints on ellipticity than shear-only studies. \\citet{2011A&A...528A..52E} and \\citet{2011MNRAS.417.2197E} proposed a new way to determine the halo ellipticity using the ratio of tangential-to-radial flexion and studied its behaviour as a radius function. \\citet{2012MNRAS.421.1443E} concluded that flexion is more sensitive to ellipticity than shear by performing a likehood analysis of mock flexion and shear data. Additionally, \\citet{2012MNRAS.419.2215V} considered the case in which cross-terms between both shear and flexion and between intrinsic galaxy ellipticities and flexion are not ignored, concluding that these terms can cause a considerable bias in the flexion estimations. In view of the increased use of the Einasto profile in cosmological studies (e.g. see \\citet{2010JCAP...08..004C,2011MNRAS.415.3177R,2011AJ....142..109C,2011arXiv1112.3120D,2011arXiv1111.3556C,2012JCAP...05..016N}), it is natural to extend its applications to weak shear and flexion lensing studies. Previously, several authors had performed weak lensing studies using the Einasto profile. For example, \\citet{2008MNRAS.388....2H} measured the cross-correlations between halo centres and mass, and between galaxies and mass, in the Millennium Run \\citep{2005Natur.435..629S}, and found that the Einasto profile provides a close fit in the inner regions of their two-part model of the halo-mass cross-correlation function. \\citet{1475-7516-2008-08-006} analysed, using a weak statistical approach, a sample of galactic- and cluster-sized dark matter haloes from the Sloan Digital Sky Survey, and obtained very similar concentration-mass relations for the NFW and Einasto profiles. \\citet{2010A&A...520A..30M} used analytical approximations of the shear of the Einasto profile to compare it with the NFW shear. Parametric models such as the singular isothermal sphere, the NFW, and S\\'ersic profiles have been used to model the dark matter distribution in weak lensing analyses (e.g. \\citealt{2009MNRAS.400.1132H,2011ApJ...729..127U,2011A&A...534A..14V,2011MNRAS.417.2197E,2012MNRAS.419.2215V,2012A&A...540A..61S}), the properties of these models having been studied by several authors (\\citealt{2000ApJ...534...34W}; \\citetalias{2006MNRAS.365..414B}; \\citealt{2009MNRAS.396.2257L}). In the case of the Einasto profile, \\citetalias{2012arXiv1202.5242R} studied the analytical properties of the Einasto profile by applying a Mellin-transform formalism. In terms of Fox $H$ and Meijer $G$ functions, they derived analytical expressions of lensing properties for all\\emph{ }values of the Einasto index, concentrating on the surface mass density, cumulative mass, deflection angle, and deflection potential. However, by means of the Mellin-transform formalism it is possible to extensively study the weak-lensing analytical properties of the Einasto profile. This study provides analytical expressions that can used to model realistic Einasto dark matter haloes in weak lensing modelling studies. In this work, we apply Mellin-transform formalism to obtain and study in detail the analytical expressions for the weak lensing properties of the Einasto profile: the shear, and first, and second flexions. This paper is organized as follows. We summarize the weak lensing formalism in Section \\ref{sec:Section2-1}, and present the Mellin-transform technique in Section \\ref{sec:Section2-2}. In Section \\ref{sec:Section3} we derive closed expressions for the shear and first and second flexions in terms of the Fox $H$ and Meijer $G$ functions. We then use the series expansions of these expressions to investigate their asymptotic behaviour. In Section \\ref{sec:Section4}, we compare our results with those for the SIS, NFW, and S\\'ersic profiles. In Section \\ref{sec:Section5}, we summarise and discuss our main results. Finally, in the appendices \\ref{Appendix-A} and \\ref{Appendix-B} we provide series expansions of the lensing properties, and explicit expressions in terms of the generalized hypergeometric function, respectively. Throughout the paper, we adopt a cosmological model with the matter density $\\Omega_{M}=0.26$, the cosmological constant $\\Omega_{\\Lambda}=0.74$, and the Hubble constant $H_0=72\\, {\\rm km}\\,{\\rm s}^{-1}{\\rm Mpc}^{-1}$. ", "conclusions": "} We have applied the Mellin transform technique to obtain closed-form expressions for the weak lensing properties of the Einasto profile. The expressions for the shear $\\gamma\\left(x\\right)$, first flexion $\\mathcal{F}\\left(x\\right)$, and second flexion $\\mathcal{G}\\left(x\\right)$ can be written in terms of the Fox $H$ function for general values of the Einasto index $n$, and can simplified in terms of the Meijer $G$ function for integer or half-integer values of $n$. We utilized the residue theorem to calculate specific power and logarithmic-power expansions for these expressions. The expansions permit us to study the asymptotic behaviour of the weak lensing properties at small and large radii. Furthermore, we employed the Slater's theorem \\citep{marichev1983handbook} to derive an expression for the convergence $\\kappa\\left(x\\right)$ in terms of the generalized hypergeometric function, which is valid for half-integer values of $n$. This enables the other expressions for the lensing properties to be written in terms of the hypergeometric function. \\noindent We have examined in detail the convergence, shear, and first and second flexions for an Einasto profile and other profiles including the singular isothermal sphere, the NFW, and S\\'ersic profiles. We found that the Einasto profile overall has a similar behaviour to these profiles. Nonetheless, this profile is clearly different from the others, particularly at small angular separations from the lens centre, where the lensing signal is stronger. At large angular separations, the Einasto profile behaves far very similarly to the NFW profile than the other profiles. We explored the dependence of the Einasto profile on the concentration parameter, our results indicating that it has a non-linear concentration dependence and that the shear and second flexion are more effective indicators of the dependence than the convergence and the first flexion. In addition, we studied the Einasto index dependence. For this parameter, the dependence seems to be weaker than for the concentration, for which the shear, second, and first flexions seem to be more sensitive to the index dependence that the convergence. We note that the magnitude of the lensing properties of a S\\'ersic model are stronger than for the other profiles, this indicates that the profile selected to model the halo must be chosen with caution as discussed by \\citet{2009MNRAS.396.2257L}. We note that the index dependence of an Einasto halo is stronger at small angular distances from the lens centre, which is the opposite of the case for a S\\'ersic halo for which at large angular distances the dependence is stronger. This means that observationally it is easier to constrain the value of the index for the Einasto profile rather than the S\\'ersic model, because the lensing signal is stronger near the lensed image. \\noindent The availability of analytical expressions for the Einasto-profile weak-lensing properties is of foremost importance, and constitutes an effort to foster the inclusion of this density profile in weak lensing modelling studies. There are several possible applications of our results for modelling studies. For example, one of them is the generation of the shear signal in weak lensing analyses, which can provide valuable information about the description of the mass density profiles. This, in turn, places constraints on the model parameters such as mass, concentration, and particularly the Einasto index, which is known to scale with mass and redshift according to $N$-body simulations \\citep{2008MNRAS.387..536G,2008MNRAS.388....2H}, and for which our results could be used to verify this variation observationally. Likewise, weak flexion might be generated using our expressions and used to constrain the model parameters, particularly when halo substructure has to be proven, providing another scale within the haloes, where the behaviour of $n$ can be studied. \\noindent Here, we considered spherically symmetrical haloes, but haloes are far from ideal symmetric objects (see e.g. \\citet{2006ApJ...646..815S,2007MNRAS.376..215B,2010MNRAS.407..891H}). Nevertheless, weak shear has proven successfully in studying deviations in the halo shape from spherical symmetry, such as the halo ellipticity \\citep{2004ApJ...606...67H,2006MNRAS.370.1008M,2007ApJ...669...21P,2009ApJ...695.1446E,2010MNRAS.405.2215O}. Similarly, weak flexion has been proposed as a tool to investigate the halo ellipticity \\citep{2009MNRAS.400.1132H,2011A&A...528A..52E,2011MNRAS.417.2197E}. Triaxiality is another aspect of the halo shape that has been explored using weak lensing \\citep{2005ApJ...632..841O,2005A&A...443..793G,2011MNRAS.416.3187S,2012MNRAS.420..596F}; ignoring the halo triaxiality can affect the parameter estimation in lens-rich clusters \\citep{2007MNRAS.380..149C,2009MNRAS.393.1235C}, leading to the cluster appearing to be more massive and concentrated, particularly, if its major axis is aligned with the line of sight \\citep{0004-637X-654-2-714,MNR:MNR14154,A&A2010...Meneghetti}. Some cluster studies where there are apparently lensing biases in the estimated concentration and mass include that of \\citet{2008ApJ...685L...9B}, who analysed the mass and concentration of four nearly relaxed clusters, that of the gravitational lens with the largest Einstein radius detected so far MACS J0717.5+3745 \\citep{2009ApJ...707L.102Z}, and \\citet{2009ApJ...699.1038O}, who obtained the radial profile of four clusters combining lensing data from the Subaru telescope. The inclusion of our results in weak lensing ellipticity and triaxiality studies is straightforward, therefore it enables the possibility of investigating the halo ellipticity and triaxiality with the Einasto profile. \\noindent Our current knowledge of the structure of the Universe on large scales will be improved by new gravitational lensing surveys such as the Dark Energy Survey% \\footnote{http://www.darkenergysurvey.org/% } (DES), Euclid% \\footnote{http://sci.esa.int/euclid/% }, the Large Synoptic Survey Telescope (LSST)% \\footnote{http://www.lsst.org/lsst/% }, the James Webb Space Telescope% \\footnote{http://www.jwst.nasa.gov/% }, KiDS% \\footnote{http://kids.strw.leidenuniv.nl/% }, Pan-STARRS% \\footnote{http://pan-starrs.ifa.hawaii.edu/public/% }, and WFIRST% \\footnote{http://wfirst.gsfc.nasa.gov/% }. These surveys will provide more accurate measurements of weak lensing that could be modelled using the analytical expressions presented in this work. \\noindent This paper constitutes a further step in studying the properties of the Einasto profile using analytical means. In addition, it extends and complements the work of \\citetalias{2012arXiv1202.5242R}, providing additional simplified expressions for their results. With this work, we hope to encourage the use of special functions such as the Fox $H$ function, the Meijer $G$ function, and the generalized hypergeometric function in astronomy and astrophysics. \\noindent" }, "1207/1207.4248_arXiv.txt": { "abstract": "Magneto-convection can produce an active region without an initial coherent flux tube. A simulation was performed where uniform, untwisted, horizontal magnetic field of 1 kG strenght was advected into the bottom of a computational domain 48 Mm wide by 20 Mm deep. The up and down convective motions produce a hierarchy of magnetic loops with a wide range of scales, with smaller loops riding ``piggy back\" in a serpentine fashion on larger loops. When a large loop approaches the surface it produces an small active region with a compact leading spot and more diffuse following spots. ", "introduction": "The standard paradigm is that active regions form when a coherent flux tube from deep in the convection zone reaches the surface, typically modeled using the thin flux tube approximation \\citep{Parker55,Fan93,Moreno-Insertis94,Caligari95,Fan09,Weber11}. \\citet{Cheung07,Martinez-Sykora08,Weber11,Fang12} have shown how important the actual convective motions are to the rise of magnetic flux. Recent simulations of active region formation have started from a coherent semi-torus of magnetic field placed in the surface layers of a model solar convection zone \\citep{Cheung10}. Our simulations show that such a coherent structure is not necessary for the formation of an active region. The action of magneto-convection itself produces rising flux tubes, which when they reach the surface can produce an active region. ", "conclusions": "These simulations show that magneto-convection itself can produce the flux tubes that give rise to active regions. The action of up and downflows on an initial supergranule size patch of horizontal field can keep the major portion of the magnetic flux confined to emerge at the surface in a similar supergranule size region. Interesting work for the future will be to investigate the evolution of larger scale magnetic structures, such as those in the flux emergence simulations of \\citet{Abbett00} and \\citet{Fan08}, as they pass through the upper convection zone. Movies of the active region formation are available on http:steinr.pa.msu.edu/$\\sim$bob/research.html\\#AR." }, "1207/1207.1782_arXiv.txt": { "abstract": "We present the results of a long-term (1999--2010) spectral optical monitoring campaign of the active galactic nucleus (AGN) Ark 564, which shows a strong \\ion{Fe}{2} line emission in {\\bf the optical. This AGN is a} narrow line Seyfert 1 (NLS1) galaxies, a group of AGNs with specific spectral characteristics. We analyze the light curves of the {\\bf permitted} H$\\alpha$, H$\\beta$, optical \\ion{Fe}{2} line fluxes, and the continuum flux in order to search for a time lag between them. Additionally, in order to estimate the contribution of iron lines from different multiplets, we fit the H$\\beta$ and \\ion{Fe}{2} lines with a sum of Gaussian components. We found that during the monitoring period the spectral variation (F$_{\\rm max}$/F$_{\\rm min}$) of Ark 564 was between 1.5 for H$\\alpha$ to 1.8 for the \\ion{Fe}{2} lines. The correlation between the \\ion{Fe}{2} and H$\\beta$ flux variations is of higher significance than that of H$\\alpha$ and H$\\beta$ (whose correlation is almost absent). The {\\bf permitted-line} profiles are Lorentzian-like, and did not change shape during the monitoring period. We investigated, in detail, the optical \\ion{Fe}{2} emission and found different degrees of correlation between the \\ion{Fe}{2} emission arising from different spectral multiplets and the continuum flux. The relatively weak and different degrees of correlations between {\\bf permitted} lines and continuum fluxes indicate a rather complex source of ionization of the broad line emission region. ", "introduction": "Narrow-line Seyfert 1 (NLS1) galaxies were first introduced as a class of active galactic nuclei (AGNs) by \\cite{ost85}. Their optical spectra show relatively narrow (FWHM $\\leq$ 2000 km s$^{-1}$) permitted lines, which are narrower than in a typical Seyfert 1 galaxy. In particular, \\cite{ost85} did show that, the permitted lines are only slightly broader than the forbidden ones, and that a strong \\ion{Fe}{2} emission is present in the optical region of the spectrum. In addition, the [O III] $\\lambda$5007/H$\\beta$ ratio, emitted in the narrow line region (NLR), varies from 1 to 5 {\\bf \\citep[][]{rod00}}, instead of the universally adopted observed value for Seyfert 1s of around 10 \\citep[][]{rod00}, indicative of the presence of high-density gas. \\cite{ost85} pointed out that the H$\\beta$ equivalent widths in NLS1s are smaller than typical values for normal Seyfert 1s, suggesting that they are not just normal Seyfert 1s seen at a particular viewing angle. Renewed interest in NLS1s arises from the discovery of their distinctive X-ray properties: they show a steep X-ray excess with a photon index of 3 below 100 keV, a steep hard X-ray continuum, and a rapid large-amplitude X-ray variability on timescales of minutes to hours \\citep[see][and references therein]{le99a,le99b,le00,pa11}. Moreover, optical studies have established that NLS1s lie at one end of the \\cite{bo92} eigenvector 1 (EV1) and that they show a relatively strong \\ion{Fe}{2} emission and a weak [O III] emission \\citep[][]{bol96}. They also represent the \"extreme Population A\" objects (FWHM H$\\beta <$ 4000 km/s) as defined by the four dimensional eigenvector 1 (4DE1) in \\cite{su07,mar10}. 4DE1 involves four parameters and NLS1 are \"extreme\" in all of them: they have the narrowest broad H$\\beta$, strongest \\ion{Fe}{2} emission, strongest X-ray excess and largest \\ion{C}{4} blueshifts. Arakelian 564 (Ark 564, IRAS 22403+2927, MGC +05-53-012) is a bright $V=14.6$ mag. \\citep{dev91}, nearby narrow-line Seyfert 1 galaxy ($z=0.02467$), with an X-ray luminosity $L_{2-10 \\rm keV}=2.4\\times 10^{43},\\rm{erg \\ s^{-1}}$ \\citep[][]{tu01}. This AGN is one of the brightest NLS1s in the X-ray band \\citep{bol96,co01,sm08}, and it shows a soft excess below $\\sim$1.5 keV and a peculiar emission-line-like feature at 0.712 keV in the source rest frame \\citep[][]{ch04}. The variations of the X-ray amplitude in the short-timescale light curve is very similar to those in the long-timescale light curve \\citep[][]{po01}, that is in contrast to the stronger amplitude variability on longer timescales, which is a characteristic of broad-line Seyfert 1 (BLS1) galaxies. In the UV part of the spectrum, this galaxy shows intrinsic UV absorption lines \\citep[][]{cr99}. In order to explore the variability characteristics of a NLS1 in different wavelength bands, a multiwavelength monitoring campaign of Ark 564 was conducted \\citep{sh01}. {\\bf The optical campaign covered the periods 1998 November -- 1999 November and 2000 May -- 2001 January, where the object was observed both photometrically (UBVRI filters) and spectrophotometrically (spectral coverage 4800--7300 \\AA) \\citep{sh01}. The data set and analysis is described in details and compared with the simultaneous X-ray and UV campaigns \\citep{sh01}.} The results of this intensive variability multiwavelength campaign show that the optical continuum is not significantly correlated with the X-ray emission \\citep[][]{sh01}. The UV campaign, carried out with the HST on 2000 May 9 and 2000 July 8, is described in \\cite{co01}. These authors found a small fractional variability amplitude of the continuum between 1365 \\AA\\ and 3000 \\AA\\ (around ~6\\%), but reported that large-amplitude short-timescale flaring behavior is present, with trough-to-peak flux changes of about 18\\% in approximately 3 days \\citep[][]{co01}. The wavelength-dependent continuum time delays in Ark 564 have been detected and these delays may indicate a stratified continuum reprocessing region \\citep[][]{co01}. Here we present the long-term monitoring of Ark 564 in the optical part of the spectrum. We analyzed the variability in the {\\bf permitted} emission lines and continuum in order to determin the size and structure of emitting regions of the {\\bf permitted} Balmer and \\ion{Fe}{2} lines. We placed particular interest in the strong \\ion{Fe}{2} lines of the H$\\beta$ spectral region whose behavior is investigated in details and discussed in this paper. The paper is organized as follow: in section 2 we describe the observations and data reduction procedures, in section 3 we give an analysis of the spectral data, in section 4 we explore the correlations between different lines and the continuum, as well as between different lines, in section 5 we investigate in more details the \\ion{Fe}{2} variation, in section 6 we discuss our results, and in section 7 we provide our conclusions. ", "conclusions": " 1) In Ark 564 during the monitoring period (1999--2010) the mean continuum and lines fluxes decreased for $\\sim$20\\%-30\\% (see Fig.~\\ref{fig05}) from the beginning (1999) to the end of the monitoring (2010). The total flux of \\ion{Fe}{2} evidently increases with the continuum flux. 2) We registered five flare-like events (two prominent and three possible) lasting $\\sim$1-3 days, when fluxes in continuum and lines changed for $\\sim$20\\% (continuum and \\ion{Fe}{2} emission) and $\\sim$10\\% for Balmer lines. 3) The flux-flux correlations between the continuum and lines are weak, where the correlation between the \\ion{Fe}{2} lines (in the red shelf of the \\ion{Fe}{2}) and continuum is slightly higher (and more significant) than between the Balmer lines and continuum. There is almost lack of correlation between the H$\\alpha$ and H$\\beta$ line fluxes. Such behavior indicate very complex physical processes in the line forming region, i.e. beside the photoionization some additional physical processes may be present. 4) We roughly estimated a lag of 2--6 days, but with large errorbars. Taking that the photoionization is probably not the only source of line excitation, the obtained results should be taken with caution. 5) We investigated in detail, the \\ion{Fe}{2} emission variability. We divided the \\ion{Fe}{2} emission in six groups according to the atomic transitions. We found that correlation between the continuum flux and emission of groups depends on the type of transition, i.e. in some case there is relatively good correlation level between the \\ion{Fe}{2} group emission ($^4$G, $^4$F group), but for $^2$H and $^6$S there is no correlation at all. 6) The Gaussian multicomponent analysis indicates that the emission of the \\ion{Fe}{2} lines is probably coming from the intermidiate line region, having velocities around 1500 km/s. The spectral variability of Ark 564 seems to be complex and different from the one observed in BLS1 \\citep[see e.g.][]{sh09}. The observed flare-like events, the small (or even lack of) correlation between H$\\alpha$ and H$\\beta$ fluxes, the different correlation degree for \\ion{Fe}{2} group emission and continuum light level, may indicate complex physics in the emitting regions, as e.g. there may be, beside the AGN, contribution of star explosions and internal shock waves." }, "1207/1207.6507_arXiv.txt": { "abstract": "We study the impulsively generated non-linear Alfv\\'en waves in the solar atmosphere, and describe their most likely role in the observed non-thermal broadening of some spectral lines in solar coronal holes. We solve numerically the time-dependent magnetohydrodynamic equations to find temporal signatures of large-amplitude Alfv\\'en waves in the model atmosphere of open and expanding magnetic field configuration, with a realistic temperature distribution. We calculate the temporally and spatially averaged, instantaneous transversal velocity of non-linear Alfv\\'en waves at different heights of the model atmosphere, and estimate its contribution to the unresolved non-thermal motions caused by the waves. We find that the pulse-driven nonlinear Alfv\\'en waves with the amplitude $A_{\\rm v}$=50 km s$^{-1}$ are the most likely candidates for the non-thermal broadening of Si VIII $\\lambda$1445.75 \\AA\\ line profiles in the polar coronal hole as reported by Banerjee et al. (1998). We also demonstrate that the Alfv\\'en waves driven by comparatively smaller velocity pulse with its amplitude $A_{\\rm v}$=25 km s$^{-1}$ may contribute to the spectral line width of the same line at various heights in coronal hole without any significant broadening. The main conclusion of this paper is that non-linear Alfv\\'en waves excited impulsively in the lower solar atmosphere are responsible for the observed spectral line broadening in polar coronal holes. This is an important result as it allows us to conclude that such large amplitude and pulse-driven Alfv\\'en waves do indeed exist in solar coronal holes. The existence of these waves and their undamped growth may impart the required momentum to accelerate the solar wind. ", "introduction": "Alfv\\'en waves are difficult to observe and yet several indications of their presence in the solar atmosphere were found during SOHO and TRACE space missions. More direct evidence for the existence of these waves in different regions of the solar atmosphere was given by high-resolution observations carried out by the Solar Optical Telescope (SOT) and the X-Ray Telescope (XRT) onboard the Hinode Solar Observatory. According to Okamoto \\& De Pontieu (2011), De Pontieu et al. (2007), and Cirtain et al. (2007), the signature of Alfv\\'en waves were observed respectively in spicules and X-ray jets using SOT and XRT instruments. Interpretations of these observational results were discussed by Erd\\'elyi \\& Fedun (2007), Tomczyk et al. (2007), and Antolin et al. (2009). Observational evidence for the existence of torsional Alfv\\'en waves in the solar atmosphere was also reported by Jess et al. (2009), who analyzed H$\\alpha$ observations obtained with high spatial resolution by the Swedish Solar Telescope (SST). They interpreted the data in terms of Alfv\\'en waves in the solar chromosphere, with periods from 12 min down to the sampling limit of the observations near 2 min, with maximum power near 6-7 min. The authors concluded that the amount of energy carried by such transversal waves was sufficient to heat the solar corona (Dwivedi \\& Srivastava 2006). These recent discoveries of Alfv\\'en waves in the solar atmosphere well-justified extensive studies of these waves that have been performed by numerous investigators in the last four decades (e.g., Hollweg 1985; Roberts 1991; Musielak \\& Moore 1995; Narain \\& Ulmschneider 1996; Roberts 2004; Antolin \\& Shibata 2010, and references cited there). The studies have covered both linear (e.g., Hollweg \\& Isenberg 2007; Murawski \\& Musielak 2010) and non-linear (e.g., Verdini \\& Velli 2007; Verdini et al. 2009; Matsumoto \\& Shibata 2010) Alfv\\'en waves, and different aspects of their generation, propagation and dissipation have been investigated. The specific objectives of these studies were to understand the role of Alfv\\'en waves in the atmospheric heating and in the acceleration of supersonic solar wind. The fast component of the solar wind originates in solar polar coronal holes (e.g., Hassler et al. 1999; Tu et al. 2005, and references cited there), which are the regions where the non-thermal broadening of spectral lines has also been observed (e.g., Hassler et al. 1990; Banerjee et al. 1998; Moran 2003; O'Shea et al., 2005; Dolla \\& Solomon 2008). Similar observations have also been done in the equatorial corona (Harrison et al., 2002). These authors proposed that the radially propagating Alfv\\'en waves may result in the non-thermal broadening of spectral line widths. These observations were also investigated analytically by Pek{\\\"u}nl{\\\"u} et al. (2002) and Dwivedi \\& Srivastava (2006). Recently, the observations from Bemporad \\& Abbo have given the first signature of the dissipation of Alfv\\'en waves neat the solar limb, which confirm the theoretical models of Dwivedi \\& Srivastava (2006), and Srivastava et al. (2007). Moreover, Zaqarashvili et al. (2006) have reported that the resonant energy conversion from Alfv\\'en to acoustic waves in the region where plasma $\\beta$ approaches unity in the solar atmosphere. This conversion can be responsible for the spectral line width variation. However, this theory only explains the most probable cause of the line-width reduction, which was observed only by O'Shea et al. (2005) in solar coronal hole. Additional problem is the fact that there is not enough observational evidence for the resonant energy conversion in the solar corona (e.g., Srivastava \\& Dwivedi 2010; McAteer et al. 2003). Hence, new studies of the role played by Alfv\\'en waves in the observed spectral line broadening were necessary. In this paper, we numerically study the behavior of large-amplitude (non-linear) Alfv\\'en waves in a model that resembles a solar coronal hole. Our main objective is to determine the role played by these waves in the spectral line broadening as observed in the coronal holes (e.g., Banerjee et al. 1998; Dolla \\& Solomon 2008). We find an agreement between our numerical results of non-linear Alfv\\'en waves and the computed line broadening of constituted synthetic spectra at different heights in model coronal hole and the observational data. This allows us to conclude that large-amplitude Alfv\\'en waves are responsible for the observed non-thermal broadening of the spectral lines in the coronal holes. Our result is important because it is an indirect evidence for the existence of non-linear Alfv\\'en waves in solar coronal holes. The outline of the paper is as follows: our numerical model is described in Sec. 2; a brief description of the used numerical code and the form of initial perturbations are given in Sec. 3; the results of our numerical simulations are presented in Sec. 4; comparison of our results to the observational data is given in Sec. 5; the obtained results are discussed in Sec. 6; and our conclusions are given in Sec. 7. ", "conclusions": "In this paper, we simulated numerically the behavior of impulsively excited non-linear Alfv\\'en waves in solar coronal holes, and studied for the first time their role in the observed broadening of spectral lines caused by these waves. We compared the obtained numerical results to the spectral line broadenings observed by Banerjee et al. (1998); Dolla \\& Solomon (2008). We found that the large-amplitude, non-linear, pulse-driven ($A_{\\rm v}$=50 km s$^{-1}$) Alfv\\'en waves are the most likely candidates for the non-thermal broadening of Si VIII $\\lambda$1445.75 \\AA\\ line profiles in the polar coronal hole (Banerjee et al. 1998). Our results also show that Alfv\\'en waves driven by comparatively smaller velocity pulse, $A_{\\rm v}$=25 km s$^{-1}$, only approximately contribute to the observed spectral line width of Si VIII $\\lambda$1445 \\AA\\ at various heights in coronal hole as reported by Dolla \\& Solomon (2008), however, without significant evidence of broadening. % The results of our numerical simulation and their comparison to the observations of line broadening are important as they become an indirect evidence for the existence of larger amplitude, pulse-driven Alfv\\'en waves in solar coronal holes. Finally, we would like to suggest that more spectroscopic observations should be carried out in future using high-resolution observations (e.g., Hinode/EIS, and also with upcoming Solar-C instruments) to search for a direct evidence of such pulse-excited non-linear Alfv\\'en waves in the solar atmosphere. Obviously, more detailed theoretical studies of a variety of pulse-driven non-linear Alfv\\'en waves excited by a range of pulse amplitudes are also needed in order to better understand the role played by these waves in the observed line broadening and the solar wind acceleration. {\\bf Acknowledgments.} The software used in this work was in part developed by the DOE-supported ASC/Alliance Center for Astrophysical Thermonuclear Flashes at the University of Chicago. This work has been supported by the Alexander von Humboldt Foundation (Z.E.M.) and by a Marie Curie International Research Staff Exchange Scheme Fellowship within the 7th European Community Framework Program (P.Ch. \\& K.M.) as well as by the \"HPC Infrastructure for Grand Challenges of Science and Engineering\" Project, co-financed by the European Regional Development Fund under the Innovative Economy Operational Program (P.Ch. \\& K.M.). We also acknowledge the CHIANTI, which is a collaborative project involving researchers at NRL (USA) RAL (UK), and the Universities of Cambridge (UK), George Mason (USA), and Florence (Italy). A.K.S. thanks Shobhna Srivastava for patient encouragement, and also to Prof. D. Tripathi, IUCAA, Pune, India for valuable discussions on spectral line-profiles and their implications." }, "1207/1207.1301.txt": { "abstract": "Taking advantage of the unprecedented combination of sensitivity and angular resolution afforded by the \\hersc\\ Space Observatory at far-infrared and submillimeter wavelengths, we aim to characterize the physical properties of cold dust within nearby galaxies, as well as the associated uncertainties, namely the robustness of the parameters we derive using different modified blackbody models. For a pilot subsample of the KINGFISH key program, we perform two-temperature fits of the \\spitz\\ and \\hersc\\ photometric data (from 24 to 500 \\mic), with a warm and a cold component, both globally and in each resolution element. We compare the results obtained from different analysis strategies. At global scales, we observe a range of values of the modified blackbody fit parameter $\\beta$$_c$ (0.8 to 2.5) and T$_c$ (19.1 to 25.1K). We compute maps of our modelling parameters with $\\beta$$_c$ fixed or treated as a free parameter to test the robustness of the temperature and dust surface density maps we deduce. When the emissivity is fixed, we observe steeper temperature gradients as a function of radius than when it is allowed to vary. When the emissivity is fitted as a free parameter, barred galaxies tend to have uniform fitted emissivities. Gathering the parameters obtained each resolution element in a T$_c$-$\\beta$$_c$ diagram underlines an anti-correlation between the two parameters. It remains difficult to assess whether the dominant effect is the physics of dust grains, noise, or mixing along the line of sight and in the beam. We finally observe in both cases that the dust column density peaks in central regions of galaxies and bar ends (coinciding with molecular gas density enhancements usually found in these locations). We also quantify how the total dust mass varies with our assumptions about the emissivity index as well as the influence of the wavelength coverage used in the fits. We show that modified blackbody fits using a shallow emissivity ($\\beta$ $<$ 2.0) lead to significantly lower dust masses compared to the $\\beta$ $<$ 2.0 case, with dust masses lower by up to 50 $\\%$ if $\\beta$$_c$ =1.5 for instance. The working resolution affects our total dust mass estimates: masses increase from global fits to spatially-resolved fits. ", "introduction": "The interstellar medium (ISM) plays a key role in the star formation history of galaxies. While stars evolve and die, they re-inject dust and gas in the ISM that will then be processed to form a new generation of stars. Dust is formed in the envelopes of late-evolved stars and is transformed in supernova shock waves or stellar winds by violent processes such as destruction, vaporization, sputtering or erosion \\citep{Jones1994,Jones1996,Serra_Dias_Cano2008}. The size of dust grains also increases through processes like accretion or coagulation accretion \\citep{Stepnik2003,Dominik1997,Dominik2008,Hirashita2009}. They participate in the gas thermal balance and screen parts of the ISM through the absorption of starlight and subsequent emission of infrared (IR) photons. Dust grains contribute to the chemistry of the ISM. Molecular hydrogen predominantly forms on grains, and molecules can freeze out on grain surfaces and undergo additional transformation \\citep{Hasegawa1993,Vidali2004,Cazaux2005}. Atomic and molecular gas are commonly traced by H{\\sc i} and CO observations. CO can be a poor tracer of molecular gas, especially in low-metallicity environments, due to the dependency of the CO-to-H$_2$ conversion factor with density and metallicity \\citep{Wilson1995,Taylor1998,Israel2000,Leroy2005,Wolfire2010}. We know that dust and gas are closely tied in the ISM. This suggests that used with H{\\sc i} and CO, dust could also be used as an additional tracer of the gas. %---------------------------- Galaxy Data --------------------------------------------------------------------------------------------- \\begin{table*} \\caption{Galaxy Data} \\label{Galaxy_data} \\centering \\begin{tabular}{ccccccccccc} \\hline \\hline &&&&&&&&&&\\\\ Galaxy & Optical & $\\alpha$$_0$ & $\\delta$$_0$ & Major diam. & Minor diam. &Distance & \\multicolumn{2}{c}{12+log(O/H)} & SFR & L$_{TIR}$\\\\ & Morphology & (J2000) & (J2000) & (arcmin) & (arcmin)& (Mpc) & (PT) & (KK) & (\\msun\\ yr$^{-1}$) & ($\\times$ 10$^{9}$ \\lsun) \\\\ (1) & (2) & (3) & (4) & (5) & (6) & (7) & \\multicolumn{2}{c}{(8)} & (9) & (10) \\\\ &&&&&&&&&&\\\\ \\hline &&&&&&&&&&\\\\ NGC~337 & SBd & 00h 59m 50.7s & - 07d 34' 44\" & 2.9 & 1.8 & 19.3 & 8.18 & 8.84 & 1.30 & 12.0\\\\ NGC~628 & SAc & 01h 36m 41.8s & ~15d 47' 17\" & 10.5 & 9.5 & 7.2 & 8.35 & 9.02 & 0.68 & 8.0 \\\\ NGC~1097 & SBb & 02h 46m 18.0s & - 30d 16' 42\" & 9.3 & 6.3 & 14.2 & 8.47 & 9.09 & 4.17 &45.0 \\\\ NGC~1291 & SB0/a & 03h 17m 19.1s & - 41d 06' 32\" & 9.8 & 8.1 & 10.4 & 8.52 & 9.20 & 0.35 & 3.5\\\\ NGC~1316 & SAB0 & 03h 22m 41.2s & - 37d 12' 10\" & 12.0 & 8.5 & 21.0 & 8.77 & 9.52 & - & 8.0 \\\\ NGC~1512 & SBab & 04h 03m 55.0s & - 43d 20' 44\" & 8.9 & 5.6 & 11.6 & 8.56 & 9.11 & 0.36 & 3.8\\\\ NGC~3351 & SBb & 10h 43m 57.5s & ~11d 42' 19\" & 7.4 & 5.0 & 9.3 & 8.60 & 9.19 & 0.58 & 8.1\\\\ NGC~3621 & SAd & 11h 18m 18.3s & - 32d 48' 55\" & 12.3 & 7.1 & 6.5 & 8.27 & 8.80 & 0.51 & 7.9\\\\ NGC~3627 & SABb & 11h 20m 13.4s & ~12d 59' 27\" & 9.1 & 4.2 & 9.4 & 8.34 & 8.99 & 1.7 & 28.0\\\\ NGC~4826 & SAab & 12h 56m 42.8s & ~21d 40' 50\" & 10.0 & 5.4 & 5.3 & 8.54 & 9.20 & 0.26 & 4.2\\\\ NGC~7793 & SAd & 23h 57m 50.4s & - 32d 35' 30\" & 9.3 & 6.3 & 3.9 & 8.31 & 8.88 & 0.26 & 2.3\\\\ &&&&&&&&&&\\\\ \\hline \\end{tabular} \\begin{list}{}{} \\item[$^{(1)}$] {\\small Galaxy Name} \\item[$^{(2)}$] {\\small Morphological type} \\item[$^{(3)}$] J2000.0 Right ascension \\item[$^{(4)}$] J2000.0 Declination \\item[$^{(5)}$] Approximate major diameter in arcminutes \\item[$^{(6)}$] Approximate minor diameter in arcminutes \\item[$^{(7)}$] Distance in megaparsecs \\item[$^{(8)}$] Mean disk oxygen abundance from \\citet{Moustakas2010} \\item[$^{(9)}$] Star formation rate in solar masses per year from Table 1 of \\citet{Calzetti2010_2} \\item[$^{(10)}$] Total infrared luminosities in the 3-to-1100 \\mic\\ in units of 10$^{9}$ solar from \\citet{Kennicutt2011} \\end{list} \\end{table*} %---------------------------------------------------------------------------------------------------------------------------- It is crucial to yield an exhaustive inventory on the range and distribution of dust temperatures as well as the radiative heating sources of dust grains. IRAS data combined with ground-based observations (SCUBA on JCMT for instance, observing at 450 and 850 \\mic), have first led to the detection of the cool phases of dust in nearby galaxies \\citep{Dunne2000,Dunne_Eales_2001,James2002}. The Infrared Space Observatory (ISO) and the \\spitz\\ {\\it Space Telescope} (with a better sensitivity) then enabled us to resolve in detail the warm and cool dust reservoirs and more accurately model the Spectral Energy Distributions (SED) of nearby objects up to 160 \\mic\\ \\citep{Draine2007}. Several issues still have to be investigated. Numerous studies thus highlighted the necessity of observing galaxies longward of 160 \\mic\\ at high resolution to properly sample the submillimeter (submm) regime of these SEDs and better characterise the total temperature range and dust masses \\citep[e.g.][]{Gordon2010}. Moreover, IR emission is known to be a good tracer of the embedded star formation in galaxies \\citep[][]{Kennicutt1998,Calzetti2007,Calzetti2010_2}. Nevertheless, the amount of IR emission not directly linked with the ongoing star formation is still poorly quantified. To address this issue, we need to understand the heating sources of the different grain populations: dust heating by the young stellar populations in H{\\sc ii} and photodissociation regions (PDR), UV photons escaping from those regions or old stellar population. The \\hersc\\ {\\it Space Observatory} \\citep{Pilbratt2010} is currently observing the FIR/submm wavelength range to probe the coldest phases of dust that were poorly unexplored. Its onboard photometers PACS (from 70 to 160 \\mic) and SPIRE (from 250 to 500 \\mic) enable the imaging of nearby galaxies at a resolution and sensitivity never reached to date by a space telescope. The \\pck\\ satellite, launched by the same rocket, performs an all-sky survey at nine wavelengths between 350 \\mic\\ and 1 cm, complementing \\hersc\\ coverage, at a lower resolution. \\citet{Planck_collabo_2011_NearbyGalaxies} studies have suggested that some nearby galaxies are both luminous and cool (presence of dust with T$<$20K), with properties comparable to those of galaxies observed with \\hersc\\ at higher redshifts \\citep{Rowan-Robinson2010,Magnelli2012}. Using \\hersc\\ observations, \\citet{Engelbracht2010} compute separate SEDs for the center and the disks of nearby galaxies and find that the cool dust is on average 15$\\%$ hotter in the center than in the disk and correlate this effect with morphological type and bar strength. Using a radial approach, \\citet{Pohlen2010} show that SPIRE surface brightness ratios - used as a proxy for cold dust temperatures - seem to decrease with radius in spirals, and conclude that dust in the outer regions is colder than in the centre of the galaxies. \\citet{Skibba2011} also investigate the relation between dust heating sources and the star formation activity with metallicity and the morphological types of their galaxies. On a more resolved scale, \\citet{Bendo2011} derive PACS and SPIRE ratio maps of spiral galaxies. They find that those ratios seem to be preferentially correlated with the global starlight, including evolved stars, than with emission from star forming regions, results consistent with those obtained by \\citet{Boquien2011} for the spiral local group galaxy M33 or by \\citet{Auld2012} for extragalactic dust clouds of Centaurus A, among others. This supports the idea of a different source of heating for the cold (15-30 K) dust than that of the warm dust. Using \\hersc\\ observations of the dwarf irregular galaxy NGC~6822, \\citet{Galametz2010} show on the contrary that SPIRE surface brightness ratios are strongly correlated with the 24 \\mic\\ surface brightness, a commonly used tracer of the star formation activity. This might indicate that star forming regions are non negligible heating sources of cold dust in that object and suggest that the source of the dust heating probably differ with the environment and vary from normal spirals to low metallicity starbursts. Observations of galaxies with \\hersc\\ and from the ground, have also suggested that the properties of cold dust grains could differ from those often used in SED models such as blackbody modified with a beta-2 emissivity law or use of graphite grains to model carbon dust. In metal-poor objects in particular, using graphite grains can lead to unphysical gas-to-dust mass ratios (G/D) compared to those expected from the metallicity of the galaxy \\citep{Meixner2010,Galametz2010}. Moreover, observations of dwarf galaxies in submm often lead to the detection of excess emission compared to extrapolations from \\spitz-based fits \\citep[]{Dumke2004, Galliano2003, Galliano2005, Bendo2006, Marleau2006, Galametz2009,OHalloran2010, Bot2010_2}. The origin of this excess is still an open issue and several explanations have been investigated so far: variations of the emissivity of the dust grains with temperature, ``spinning dust\" emission, cold dust etc. Many investigations still need to be carried out to disentangle between those different hypotheses. We have obtained photometric \\hersc/PACS (70 to 160 \\mic) and \\hersc/SPIRE (250 to 500 \\mic) observations of nearby galaxies as part of the open time key programme KINGFISH \\citep[Key Insights on Nearby Galaxies: A Far-Infrared Survey with \\hersc,][]{Kennicutt2011}. The KINGFISH sample comprises 61 galaxies probing a wide range of galaxy types (from elliptical to irregular galaxies), metallicity (7.3 $\\le$ 12+log(O/H) $\\le$ 9.3) and star-forming activity (10$^{-3}$ $\\le$ SFR $\\le$ 7 \\msun\\ yr$^{-1}$). The KINGFISH galaxies are located between 3 and 31 Mpc, leading to ISM resolution elements of 0.5 to 5.4 kpc at SPIRE 500 \\mic\\ resolution (36\\arcsec). We refer to \\citet{Kennicutt2011} for a detailed description of the sample and the observation strategy. We aim to benefit from the good coverage and resolution of \\hersc\\ to sample the complete thermal dust emission spectrum and probe the spatially-resolved properties of the cold dust grain population. We refer to the work of \\citet{Aniano2012a} (hereafter [A12a]) and Aniano et al. 2012b (in prep, full KINGFISH sample, hereafter [A12b]) for an analysis of other dust properties (PAH contribution to the total dust mass, distribution of the starlight intensity, dust mass surface density, among others) in the KINGFISH sample using the \\citet{Draine_Li_2007} dust models. The two studies are complementary. The following paper particularly focuses on investigating the distribution of cold dust temperatures and dust mass surface densities in our objects. The SED model applied in this work (two modified-blackbody fit) allows for variations in the emissivity index of the dust grains and enables us to test the robustness of applying different assumptions of emissivity index to our extragalactic objects and quantify how this affects the temperatures and masses derived. This paper is structured as follows. In $\\S$2, we describe the \\hersc\\ observations and the convolution techniques we apply to degrade images to the lowest resolution (that of SPIRE 500 \\mic). In $\\S$3, we discuss the integrated properties of the cold dust (temperature and emissivity). We also study how uncertainties in the flux measurements can lead to degeneracies between parameters. In $\\S$4, we produce maps of the dust properties of our galaxies and discuss the choice of a fixed or free emissivity parameter on the temperature and emissivity maps obtained. We also study the distribution of the dust mass surface densities in our objects. We finally analyze how the total dust mass depends on the type of modified blackbody model we apply, the wavelength coverage and the working resolution. %=============================================================== %Observations and data reduction %=============================================================== ", "conclusions": "We combine \\spitz/MIPS and \\hersc\\ data to derive global properties of the cold dust in a sample of 11 nearby galaxies using a two-temperature fit to model their SEDs. At global scales, \\begin{itemize} \\item Our objects show a wide range of SEDs, with a clear evolution of the 100-to-24 \\mic\\ flux density ratios from actively to quiescent star-forming objects and variations on the global cold dust parameter $\\beta$$_c$ we derive, from $\\beta$$_c$ = 0.8 to 2.5. \\item We perform a Monte Carlo simulation in order to quantify the effect of uncertainties on the parameters derived from our fitting technique. We show that even with an exhaustive sampling of the global thermal dust emission, an artificial T-$\\beta$ anti-correlation can be created by uncertainties in the flux measurements. \\end{itemize} We use the unprecedented spatial resolution of \\hersc\\ to derive resolved dust properties within our objects and study the robustness of our result with our SED model assumptions (emissivity index, SED coverage, resolution). Our study at local scales shows that: \\begin{itemize} \\item When we fix the cold emissivity index to a standard $\\beta$$_c$=2.0, we obtain a smooth distribution of the cold temperatures, with a radial decrease toward the outer parts of galaxies. In some objects, the temperature distribution follows the distribution of star-forming regions (NGC~3627 for instance). Fixing $\\beta$$_c$ to 1.5 or 1.0 influences the temperature range derived but does not strongly modify the temperature relative distribution. \\item When $\\beta$$_c$ is free, barred spirals (e.g. NGC~1097, NGC~1316, NGC~1512, NGC~3351 or NGC~3627) show temperature distributions that are similar to those obtained with a fixed $\\beta$$_c$ and homogeneous emissivity index values across the galaxies. For non-barred galaxies, $\\beta$$_c$ seems to decrease with radius. We nevertheless obtain homogeneous cold temperature distributions for those objects. The fact that these temperature maps do not seem to relate to any dust heating sources questions the use of a free $\\beta$ in those objects. Applying a $\\beta$-free model to deduce an average emissivity index across the galaxy and derive the temperature map using this median value could help to avoid potential degeneracies between T and $\\beta$ in those objects. \\item Gathering our local results in a T$_c$-$\\beta$$_c$ diagram clearly leads to a correlation between the two parameters. This correlation could be linked with {\\it 1)} a real emissivity variation of dust grains with temperature {\\it 2)} noise effects and uncertainties on flux measurements {\\it 3)} temperature mixing along the line of sight. \\item The dust mass surface density usually peaks in the center of galaxies but also seems to coincide with extra-nuclear peaks of star formation at the end of the bars (usually corresponding to molecular gas reservoirs). \\item The dust mass estimates seem to be affected by resolution, with a systematic increase of total dust masses in the ``resolved\" case compared to dust masses globally determined. \\item The submm coverage of our SEDs with \\hersc\\ observations modifies the dust masses compared to those derived from \\spitz-based fits. Two MBB fits performed with data up to 250 \\mic\\ differ from those derived using the complete SPIRE coverage (data up to 500 \\mic) by up to 17$\\%$ if $\\beta$$_c$=2.0, up to 30$\\%$ if $\\beta$$_c$=1.0. \\item Dust masses derived globally or on a local basis using the \\citet{Draine_Li_2007} model only differ from models using modified blackbodies and $\\beta$$_c$=2.0 by less than 30$\\%$. Dust masses derived from modified blackbody fits using a shallower parameter $\\beta$$_c$ are systematically lower that those derived with $\\beta$$_c$=2.0 (lower by 25 to 46 $\\%$ if $\\beta$=1.5 for instance in our sample). Our results highlight the necessity to better investigate the variations of the cold dust properties to properly quantify the total dust mass if we aim for instance, to use it as a gas tracer. \\end{itemize} This study aimed to investigate the properties of cold dust unveiled by \\hersc\\ observations of nearby galaxies. Questions remain due to possible degeneracies between the grain emissivities and temperatures. Moreover, this paper uses isothermal fits, which implies that the intrinsic emissivity index of the dust grains and the emissivity index directly obtained from our fits are one and the same. The presence of several dust grain temperatures in our objects both locally and at global scales also makes a direct extrapolation of dust properties difficult. Nevertheless, combining those results with other derived quantities such as the elemental abundances in the ISM or reliable gas-to-dust mass ratios as well as further investigations on closer objects will help us to disentangle between the different hypotheses. %=============================================================== % Acknowledgements %===============================================================" }, "1207/1207.5732_arXiv.txt": { "abstract": "Studies of the Galactic population of radio pulsars have shown that their luminosity distribution appears to be log-normal in form. We investigate some of the consequences that occur when one applies this functional form to populations of pulsars in globular clusters. We use Bayesian methods to explore constraints on the mean and standard deviation of the luminosity function, as well as the total number of pulsars, given an observed sample of pulsars down to some limiting flux density, accounting for measurements of flux densities of individual pulsars as well as diffuse emission from the direction of the cluster. We apply our analysis to Terzan~5, 47~Tucanae and M~28, and demonstrate, under reasonable assumptions, that the number of potentially observable pulsars should be within 95\\% credible intervals of $147^{+112}_{-65}$, $83^{+54}_{-35}$ and $100^{+91}_{-52}$, respectively. Beaming considerations would increase the true population size by approximately a factor of two. Using non-informative priors, however, the constraints are not tight due to the paucity and quality of flux density measurements. Future cluster pulsar discoveries and improved flux density measurements would allow this method to be used to more accurately constrain the luminosity function, and to compare the luminosity function between different clusters. ", "introduction": "Globular clusters are spherical collections of stars located throughout the haloes of galaxies. Once thought to be composed entirely of old metal-poor population II stars, they are now also believed to form during interactions or collisions of galaxies, therefore containing younger stars having higher metallicities \\citep[see][]{zep03}. The total masses of globular clusters range up to the order of $10^6 M_{\\sun}$ \\citep[see][]{mey97}, and core stellar number densities reach $10^6$ pc$^{-3}$. The high core densities lead to dynamical interactions between stellar systems that are found less commonly in the Galactic plane. For example, globular clusters favour the formation of low-mass X-ray binaries (LMXBs) that are believed to be the progenitors of millisecond pulsars \\citep[MSPs;][]{alp82,arc09}, and hence, the fraction of MSPs among all pulsars in globular clusters is much larger than that in the Galactic field ($\\sim$97\\% versus $\\sim$11\\%). In addition, the binary MSPs in globular clusters tend to have higher eccentricities compared to their field counterparts, due to exchange or fly-by encounters. MSPs, due to their formation history, can be considered long-lived tracers of LMXBs, and therefore, constraints on the MSP content of globular clusters provide unique insights into binary evolution and the integrated dynamical history of globular clusters, while determining the radio luminosity function of these pulsars helps shed light on the radio emission mechanism in action in these compact objects. Pulsar searches of globular clusters have yielded impressive returns in recent years \\citep[see][]{cam05}, with currently 144 pulsars known in 28 clusters\\footnote{See Paulo Freire's globular cluster pulsar catalogue at http://www.naic.edu/$\\sim$pfreire/GCpsr.html}. Of these, three clusters, Terzan 5, 47 Tucanae and M 28 are known to harbour more than 10 pulsars each, the most populous being Terzan 5 with 34 (Ransom S. M., private communication). In this paper, we describe a Bayesian method that we have developed to compute an estimate of the true number of pulsars in a given cluster, given an observed population. There have been many attempts to constrain the population size of pulsars in all globular clusters in the Galaxy \\citep*[see, for example,][]{kul90b,bag11}. This work is different in that it treats clusters individually instead of dealing with the total population. Bayesian approaches to constrain the pulsar population of individual globular clusters have been used specifically for the case of young (non-recycled) pulsars by \\citet{boy11}. This work focuses on the entire radio pulsar content of the cluster -- the majority of which is made up of old (recycled) pulsars -- and additionally, attempts to constrain luminosity function parameters jointly with population size. \\citet{fau06} have shown that the luminosity distribution of non-recycled pulsars in the Galactic field appears to be log-normal in form. More recently, \\citet{bag11} have verified that the observed luminosities of recycled pulsars in globular clusters are consistent with this result. Assuming, therefore, that there is no significant difference between the nature of Galactic and cluster populations, we investigate some of the consequences that occur when one applies this functional form to populations of pulsars in individual globular clusters. For a log-normal (base-10) distribution of pulsar luminosities, the luminosity function is given by \\begin{equation} f(\\log L) = \\frac{1}{\\sigma \\sqrt{2\\pi}} e^{-\\frac{(\\log L - \\mu)^2}{2 \\sigma^2}}, \\end{equation} where $L$ is the luminosity in mJy kpc$^2$, $\\mu$ is the mean and $\\sigma$ is the standard deviation of the distribution. We are interested in the situation where we observe $n$ pulsars with luminosities above some limiting luminosity $L_{\\rm min}$. Given this sample of pulsars, we ask what constraints we can place on their luminosity function parameters, in addition to the potentially observable population size $N$ (that is, the population of pulsars beaming towards the Earth). Another way of thinking about this problem is that there is a family of luminosity function parameters and population sizes that is consistent with an observation of $n$ pulsars above the luminosity limit of the survey, and we are analyzing the posterior probabilities of different members of this family given the observations. This paper is organized as follows: In \\S\\ref{sec_bayesian}, we describe our technique. In \\S\\ref{sec_apps}, we apply the technique to observations of a few globular clusters to determine the constraints on the luminosity function parameters and population size. Later, we refine our results using \\emph{a priori} information on the luminosity function parameters to get a better estimate of the number of pulsars in those clusters. A summary and our conclusions are presented in \\S\\ref{sec_conclusions}. ", "conclusions": "\\label{sec_conclusions} We have developed a Bayesian technique to constrain the luminosity function parameters and population size of pulsars in individual globular clusters, given a data set that consists of the number of observed pulsars, the flux densities of the individual pulsars in the cluster and the total diffuse flux emission from the direction of the globular cluster, assuming a log-normal luminosity function. We have applied our analysis to a few globular clusters and have demonstrated the utility of this technique in constraining the aforementioned parameters. Our technique is applied in two different ways -- first, with no prior information, and second, assuming prior knowledge of the possible ranges of $\\mu$ and $\\sigma$. As shown for Terzan 5, the results for the first approach do not constrain $N$, $\\mu$ or $\\sigma$ very well due to paucity of data, but the latter two do exhibit consistency with the values found by \\citet{fau06} and \\citet{bag11}. For the second approach in which we assume prior information to bound $\\mu$ and $\\sigma$, the priors help better constrain the total number of pulsars in the cluster. The technique we have developed here should prove useful in future studies of the globular cluster luminosity function where ongoing and future pulsar surveys are expected to provide a substantial increase in the observed populations of pulsars in many clusters. In particular, we anticipate that the increased amount of data would enable us to constrain the distributions of $\\mu$ and $\\sigma$ independently (i.e. without the need to assume prior information from the Galactic pulsar population). Further interferometric measurements of the diffuse radio flux in many globular clusters could provide improved constraints on $\\mu$ and $\\sigma$ by measuring the flux contribution from the individually unresolvable population of pulsars." }, "1207/1207.6111_arXiv.txt": { "abstract": "\\noindent We discuss a new class of tribrid inflation models in supergravity, where the shape of the inflaton potential is dominated by effects from the K\\\"{a}hler potential. Tribrid inflation is a variant of hybrid inflation which is particularly suited for connecting inflation with particle physics, since the inflaton can be a D-flat combination of charged fields from the matter sector. In models of tribrid inflation studied so far, the inflaton potential was dominated by either loop corrections or by mixing effects with the waterfall field (as in ``pseudosmooth'' tribrid inflation). Here we investigate the third possibility, namely that tribrid inflation is dominantly driven by effects from higher-dimensional operators of the K\\\"{a}hler potential. We specify for which superpotential parameters the new regime is realized and show how it can be experimentally distinguished from the other two (loop-driven and ``pseudosmooth'') regimes. ", "introduction": "Inflation is a very successful paradigm for solving the horizon and flatness problems and for creating the almost scale-invariant, adiabatic perturbations that have been observed in the cosmic microwave background (CMB) \\cite{Guth:1980zm,inf2,inf3,inf4}. However, it is not clear how cosmic inflation is realized within particle theory. Most importantly, the identity of the inflaton field, which drives the rapid expansion, is still unknown. Various models of inflation have been proposed in the literature, however often the inflaton is just an additional gauge singlet, rather disconnected from the rest of the theory.\\footnote{Some notable exceptions are, for example, inflection point inflation in the Minimal Supersymmetric Standard Model \\cite{Allahverdi:2006iq,mssmInflation}, Standard Model Higgs inflation \\cite{smHiggsInflation, smHiggsInflation3} and GUT Higgs inflation \\cite{gutHiggsInflation,gutHiggsInflation2}.} A proposed framework for connecting inflation with particle physics is supersymmetric hybrid inflation \\cite{susyHybrid,Dvali:1994ms,Linde:1997sj,RehmanHybrid} in which inflation is ended by a particle physics phase transition. The energy scale of this phase transition is around the Grand Unification (GUT) scale: $\\Lambda_{\\text{inflation}} \\sim M_{\\text{GUT}} \\sim 10^{16}$ GeV. This inspires hope that hybrid inflation may be connected to the spontaneous breaking of a GUT symmetry \\cite{Dvali:1994ms}. However, the inflaton itself must be a singlet in such models, and the predictions for the CMB fluctuations -- which depend on the properties of the inflaton -- usually cannot be related to other observables. This issue was improved by the development of supersymmetric tribrid inflation \\cite{sneutrinoHybrid,Antusch:2008pn,tribridBasic,Antusch:2009vg}, which is a variant of hybrid inflation where the inflaton itself can be charged under the symmetries of the particle theory \\cite{Antusch:2010va,Antusch:2011ei}. In particular, it can be a charged matter particle and might have observable effects in the low-energy theory. This could make it possible to determine the inflaton couplings both from particle physics constraints and from measurements of the CMB at the same time. In models of tribrid inflation studied so far, the inflaton potential was dominantly generated by either one-loop corrections from the Coleman-Weinberg potential, or by small vacuum expectation values of non-inflaton fields like the waterfall field. We refer to these regimes as loop-driven and (pseudo)smooth \\cite{pseudosmooth} tribrid inflation. The goal of this paper is to complete the discussion of variants of tribrid inflation by analyzing the third possibility, namely the case that tribrid inflation is driven dominantly by effects from higher-dimensional operators of the \\kahler potential. Based on a generalized superpotential for tribrid inflation we specify for which parameters the \\kahlerdriven regime is realized, calculate the slow-roll predictions and show how it can be experimentally distinguished from the other two regimes by a measurement of the running of the spectral index. We also discuss how the new class of inflation models might be embedded into supersymmetric GUT and/or flavour models. The rest of the paper is structured as follows: We first motivate and introduce tribrid inflation in section \\ref{sec:generalizedTribrid}. In sections \\ref{sec:scalarPotential} -- \\ref{sec:loops} we discuss the \\kahlerdriven regime and derive the predictions using a power series expansion of the \\kahler potential. Section \\ref{sec:summary} summarizes our findings and compares our results for \\kahlerdriven with loop-driven and pseudosmooth tribrid inflation. There we also discuss how our results can be used for model building, using tribrid inflation as a framework for building explicit particle physics models of inflation. ", "conclusions": "\\label{sec:summary} \\subsubsection*{The three regimes of tribrid inflation} The generalized tribrid superpotential of eq.~\\eqref{eq:generalizedTribridW} can account for tribrid inflation in various ways. Apart from the formerly known regime where the inflaton potential is generated by loop corrections (see e.g.\\ \\cite{valerieInflation}) and the recently discussed pseudosmooth regime \\cite{pseudosmooth}, we have shown that tribrid inflation can also happen in a \\kahlerdriven regime where the inflaton potential is mostly generated by Planck-suppressed operators in the \\kahler potential. The three possible regimes, along with the required superpotential parameters $l$ and $m$, are shown in table~\\ref{tab:summaryFinalML}. The K\\\"{a}hler- and loop-driven regimes require $l \\geq m = 2$ to correctly end with a waterfall transition, whereas the pseudosmooth regime requires $l \\geq m > 2$ to have suitable small-field trajectories. \\begin{table}[hbt] \\centering \\begin{tabular}{ | l | c | c | c | c | } \\hline & $m=1$ & $m=2$ & $m=3$ & $m = 4$ \\\\ \\hline \\multirow{2}{*}{$l=2$} & \\multirow{2}{*}{\\textcolor{red}{not possible}} & \\textcolor[rgb]{0,0.55,0}{K\\\"{a}hler-driven} & \\multirow{2}{*}{\\textcolor{red}{not possible}} & \\multirow{2}{*}{\\textcolor{red}{not possible}} \\\\ & & \\textcolor[rgb]{0,0.55,0}{or loop-driven} & & \\\\ \\hline \\multirow{2}{*}{$l=3$} & \\multirow{2}{*}{\\textcolor{red}{not possible}} & \\textcolor[rgb]{0,0.55,0}{K\\\"{a}hler-driven} & \\multirow{2}{*}{\\textcolor[rgb]{0,0.55,0}{pseudosmooth}} & \\multirow{2}{*}{\\textcolor{red}{not possible}} \\\\ & & \\textcolor[rgb]{1,0.5,0}{or loop-driven} & & \\\\ \\hline \\multirow{2}{*}{$l=4$} & \\multirow{2}{*}{\\textcolor{red}{not possible}} & \\textcolor[rgb]{0,0.55,0}{K\\\"{a}hler-driven} & \\textcolor[rgb]{1,0.5,0}{pseudosmooth} & \\multirow{2}{*}{\\textcolor[rgb]{0,0.55,0}{pseudosmooth}} \\\\ & & \\textcolor[rgb]{1,0.5,0}{or loop-driven} & \\textcolor[rgb]{1,0.5,0}{or smooth} & \\\\ \\hline \\end{tabular} \\caption{Viable (green) and dysfunctional (red) superpotential choices for supersymmetric tribrid inflation. The text indicates which regimes of tribrid inflation are possible for a given combination of $l$ and $m$. For the green entries, it has been explicitly shown that slow-roll inflation in agreement with observations is possible. The orange entries satisfy the necessary conditions and may or may not feature suitable slow-roll trajectories.} \\label{tab:summaryFinalML} \\end{table} It is interesting to note that although all three regimes make similar predictions for $r \\lesssim 0.01$ and $\\braket{H} \\gtrsim O(10^{16} \\, \\text{GeV})$, the \\kahlerdriven regime can be distinguished from the other two regimes by a measurement of $\\alpha_s$, which is predicted to be $\\lvert \\alpha_s \\rvert < 10^{-3}$ \\cite{pseudosmooth,valerieInflation} for the loop-driven and pseudosmooth regimes and $\\alpha_s > 0$ for the \\kahlerdriven regime. \\subsubsection*{\\kahlerdriven tribrid inflation} In this paper, the \\kahlerdriven regime of tribrid inflation was analyzed for the first time in detail, using a generic expansion of the \\kahler potential. We noticed that the large number of model parameters can be mapped to an effective model with only three free parameters: $a$ and $b$, which can be calculated from the \\kahler potential using eqs.~\\eqref{eq:kahlerAK}--\\eqref{eq:kahlerBK}, and the spectral index $n_s$, which will be measured with improved precision by the Planck satellite. Using these parameters, we analytically derived the slow-roll predictions for the tensor-to-scalar ratio $r$ and for the running of the spectral index $\\alpha_s$ (fig.~\\ref{fig:kahlerRAlphaSResult}). We also calculated constraints on the model parameters $\\braket{H}$ (fig.~\\ref{fig:kahlerHiggsVev}) and $\\lambda$ (fig.~\\ref{fig:kahlerLambda}) depending on $a$, $b$ and $n_s$ for the different choices of $l$ and $n$. Successful inflation can occur for a large range of parameters. The required tuning of the \\kahler potential, which is typical for supergravity theories, is pretty mild: it is sufficient to tune a single parameter $\\kappa_{101} \\simeq 0.94$--$0.99$, while all other parameters in the \\kahler potential can take values of $O(1)$. \\emptyline Our results can be used for model-building in various ways. If one knows the \\kahler potential from some UV completion, one can calculate $a$ and $b$ using eqs.~\\eqref{eq:kahlerAK}--\\eqref{eq:kahlerBK}. The spectral index can be fixed by the upcoming measurement by the Planck satellite. One can then read off the constraints on the superpotential parameters $\\lambda$ and $\\braket{H}$ and the CMB observable $\\alpha_s$ from our plots or calculate the numerical value from our analytical slow-roll results (eqs.~\\eqref{eq:phi0}, \\eqref{eq:kahlerAlphaSResult}, \\eqref{eq:V0}, \\eqref{eq:phic}, \\eqref{eq:lambdaKahler3} and \\eqref{eq:lambdaKahler2}). If one does not know the \\kahler potential, then $a$ could be fixed by a future measurement of $\\alpha_s$. This suffices for an estimate of $\\braket{H}$ and $\\lambda$, as they do not strongly depend on $b$. One can also estimate the parameter $a$ by a measurement of the superpotential coupling $\\lambda$; when tribrid inflation is embedded into a full particle physics model, this coupling may be observable in the low-energy theory. In sneutrino tribrid inflation, as discussed for example in the loop-driven regime (with $l = m = 2$) in \\cite{valerieInflation}, $\\lambda$ is related to the observable neutrino mass. \\emptyline We took particular care to check our approximations: We explicitly showed that the slow-roll parameters remain small throughout inflation and that our semiclassical approximation is not spoiled by quantum fluctuations. We estimated the effects of higher-dimensional operators, which turn out to be small, and proved that loop effects are negligible for a large part of parameter space, especially for $l > 2$ (fig.~\\ref{fig:kahlerLoopPotential}). We also compared our analytical approximations with numerical computations, and found excellent agreement. We therefore believe that our treatment covers the full range of K\\\"{a}hler-driven slow-roll tribrid inflation models with good accuracy. \\subsubsection*{Outlook} Together with the papers on loop-driven (see e.g.\\ \\cite{valerieInflation}) and pseudosmooth \\cite{pseudosmooth} tribrid inflation, this paper provides a blueprint for applying tribrid inflation to supersymmetric model-building. It has been shown that a variety of superpotentials can feature tribrid inflation, in particular all superpotentials of the form of eq.~\\eqref{eq:generalizedTribridW} with $l \\geq m = 2$ or $l=m>2$. It should now be possible to identify particle physics models which can contain the required superpotential terms. The inflaton can be a D-flat combination of (charged) matter fields \\cite{Antusch:2010va}, and the waterfall phase transition might be identified with the spontaneous breaking of, e.g., a GUT symmetry \\cite{Dvali:1994ms} or family symmetry \\cite{Antusch:2008gw}. Once a particle theory with a suitable superpotential has been found, the regime of tribrid inflation can be determined from $l$ and $m$ using table~\\ref{tab:summaryFinalML}, and the inflationary predictions can be read off the graphs in the corresponding paper. As we have explained, this works particularly well for the K\\\"{a}hler-driven regime, where we have derived strong correlations between several model parameters and the cosmological observable $\\alpha_s$. We conclude that models of supersymmetric tribrid inflation are indeed promising candidates for realizing inflation in close contact with particle physics." }, "1207/1207.5112_arXiv.txt": { "abstract": "We report on the discovery of large-amplitude flickering from V648 Car (= SS73-17), a poorly studied object listed amongst the very few hard X-ray emitting symbiotic stars. We performed milli-magnitude precision optical photometry with the Swope Telescope at the Las Campanas Observatory, Chile, and found that V648 Car shows large U-band variability over time scales of minutes. To our knowledge, it is amongst the largest flickering of a symbiotic star ever reported. Our finding supports the hypothesis that symbiotic WDs producing hard X-rays are predominantly powered by accretion, rather than quasi-steady nuclear burning, and have masses close to the Chandrasekhar limit. No significant periodicity is evident from the flickering light curve. The ASAS long-term V light curve suggests the presence of a tidally distorted giant accreting via Roche Lobe overflow, and a binary period of $\\sim520$ days. On the basis of the outstanding physical properties of V648 Car as hinted by its fast and long-term optical variability, as well as by its nature as hard X-ray emitter, we therefore call for simultaneous follow-up observations in different bands, ideally combined with time-resolved optical spectroscopy. ", "introduction": "V648 Car was a forgotten red variable star of magnitude $V \\sim$ 10 until it was realized its nature as a hard X--ray emitter, i.e. when source IGR J10109-5746 was discovered at hard X--rays by the {\\it INTEGRAL} satellite (Kuiper et al. 2006; Bird et al. 2010). Indeed, V648 Car lies within the error circle of this high-energy source, and Masetti et al. (2006) first suggested that the red star could be the optical counterpart of the hard X--ray emitter, mainly on the basis of the positional correlation, subsequently confirmed by the detection of a soft X--ray counterpart with arcsecond-sized error on its position (Tueller et al. 2005).\\\\ The high-energy behavior of this object was explored with several X--ray studies with {\\it Suzaku} (Smith et al. 2008), {\\it Swift} (Kennea et al. 2009) and {\\it Chandra} (Eze et al. 2010). These showed that the spectrum of the source is described with a highly absorbed thermal plasma model of temperature kT $\\approx$ 10 keV with superimposed emission lines of ionized iron. Concerning the light curve, the {\\it X-ray Telescope (XRT)} on board of Swift did not detect any rapid (100 to 5000 s) periodic variability, which led Kennea et al. (2009) to state that the X--rays come more likely from the boundary layer around a massive (estimated between 1.1 and 1.35 M$_\\odot$) white dwarf (WD). At these extensive studies of V648 Car in the X--ray domain it has not yet corresponded a parallel detailed investigation in the optical, if we exclude its previous inclusion in both the Sanduleak \\& Stephenson's catalog of emission line objects of the southern Milky Way (where it appears as SS73-17 and is classified as a M3ep+OB star -- Sanduleak \\& Stephenson 1973) and in the similar catalog by Henize (who lists SS73-17 as Hen 3-380 and describes it as having Balmer lines as well as single ionized calcium in emission -- Henize 1976). Later spectroscopic observations of 33 unclassified emission line stars from the original Sanduleak \\& Stephenson's list briefly reported on V648 Car being an ``example of Mira type star [...] having H$\\alpha$ and H$\\beta$ in emission'' (Pereira et al. 2003).\\\\ The M4\\,III spectral type as derived by Pereira et al. from the relative strength of the TiO spectral bands, and an apparent blue excess (see their Fig. 1), eventually brought Masetti et al. (2006) to candidate V648 Car as symbiotic star, classification further advocated by the analogy with RT Cru, a previous case of association between a symbiotic star and a newly-discovered {\\it INTEGRAL} hard X--ray source (Chernyakova et al. 2005; Masetti et al. 2005; Tueller et al. 2005; Luna \\& Sokoloski 2007).\\\\ Symbiotic stars are long-period interacting binaries composed of a hot compact star -- generally a WD -- and an evolved giant star, whose mutual interaction via accretion processes triggers the extended emission recorded from radio to X-rays. Out of $\\sim200$ symbiotic stars known today (Belczy{\\'n}ski et al. 2000), only 5 have been reported to emit significantly above 5 keV (V648 Car, RT Cru, T CrB, CH Cyg and MWC 560). It has been argued that the number of hard X--ray emitting symbiotic stars in the Galaxy may actually be much larger, and they have also been proposed as candidates for Ia supernovae progenitors because all the systems known to date seem to host WDs whose masses approach the Chandrasekhar limit (Eze 2011). In the framework of an observing campaign aimed at characterizing the WD properties in different type of accreting binary stars, we have started a systematic search for flickering (i.e., stochastic photometric variations on time-scales of a few minutes) in southern symbiotic stars (Angeloni et al. 2012, in preparation). This survey has been conceived as the complementary extension of the survey by Sokoloski et al. (2001) to the southern hemisphere and to dimmer objects. For our first ``test-run'', we have recognized V648 Car as a priority target by virtue of its strong similarity with CH Cyg and MWC560, two other hard X--ray symbiotic stars known to feed powerful jets and to show optical flickering (Leedj{\\\"a}rv 2004). In this letter, we present the outcome of our fast photometric monitoring of V648 Car conducted with the Swope Telescope at the Las Campanas Observatory, Chile. We first describe the observations and the data reduction/analysis process in Sect. 2. The results are presented in Sect. 3, then discussed in Sect. 4. Concluding remarks follow in Sect. 5. ", "conclusions": "\\label{discussion} We have discovered fast, large-amplitude flickering of V648 car, a poorly studied object listed amongst the very few hard X-ray emitting symbiotic stars. With an overall U-band amplitude of $\\sim0.6$ mag over time scales of minutes, it is one of the largest flickering ever reported from a symbiotic star. This finding is a further confirmation that symbiotic WDs producing hard X-rays are predominantly powered by accretion (rather than quasi-steady nuclear burning on the surface of the WD), and therefore that the insights about the masses of these WDs based on this interpretation are fundamentally correct. The ASAS light curve suggests the presence of a tidally distorted giant accreting via RLO, and a binary period of $\\sim520$ days. The outstanding character of V648 Car, as reconstructed from the puzzle of its fast and long-term optical variability and by its hard X-ray emitting properties, clearly demands a coordinated follow-up: e.g., simultaneous X/UV/optical observations in different bands, ideally combined with time-resolved optical spectroscopy." }, "1207/1207.7117_arXiv.txt": { "abstract": "Halos are biased tracers of the dark matter distribution. It is often assumed that the initial patches from which halos formed are locally biased with respect to the initial fluctuation field, meaning that the halo-patch fluctuation field can be written as a Taylor series in the dark matter density fluctuation field. If quantities other than the local density influence halo formation, then this Lagrangian bias will generically be nonlocal; the Taylor series must be performed with respect to these other variables as well. We illustrate the effect with Monte-Carlo simulations of a model in which halo formation depends on the local shear (the quadrupole of perturbation theory), and provide an analytic model which provides a good description of our results. Our model, which extends the excursion set approach to walks in more than one dimension, works both when steps in the walk are uncorrelated, as well as when there are correlations between steps. For walks with correlated steps, our model includes two distinct types of nonlocality: one is due to the fact that the initial density profile around a patch which is destined to form a halo must fall sufficiently steeply around it -- this introduces $k$-dependence to even the linear bias factor, but otherwise only affects the monopole of the clustering signal. The other is due to the surrounding shear field; this affects the quadratic and higher order bias factors, and introduces an angular dependence to the clustering signal. In both cases, our analysis shows that these nonlocal Lagrangian bias terms can be significant, particularly for massive halos; they must be accounted for in, e.g., analyses of higher-order clustering in Lagrangian or Eulerian space. Comparison of our predictions with measurements of the halo bispectrum in simulations is encouraging. Although we illustrate these effects using halos, our analysis and conclusions also apply to the other constituents of the cosmic web -- filaments, sheets and voids. ", "introduction": "The virialized halos which are identified in simulations of gravitational clustering are biased tracers of the underlying matter field. Typically, this bias is described in two ways, either by relating the halo and mass fields at the time the halos were identified (e.g., the present), or by identifying the patches in the initial conditions which are destined for form halos, and describing the bias between these patches and the initial mass fluctuation field \\cite{mw1996}. These are known as Eulerian and Lagrangian bias, respectively. In either case, the simplest models assume that this bias is local, meaning that the biased field can be written as a (deterministic) function of the mass field. However, the nonlinearly evolved mass field is a nonlocal function of the initial one, so Lagrangian and Eulerian bias cannot both be local \\cite{clmp1998, css2012, bsdm2012}. It has recently been noted that neither of the two best studied models of Lagrangian bias, peaks theory \\cite{bbks1986} and the excursion set approach \\cite{bcek1991}, are local. This is because both approaches predict that the abundance of biased tracers (peaks or halos) should depend, not just on the local values of the overdensity field, but on derivatives of the field as well \\cite{ms2012}. This gives rise to a rather specific form for nonlocal bias, in which the bias is most naturally described in Fourier space, where it is $k$-dependent, even at the linear level \\cite{dcss2010, mps2012, ps2012b}. The main goal of the present work is to explore models in which the nonlocality of bias is qualitatively different, arising only at second order, and associated with anisotropies in the initial field. This source of nonlocality is generic to models in which halos form from an anisotropic collapse \\cite{bm1996}, for which there is considerable evidence \\cite{smt2001}. Section~\\ref{biasmodels} describes the relation between Lagrangian and Eulerian bias in nonlocal models, and Section~\\ref{x6d} describes our excursion set based treatment of the origin of nonlocal Lagrangian bias, providing an analytic description of the effect. The technical problem which is the subject of this section is to provide an accurate formula for the first crossing distribution of a barrier by $n$-dimensional walks with correlated steps; we use Monte-Carlo simulations of such walks to illustrate the accuracy of our analytic formulae. Section~\\ref{carmen} compares our predictions with estimates of the nonlocal halo bias term in numerical simulations of hierarchical clustering, and a final section summarizes our findings. ", "conclusions": "If halo bias is nonlocal in Lagrangian space (Eq.~\\ref{dhnonlocal}), then this will add nonlocality in Eulerian space bias (Eq.~\\ref{nonlocalbEul}). We provided an explicit calculation of this effect, in which halo formation depends on the initial local density $\\delta_0$ and shear field $q_0^2$. In the excursion set approach, the problem of estimating halo abundances reduces to solving for the first crossing distribution of a suitably chosen barrier (Eq.~\\ref{Bq}) by $6$-dimensional random walks. In particular, we argued that a barrier which is linear in $q_0$ should provide a good first approximation to the physics of halo collapse (Figure~\\ref{gif2Bq}). And we provided a simple but accurate analytic approximation for the first crossing distribution for the case in which walks have uncorrelated steps (Eq.~\\ref{vfv}); the approximation follows from treating the full 6-dimensional problem as an effective one-dimensional one (Eq.~\\ref{Bqapprox}). Predictions for halo bias come from studying walks which do not start from the origin. We argued that the associated first crossing distribution is best thought of in terms of a shifted barrier (Eq.~\\ref{Bshifted}), from which it is straightforward to derive formulae for halo bias formulae (Eqs.~\\ref{b1b2} and~\\ref{c2}). These formulae, which quantify how the large scale density and shear fields affect halo abundances as a function of mass and time, are quite accurate (Fig.~\\ref{b6rnot8}). For walks with correlated steps, the predicted first crossing distributions and bias formulae can be written in units in which they are universal (independent of power spectrum and cosmology). We argued that this universality should be weakly broken if steps are correlated (Fig.~\\ref{c2b1-corr}). In this case, too, we provided analytic approximations for the unconditional first crossing distribution (Eq.~\\ref{vfvMS}) and halo bias factors (Eqs.~\\ref{b1b2-corr} and~\\ref{c2-corr}), which were quite accurate (Fig.~\\ref{cdm6d}). Our results indicate that nonlocal bias effects, as quantified by the parameter $c_2 = c_2^{\\rm L} - 8b_1^{\\rm L}/21$ of Eq.~(\\ref{nonlocalbEul}), can be substantial for the most massive halos. Our analysis is easily extended to describe the more complicated case in which the barrier $\\delta_c$ depends on $(e,p)$ of Eq.~(\\ref{ep}) rather than simply $q^2$. This is because $\\delta_c(e,p)$ is a function of the combination $e\\delta$ and $p\\delta$ \\cite{smt2001}, and, like $q^2$, these combinations are actually independent of $\\delta$ \\cite{st2002}. Alternative parametrizations of this \\cite{smlw2007} have $\\delta_c(v,w)$ where $v$ and $w$ (defined in Eq.~\\ref{vw}) are also independent of $\\delta$. This means that the analog of Eq.~(\\ref{vfvavq}), in which one averages over a distribution of first crossing distributions, remains a good approximation. In addition, the idea that one can map the 6-dimensional walk problem to a 1-dimensional moving barrier problem should continue to hold when $\\delta_c(e,p)$ or $\\delta_c(v,w)$ rather than $\\delta_c(q)$, so the analog of Eq.~(\\ref{vfvMS}) should also provide a reasonable approximation. Therefore, we believe our analysis should be applicable to other parametrizations or models of the effects of nonlocality. Indeed, our analysis should also apply to cases where one places conditions on the eigenvalues of the deformation tensor (e.g. all three have the same sign), rather than the combinations $e,p$, or $v,w$. In this respect, it provides the basis for modelling not just halos, but the abundance and spatial distribution of superclusters, filaments, sheets and voids as well. In particular, our analysis predicts that all of these constituents of the cosmic web should exhibit nonlocal bias effects; it will be interesting to see if such effects are discovered in simulations. In our analysis of nonlocal halo bias, we assumed that the effect of the large scale environment was to affect the distribution of $\\delta$ and $q$ on smaller scales, but not the shape of the collapse barrier. If the halo formation process depends on the surrounding environment, as the analysis in \\cite{vd2008} suggests, then this will provide an additional contribution to nonlocal bias. It is straightforward to include such an effect in our treatment of conditional $n$-dimensional walks, but our current expressions do not do so. We stated at the start that the nonlocal effects which are the subject of this paper, and which enter at the quadratic level ($q^2$, $e\\delta$, etc.), are qualitatively different from those which enter even at the linear level, and contribute $k$-dependent terms to the bias. In principle, both effects are present in our correlated walks calculation -- the latter arise from the dependence of the first crossing distribution on the derivative of the initial density field \\cite{ms2012, mps2012}, rather than on the anisotropic distribution of the mass. Since our main goal was to illustrate the effects of $q^2$, etc., we ignored these other terms, but a more complete model would include them. Separating these terms from one another should be possible, since the terms which depend on derivatives of the field will only contribute to the monopole of the bias. This is the subject of work in progress. A first comparison with numerical simulations (Figure~\\ref{gamma2resFig}) showed that our model is fairly accurate at describing the magnitude of nonlocal Lagrangian bias, in particular after accounting for correlations between steps. A more detailed comparison with simulations may require merging the formalism here with the additional requirement that one is interested in special positions (such as peaks), rather than random positions in the field. As noted in \\citep{ps2012b}, the formalism of \\citep{ms2012} on which our analysis is based allows the inclusion of this requirement with no additional conceptual complications; for peaks in a Gaussian field, this is particularly straightforward, because the distribution of the shear field around such peaks is known \\cite{vdWB1996}. Our analysis also suggests that a fruitful extension of the peaks model to smaller masses than that on which it usually breaks down is to look for peaks in the field defined by $\\delta - \\sqrt{q^2/q_c^2}$." }, "1207/1207.7267_arXiv.txt": { "abstract": "Over the last few years, the fruitful data provided by the Large Area Telescope aboard the Fermi Gamma-ray Space Telescope has revolutionized our understanding of high-energy processes in \\gc, particularly those involving compact objects like millisecond pulsars (MSPs). Gamma-ray emission between 100~MeV to 10~GeV has been detected from more than a dozen \\gc~in our Galaxy, most notably 47~Tucanae and Terzan~5. Based on a sample of known gamma-ray \\gc, empirical relations between the gamma-ray luminosity and properties of \\gc~such as stellar encounter rate, metallicity, as well as optical and infrared photon energy density in the cluster, have been derived. The gamma-ray spectra are generally described by a power law with a cut-off at a few GeV. Together with the detection of pulsed \\grs~from a millisecond pulsar in a globular cluster, such spectral signature gives support that \\grs~from \\gc~are collective curvature emission from magnetospheres of MSPs within the cluster. Alternative models in which the inverse-Compton emission of relativistic electrons accelerated close to MSPs or the pulsar wind nebula shocks have also been suggested. Observations at $>$10~GeV by Fermi/LAT and atmospheric Cherenkov telescopes like H.E.S.S.-II, MAGIC-II, VERITAS, and CTA will help to settle some questions unanswered by current data. We also discuss TeV observations of \\gc, as well as observational prospects of gravitational waves from double neutron stars in \\gc. ", "introduction": "Globular clusters are the oldest gravitationally-bounded stellar systems in the Galaxy. Nearly 160 \\gc~are known nowadays~\\citep[][2010 edition]{harris_catalog}. They fill a spherical halo around the Galaxy, many of them are located within the Galactic bulge. Due to the high concentration of stars within the \\gc, they host a large number of compact objects including neutron stars and white dwarfs; many of them are found in binary systems, forming for example low-mass X-ray binaries (LMXBs) and Cataclysmic variables. Since 1970s, it has been known that the formation rate per unit mass of LMXBs (Alpar et al. 1982) is orders of magnitude greater in globular clusters than in the rest of the Galaxy (Katz 1975; Clark 1975). As millisecond pulsars (MSPs) are generally believed to be descendants of LMXBs, it becomes natural that $\\sim$80\\% of the known MSPs are detected in \\gc~(cf. Manchester et al. 2005). Theoretical arguments have long asserted that the formation of LMXBs (and therefore their decedents MSPs) is made efficient through frequent stellar encounters. Using the X-ray populations in various \\gc~unveiled by the Chandra X-Ray Observatory, Pooley et al. (2003) and Gendre et al. (2003) found a positive correlation between the number of LMXBs in \\gc~and the stellar encounter rate, $\\Gamma_\\mathrm{c}$, putting the dynamical formation scenario of LMXBs in \\gc~on an observational ground. Using the cumulative luminosity distribution functions of radio millisecond pulsars (MSPs) in \\gc~as a probe of the resided MSPs in the clusters, \\citet{Hui10_metallicity} found that the number of MSPs in a globular cluster is correlated with its stellar encounter rate, as well as its metallicity. This finding provides the first observational evidence of the dynamical origin of MSPs. This is easy to understand since MSPs are descendants of LMXBs. In this paper we review main results revealed by high-energy \\gr~observations of Galactic \\gc, which were mainly contributed by data taken using the Large Area Telescope (LAT) on board the Fermi Gamma-ray Space Telescope since its launch in June 2008. Towards the end of the paper we also discuss the recent TeV observations as well as observational prospects related to gravitational wave experiments of \\gc. The readers are referred to an earlier review by~\\citet{Bednarek_review} from a more theoretical aspect. ", "conclusions": "" }, "1207/1207.7273_arXiv.txt": { "abstract": "We present the direct measurement of the Hubble constant, yielding the direct measurement of the angular-diameter distance to NGC 6264 using the H$_{2}$O megamaser technique. Our measurement is based on sensitive observations of the circumnuclear megamaser disk from four observations with the Very Long Baseline Array, the Green Bank Telescope and the Effelsberg Telescope. We also monitored the maser spectral profile for 2.3 years using the Green Bank Telescope to measure accelerations of maser lines by tracking their line-of-sight velocities as they change with time. The measured accelerations suggest that the systemic maser spots have a significantly wider radial distribution than in the archetypal megamaser in NGC 4258. We model the maser emission as arising from a circumnuclear disk with orbits dominated by the central black hole. The best fit of the data gives a Hubble constant of $H_{0} =$ 68$\\pm$9 km~s$^{-1}$~Mpc$^{-1}$, which corresponds to an angular-diameter distance of 144$\\pm$19 Mpc. In addition, the fit also gives a mass of the central black hole of (3.09$\\pm$0.42)$\\times$10$^{7}$ $M_{\\odot}$. The result demonstrates the feasibility of measuring distances to galaxies located well into the Hubble flow by using circumnuclear megamaser disks. ", "introduction": "Cosmology research has entered a new era since the discovery of the acceleration of the Universe (Perlmutter et al. 1999; Riess et al. 1998). ``Dark Energy\" (DE), which has negative pressure and accounts for 73\\% of the total energy density of the Universe, is currently the best candidate to explain the cosmic acceleration, and understanding its nature has become one of the most important problems in modern astronomy and physics. There are several promising methods to explore dark energy with high accuracy via its equation-of-state parameter $w$ (Frieman, Turner \\& Huterer 2008). As pointed out by Hu (2005), among all observables for probing DE in light of the Microwave Background Radiation (CMB), $w$ is most sensitive to variations in $H_{0}$. Hu (2005) concluded that the single most important complement to the CMB for measuring $w$ at $z$$\\sim$0.5 is a determination of the Hubble constant (at $z$$\\sim$0) to better than a few percent. This insight forms the fundamental motivation for the Megamaser Cosmology project (MCP; Reid et al. 2009 (Paper I); Braatz et al. 2010 (Paper II); Kuo et al. 2011 (Paper III); and Reid et al. 2012 (Paper IV)), which aims to determine $H_{0}$ to 3\\% accuracy. The key to a precise and direct determination of the Hubble constant is to measure accurate distances to galaxies well into the Hubble flow (i.e. $\\geq$ 30 Mpc). The galaxies must be distant to reduce the contribution of the uncertainty coming from peculiar velocities relative to a smooth Hubble flow. Among all the approaches to measure accurate distances directly, the megamaser method pioneered by the study of NGC 4258 (Herrnstein et al. 1999) has proven to yield accurate and direct distance measurements to galaxies beyond our Local Group. The methodology of the megamaser technique for distance determination can be found in Paper I, Paper II, and the modified version of the techinique that directly measure the Hubble constant can be found Paper IV. The application of the megamaser technique for measuring black hole mass can be found in Paper III. In the MCP, we aim to measure the Hubble constant with $\\sim$10\\% or better accuracy from each of about 10 galaxies, thereby determining $H_{0}$ to $\\sim$3\\% after averaging the results. The first galaxy measured by the MCP was UGC 3789 (Paper I, Paper II \\& Paper IV). A preliminary analysis using a simple model of the maser disk in Paper II gave a Hubble constant of 69$\\pm$11 Mpc (16\\% accuracy), and the overall uncertainty of the Hubble constant has recently been further reduced to $\\sim$10\\% when new data and better modeling were included (Paper IV). In this paper, we present the direct Hubble constant measurement from NGC 6264, the first galaxy beyond 100 Mpc measured with the megamaser method. In section 2, we present our VLBI and single-dish observations. In section 3, we show the analysis of the centripetal accelerations of the masers in NGC 6264. The Hubble constant and distance determinations are presented in section 4. In section 5, we discuss the challenges of applying the maser technique to distant galaxies. Finally, we summarize the results in section 6. ", "conclusions": "Our main conclusions are the following: \\begin{itemize} \\item[1.] The application of the megamaser technique to galaxies deep in the Hubble flow is intrinsically more difficult because of significantly lower flux densities and smaller disk angular size. To image the maser disks in distant galaxies more efficiently, we developed a method to use multiple maser lines at different locations in the maser disk for performing self-calibration on weaker sources. \\item[2.] The systemic masers in NGC 6264 are located at four or more radii from the black hole. 3-dimensional modeling that utilizes all information of maser spots including their positions, velocities, and accelerations is essential to measure the distance to NGC 6264 because of the complexity of its maser disks. \\item[3.] We modeled the H$_{2}$O maser disk in NGC 6264 with a Bayesian approach and obtained a Hubble constant of $H_{0} =$ 68 $\\pm$9 km~s$^{-1}$~Mpc$^{-1}$, corresponding to an angular-diameter distance of 144$\\pm$19 Mpc. The Hubble constant is consistent with the value obtained from UGC 3789 in Paper IV ($H_{0} = 69\\pm7$). This is the first time the megamaser technique has been successfully applied to a galaxy beyond 100 Mpc. \\end{itemize} We expect that including new observations will enable a 10\\% accurate estimate of the Hubble constant from NGC 6264, as well as constraining the magnitude of orbital eccentricity in the maser disk. \\appendix" }, "1207/1207.3720_arXiv.txt": { "abstract": "{ Since the first discovery of microlensing events nearly two decades ago, gravitational microlensing has accumulated tens of TBytes of data and developed into a powerful astrophysical technique with diverse applications. The review starts with a theoretical overview of the field and then proceeds to discuss the scientific highlights. (1) Microlensing observations toward the Magellanic Clouds rule out the Milky Way halo being dominated by MAssive Compact Halo Objects (MACHOs). This confirms most dark matter is non-baryonic, consistent with other observations. (2) Microlensing has discovered about 20 extrasolar planets (16 published), including the first two Jupiter-Saturn like systems and the only ``cold Neptunes\" yet detected. They probe a different part of the parameter space and will likely provide the most stringent test of core accretion theory of planet formation. (3) Microlensing provides a unique way to measure the mass of isolated stars, including brown dwarfs to normal stars. Half a dozen or so stellar mass black hole candidates have also been proposed. (4) High-resolution, target-of-opportunity spectra of highly-magnified dwarf stars provide intriguing ``age'' determinations which may either hint at enhanced helium enrichment or unusual bulge formation theories. (5) Microlensing also measured limb-darkening profiles for close to ten giant stars, which challenges stellar atmosphere models. (6) Data from surveys also provide strong constraints on the geometry and kinematics of the Milky Way bar (through proper motions); the latter indicates predictions from current models appear to be too anisotropic compared with observations. The future of microlensing is bright given the new capabilities of current surveys and forthcoming new telescope networks from the ground and from space. Some open issues in the field are identified and briefly discussed. } ", "introduction": "\\label{sec:intro} Gravitational microlensing in the local group refers to the temporal brightening of a background star due to intervening objects. \\cite{Ein36} first studied (micro)lensing by a single star, and concluded that ``there is no great chance of observing this phenomenon.'' Although there were some works in intervening years by \\cite{Ref64} and \\cite{Lie64}, the field was revitalized by \\cite{Pac86} who proposed it as a method to detect MAssive Compact Halo Objects (MACHOs) in the Galactic halo. From observations of microwave background radiation and nucleosynthesis (see, e.g. \\citealt{Kom11, Ste07}), it is clear that most of the dark matter must be non-baryonic, and so the original goal of microlensing is now obsolete. Nevertheless, microlensing has developed into a powerful technique with diverse applications in astrophysics, including constraints on MACHOs, the study of the structure of the Milky Way, stellar atmospheres and the detection of extrasolar planets and stellar-mass black hole candidates. Since the first discoveries of microlensing events in 1993 \\citep{Alc93, Uda93}, the field has made enormous progress in the last two decades. A number of reviews have been written on this topic (e.g. \\citealt{Pac96, Mao01, Eva03, Wam06}), with the most recent highlights given in \\cite{Mao08a}, \\cite{Gou09a} and \\cite{Gau10}. The readers will also greatly benefit from two recent, comprehensive conference proceedings: the Manchester Microlensing Conference\\footnote{http://pos.sissa.it/cgi-bin/reader/conf.cgi?confid=54} and the 2011 Sagan Exoplanet Summer Workshop: Exploring Exoplanets with Microlensing\\footnote{http://nexsci.caltech.edu/workshop/2011/}. The workshop materials contain not only recent scientific highlights but also hands-on exercises for data reduction and modelling. The structure of this review is as follows. \\sect\\ref{sec:basics} introduces the basics of gravitational microlensing, which reproduces \\cite{Mao08b} in a slightly modified form; \\sect\\ref{sec:apps} builds on the introduction and discusses the applications of gravitational microlensing. We finish this review with an outlook for the field in \\sect\\ref{sec:future}. Due to the rapid expansion of the field, it is unavoidable that the reference list is incomplete (and somewhat biased). ", "conclusions": "" }, "1207/1207.6105_arXiv.txt": { "abstract": "We present a robust method to constrain average galaxy star formation rates, star formation histories, and the intracluster light as a function of halo mass. Our results are consistent with observed galaxy stellar mass functions, specific star formation rates, and cosmic star formation rates from $z=0$ to $z=8$. We consider the effects of a wide range of uncertainties on our results, including those affecting stellar masses, star formation rates, and the halo mass function at the heart of our analysis. As they are relevant to our method, we also present new calibrations of the dark matter halo mass function, halo mass accretion histories, and halo-subhalo merger rates out to $z=8$. We also provide new compilations of cosmic and specific star formation rates; more recent measurements are now consistent with the buildup of the cosmic stellar mass density at all redshifts. Implications of our work include: halos near $10^{12}\\Msun$ are the most efficient at forming stars at all redshifts, the baryon conversion efficiency of massive halos drops markedly after $z\\sim 2.5$ (consistent with theories of cold-mode accretion), the ICL for massive galaxies is expected to be significant out to at least $z\\sim 1-1.5$, and dwarf galaxies at low redshifts have higher stellar mass to halo mass ratios than previous expectations and form later than in most theoretical models. Finally, we provide new fitting formulae for star formation histories that are more accurate than the standard declining tau model. Our approach places a wide variety of observations relating to the star formation history of galaxies into a self-consistent framework based on the modern understanding of structure formation in $\\Lambda$CDM. Constraints on the stellar mass---halo mass relationship and star formation rates are available for download online. ", "introduction": "Constraining the buildup of stellar mass in galaxies provides fundamental constraints on galaxy formation models. An ever growing number of galaxy observations have been made, covering a time period from 500 Myr after the Big Bang \\citep[e.g.][]{Zheng12} to the present day. A model of galaxy formation that can match observations over this entire stretch of time would represent a significant aid to our understanding. So far, matching the evolution of the stellar masses and star formation rates of galaxies over this entire epoch with has proved difficult \\citep[see, e.g.,][]{Lu12,Borgani11, Weinmann12}. Despite challenges of reproducing these observations with detailed models of galaxy formation, significant progress has been made in recent years with empirical models that connect the evolution of galaxy properties to the evolution of dark matter halos. Over the past decade, a range of studies have associated galaxies with dark matter halos at a given epoch, using a variety of techniques, including Halo Occupation Distribution modeling \\citep[e.g.,][]{Berlind02, Bullock02} the Conditional Luminosity Function modeling \\citep[e.g.,][]{Yang03}, and variants of the abundance matching technique \\citep[e.g.,][]{Colin99, kravtsov_klypin:99, Neyrinck04, Kravtsov04, Vale04, Vale06, Tasitsiomi04, conroy:06, Shankar06, Berrier06, Marin08, Guo-09, moster-09, Behroozi10}. The simplest abundance matching models assign the most massive observed galaxies in rank order to the largest halos in an equal simulation volume. With only slight modifications (e.g., using the peak mass for satellite halos), this technique accurately reproduces the redshift-- and scale--dependent clustering of galaxies \\citep[e.g.,][]{conroy:06,Reddick12}. Because the cosmological model provides a prediction for the buildup of dark matter halos, one can combine knowledge of the galaxy-halo connection at different epochs with knowledge of the mass accretion and merger histories of halos to constrain the buildup of stellar mass in galaxies over cosmic time. This was first done in a comprehensive way by \\cite{cw-08}, who provided an empirical constraint on the star formation histories and stellar mass growth in galaxies from $z=2$ to the present. This approach has been explored in several studies using variants of the techniques \\citep{zheng-07,White07,Firmani10,Leitner11,Bethermin12,Wang12, Moster12}, which together provide important constraints on the basic picture for the buildup of stars in galaxies and their connection to dark matter halos. These studies represented important advances, albeit with shortcomings. For example, most of these studies only modeled galaxy evolution from $z\\sim 2$ to the present due to perceived conflicts between the integrated cosmic star formation and the cosmic stellar mass density at $z>1$ \\citep{Wilkins08,Hopkins06b}. In addition, because passive stellar mass loss (e.g., supernovae of massive stars) depends on galaxies' star formation histories, these studies had to make assumptions about the ages of stars already formed at $z>2$. Finally, these studies paid limited attention to the vast array of observational uncertainties as well as modeling uncertainties affecting their derived constraints; indeed, many results are presented without error bars \\citep{cw-08,Firmani10,Bethermin12,Wang12}. Here we present a new technique and a new compilation of observations, aimed squarely at resolving these issues. On the observational side, new observations of galaxies at very high redshifts \\citep[e.g.][]{Bouwens11,BORG12,Bouwens11b,McLure11} have made it possible to place constraints on galaxy formation all the way to $z\\sim8$. At lower redshifts, we present the first constraints from stellar mass functions based on the PRIMUS survey, as well as from a new analysis at $z=0$ based on SDSS and GALEX data \\citep{Moustakas12}. These are obtained from a consistent methodology from $z=0$ to $z=1$, and thus alleviate many concerns about matching inconsistently-identified galaxies at different epochs. Using these combined data sets, we find that observations of the evolution of galaxy star formation rates and stellar masses can now be reconciled (see also \\citealt{Bernardi10} and \\citealt{Moster12}). Our new method constrains the galaxy--halo relation using observed galaxy stellar mass functions \\textit{as well as} specific star formation rates and cosmic star formation rates. As with previous studies, we match observed galaxies to halos, but the additional information on star formation rates allows us to break degeneracies and directly constrain the buildup of the intracluster light, as well as the amount of stars that can transfer from satellites to the central galaxy during mergers. We account for a number of statistical and systematic effects, including uncertainties from stellar population synthesis models, dust models, star formation history models, the faint-end slope of the stellar mass function, observational completeness, scatter between stellar mass and halo mass. Given a parametrization for the intrinsic stellar mass---halo mass relationship as well as for these observational effects, we use a Markov Chain Monte Carlo (MCMC) method to find the allowed posterior distribution. This is important not only for addressing concerns that our results may be biased by limited observational constraints at high redshifts, but it also gives us the power to determine which observations would best improve the resulting constraints on the relationship between stellar mass, star formation, and halo mass. The direct results of this analysis are constraints on the distribution of galaxy stellar masses as a function of halo mass and redshift. From these, many other constraints relevant to galaxy formation are derived, including the average star formation rate as a function of halo mass and redshift, the average star formation history in galaxies at a given epoch as a function of galaxy stellar mass or host halo mass, the instantaneous baryon conversion efficiency of galaxies as a function of mass and redshift, the buildup of the intracluster light, and the evolution of the stellar mass to halo mass ratio for progenitors of today's galaxies, all including full uncertainties and covering a redshift range from $z=0$ to $z=8$. We provide a broad overview of our methodology in \\S \\ref{s:methodology}, including our parametrization of the stellar mass -- halo mass relation and the relevant uncertainties, with additional details in Appendices \\ref{a:av_afh}-\\ref{a:mergers}. We discuss the observational data sets relevant to our method in \\S \\ref{s:data}, with special attention to new measurements of cosmic star formation rates. We discuss the simulations that we use in \\S \\ref{s:sim}, along with recalibrations of the halo mass function (Appendix \\ref{a:tinker}), halo mass accretion rates (Appendix \\ref{a:mah}), and subhalo merger/disruption rates (Appendix \\ref{a:disruption}). We present our main results in \\S \\ref{s:results}, with discussion in \\S \\ref{s:discussion} and a summary of our conclusions in \\S \\ref{s:conclusions}. Throughout this work, we assume a \\cite{chabrier-2003-115} initial mass function, the stellar population synthesis model of \\cite{bc-03}, and the dust model employed in \\cite{blanton-roweis-07}. We convert data sets from other papers to these models as necessary. We additionally assume a flat, $\\Lambda$CDM cosmology with parameters $\\Omega_M = 0.27$, $\\Omega_\\Lambda = 0.73$, $h=0.7$, $n_s = 0.95$, and $\\sigma_8 = 0.82$. ", "conclusions": "\\end{figure*} \\begin{table*} \\vspace{-4ex} \\begin{center} \\caption{Table of Parameters} \\label{t:parameters} \\begin{tabular}{rccc} \\hline \\hline Symbol & Description & Equation & Section\\\\ \\hline $M_1$ & Characteristic halo mass & \\ref{e:cosmic_sfh} & \\ref{s:smhm}\\\\ $\\epsilon$ & Characteristic stellar mass to halo mass ratio & \\ref{e:cosmic_sfh} & \\ref{s:smhm}\\\\ $\\alpha$ & Faint-end slope of SMHM relation & \\ref{e:cosmic_sfh} & \\ref{s:smhm}\\\\ $\\delta$ & Strength of subpower law at massive end of SMHM relation & \\ref{e:cosmic_sfh}& \\ref{s:smhm}\\\\ $\\gamma$ & Index of subpower law at massive end of SMHM relation & \\ref{e:cosmic_sfh}& \\ref{s:smhm}\\\\ $\\nu$ & Exponential cutoff of evolution of $M_\\ast(M_h)$ with scale factor & \\ref{e:redshift_scaling} & \\ref{s:smhm}\\\\ $\\xi$ & Scatter in dex of true stellar mass at fixed halo mass & \\ref{e:scatter} & \\ref{s:smhm}\\\\ $M_{h,ICL}$ & Characteristic halo mass at which half of stellar mass growth is due to mergers & \\ref{e:f_sfr} & \\ref{s:icl}\\\\ \\hline $\\mu$ & Systematic offset in stellar masses and SFRs & \\ref{e:syst} & \\ref{s:obs_syst}\\\\ $\\kappa$ & Systematic offset in stellar masses for active galaxies & \\ref{e:syst2} & \\ref{s:obs_syst}\\\\ $\\sigma$ & Scatter in measured stellar mass at fixed true stellar mass & \\ref{e:psf} & \\ref{s:obs_syst}\\\\ $c$ & Galaxy detection completeness & \\ref{e:completeness} & \\ref{s:obs_syst}\\\\ $b$ & Fraction of incompleteness due to burstiness (as opposed to dustiness) & \\ref{e:burstiness_correction} & \\ref{s:obs_syst}\\\\ $\\rho$ & Correlation of SFR to stellar mass at fixed halo mass & \\ref{e:rho_def} & \\ref{s:obs_ssfr_csfr}\\\\ \\hline \\end{tabular} \\end{center} \\end{table*} \\begin{table*} \\begin{center} \\caption{Table of Priors} \\label{t:priors} \\begin{tabular}{rccc} \\hline \\hline Symbol & Description & Equation & Prior\\\\ \\hline $\\xi_0$ & Value of $\\xi$ at $z=0$, in dex & \\ref{e:scatter} & $G(0.16,0.04)$\\\\ $\\xi_a$ & Redshift scaling of $\\xi$, in dex & \\ref{e:scatter} & $G(0,0.16)$\\\\ $\\mu_0$ & Value of $\\mu$ at $z=0$, in dex & \\ref{e:syst} & $G(0,0.14)$\\\\ $\\mu_a$ & Redshift scaling of $\\mu$, in dex & \\ref{e:syst_mu} & $G(0,0.22)$\\\\ $\\kappa_0$ & Value of $\\kappa$ at $z=0$, in dex & \\ref{e:syst2} & $G(0,0.24)$\\\\ $\\kappa_a$ & Redshift scaling of $\\kappa$, in dex & \\ref{e:syst_kappa} & $G(0,0.3)$\\\\ $\\sigma_z$ & Redshift scaling of $\\sigma$, in dex & \\ref{e:psf} & $G(0.05,0.015)$\\\\ $A$ & Amplitude of galaxy detection completeness & \\ref{e:completeness_bare} & $U(0,1)$\\\\ $z_c$ & Onset redshift of galaxy detection completeness & \\ref{e:completeness_bare} & $z_c>0.8$\\\\ $b$ & Fraction of incompleteness due to burstiness (as opposed to dustiness) & - & $U(0,1)$\\\\ $\\rho_{0.5}$ & Correlation of SFR to stellar mass at fixed halo mass at $a=0.5$ & \\ref{e:rho_def} & $U(0.23,1.0)$\\\\ \\hline \\end{tabular} \\end{center} \\tablecomments{$G(x,y)$ denotes a Gaussian distribution with center $x$ and width $y$. $U(x_1,x_2)$ denotes a uniform distribution from $x_1$ to $x_2$. The remaining parameters have no explicit priors; $M_1$ and $\\epsilon$ are explored in logarithmic space, whereas $\\alpha,\\delta$, and $\\gamma$ are explored in linear space.} \\end{table*} \\subsection{Methodology Summary} \\label{s:methodology_summary} Although the equations involved are somewhat complicated, the logical steps involved in our method are straightforward and are shown visually in Fig.\\ \\ref{f:methodology_summary}. These steps include: \\begin{enumerate} \\item Choose parameters for the stellar mass -- halo mass relation as well as observational and methodology uncertainties (see Table \\ref{t:parameters} for full list; see Table \\ref{t:priors} for adopted priors) using a Markov Chain Monte Carlo method. \\item Use the chosen stellar mass -- halo mass relation to populate halos with galaxies, and use merger rates and mass accretion histories calculated from dark matter simulations to infer galaxy growth rates. \\item From the previous step, calculate the average star formation rates as a function of halo mass and redshift, and use those to calculate the average specific star formation rate and the cosmic star formation rate. In addition, using the abundances of halos, calculate the stellar mass function. \\item Apply corrections for observational errors and biases to the derived stellar mass function and star formation rates. \\item Compare with observational measures of the stellar mass function and star formation rates (i.e., sum $\\chi^2$ errors for each data point) to calculate the likelihood that the chosen stellar mass -- halo mass relation matches observed results (i.e., $\\exp(-0.5\\chi^2)$). \\item Return to step \\#1 until the MCMC algorithm has converged. \\end{enumerate} To ensure convergence on a space with such a large number of parameters, we use the Adaptive Metropolis exploration method \\citep{Haario01}. Although the algorithm of \\cite{dunkley-2005} indicates that we converge after only $5\\times10^{5}$ points, we continue to run for $4\\times10^6$ total points to minimize the chance of unexplored regions and to ensure that the adaptive updates of the step covariance matrix converge and do not bias our final results." }, "1207/1207.6275_arXiv.txt": { "abstract": "We report an analysis of the interstellar $\\gamma$-ray emission from the Chamaeleon, R~Coronae Australis (R CrA), and Cepheus and Polaris flare regions with the {\\it Fermi} Large Area Telescope. They are among the nearest molecular cloud complexes, within $\\sim$ 300 pc from the solar system. The $\\gamma$-ray emission produced by interactions of cosmic-rays (CRs) and interstellar gas in those molecular clouds is useful to study the CR densities and distributions of molecular gas close to the solar system. The obtained $\\gamma$-ray emissivities above 250 MeV are (5.9 $\\pm$ 0.1$_{\\rm stat}$ $^{+0.9}_{-1.0}$$_{\\rm sys}$) $\\times$ 10$^{-27}$ photons s$^{-1}$ sr$^{-1}$ H-atom$^{-1}$, (10.2 $\\pm$ 0.4$_{\\rm stat}$ $^{+1.2}_{-1.7}$$_{\\rm sys}$) $\\times$ 10$^{-27}$ photons s$^{-1}$ sr$^{-1}$ H-atom$^{-1}$, and (9.1 $\\pm$ 0.3$_{\\rm stat}$ $^{+1.5}_{-0.6}$$_{\\rm sys}$) $\\times$ 10$^{-27}$ photons s$^{-1}$ sr$^{-1}$ H-atom$^{-1}$ for the Chamaeleon, R~CrA, and Cepheus and Polaris flare regions, respectively. Whereas the energy dependences of the emissivities agree well with that predicted from direct CR observations at the Earth, the measured emissivities from 250 MeV to 10 GeV indicate a variation of the CR density by $\\sim$ 20 \\% in the neighborhood of the solar system, even if we consider systematic uncertainties. The molecular mass calibrating ratio, $X_{\\rm CO} = N({\\rm H_{2}})/W_{\\rm CO}$, is found to be (0.96 $\\pm$ 0.06$_{\\rm stat}$ $^{+0.15}_{-0.12}$$_{\\rm sys}$) $\\times$10$^{20}$ H$_2$-molecule cm$^{-2}$ (K km s$^{-1}$)$^{-1}$, (0.99 $\\pm$ 0.08$_{\\rm stat}$ $^{+0.18}_{-0.10}$$_{\\rm sys}$) $\\times$10$^{20}$ H$_2$-molecule cm$^{-2}$ (K km s$^{-1}$)$^{-1}$, and (0.63 $\\pm$ 0.02$_{\\rm stat}$ $^{+0.09}_{-0.07}$$_{\\rm sys}$) $\\times$10$^{20}$ H$_2$-molecule cm$^{-2}$ (K km s$^{-1}$)$^{-1}$ for the Chamaeleon, R~CrA, and Cepheus and Polaris flare regions, respectively, suggesting a variation of $X_{\\rm CO}$ in the vicinity of the solar system. From the obtained values of $X_{\\rm CO}$, the masses of molecular gas traced by $W_{\\rm CO}$ in the Chamaeleon, R~CrA, and Cepheus and Polaris flare regions are estimated to be $\\sim$ 5$\\times$10$^{3}$ $M_{\\odot}$, $\\sim$ 10$^{3}$ $M_{\\odot}$, and $\\sim$ 3.3$\\times$$10^{4}$ $M_{\\odot}$, respectively. A comparable amount of gas not traced well by standard \\HI\\ and CO surveys is found in the regions investigated. ", "introduction": "Observations of high-energy $\\gamma$-ray emission ({\\it E} $\\gtrsim$ 30 MeV) from molecular clouds can be used to study the cosmic-ray (CR) production, the CR density, and the distribution of the interstellar medium (ISM) in such systems. $\\gamma$-rays are produced in the ISM by interactions of high-energy CR protons and electrons with the interstellar gas, via nucleon-nucleon collisions, electron Bremsstrahlung, and inverse Compton (IC) scattering. Since the $\\gamma$-ray production cross section is almost independent of the chemical or thermodynamic state of the ISM, and the interstellar gas is essentially transparent to those high-energy photons, observations in $\\gamma$-rays have been recognized as a powerful probe of the distribution of interstellar matter. If the gas column densities are estimated with good accuracy by observations in other wavebands such as radio, infrared, and optical, the CR spectrum and density distributions can be examined as well. Molecular clouds that are within 1 kpc from the solar system (namely nearby molecular clouds) and have masses greater than a few 10$^3$ {\\it M}$_{\\odot}$ are well suited for an analysis of their $\\gamma$-ray emission to investigate the distribution of CR densities and interstellar gas since they are observed at high latitudes and therefore largely free from confusion with the strong emission from the Galactic plane. Study of such nearby molecular clouds in $\\gamma$-rays can be dated back to the COS-B era (e.g., Bloemen et al. 1984) and was advanced by the EGRET on board {\\it Compton Gamma-Ray Observatory} (e.g., Hunter et al. 1994). Although some important information has been obtained on properties of CRs and the ISM by these early observations, detailed studies have only been performed on giant molecular clouds with masses greater than $\\sim$ 10$^5$ {\\it M}$_{\\odot}$ such as the Orion complex (e.g., Digel et al. 1999). The data above 1 GeV, which are crucial to study CR nuclei spectra, suffered from the limited photon statistics, angular resolution, and energy coverage of these early missions. The advent of the {\\it Fermi} Gamma-ray Space Telescope launched in 2008 has improved the situation significantly. The sensitivity of the LAT (Large Area Telescope) on board {\\it Fermi} % is more than an order of magnitude better than that of the EGRET, and enables resolving more point sources and studying the diffuse $\\gamma$-ray emission with unprecedented sensitivity. In addition, newer surveys of the ISM (e.g., Dame et al. 2001, Kalberla et al. 2005, and Grenier et al. 2005) allow us to investigate the CR spectral and density distributions with better accuracy. Here we report a {\\it Fermi} LAT study of diffuse $\\gamma$-rays from the Chamaeleon, R Coronae Australis (R CrA), and Cepheus and Polaris flare molecular clouds. They are among the nearest ($\\lesssim$ 300 pc from the solar system) molecular clouds exhibiting star formation activity. Although EGRET observed $\\gamma$-ray emission associated with the molecular gas in the Chamaeleon region (Grenier et al. 2005), no detailed study of CR and matter distributions for the Chamaeleon and R CrA regions has been performed yet since they have rather small masses ($\\lesssim$ 10$^{4}$ {\\it M}$_{\\odot}$, about 1/10 of that of the Orion molecular cloud) and consequently small $\\gamma$-ray fluxes. We also analyzed in detail the region of the Cepheus and Polaris flares which was included in the {\\it Fermi} LAT study of the second Galactic quadrant (Abdo et al. 2010b). It is located in the direction almost opposite to the Chamaeleon region in the Gould Belt (see, e.g., Perrot and Grenier 2003), therefore we can investigate the distribution of the CR density over several hundred pc but still inside the coherent environment of the Gould Belt. This paper is organized as follows. We first describe the observations as well as the sky model preparation and the data analysis in Section \\ref{sec:Data_analysis}, and show the obtained results in Section \\ref{sec:Results}. We then discuss the CR and matter distributions in Section \\ref{sec:Discussion} and give conclusions in Section \\ref{sec:Summary_and_conclusions}. ", "conclusions": "\\label{sec:Summary_and_conclusions} We have studied the $\\gamma$-ray emission from the Chamaeleon, R CrA, and Cepheus and Polaris flare molecular clouds close to the solar system ($\\lesssim$ 300 pc) using the first 21 months of {\\it Fermi} LAT data. Thanks to the excellent performance of the LAT, we have obtained unprecedentedly high-quality emissivity spectra of the atomic and molecular gas in these regions in the 250 MeV -- 10 GeV range. The $\\gamma$-ray emissivity spectral shapes in three regions agree well with the model for the LIS (a model based on local CR measurement), thus indicating a similar spectral distribution of CRs in these regions. The emissivities, however, indicate a variation of the CR density of $\\sim$ 20 \\% within $\\sim$ 300 pc around the solar system, even if we consider the systematic uncertainties. We consider possible origins of the variation are non-uniform supernova rate and anisotropy of CRs depending on the propagation conditions. The molecular mass calibration ratio $X_{\\rm CO}$ for the Chamaeleon cloud and the R CrA cloud are comparable, whereas that of the Cepheus and Polaris flare region is $\\sim$ 2/3 of the others, suggesting a variation of $X_{\\rm CO}$ in the vicinity of the solar system. From the obtained values of $X_{\\rm CO}$, the masses of gas traced by $W_{\\rm CO}$ in the Chamaeleon, R CrA, and Cepheus and Polaris flare regions are estimated to be $\\sim$ 5 $\\times$ $10^{3}$ $M_{\\odot}$, $\\sim10^{3}$ $M_{\\odot}$, and $\\sim$ 3.3 $\\times$ $10^{4}$ $M_{\\odot}$ respectively. Similar amounts of gas are inferred to be in the phase not well traced by the {\\HI} or CO lines. Accumulation of more $\\gamma$-ray data, particularly at high energies, and progress in ISM studies, will reveal the CR and matter distribution in greater detail. The {\\it Fermi} LAT Collaboration acknowledges generous ongoing support from a number of agencies and institutes that have supported both the development and the operation of the LAT as well as scientific data analysis. These include the National Aeronautics and Space Administration and the Department of Energy in the United States, the Commissariat $\\grave{\\rm a}$ l'Energie Atomique and the Centre National de la Recherche Scientifique/Institut National de Physique Nucl$\\acute{\\rm e}$aire et de Physique des Particules in France, the Agenzia Spaziale Italiana and the Istituto Nazionale di Fisica Nucleare in Italy, the Ministry of Education, Culture, Sports, Science and Technology (MEXT), High Energy Accelerator Research Organization (KEK), and Japan Aerospace Exploration Agency (JAXA) in Japan, and the K. A. Wallenberg Foundation, the Swedish Research Council, and the Swedish National Space Board in Sweden. Additional support for science analysis during the operations phase is gratefully acknowledged from the Istituto Nazionale di Astrofisica in Italy and the Centre National d'$\\acute{\\rm E}$tudes Spatiales in France. We thank the GALPROP team for providing a development version of GALPROP model adjusted to the measurement data by the LAT. GALPROP development is supported by NASA Grant NNX09AC15G and by the Max Planck Society." }, "1207/1207.1987_arXiv.txt": { "abstract": "{ HR4796 is a young, early A-type star harbouring a well structured debris disk, shaped as a ring with sharp inner edges. The inner edge might be shaped by a yet unseen planet inside the ring; the outer one is not well understood. The star forms together with the M-type star HR4796B, a binary system, with a projected separation of $\\simeq$ 560 AU.} {Our aim is to explore the surroundings of HR4796A and B, both in terms of extended or point-like structures.} {Adaptive optics images at L'-band were obtained with NaCo in Angular Differential Mode and with Sparse Aperture Masking (SAM). We analyse the data as well as the artefacts that can be produced by ADI reduction on an extended structure with a shape similar to that of HR4796A dust ring. We determine constraints on the presence of companions using SAM and ADI on HR4796A, and ADI on HR4796B. We also performed dynamical simulations of a disk of planetesimals and dust produced by collisions, perturbed by a planet located close to the disk outer edge.} {The disk ring around HR4796A is well resolved. We highlight the potential effects of ADI reduction of the observed disk shape and surface brightness distribution, and side-to-side asymmetries. We produce 2D maps of planet detection limits. No planet is detected around the star, with masses as low as 3.5 \\mjup at 0.5\" (58 AU) and less than 3 \\mjup in the 0.8-1\" range along the semi-major axis. We exclude massive brown dwarfs at separations as close as 60 mas (4.5 AU) from the star thanks to SAM data. The detection limits obtained allow us to exclude a possible close companion to HR4796A as the origin of the offset of the ring center with respect to the star; they also allow to put interesting constraints on the (mass, separation) of any planet possibly responsible for the inner disk steep edge. Using detailed dynamical simulations, we show that a giant planet orbiting outside the ring could sharpen the disk outer edge and reproduce the STIS images published by Schneider et al. (2009). Finally, no planets are detected around HR4796B with limits well below 1 \\mjup at 0.5\" (35 AU).} {} ", "introduction": "Understanding planetary systems formation and evolution has become one of the biggest challenges of astronomy, since the imaging of a debris disk around $\\beta$ Pictoris in the 80's (\\cite{smith84}) and the discovery of the first exoplanet around the solar-like star 51 Pegasi during the 90's (\\cite{mayor95}). Today, about 25 debris disks have been imaged at optical, infrared, or submillimetric wavelengths (http://astro.berkeley.edu/kalas/disksite). Debris disks trace stages of system evolution where solid bodies with sizes significantly larger than the primordial dust size (larger than meters or km sized) are present to account for through collisions, the presence of short lived dust. They are thought to be privileged places to search for planets. This is particularly true for those showing peculiar structures (e.g. rings with sharp edges) or asymmetries, spirals even though other physical effects not involving planets could also lead to the formation of similar structures. \\cite{takeuchi01} for instance showed that relatively small amounts (typ. 1-a few Earth masses) of gas can shape the dust disk through gas-dust interactions into rings (see below). It is remarkable however that all the stars around which relatively close (separations less than 120 AU) planets have been imaged are surrounded by debris disks: a $\\leq$ 3-\\mjup planetary companion was detected in the outskirts of Fomalhaut's debris disk (119 AU from the star; \\cite{kalas08}), four planetary companions of 7-10 \\mjup were imaged at 15, 24, 38 and 68 AU (projected separations) from HR 8799 (\\cite{marois08}, \\cite{marois10}). Using VLT/NaCo L'-band saturated images, we detected a 9\\pm 3 \\mjup planet in the disk of $\\beta$ Pictoris ($\\simeq$ 12 Myr) with an orbital radius of 8-12 AU from the star (\\cite{lagrange10}; \\cite{chauvin12}). More recent studies at Ks show that $\\beta$ Pic b is located in the inclined part of the disk (\\cite{lagrange11}), conforting the link between disk morphology and the presence of a planet (\\cite{mouillet97}; \\cite{jca01}). $\\beta$ Pic b also confirms that giant planets form in a timescale of 10 Myr or less. Interestingly, $\\beta$ Pic b and maybe also HR8799e could have formed in situ via core accretion, in contrast with the other, more remote, young companions detected with high-contrast imaging. If formed in situ, the latter probably formed through gravitational instabilities within a disk, or through the fragmentation and collapse of a molecular cloud. There are many exciting questions regarding disks and planets: could different planet formation processes be at work within a given disk? Disks and planets are known to exist in binary systems; (how) do massive companions impact on the dynamical evolution of inner planets and disks? In a recent study, \\cite{rodriguez11} showed that for a binary system to have a disk, it must either be a very wide binary system with disk particles orbiting a single star or a small separation binary with a circumbinary disk. Such results can help the search for planets if one relates debris disks with planet formation. However, another question, already mentioned, is to which extent and how can debris disks indicate the presence of already formed planets? A particularly interesting system in the present context is HR4796A, consisting of an early-type (A0), young, close-by star (see Table~\\ref{star}) surrounded by dust, identified in the early 90's (\\cite{jura91}) and resolved by \\cite{koerner98} and \\cite{jaya98} at mid-IR from the ground, and at near-IR with NICMOS on the HST (\\cite{schneider99}) as well as from the ground, coupling coronagraphy with adaptive optics (\\cite{jca99}). The resolved dust shapes as a narrow ring, with steep inner and outer edges. The steepness of the inner edge of the dust ring has been tentatively attributed to an unseen planet (\\cite{wyatt99}); however, none has been detected so far. The disk images + SED modeling required at least two populations of grains, one, narrow (a few tens AU) cold ring, located at $\\simeq$ 70 AU, and a second one, hotter and much closer to the star (\\cite{jca99}), although the existence of an exozodiacal dust component is debated (\\cite{li03}). \\cite{wahhaj05}, argued that in addition to the dust responsible for the ring-like structure observed at optical/near IR wavelengths, a wider, low-density component should be present at similar separations to account for the thermal IR images. Recently, higher quality (SN/angular resolution) data were obtained with STIS (\\cite{schneider09}) and from the ground (\\cite{thalmann11}), the latter using performant AO system on a 8 m class telescope, as well as Angular Differential Imaging (ADI, see below). With the revised distance of HR4796 with respect to earlier results, the ring radius is at about 79 AU, and has a width of 13 AU (STIS 0.2-1 $\\mu$m data). Furthermore, these authors show a 1.2-1.6 AU physical shift of the projected center of the disk wrt the star position along the major axis and \\cite{thalmann11} moreover measures a 0.5 AU shift along the disk minor axis. Finally, and very interestingly in the context of planetary system formation, HR4796 has also a close-by (7.7 arcsec, i.e. about 560 AU projected separation), M-type companion (\\cite{jura93}), plus a tertiary one, located much further away (13500 AU in projection; \\cite{kastner08}). The closest companion may have played a role in the outer truncature of the disk, even though, according to \\cite{thebault10}, it alone cannot account for the sharp outer edge of HR4796 as observed by STIS. \\begin{table}[t!] \\caption{HR 4796 stellar properties. (1)\\cite{vanleeuwen07}, (2) \\cite{stauffer95}. } \\label{star} \\begin{center} \\begin{tabular}{l l} \\hline Spectral Type & A0V \\\\ $V$ & 5.78 \\\\ $B-V$ & 0.009\\\\ $\\pi$ & 3.74 \\pm 0.33 mas (1)\\\\ Age & 8\\pm 2 Myr (2) \\\\ \\hline \\end{tabular} \\end{center} \\label{star} \\end{table} ADI (\\cite{marois06}) is a technics that has proved to be very efficient in reaching very high contrast from the ground on point-like objects. It has also been used to image disks around HD61005 (the Moth disk, \\cite{buenzli10}), HD32297 (\\cite{bocca11}), and \\bp (\\cite{lagrange11}), but it should be used with care when observing extended structures, as the morphology of these structures may be strongly impacted by this method (Milli et al, 2012, in prep). We present here new high contrast images of HR4796 obtained with NaCo on the VLT at L' band, both in ADI and Sparse Aperture Mode (SAM, \\cite{tuthill10}) aiming at exploring the disk around HR4796A, and at searching for possible companions around HR4796A as well as around HR4796B. SAM and ADI are complementary as the first give access to regions in the 40-400 mas range from the star, and ADI further than typically 300 mas. Pedagogical examples of SAM performances and results are reported in \\cite{lacour11}. ", "conclusions": "\\subsection{The inner disk sharp edge} One of the most remarkable features of the HR4796 disk is certainly the offset of the disk center with respect to the star (also observed in the case of HD141569, and Fomalhaut). Two explanations are 1) the presence of a close, fainter companion (in such a case, the disk would be a circumbinary disk and orbit around the binary center of mass), and 2) the presence of a companion close to the disk inner edge on an eccentric orbit that induces a forced eccentricity to the disk ring by secular gravitational interaction, an explanation which was proposed to explain the eccentricity of the Fomalhaut disk (\\cite{quillen06}, \\cite{kalas05}, \\cite{kalas08}, \\cite{chiang09}). We first investigate whether this offset could be due to the presence of a close companion. In such a case, the ellipse center would mark the center of mass of the binary system. Using the center of mass definition, it appears that the mass of a body necessary to shift the center of mass at the observed position of the ellipse center would be much larger than the detection limit obtained with SAM between 40 and 400 mas or between 400 mas and 1 arcsec (ring position) with the ADI data. A companion located between 23 and 40 mas would have a mass larger or comparable to that of HR4696A; such a scenario must be excluded as under such conditions, the photometric center of the system would also be shifted. Then, the most plausible explanation to the offset is a light eccentric planet close to the inner edge of the disk. We now try to use the detection limits found in this paper to constrain the properties of an inner planet that could be responsible for the steep inner edge observed with HST/STIS data, which, conversely to ADI data, is not impacted by ADI reduction effects. \\cite{wisdom80} showed that in the case of a planet and particles on circular orbits, we have the relation $\\delta$a/a = 1.3.(Mp/Ms)$^{2/7}$ where Mp and Ms are the planet and star masses, a is the orbital radius of the planet and $\\delta$ a the distance between the planet and the disk inner edge. Hence if a planet sculpts the inner edge, its mass and distance from the inner edge must satisfy this relation. Assuming an inner edge located at 77 AU, we can derive the mass of the planet necessary to produce this sharp edge, as a function of its distance to the edge, and test whether such a planet would have been detected or not. This is done in Figure\\ref{map_planete} where we show the region, inside the yellow ellipse, that, given the present detection limits, have to be excluded. Hence the only possible location of the planets responsible for the inner edge is between the yellow ellipse and the red one (which traces the inner edge of the disk). We see that along the major axis, only the planets closest to the inner edge (less than $\\simeq$ 10 AU) remain out of the present detection capabilities. Hence if a planet is responsible for the inner edge sculpting and is located along or close to the major axis, then it should be a low mass planet, and located further than 63 AU, ie within about 15 AU from the edge. Along the minor axis, due to the projection effects, the presence of planets is much less constrained: only planets at more than 26 AU from the edge would have been detected. The previous constrains were obtained assuming the planet and the perturbated bodies are both on circular orbits. The actual disk eccentricity beeing very small, about 0.02, this assumption is reasonable. As an exercice, we investigate the impact of a higher eccentricity, using the results of \\cite{mustill11}, who revisited this scenario, assuming the perturbating bodies were on an eccentric orbit; the relation becomes : $\\delta$ a/a = 1.8.e$^{1/5}$.(Mp/Ms)$^{1/5}$. With the same reasonning, and assuming an eccentricity of 0.1, we provide in Figure\\ref{map_planete} the possible locations of a planet responsible for the inner edge, with the same color conventions as in the circular case. Again, comparison with Figure\\ref{limdet2D_hr4796A} shows that if a planet was responsible for the inner edge sculpting, and located along or close to the semi-major axis, then it would have to be located less than $\\simeq$ 25 AU from the edge of the disk. The location of planets along or close to the minor axis would not be significantly constrained. This case, even though not adapted to the present case as the disk eccentricity is very small, illustrates the impact of this parameter on the planet detection capabilities. \\begin{figure} \\centering \\includegraphics[angle=0,width=.9\\hsize]{figures/sketch_new.ps} \\caption{ Possible remaining locations of a planet inner to the disk that could produce an inner sharp edge at 77 AU given the present detections limits (Left: case of a circular orbit; Right: case of an elliptical orbit with e=0.1). The possible locations are restricted to the region between the yellow curve and the red ellipse that mimics the disk itself. North is up, East to the left. } \\label{map_planete} \\end{figure} \\subsection{The outer disk sharp edge} Another striking feature of the system is the very steep disk outer edge. Radiation pressure from A-type stars induces surface brightness distributions in the outer part of disks with slopes of -3.5 to -5, depending on the assumptions related to the production laws of small grains (see for instance \\cite{leca96}, \\cite{thebault08} and ref. there-in). In particular, \\cite{thebault08} modeled the outer parts of collision rings with an initially steep outer edge, following the motion of the small grains produced through collisions and submitted to radiation pressure and showed that, as already proposed, the profile of the resulting SBD in the outer part of the disk followed a r$^{-3.5}$ law. They showed that unless the disks are extremely and unrealistically dense, and prevent the small grains from escaping, the disks have to be extremely \"cold\" (with an average free eccentricity of $\\leq$ 0.0035) to explain an outer power law of r$^{-6}$ (which was the value adopted at this time for the HR4796 profile). AO data suggest that the situation could be even more radical with an even steeper outer edge than previously thought. Also, if confirmed, the fact that we possibly find at 3.8 $\\mu$m a disk width different from that found in the optical (0.2-1$\\mu$m) with STIS would argue against an extremely dense disk, as, in such a case, large bodies and small grains would be in the same regions. The cold disk scenario can nonetheless be also problematic as, in such a case small grains should be underabundant (\\cite{thebault08}) and the optical/near-IR fluxes would be produced by large particules (typ. sizes 50 $\\mu$m). We should expect then a color index different from that observed (see \\cite{debes08}). The latter rather predict that the scattered light flux is dominated by 1.4 $\\mu$m dust, which seems difficult to explain within the dynamically-cold-disk scenario. Could gas be responsible for the observed steep outer edge? \\cite{takeuchi01} investigated the impact of the gas in such debris disks, and showed that even small amount of gas, 1- a few Earth masses, could partly balance the effects of gravitation, radiation pressure and Poynting-Robertson drag and alter the grains dynamics differentially, and lead to grains spatial distributions, that, depending on the grain sizes, could be different from those expected in a disk-free gas. Under such processes, the gas could be responsible for ring-like structures at distances depending on the dust size considered. In their attempt to investigate disks roughly similar to HR4796 and HD141569, assuming 1 MEarth gas, they showed that, conversely to large grains which occupy the whole gas disk, grains with sizes ($\\simeq$ 10-200 $\\mu$m) tend to concentrate in the outer gas disk, where the gaseous density sharply decreases. Hence these grains would form a narrow ring which position traces the change in the radial distribution the gaseous disk. Under this {\\it a priori} attractive scenario, grains with sizes 1-10 $\\mu$m would still be blown away. The main problem with this hypothesis is that it requires a gas disk with a relatively sharp outer edge, and thus an explanation for such an edge. Another issue is that \\cite{takeuchi01} did not explore the SBD profile beyond the main dust ring, so that it remains to be see whether slopes in the -10 range are possible. Finally, it is worth noting that so far no circumstellar gas has been detected either in atomic species through absorption spectroscopy (but the system being inclined, the non detection is not a strong constraint) or molecular species, either CO (\\cite{liseau99}) or H2 (N(H2)$\\leq 10^{15} cm^{-2}$ \\cite{martin08}). In any case, such a scenario can be tested in the forthcoming years with high angular resolution observations on a wide range of wavelengths. We finally study the possibility that the outer disk is sculpted by massive bodies. The first candidate we might think of is HR4796B. \\cite{thebault10} investigated the possibility that the disk could be sculpted by HR4796B, if orbiting on a rather eccentric orbit (e$\\geq$ 0.45) but again showed that, even under such conditions the outer profile would not be so steep. This is mainly because the companion star is not able to dynamically remove small grains from the outer regions at a pace that can compensate for their steady collisional production in the parent body ring. An alternative explanation could be the presence of a close, unseen outer planet. We investigate this scenario using the new code developed by \\cite{thebault2012} to study perturbed collisionally active debris discs. The code computes the motion of planetesimals submitted to the gravitational perturbation of a planet; and follows the evolution of small dust realeased through collisions among the planetesimals and submitted to radiation pressure and Poynting-Robertson effect (note that in the present case, radiation pressure largely drives the grains dynamics once produced). The configuration we consider is a narrow ring of large parent bodies, a birth ring of width $\\sim 8$AU centered on 71 AU. The collisional production and destruction rates of small grains (those which contribute to the luminosity beyond the main ring) in this parent body ring is parameterized by the average vertical optical depth within the main ring, taken to be $\\tau = 5\\times 10^{-3}$ \\footnote{for more detail on the procedure, see \\cite{thebault2012}}. The resulting SBD (case face-on) is derived assuming grey scattering. It can be directly compared to the curve presented in Figure 6 of Schneider et al (2009) paper (de-projected curve\\footnote{the FWHM of the STIS PSF being quite narrow compared to the disk width, its impact is very limited.}). For the perturbing planet's mass, we consider 3 different values: $8M_{Jup}$, $5M_{Jup}$ and $3M_{Jup}$, which are consistent with the constraints imposed by our observational non-detection. Note that $8M_{Jup}$ is only marginally possible in a very narrow region along the disc's semi-minor axis, but we have to keep in mind that the detection limits are derived from masses-brightness relationships that are debated at young ages. We assume a circular orbit for the planet (the most favourable case for cleaning out the region beyond the ring, see \\cite{thebault2012}) and place it as close as possible to the parent body ring, i.e., so that the outer edge of the observed ring (around 75 AU) corresponds to the outer limit for orbital stability imposed by planetary perturbations. This places the planet at a distance to the central star comprised between $92$ and $\\sim 99$ AU depending on its mass. In Figure~\\ref{dyn_simu}, we show the SBD obtained for such a configuration. Note that, for each planet mass, we are not showing an azimutal average but the \"best\" radial cut, i.e. the one that gives the closest match to the deprojected NE side SBD obtained in Figure 6 of \\cite{schneider09}. As can be clearly seen, the $8M_{Jup}$ case provides a good fit to the observed profile: the maximum of the SBD roughly corresponds to the outer edge of the parent body disk and is followed by a very sharp brightness decrease, with a slope $\\leq$ -10 between 75 and 95 AU, i.e. between brightnesses of 1 and 0.1, a range corresponding approximately to the dynamical range accessible to the available images. This is significantly steeper than the one that would be expected if no planet was present (-3.5 according to \\cite{thebault08}) and is fully compatible with the observed sharp luminosity decrease. Longwards 95 AU, the flux level is lower ($\\leq$ 0.1 ADU) and the SBD is flatter (slope $\\simeq$ -3.8). We also note a plateau inside the parent body ring at a level of $\\sim 0.2$, due to both the inward drift of small grains because of the Poynting-Robertson effect and to the dynamical injection of particles after close encounters with the planet. Of course, not too much significance should be given to the SBD obtained inwards of the disc since our simulations (focused on the outer regions) do not consider any inner planet shaping the inner edge of the disc. Nevertheless, they show that, should \"something\" have truncated the disc at around 67\\,AU in the past, then the effect of one external planet on such a truncated disc could lead to an SBD compatible with observations in the inner regions. For the $5M_{Jup}$ case, the fit of the observed SBD is slightly degraded, but mostly in the region beyond 90\\,AU where flux levels are close to the 0.1 threshold. For the $3M_{Jup}$ case, however, the fit gets very poor for almost the whole outer region (keeping in mind that we are here showing the best radial cut). We conclude that a $8M_{Jup}$ planet located on a circular orbit at $\\sim 25\\,$AU from the main ring provides a satisfying fit (especially considering the non-negligeable uncertainties regarding flux values far from the main ring) to the observed SBD. According to the derived detection limits, such a massive planet would have been detected almost everywhere except in a very narrow region along the disc's semi-minor axis; however, we remind the uncertainties inherent to the models used to link planet masses and luminosities as a function of the system's age. In any case, even a less massive perturber of e.g. $5_{MJup}$ would still give an acceptable fit of the observed luminosity profile. The external planet scenario thus seems the most likely one for shaping the outer regions of the disc. Of course, these results are still preliminary and should be taken with caution. A more thorough numerical investigation should be carried out, exploring a much wider parameter space for planet masses and orbit, as well as deriving other outputs that can be compared to observations, such as 2-D synthetic images. Such a large scale numerical study exceeds the scope of the present work and will be the purpose of a forthcoming paper. Note also that the more general issue of how planets shape collisionally active debris disks will be thoroughly investigated in a forthcoming paper (Thebault, 2012b, in prep.). \\begin{figure*} \\centering \\includegraphics[angle=0,width=.45\\hsize]{figures/radhrsimu.ps} \\caption{ Synthetic surface brightness profiles obtained, using Thebault (2012)'s numerical model, for 3 different masses of a putative perturbing outer planet: $8M_{Jup}$, $5M_{Jup}$ and $3M_{Jup}$. For each case, the planet has a circular orbit and is placed as close as possible to the main ring of large parent bodies in order to truncate it at about 75AU without destroying it (see text for more details). The observed profile, derived from \\cite{schneider09}, is shown for comparison. For each planet mass, we show the radial cut (i.e., for one position angle along the disc) that provides the best fit to this observed profile. The horizontal line delineates approximately the part (above this line) of the SBD accessible to the observations assuming a dynamical range of 10. } \\label{dyn_simu} \\end{figure*}" }, "1207/1207.2089_arXiv.txt": { "abstract": "{We discuss the basic features of the propagation of Ultra High Energy Cosmic Rays in astrophysical backgrounds, comparing two alternative computation schemes to compute the expected fluxes. We also discuss the issue of the transition among galactic and extra-galactic cosmic rays using theoretical results on fluxes to compare different models. } ", "introduction": "\\label{intro} Ultra High Energy Cosmic Rays (UHECR) are the most energetic particles observed in nature with (detected) energies up to $3\\div 5 \\times 10^{20}$ eV. The experimental observations of UHECR are performed nowadays by the Auger observatory in Argentina \\cite{Auger} and by the HiRes \\cite{HiRes} and Telescope Array (TA) \\cite{TA} observatories in the USA. The observation of these particles poses many interesting questions mainly on their nature and origin. The study of the propagation of UHECR through astrophysical backgrounds can be very useful to interpret the observations and, may be, to have important insights on the possible sources of these particles. The propagation of UHECR from the sources to the observer is mainly conditioned by the interactions with the intervening astrophysical backgrounds such as the Cosmic Microwave Background (CMB) and the Extragalactic Background Light (EBL). While the propagation of protons is conditioned only by the CMB radiation field \\cite{Nuclei}, nuclei propagation is also affected by the EBL radiation \\cite{Nuclei}. Several propagation dependent features in the spectrum can be directly linked to the chemical composition of UHECR and/or to the distribution of their sources \\cite{Boncioli}. Among such features particularly important is the Greisin, Zatsepin and Kuzmin (GZK) suppression of the flux, an abrupt depletion of the observed proton spectrum due to the interaction of the UHE protons with the CMB radiation field \\cite{GZK}. The GZK suppression, as follows from the original papers, is referred to protons and it is due to the photo-pion production process on the CMB radiation field ($p+\\gamma_{CMB} \\to \\pi + p$). In the case of nuclei the expected flux also shows a suppression at the highest energies that, depending on the nuclei specie, is due to the photo-disintegration process on the CMB and EBL radiation fields ($A + \\gamma_{CMB,EBL} \\to (A - nN) + nN$) \\cite{Nuclei}. In any case, the interaction processes between UHE particles and astrophysical backgrounds will condition the end of the CR spectrum at the highest energies and the high energy behavior of the flux can be used as a diagnostic tool for the chemical composition of the observed particles. Another important feature in the spectrum that can be directly linked with the nature of the primary particles and their origin (galactic/extra-galactic) is the pair-production dip \\cite{dip}. This feature is present only in the spectrum of UHE extragalactic protons and, as the GZK, is a direct consequence of the proton interaction with the CMB radiation field, in particular the dip brings a direct imprint of the pair production process $p+\\gamma_{CMB} \\to p+e^{+} + e^{-}$ suffered by protons in their interaction with CMB radiation. The experimental observations show contradictory results mainly on the chemical composition of UHECR. While Auger favors a light (proton) composition at low energies and an heavy (nuclei) composition at the highest energies \\cite{Chem}, HiRes and TA show a proton dominated spectrum at all energies \\cite{Chem}. A clear understanding of the chemical composition of UHECR is of paramount importance in the study of such particles in particular in the determination of their possible sources and to tag the transition among galactic and extra-galactic CR \\cite{TransReview}. Apart from the observations on chemical composition, also the observations on fluxes show a certain level of disagreement among HiRes and TA on one side and Auger on the other side. HiRes and TA experiments show a flux with a clear observation of the protons GZK suppression and the pair-production dip, coherently with their composition observations \\cite{Flux}. The situation changes if the Auger results are taken into account. The latest release of the Auger data on flux shows a spectrum not compatible with the pair production dip with an high energy suppression not compatible with the GZK suppression as expected for protons. Signaling a possible deviation from a proton dominated spectrum to an heavier composition at the highest energies, confirming with the flux behavior the observations on chemical composition \\cite{Flux}. This puzzling situation, with different experiments favoring different scenarios, shows the importance of a systematic study of the propagation of UHECR in astrophysical backgrounds. In the present paper we will present two alternative ways to study the propagation of UHECR in astrophysical backgrounds: through a Monte Carlo (MC) approach and through an analytic computation scheme based on the kinetic theory of particles propagation. We will than compare theoretical results with observations, considering in particular the issue of the transition from galactic to extra-galactic CR. The paper is organized as follows: in section \\ref{prop} we will briefly review theoretical models to study UHECR propagation, in section \\ref{trans} we will use the results of the propagation models to determine the features of the transition among galactic and extra-galactic CR, in this section we will use the results of \\cite{AmatoBlasi} for the galactic component, finally we will conclude in section \\ref{conclude}. ", "conclusions": "\\label{conclude} The theoretical study of the propagation of UHECR through astrophysical backgrounds has reached a remarkable level of refinement with several different approaches providing very reliable results. Here, in particular, we have compared the kinetic approach of \\cite{Nuclei} with the MC approach presented in \\cite{SimProp} showing the very good agreement of their results (figure \\ref{fig1}). Using these theoretical tools it is possible to study with good accuracy different models for the production of UHECR by comparing experimental data with theoretical expectations. In the present paper we have considered the dip model \\cite{dip} and the disappointing model \\cite{disapp}. The dip model is based on the assumption of a pure proton composition of UHECR with an injection power law index $\\gamma_g>2.5$, the experimental data of HiRes and TA confirm with very good accuracy this model (figure \\ref{fig2}, left panel) while the Auger data disfavor it with an heavier composition at the highest energies. The disappointing model, on the other hand, is based on the assumption of a mixed composition at the sources with an injection power law index $\\gamma_g<2.5$ and a relatively low maximum energy for protons ($E_{max}^p=4\\times 10^{18}$ eV). This model provides a very accurate description of the Auger data (figure \\ref{fig2}, right panel), while it does not reproduce the HiRes and TA observations. To better characterize these theoretical models, we have also focused the attention on the transition among galactic and extra-galactic CR assuming the galactic component as computed in \\cite{AmatoBlasi}. In the case of the dip model the transition is a sharp change from heavy (galactic) to light (extra-galactic) dominance, starting at energies around $10^{17}$ eV. The total flux in the case of the dip model shows a very good agreement with the all particle spectrum observed by different experiments (figure \\ref{fig3}, left panel). In the case of the disappointing model the transition is placed at energies around $10^{18}$ eV and the total flux seems not in perfect agreement with observations in particular at energies around the transition region (figure \\ref{fig3}, right panel)." }, "1207/1207.2795_arXiv.txt": { "abstract": "We report results from a deep Jansky Very Large Array (JVLA) search for $^{12}$CO $J=1-0$ line emission from galaxies in a candidate galaxy cluster at $z\\sim1.55$ in the COSMOS field. We target four galaxies with optical spectroscopic redshifts in the range $z=1.47-1.59$, consistent with the likely redshift for the candidate galaxy cluster. Two of these four galaxies, ID 51613 and ID 51813, are nominally detected in CO $1-0$ line emission at the $3-4\\sigma$ level. We find CO luminosities of $(2.42\\pm0.58)\\times10^{10}$ K km s$^{-1}$ pc$^2$ and $(1.26\\pm0.38)\\times10^{10}$ K km s$^{-1}$ pc$^2$, respectively. Taking advantage from the clustering and expanded 2-GHz bandwidth of the JVLA, we perform a search for emission lines in the proximity of optical sources within the field of view of our observations ($60''$). We limit our search to galaxies with $K_\\mathrm{S}<23.5$ (AB) and $z_\\mathrm{phot}=1.2-1.8$. We find 2 bright optical galaxies, ID 51207 and ID 51380, to be associated with significant emission line peaks ($>4\\sigma$) in the data cube, which we identify with the CO $1-0$ line emission. To test the reliability of the line peaks found, we performed a parallel search for line peaks using a Bayesian inference method. Both CO line emitting candidates are identified with probabilities of 13\\% and 72\\% that there are line emitting sources in each case, respectively. Monte Carlo simulations show that such associations are statistically significant, with probabilities of chance association of 3.5\\% and 10.7\\% for ID 51207 and ID 51380, respectively. Modeling of their optical/IR spectral energy distributions (SED) indicates that the CO detected galaxies and candidates have stellar masses and star formation rates (SFRs) in the range $(0.3-1.1)\\times10^{11}\\ M_\\odot$ and $60-160\\ M_\\odot$ yr$^{-1}$, with star formation efficiencies (SFEs) comparable to that found in other star-forming galaxies at similar redshifts. By comparing the space density of CO emitters derived from our observations with the space density derived from previous CO detections at $z\\sim1.5$, and with semi-analytic predictions for the CO luminosity function, we suggest that the latter tend to underestimate the number of CO galaxies detected at high-redshift. Finally, we argue about the benefits of future searches for molecular gas line emission in clustered fields with upcoming submillimeter/radio facilities. ", "introduction": "\\begin{figure*} \\centering \\includegraphics[scale=0.6]{fig1.ps} \\includegraphics[scale=0.6]{fig2.ps} \\caption{{\\it (Left:)} $Bi'K$ color composite of a $3'\\times3'$ region around the candidate cluster center. Cyan contours show the XMM {\\it Newton} X-ray emission at $2, 3$ and $4\\sigma$ significance. White circles show spectroscopically confirmed galaxies with redshifts $z_\\mathrm{spec}=1.470-1.595$. Yellow circles represent spectroscopically confirmed galaxies with $z_\\mathrm{spec}=1.2-1.8$. A large $30''$ radius circle indicates the location and field of view of our deep JVLA observations of the molecular gas. {\\it (Right:)} Close-up to the central $80''\\times80''$ around the JVLA pointing position. Labels indicate the sources ID tags.\\label{fig:field}} \\end{figure*} One of the major goals of modern observational cosmology is to understand how the gas in the diffuse interstellar medium (ISM) of galaxies converts into stars and how both phases evolve with cosmic time. A major advance has been the determination of the evolution of the SFR density of the Universe. Deep optical and radio surveys indicated that the SFR density steadily increases from $z=5$, with a peak at $z\\sim3-1$, and steeply declines from $z=1$ to the present \\citep[e.g., ][]{Lilly1996, Madau1996, Cowie1999, Steidel1999, Giavalisco2004, Ouchi2004, Bouwens2007, Smolcic2009}. Although the contribution from luminous, merger-driven starburst galaxies at high-redshift to the SFR density appears to be significant, the population that drives such evolution appears to be formed by more quiescent massive star-forming galaxies \\citep{Daddi2007, Rodighiero2011}. The close relationship between the SFR and the amount of molecular gas in galaxies, from which stars form, suggest that the evolution of the SFR density is the result of the evolution of the molecular gas density in galaxies across cosmic times. Measuring the evolution of this molecular gas density is thus critical to understand the formation of molecular gas and stars in galaxies. For this, it is necessary to perform blind surveys of the sky that allow us to measure the molecular gas content in a CO flux limited sample as a function of redshift and luminosity \\citep[e.g. observations in the local Universe and simulations by ][]{Keres2003, Obreschkow2009a, Obreschkow2009b}. The technical limitations (e.g. bandwidth, sensitivity) of current submillimeter and radio facilities, along with the intrinsic faintness of the CO emission lines used to measure the amount of molecular gas in galaxies have, however, precluded such studies. The advent of sensitive interferometers such as the Jansky Very Large Array \\citep[JVLA; ][]{Perley2011} and the Atacama Large Millimeter Array \\citep[ALMA; ][]{Wootten2009} will make these studies possible. A huge step toward this goal has been the recent detection of significant CO line emission in massive star-forming galaxies at high redshift \\citep{Daddi2008, Daddi2010a, Tacconi2010, Geach2011}. The large molecular gas masses derived are comparable to that observed in submillimeter galaxies (SMGs), in the range $\\sim10^{10-11}\\ M_\\odot$, however these galaxies are forming stars at $\\sim5-10$ times lower rates, with typical SFRs in the range $\\sim50-300$ M$_\\odot$ yr$^{-1}$. These findings imply that they have low SFEs and thus consume the gas at longer timescales, in agreement with their long duty cycle times of $>0.5$ Gyr \\citep{Daddi2007}. From the analysis of the CO molecular gas and H$\\alpha$ ionized gas kinematics, it is expected that between 1/3 to 2/3 of the star forming galaxies at $z\\sim2$ are consistent with rotating disks \\citep{ForsterSchreiber2009, Tacconi2010}, and suggest that the dominant mechanism to cool gas into molecular clouds is driven by cold gas accretion rather than galaxy interactions \\citep{Keres2005, Dekel2009}. Follow-up multi-transition CO observations recently discovered that a cold gas component, with gas excitation conditions similar to that seen in local spiral galaxies, is present in these objects \\citep{Dannerbauer2009, Aravena2010}. Remarkably, observations at $z=0, z=0.5$ and $z=1-2$ suggest that there has been a strong evolution in the molecular gas fractions in star forming galaxies over the past $10$ Gyr \\citep{Tacconi2010, Lagos2011, Geach2011}. In this paper, we present results of a deep search for CO molecular gas from galaxies that are part of a candidate galaxy cluster at $z\\sim1.5$ with the JVLA. The scope of this paper is to perform the first attempt to conduct an efficient blind search for CO emission from galaxies located in a part of the sky that likely contains a higher than average number of CO emitters. The clustering of galaxies in angular scales and along the line of sight makes it possible to effectively increase the number of galaxies observed within a single pointing and frequency tuning with the JVLA. We use the extensive multi-wavelength dataset available for the COSMOS field, including the accurate photometric redshifts, to inform the location of optical galaxies and to guide the search for CO emission. In section 2, we present the properties of the (overdense) target field and the details of the JVLA observations. In section 3, we present the main results from this paper. We report the detection of CO $1-0$ emission in two spectroscopically confirmed galaxies and analyze the 45.5 GHz continuum emission and properties of a spectroscopically confirmed radio loud galaxy. In sections 3.3 and 3.4, we present the search for line emission peaks from optical galaxies in the field. Significant line peaks associated with optical sources that have compatible photometric redshifts are identified as CO line candidates. The line search is performed using two different algorithms. From sections 3.5 to 3.7, we analyze the likelihood that such associations are real based on both a statistical approach and on the detection of the CO candidates in the IR wavelength regime. In section 4, we discuss the derived SFEs for the two CO detected spectroscopically confirmed galaxies and for the two CO candidate emitters, as well as the implications of these results for future experiments to measure the CO source counts at high redshift. We adopt a concordance $\\Lambda$CDM cosmology throughout, with $H_0=71$ km s$^{-1}$ Mpc$^{-1}$, $\\Omega_M=0.27$, and $\\Omega_\\Lambda = 0.73$ \\citep{Spergel2007}. \\begin{figure} \\centering \\includegraphics[scale=0.5]{fig3.ps} \\caption{Projected number density of galaxies with $z_\\mathrm{phot}=1.5-1.6$ and SFR $>5$ M$_\\odot$ yr$^{-1}$ in a $3'\\times3'$ region around the cluster center. The grayscale and contours represent the density of galaxies in this field, $\\delta_C$, given in terms of the average density $\\delta_0$ of similarly selected galaxies in the COSMOS field, $\\delta = \\delta_C/\\delta_0$. Contours range from 2 to 7 in steps of +1. The red circle represents the location of the JVLA pointing and primary beam (PB) FWHM. \\label{fig:densitymap}} \\end{figure} \\begin{figure} \\centering \\includegraphics[scale=0.55]{fig4.ps} \\caption{Color-magnitude diagram ($z-[3.6]$ vs. $[3.6]$) of galaxies located within $r_{200}$ ($0.9'$) from the center of the cluster candidate. The gray dashed line shows a model red sequence for a cluster formation redshift of $z_f=3$ from \\citet{Lidman2008}. The gray star symbol represents the the characteristic magnitude $m^\\star$ for passively evolving galaxies. Blue filled squares correspond to galaxies in the redshift range $z=1.5-1.7$. The red filled circles represent galaxies located in the same redshift range but identified to belong to the forming red-sequence. Small black circles show background/foreground galaxies in the field. \\label{fig:cm_diagram}} \\end{figure} \\begin{figure} \\includegraphics[scale=0.55]{fig5.ps} \\caption{The distribution of photometric redshifts of galaxies located within $r_{200}$ from the center of the cluster candidate. The open histogram shows all the galaxies in the field. The red histogram shows the redshift of the red-sequence identified galaxies, and the dashed histogram indicate the redshifts of galaxies shown in blue in Fig. \\ref{fig:cm_diagram}. \\label{fig:zdistr_cm}} \\end{figure} \\begin{table*} \\centering \\caption{Observed CO properties for galaxies with spectroscopic redshift\\label{tab:1}} \\begin{tabular}{cccccccc} \\hline ID$^a$ & $\\alpha_{J2000}$ $^b$& $\\delta_{J2000}$ $^b$& $z_\\mathrm{opt}$ $^c$& $z_\\mathrm{CO}$ $^d$& $S_\\mathrm{CO}dv$ $^e$& $L'_\\mathrm{CO}$ $^f$& $M(\\mathrm{gas})$ $^g$ \\\\ & & & & & (Jy km s$^{-1}$) & ($10^{10}$ K km s$^{-1}$ pc$^2$) & ($10^{10}\\ M_\\odot$) \\\\ \\hline\\hline 50480 & 10\\ 02\\ 40.47 & +01\\ 34\\ 41.2 & 1.523 & - & $<0.053$ & $<0.7$& $\\ldots$ \\\\ 51130$^{\\dagger}$ & 10\\ 02\\ 42.25 & +01\\ 34\\ 32.5 & 1.519 & - & $<0.044$ & $<0.6$& $\\ldots$ \\\\ 51613 & 10\\ 02\\ 43.36 & +01\\ 34\\ 20.9 & 1.516 & 1.517 & $0.20\\pm0.05$ & $2.42\\pm0.58$ & $8.7\\pm2.1$ \\\\ 51858 & 10\\ 02\\ 40.43 & +01\\ 34\\ 13.1 & 1.560 & 1.556 & $0.10\\pm0.03$ & $1.26\\pm0.38$ & $4.5\\pm1.4$ \\\\ \\hline \\end{tabular} \\begin{flushleft} \\begin{footnotesize} \\noindent $^{\\dagger}$ Radio galaxy; $^a$ COSMOS ID; $^b$ Position of the optical source; $^c$ Optical redshift; $^d$ CO redshift; $^e$ Spatially and velocity integrated line flux; $^f$ CO luminosity; $^g$ Gas mass obtained using $\\alpha_\\mathrm{CO}=3.6\\ M_\\odot$ (K km s$^{-1}$ pc$^2$)$^{-1}$ \\end{footnotesize} \\end{flushleft} \\end{table*} ", "conclusions": "In this work, we have presented deep CO $1-0$ line observations of galaxies located in a galaxy cluster candidate at $z=1.5$ using the JVLA. The candidate cluster was identified using the red-sequence technique and is associated to an overdensity of $\\sim\\times7$ of star-forming galaxies. We use the spatial clustering and expanded bandwidth of the JVLA to simultaneously observe the CO emission from four galaxies with available optical spectroscopic redshifts in a single pointing and frequency setting. We detect, at the $\\sim3\\sigma$ and $\\sim4\\sigma$ level, the CO $1-0$ emission line from two of the galaxies with available spectroscopic redshifts: ID 51613 and ID 51858 at $z=1.516$ and $z=1.556$, respectively. We find $L'_\\mathrm{CO}=(3.9\\pm1.1)\\times10^{10}$ K km s$^{-1}$ pc$^2$ and $L'_\\mathrm{CO}=(1.3\\pm0.4)\\times10^{10}$ K km s$^{-1}$ pc$^2$, respectively. Both galaxies are also detected with {\\it Spitzer} at 24 $\\mu$m, and with {\\it Herschel} at 250 $\\mu$m and 350 $\\mu$m. While modeling of their dust properties indicates total IR luminosities of $5.7\\times10^{11}\\ L_\\odot$ and $9.5\\times10^{11}\\ L_\\odot$, respectively, characterization of their multi-wavelength SEDs suggests both are young star-forming galaxies that have formed a significant fraction of their stellar content, with stellar masses of $4.5\\times10^{10}\\ M_\\odot$ and $6.0\\times10^{10}\\ M_\\odot$, respectively. We performed a blind search for significant ($>4\\sigma$) emission line peaks in the JVLA data cube around the position of optical sources that were selected to have photometric redshift in the range $z=1.2-1.8$ and limited to $K_S<23.5$. This selection criteria ensures an accurate photometric redshift, while selecting galaxies with stellar masses $>2\\times10^9\\ M_\\odot$. We find that 2 of these selected optical galaxies are associated with significant emission line peaks ($>4\\sigma$), which are thus identified as CO $1-0$ line emission candidates. Both galaxies, ID 51207 and ID 51380, are found to have photometric redshifts in the range $1.4-1.6$. Based on Monte Carlo simulations of the distribution of such emission line peaks in the field, we find that any optical galaxy that follows our selection criteria and is located within $1''$ from a candidate emission line peak is statistically unlikely to be associated to such peak by chance. In the case of ID 51207, such probability is $3.7\\%$ while for ID 51380 we find 10\\%. Using in parallel a Bayesian inference approach to select emission line peaks in the data cube, we recover both CO-emitting candidate sources. Here, the line identified with ID 51380 is found to have a probability of $70\\%$ of being a CO emission line, however the line identified with ID 51207 is found to have a rather low probability of 13\\%. We find CO based spectroscopic redshifts of 1.530 and 1.551 with $L'_\\mathrm{CO}=(1.03\\pm0.25)\\times10^{10}$ K km s$^{-1}$ pc$^2$ and $L'_\\mathrm{CO}=(1.12\\pm0.26)\\times10^{10}$ K km s$^{-1}$ pc$^2$ for ID 51207 and ID 51380, respectively. These galaxies, ID 51207 and ID 51380, are detected at 24 $\\mu$m and 250 $\\mu$m with {\\it Spitzer} and {\\it Herschel}. Such measurements lead to IR luminosities of $4.9\\times10^{11}\\ L_\\odot$ and $3.3\\times10^{11}\\ L_\\odot$ for ID 51207 and ID 51380, respectively. In the former case, its optical/IR SED suggest a $\\sim$1 Gyr old, very massive star forming galaxy, with a stellar mass of $1.1\\times10^{11}\\ M_\\odot$, while for ID 51380 the SED indicates that it corresponds to a young star-forming galaxy, with a stellar mass of $2.8\\times10^{10}\\ M_\\odot$. In one case, we associated a line peak candidate with an optical source that had a previous optical spectroscopic redshift. This secure optical redshift of $z=1.240$ is incompatible with the one that would be derived from the identification of the line peak with the CO $1-0$ emission line. This optical source fits well within our source selection criteria, $K<23.5$ (AB) and $z_\\mathrm{phot}=1.2-1.8$, however if we had identified it with the CO line emission, the non-detection in the IR bands and its implied SFR from the SED fitting method would have implied a very low SFE. We argue that in the absence of high-significance line detections, tight restrictions should be imposed on sources to be identified with CO line emission, such as detection in the IR bands or a lower limit in SFRs. We measured the space density of CO galaxies compared to the space density of CO emitters estimated from the 6 BzK galaxies significantly detected in CO emission by Daddi et al., and compared to predictions from semi-analytic simulations. Overall, we find that all observations are only roughly consistent with the simulations, despite the low number of detections of typical star-forming galaxies at high-redshift. Clearly, observations of molecular gas from a statistically significant sample of these galaxies is necessary to measure the evolution of CO luminosity function with redshift, and thus constrain models of galaxy formation and evolution. We conclude by discussing the advantages of performing deep CO observations of star-forming galaxies in clustered fields compared to blank-fields in the sky." }, "1207/1207.5510_arXiv.txt": { "abstract": "The Extrasolar Planet Search with PRIMA project (ESPRI) aims at characterising and detecting extrasolar planets by measuring the host star's reflex motion using the narrow-angle astrometry capability of the PRIMA facility at the Very Large Telescope Interferometer. A first functional demonstration of the astrometric mode was achieved in early 2011. This marked the start of the astrometric commissioning phase with the purpose of characterising the instrument's performance, which ultimately has to be sufficient for exoplanet detection. We show results obtained from the observation of bright visual binary stars, which serve as test objects to determine the instrument's astrometric precision, its accuracy, and the plate scale. Finally, we report on the current status of the ESPRI project, in view of starting its scientific programme. ", "introduction": "\\label{sec:intro} High-precision astrometry with an accuracy in the 0.01-1 milli-arcsecond (mas) range is a powerful technique to detect and characterise extrasolar planets and substellar companions of dwarf stars\\cite{Pravdo:2005fu, Casertano:2008th, Sahlmann:2011fk, Lazorenko:2011lr}, especially when complemented with radial velocity measurements. With its sensitivity to the true companions mass, the astrometry technique is particularly suited to constrain the mass distributions of exoplanet populations\\cite{Sahlmann2012PhD}. So far, only few instruments are capable of realising the required precision and long-term stability $<1$~mas, but it has been demonstrated e.g. with VLT optical imaging\\cite{Lazorenko:2009ph}, the HST\\cite{Benedict:2010ph}, and infrared interferometry\\cite{Muterspaugh:2010lr}.\\\\ Using the new capabilities offered by the PRIMA facility\\cite{Delplancke:2008xr, Belle2008} at the Very Large Telescope Interferometer (VLTI) of the European Southern Observatory, we plan to carry out an astrometric planet search and characterisation programme, which requires an accuracy of 0.01-0.1 mas for the relative position measurement of two stars in a narrow field of view $\\lesssim30\\arcsec$. We are currently commissioning this mode of operation and once completed, the ESPRI scientific observing programme will begin and the astrometric mode of PRIMA will be offered to the community. PRIMA will then be the unique public facility capable of delivering high-precision relative astrometry to the general user.\\\\ % In Section \\ref{sec:principle}, we describe the measurement principle and its implementation with PRIMA at the VLTI. Section \\ref{sec:obs} discusses the test observations we have carried out on bright visual binary systems and Section \\ref{sec:res} presents the results. The status of the science programme of ESPRI is summarised in Section \\ref{sec:espri}. We conclude in Section \\ref{sec:conclusions} and indicate directions for the immediate future of the project in Section \\ref{sec:outlook}. ", "conclusions": "\\label{sec:conclusions} \\begin{enumerate} \\item The PRIMA dual-feed facility has successfully been integrated in the VLTI and is operational for astrometric observations. Several technical runs and three astrometric commissionings have taken place in 2011/12 with the goal of establishing the instrument's astrometric capabilities. \\item The astrometric precision of PRIMA lies in the expected range and a performance of 0.03 mas on bright small-separation ($\\lesssim$10\\arcsec) binaries has been achieved. % \\item For wide-separation binaries ($\\gtrsim$10\\arcsec) and long ($>$2 h) observing sequences, large systematic errors that are correlated with the field rotation are observed. Those errors can amount to several tens of micro-meters corresponding to tens of milli-arcseconds in astrometry. \\item By comparison with NACO observations, we find that the plate scale of PRIMA is accurate at better than $10^{-3}$. \\item The astrometric accuracy of PRIMA in the current (February 2012) state is limited to a few milli-arcseconds by systematic errors. There is evidence that those errors originate in the telescope beam train between the primary mirror and the M9 mirror, which is not monitored by the laser metrology system. The decision to place the metrology endpoint at M9 was taken during the design phase of PRIMA to limit the impact on the operational infrastructure of the VLTI. A design with the endpoints at the secondary mirror was proposed at the time, but deferred until proven necessary to reach the required astrometric accuracy\\cite{Delplancke2006}. A hardware modification to place the endpoints at the level of the secondary mirror is currently being investigated. \\item The astrometric performance of PRIMA as of February 2012 is insufficient to allow us to begin with the ESPRI programme for extrasolar planet search and characterisation. \\item The preparation of the ESPRI target list has been concluded and we will begin with scientific observations of exoplanet targets when the long-term astrometric accuracy of PRIMA will have reached the $\\sim$0.1 mas level on targets with $m_K=7-8$ and $\\Delta m_K=2-6$. \\end{enumerate}" }, "1207/1207.0429_arXiv.txt": { "abstract": "The half-skyrmions that appear in dense baryonic matter when skyrmions are put on crystals modify drastically hadron properties in dense medium and affect strongly the nuclear tensor forces, thereby influencing the equation of state (EoS) of dense nuclear and asymmetric nuclear matter. The matter comprised of half skyrmions has vanishing quark condensate but non-vanishing pion decay constant and could be interpreted as a hadronic dual of strong-coupled quark matter. We infer from this observation {combined with certain predictions of hidden local symmetry in low-energy hadronic interactions } a set of new scaling laws -- called ``new-BR\" -- for the parameters in nuclear effective field theory controlled by renormalization-group flow. They are subjected to the EoS of symmetric and asymmetric nuclear matter, and are then applied to nuclear symmetry energies and properties of compact stars. The changeover from the skyrmion matter to a half-skyrmion matter that takes place after the cross-over density $n_{1/2}$ provides a simple and natural field theoretic explanation for the change of the EoS from soft to stiff at a density above that of nuclear matter required for compact stars as massive as $\\sim 2.4M_\\odot$. Cross-over density in the range $ 1.5n_0 \\lsim n_{1/2} \\lsim 2.0 n_0$ has been employed, and the possible skyrmion half-skyrmion coexistence {or cross-over} near $n_{1/2}$ is discussed. The novel structure of {the tensor forces and} the EoS obtained with the new-BR scaling is relevant for neutron-rich nuclei and compact star matter and could be studied in RIB (rare isotope beam) machines. ", "introduction": "The topological soliton called skyrmion~\\cite{skyrme} has turned out to be exceedingly pervasive in a variety of space-time dimensions ranging from 3 to 5 in many areas of physics~\\cite{multifacet} and has been beautifully observed in such systems as quantum Hall or cold atoms and more recently in a monoatomic magnetic film (see e.g., \\cite{heinze}). In contrast, the situation with its role in nuclear physics has been much less clear and with rather limited success. In this note, we make an attempt to uncover the power, hitherto unexploited, of skyrmion in strong interaction physics focusing on nuclear and dense matter. In contrast to condensed matter, the effect of skyrmion structure in strong interactions turns out to be indirect and hence much less transparent. In this work we show with simple plausible arguments that the skyrmion picture can indeed make a novel prediction on the properties of compact stars that has not been made thus far by other approaches. The arguments made in formulating the theoretical framework are neither rigorous nor completely unambiguous. { Although the crystal structure which is valid in the large $N_c$ limit could be applicable at very large density, it is not clear that it can be used in the density regime that we are concerned with, which will be a few times the normal nuclear matter density. What we will be exploiting is, however, the topological structure provided by the skyrmion configuration, which is insensitive to spatial symmetry.} In proceeding we will rely on what Nature indicates at normal densities and then extrapolate to high densities using a hidden local symmetry (HLS) structure with well-defined degrees of freedom . The starting point of our work is that when a large number of skyrmions as baryons are put on an FCC (face-centered-cubic) crystal to simulate dense matter, the skyrmion matter undergoes a transition to a matter consisting of half-skyrmions~\\cite{goldhaber} in CC configuration at a density that we shall denote as $n_{1/2}$. This density is difficult to pin down precisely but it is more or less independent of the mass of the dilaton scalar, the only low-energy degree of freedom that is not well-known in free space. It has been estimated to lie typically at between 1.3 and 2 times the normal nuclear matter density $n_0$ ~\\cite{half}. The half-skyrmion phase, made up of fractionized baryon numbers, is characterized by the quark condensate $\\la\\bar{q}q\\ra$ that vanishes on the average in the unit cell with, however, chiral symmetry still broken, so the pion is present. It likely has an inhomogeneous spatial distribution of baryon density. There is no obvious order parameter for the ``transition\" although there can be higher-dimension field operators representing an emergent symmetry that could be identified at quantum level. What can distinguish the two ``phases\" are the different degrees of freedom with different topological charges. Among the predictions made so far with the half-skyrmion phase, the most striking one -- which is the main object of this article and has not been made by other approaches -- was that the presence of $n_{1/2}$ strongly modifies the tensor forces in nuclear interactions and in particular the symmetry energy at densities $n> n_{1/2}$~\\cite{LPR,LR12}. In this note, we confront our predictions with nature by translating the (semi-)classical results of \\cite{LPR} into the parameters of an effective Lagrangian having chiral symmetry and conformal symmetry, and then do a quantum-EFT calculation for nuclear matter and compact-star matter using a renormalization-group (RG) based formalism~\\cite{kuo}. For the range $n_{1/2}= (1.5-2.0) n_0$ considered in \\cite{LPR}, we have applied our formalism to neutron star calculations and a comparison of our results with the recently discovered two-solar-mass neutron star \\cite{demorest2010} will be discussed. An interesting result of the calculation is that the skyrmion-half-skyrmion crossover makes the EoS stiffer at the crossover density $n_{1/2}$ and beyond, thereby leading to more massive stars. In a nutshell, our strategy is as follows. Up to the nuclear matter density $n_0$, our nuclear effective field theory (EFT) will be guided by symmetries of low-energy QCD (such as chiral symmetry, hidden local symmetry etc.) backed by nuclear phenomenology available up to density near $n_0$. There the effective Lagrangian will be endowed with parameters suitably scaling in the vacuum sliding with the density. We will assume that one can use the same EFT up to the density $n_{1/2}$ at which half-skyrmions appear which we take to be above but not far above $n_0$. Above $n_{1/2}$ for which there are neither experimental data nor model-independent theoretical tools available, we will take the properties indicated by the skyrmion-half-skyrmion transition based on hidden local symmetric Lagrangian and certain predicted property of hidden local fields as the chiral critical point is approached. The effective Lagrangian so given is then translated into an effective nuclear field theory that is subject to many-body techniques that account for high order quantum effects. ", "conclusions": "In this paper, we subjected the nuclear effective field theory anchored on RG flow, with the parameters of the Lagrangian sliding with density, to normal nuclear matter and dense compact-star matter. The scaling behavior used here differs from the old BR scaling~\\cite{BR91}, in that at a density $n_{1/2}>n_0$, a topological change takes place from skyrmion matter to half-skyrmion matter, giving rise to a modified scaling new-BR. The changeover from skyrmion matter to half-skyrmion matter is characterized by a vanishing quark condensate $\\la\\bar{q}q\\ra=0$ but a nonvanishing pion decay constat $f_\\pi\\neq 0$. Thus it is not a standard phase transition \\`a la Ginzburg-Landau-Wilson paradigm although two different phases are involved; it appears to involve an emergent symmetry not present in the fundamental theory, QCD. At the semi-classical approximation made in the calculation, the half-skrymions are not deconfined in contrast to what happens in certain condensed matter systems~\\cite{deconfined}. They are bound or confined, so they are not propagating degrees of freedom. What characterizes the system is that the mass of the baryon made up of two `bound' half-skyrmions remains more or less unscaled, not going to zero up to the density $n_c$ at which the quarks get deconfined, whereas the $\\rho$-meson mass is expected to drop faster in the half-skyrmion phase than in the skyrmion phase. This means that the origin of the most, if not all, of the nucleon mass is not in the dynamical symmetry breaking of chiral symmetry, in contrast to the meson mass, with a substantial mass of the nucleon coming from a hitherto unknown source. This is similar to what is described in the parity-doublet model of the nucleon~\\cite{detar,PLRS}.We should note however that this picture is clearly at odds with the constituent quark model -- which has a strong theoretical support from QCD in the large $N_c$ limit~\\cite{weinberg} -- where the ratio of the meson mass over the baryon mass is 2/3. Whether or not the constituent quark model is applicable in nuclear medium is not known, but if there were an $m_0$ for the quark which is not small, then it should be possible that the constituent quark model hold in dense medium and the ratio remain more or less the same. In this case, the scaling could be considerably different from the new-BR. It is intriguing that the two consequences of the changeover at $n_{1/2}$, namely, the drastic modification of the nuclear tensor force and the stiffening of the EoS of dense matter at $n_{1/2}$, seem to be hinting at the mechanism for the generation of $\\sim 99$ \\% of the nucleon mass in the strong interactions. See \\cite{MR-mass} for discussions on this matter. The salient features obtained in the RG-implemented effective theory approach adopted in this paper can be summarized as follows: \\begin{enumerate} \\item Without a suitable scaling in the Lagrangian that figures in $V_{low-k}$ (or incorporating many-body forces), symmetric nuclear matter cannot be stabilized at the right density and with correct binding energy. \\item Our calculations have essentially two scaling parameters: one is $c_I \\approx 0.13$ for all mesons (vector mesons and scalar meson) and the nucleon in region I, and in region II we have $c_{II}=c_{I}$ for mesons and the vector coupling and an additional parameter $y(n)\\approx 0.8$ for the nucleon. With these two parameters, one can explain satisfactorily the saturation density, the binding energy and the compression modulus of symmetric nuclear matter as well as the nuclear symmetry energy, and predict the EoSs for symmetric and asymmetric nuclear matter at high density and compact-star matter. Our results give a good fit to all quantities that are available experimentally at densities up to $n\\sim 4n_0$. \\item The topology change from skyrmion to half-skyrmion at $n_{1/2}$ changes the slope of the EoS, making it stiffer in the half-skyrmion phase and raises the maximum mass of compact stars to $\\sim 2.4M_\\odot$. Verifying the presence and the role of the topology change at $n_{1/2}$ should be feasible at RIB machines. \\end{enumerate} In our treatment, $n$-body forces for $n>2$ have not been taken into account. As mentioned, 3-body forces -- in place of BR -- could equally well provide the repulsion needed to stabilize nuclear matter. This does not mean that the many-body forces and the BR are alternatives. They should both come in together. In principle, there should be no problem in including both BR and many-body forces in a way consistent with the tenet of chiral expansion. What one has to do in the presence of such n-body potentials is then to suitably modify the scaling properties of the Lagrangian, since direct and indirect chiral symmetry effects are compounded in physical quantities in a variety of different chiral expansion schemes as illustrated in \\cite{FR}. A fully consistent way of doing the calculation would be to have {\\em both} the scaling and many-body potentials treated together with certain constraints, such as thermodynamic consistency, taken into account. We also note that our EoS is very close to the EoS found in Ref.[40] with a similar stiffening throughout the range of density considered, where the sound velocity never exceeds 0.9. We have not taken into account strangeness degrees of freedom -- such as kaons, hyperons, strange quarks etc. -- into the EoS for neutron-rich matter. In our formulation anchored on dense skyrmion matter, as described in \\cite{LR12}, hyperons can enter only {\\em after} kaons condense. Therefore the issue here is how kaon condensation can take place after changing from skyrmion matter to half-skyrmion matter. There are two opposing mechanisms to consider in the process. One is that in the presence of the topology change at $n_{1/2}$, the mass of $K^-$ has a propitious drop not present in conventional chiral perturbation treatments~\\cite{PKR}. This goes in the direction of lowering the critical density for kaon condensation. The other is the effect of stiffening the EoS. It is known for instance in phenomenological studies that the more repulsion there is in non-strange nuclear interactions, the higher the kaon condensation critical density goes up~\\cite{pandha}. What will happen in compact stars therefore will depend crucially on which one dominates. One intriguing possibility is that the stiffening postpones the drop of $m_K^*$ in a manner analogous to the stiffening at the smooth crossover at a density $\\sim (2-4) n_0$ from hadron to non-strange quark phase in the hybrid model that also yields the maximum star mass $\\sim 2.3 M_{\\odot}$~\\cite{masuda}. This will also have an important impact on the cooling of the star, since the appearance of strange flavor at higher density will prevent fast direct URCA process from setting in too precociously. It should be stressed that in our approach, strangeness in the form of condensed kaons (or equivalently hyperons) may enter at near or even before the density to which our theory with topology change can be extended, say $\\sim 4.5n_0$. Therefore the extrapolation beyond such density with polytropes, without accounting for strangeness degrees of freedom, potentially violating causality, should be taken as merely exploratory. One important aspect in our treatment that requires serious studies is the correlation between the behavior of the in-medium nucleon mass $m_N^*$ and that of the in-medium $\\omega$-N coupling $g_{\\omega NN}$ which is related to the $U(1)$ gauge coupling $g^*_\\omega$. We have adopted in our calculation the information from the skyrmion crystal calculations~\\cite{half,ma-crystal} and the parity-doubling nucleon model~\\cite{PLRS} that the nucleon mass drops only about 20\\% up to the highest density we are considering. We have taken the scaling $y$ effective in Region II to be constant as indicated in the skyrmion-crystal calculation~\\cite{ma-crystal} and in the one-loop RG analysis of HLS Lagrangian. As stated, were we to drop the $\\omega$-nucleon coupling according to $g^*_\\omega/g_\\omega\\approx g^*_\\rho/g_\\rho\\approx g^*/g=\\Phi_{II}$ as one would expect if flavor $U(2)$ symmetry held in Region II, the EoS would become much too soft above $n_0$ to be compatible with the existence of the 2-solar mass object observed in nature. We kept $g^*_\\omega/g_\\omega\\approx 1$ while letting the $\\omega$ mass scale. Now to quantify the above observation, we have examined the effect of dropping $\\omega$-NN coupling for given $m_N^*$s. Writing the $\\omega$-nucleon coupling in Region-II as $g^*_\\omega/g_\\omega=(1+c_{II,N\\omega}~n/n_0)^{-1}$, we have found at $n=2.5n_0$, $E_0/A= (-33.9, -50.9,-68.4) $ MeV for y(n)=0.77 and $E_0/A= (11.68, -1.26, -14.56)$ MeV for y(n)=0.60 for the scaling constant of the $\\omega$-NN coupling $c_{II,N\\omega} = (0.046, 0.093, 0.139)$ with all other parameters fixed to (A) of Fig. 3. One sees that the EoS is extremely sensitive to the in-medium properties of both the nucleon mass and the $\\omega$-NN coupling. There are two implications that follow from this calculation. One is that $U(2)$ symmetry can be badly broken in dense medium and as a consequence the vector manifestation of HLS~\\cite{HY:PR} does not apply to the in-medium $\\omega$ meson although its mass may approach zero as the $\\rho$ mass does \\`a la mended symmetry. The other is that the in-medium nucleon mass and $\\omega$-NN coupling must be strongly correlated. One-loop renormalization group equations with the generalized hidden local symmetry Lagrangian implemented with baryons (with no dilatons) of \\cite{PLRS1} show that in the chiral limit, both $m_\\rho^*$ and $m_\\omega^*$ approach zero as the dilaton limit fixed point is approached. So does the nucleon mass $m_N^*$ in the standard (or ``naive\") assignment for the nucleon (see \\cite{PLRS1}). However while the vector manifestation of HLS~\\cite{HY:PR} requires that $g_\\rho^*/g_\\rho\\propto \\la\\bar{q}q\\ra^*/\\la\\bar{q}q\\ra\\rightarrow 0$ near chiral restoration, if $U(2)$ symmetry is violated in medium, the in-medium $\\omega$-NN is predicted to drop much more slowly than the $\\rho$-NN coupling~\\cite{PLRS}. At one-loop order the $\\omega$-nucleon coupling is found not to scale. It is only at two-loop and higher order that scaling sets in. One can see from the RGEs the interplay between the slow scalings of the coupling and nucleon mass. This behavior agrees qualitatively with what was noticed above where lowering the nucleon mass required reducing the coupling $g_\\omega$ in order to have the symmetry energy lie within the range given by heavy-ion data. \\subsection*" }, "1207/1207.1105_arXiv.txt": { "abstract": "Cosmic microwave background polarization encodes information not only on the early universe but also dark energy, neutrino mass, and gravity in the late universe through CMB lensing. Ground based surveys such as ACTpol, PolarBear, SPTpol significantly complement cosmological constraints from the Planck satellite, strengthening the CMB dark energy figure of merit and neutrino mass constraints by factors of 3-4. This changes the dark energy probe landscape. We evaluate the state of knowledge in 2017 from ongoing experiments including dark energy surveys (supernovae, weak lensing, galaxy clustering), fitting for dynamical dark energy, neutrino mass, and a modified gravitational growth index. Adding a modest strong lensing time delay survey improves those dark energy constraints by a further 32\\%, and an enhanced low redshift supernova program improves them by 26\\%. ", "introduction": "The cosmic microwave background (CMB) radiation temperature and polarization power spectra have provided a cornucopia of cosmological information, on both the early and late universe \\cite{wmap7, arbar,quad, bicep, act, spt}. In the next five years considerably more detailed information will delivered by the Planck satellite \\cite{planck} and ground based, high resolution polarization experiments such as ACTpol \\cite{actpol}, POLAR \\cite{polar}, PolarBear \\cite{polarbear} and SPTpol \\cite{sptpol}. Such high resolution experiments carry substantial information not only on the primordial perturbations but on the matter density power spectrum at redshifts $z\\approx1-5$. The matter perturbations gravitationally lens the primordial CMB and from the reconstructed deflection field one can extract the cosmological parameters that impact the matter power spectrum. These include properties of dark energy, including the time variation of its equation of state, the sum of neutrino masses, and the gravitational growth index to test general relativity. In this article we focus on the role of such near term CMB polarization information in constraining these cosmological parameters simultaneously, since all of them affect the CMB lensing. We will also bring in other dark energy survey information expected within the five year horizon to complement the CMB, to give an overall reasonable assessment of the constraints in 2017. We emphasize that the CMB is the main focus of this article and we take a fiducial combination of dark energy surveys with no intent to study all possible permutations. Section~\\ref{sec:cmb} gives a brief review of CMB lensing and describes the near term CMB data sets. In Sec.~\\ref{sec:other} we present the cosmological parameter set and the Fisher matrix constraints from the CMB, then adding in various dark energy probes. Leverage from possible further near term data is discussed in Sec.~\\ref{sec:further}. ", "conclusions": "\\label{sec:concl} The origin of cosmic acceleration is a fundamental mystery of physics, and the subject of active and varied observational efforts. We find and quantify that the newly maturing probe of CMB lensing can play a significant role. Within the next 5 years maps of the CMB lensing deflection field over wide areas will be obtained, considerably strengthening our cosmological knowledge. These experiments are underway or imminent, and combination with ongoing or near term dark energy surveys will provide a large step forward in evaluating dark energy, neutrino, gravity, and primordial perturbation properties. The high resolution, low noise, ground based CMB polarization experiments give powerful leverage. Indeed, in combination with all other survey probes in the commonly used, ``vanilla plus $w_0$, $w_a$'' parameter space the result would be a very strong dark energy figure of merit FOM[$w_0,w_a$]=406. However, it is overly restrictive to assume perfect knowledge of the sum of neutrino masses and the behavior of gravity, both of which affect the growth of structure we are using as a cosmological probe. We present results including simultaneous fits for not only the vanilla parameters but dark energy dynamics $w_0$ and $w_a$, gravity $\\gamma$, and sum of neutrino masses $\\sum m_\\nu$. Concentrating on data expected to arrive within the next 5 years, we examine constraints from realistic near term experiments measuring CMB satellite and ground based polarization, supernova distances, and the galaxy power spectrum, as the baseline combination of surveys. The complementarity between the different probes is highlighted and quantified. In addition to the cosmological parameter estimation and figures of merit in dark energy and gravity--neutrino planes, we analyze the dependence of the results on the sky area of the CMB space/ground overlap and of galaxy surveys. We further study the impact of weak lensing data. Interestingly, we find that substantial further leverage can be obtained in supplementing the canonical program with a modest strong lensing time delay survey, requiring of order $\\sim$200 orbits on the Hubble Space Telescope, and a tightly systematics controlled nearby supernova program. This overall combination of experiments would deliver knowledge of the dark energy time variation $w_a$ to 0.25 and of the sum of the neutrino masses to 0.055 eV with data taken by $\\sim$2017. CMB lensing thus acts a powerful component of a dark energy and cosmology survey program, enabling us to close in chasing down cosmic acceleration. For further improvements between 2017 and when the next generation of cosmic volume surveys deliver data, promising paths forward include the use of cross-correlations between probes, the development of new probes (just as CMB lensing and strong lensing grew viable), reduction of weak lensing systematics, and the ability to accurately use smaller scales ($k_{max}>0.125\\,h$/Mpc) in the galaxy power spectrum." }, "1207/1207.6513_arXiv.txt": { "abstract": "We describe a simple method for estimating the vertical column density in Smoothed Particle Hydrodynamics (SPH) simulations of discs. As in the method of \\cite{PolytropicCooling}, the column density is estimated using pre-computed local quantities and is then used to estimate the radiative cooling rate. The cooling rate is a quantity of considerable importance, for example, in assessing the probability of disc fragmentation. Our method has three steps: (i) the column density from the particle to the mid plane is estimated using the vertical component of the gravitational acceleration, (ii) the ``total surface density'' from the mid plane to the surface of the disc is calculated, (iii) the column density from each particle to the surface is calculated from the difference between (i) and (ii). This method is shown to greatly improve the accuracy of column density estimates in disc geometry compared with the method of Stamatellos. On the other hand, although the accuracy of our method is still acceptable in the case of high density fragments formed within discs, we find that the Stamatellos method performs better than our method in this regime. Thus, a hybrid method (where the method is switched in regions of large over-density) may be optimal. ", "introduction": "Smooth Particle Hydrodynamics (SPH) \\citep{LucySPH,MonaghanSPH} is a Lagrangian technique for simulating fluid flows using a particle representation. This technique assigns each particle a mass, position and internal energy and then interpolates state variables, such as density, by ``smoothing'' over neighbouring particles. Gravity has been incorporated into the SPH formalism, allowing SPH to be used to simulate astrophysical fluids. However, the thermal evolution of many astrophysical systems is often governed by radiative transfer effects, in addition to hydrodynamic energy transfer. As an accurate description of a system's thermal evolution is of vital importance in many astrophysical systems (e.g. accretion discs, stellar gas clouds) a lot of effort has recently been made to add radiative transfer to SPH \\citep{RadiativeSPH1,RadiativeSPH2,PolytropicCooling,HybridCooling}. Unfortunately, a full three dimensional, frequency dependent description of radiative transfer is currently computationally impossible, so simplifying assumptions have to be made. Since SPH is a Lagrangian method, the purpose of modeling radiative transfer effects is to provide a net cooling (or heating) rate {\\it per particle} as a result of radiative processes. For example, a popular choice for simulating self-gravitating discs is to use a cooling time prescription, where the radiative cooling rate $\\dot{U}$ is set equal to $\\frac{-U}{t_{cool}}$ where $U$ is the cooling rate per unit mass and $t_{cool}$ is simply parameterised as a prescribed multiple of the local dynamical timescale \\citep{BetaCooling}. Although such a description grossly oversimplifies the underlying physics, it has a very low computational cost and so simulations including approximating radiative cooling can be run without sacrificing spatial or temporal resolution. Recently, Stamatellos et al have proposed a method that improves upon the cooling time prescription without significantly increasing the computational cost \\citep{PolytropicCooling}. Forgan et al extended this method to include heat transfer between particles, by combining the Stamatellos method with the flux limited diffusion (FLD) method \\citep{HybridCooling}. The Stamatellos method estimates the optical depth by assuming that the relationship between the two locally computed variables, density and gravitational potential, and the optical depth is the same as it is for a mass-weighted average of that relationship over a self-gravitating polytropic sphere. They then estimate the local cooling rate using only the local temperature and optical depth. This estimate tends to the radiative diffusion approximation at high optical depths but does not involve calculating noisy derivatives as is required in a proper implementation of radiative diffusion \\citep{SPHFLD}. The aim of such methods is not to achieve perfect agreement with the full radiative transfer equations, but to achieve an accuracy that is comparable in magnitude to the other uncertainties such as those associated with the grey approximation and the appropriate values of the frequency averaged opacity. However, although the Stamatellos method has proved successful for modelling spherical systems, Wilkins \\& Clarke have shown that it can systematically underestimate the cooling in disc geometries by as much as a factor of four in the region of the mid plane, where column density estimation is critical to estimating the cooling \\citep{WilkinsAndClarke}. This underestimate stems from the Stamatellos method overestimating the column density, $\\Sigma$, which appears as a quadratic term in the expression for the cooling rate in the optically thick limit. An example of where radiative transfer is important to the evolution of a disc system is the study of gravitational instabilities in proto-planetary disc. The cooling time prescription has been used to study these discs in great detail, but a more accurate description of the cooling in such systems is necessary to improve our understanding of these gravitational instabilities \\citep{DataPaper}. In this paper we propose a variation on the Stamatellos method for calculating the cooling rate for the special case of disc geometries, by improving the estimate of the column density in this case. Our method still only requires the use of local quantities, already calculated by the gravitational and hydrodynamic codes and as such remains computationally inexpensive. The paper is organized as follows. In section \\ref{sec:method}, we describe our new method in detail. In section \\ref{sec:analtests}, we test our method on a series of discs for which all quantities of interest can be obtained analytically or semi-analytically. In section \\ref{sec:realdiscs} we test our method on realistic disc simulations that have been evolved long enough to either fragment or reach marginal stability. Finally, section \\ref{sec:conclusion} summarises our conclusions. ", "conclusions": "\\label{sec:conclusion} In this paper we have presented a new method for efficiently and accurately estimating the cooling in disc geometries. To achieve this, our method estimates the column density between each particle and the surface of the disc, via estimates of the column density to the mid-plane and the total surface density. We verify the accuracy of our column density estimation (and hence our cooling rates) by comparison with discs for which the analytic form of the column density can be calculated. We test our method on a realistic proto-planetary disc simulation, that has been evolved for long enough to reach marginal stability and develop a typical spiral structure. Finally, we test our method on a fragmented disc and find that our method does not perform as well as the Stamatellos method within the fragments, due to the locally spherical geometry. We suggest that this shortcoming can be resolved by using the Stamatellos method only within regions of extremely high density (i.e. the fragments). We find throughout our tests that the accuracy of our method remains high (i.e. typical errors of order a few tens of per cent) and conclude that it is ideally suited for use in problems that depend on an accurate estimate of the cooling rate in disc geometries." }, "1207/1207.3793_arXiv.txt": { "abstract": "We present the results of interferometric observations of the cool core of Abell 1795 at CO(1-0) using the Combined Array for Research in Millimeter-Wave Astronomy. In agreement with previous work, we detect a significant amount of cold molecular gas (3.9 $\\pm$ 0.4 $\\times$10$^9$ M$_{\\odot}$) in the central $\\sim$10 kpc. We report the discovery of a substantial clump of cold molecular gas at clustercentric radius of 30 kpc (2.9 $\\pm$ 0.4 $\\times$10$^9$ M$_{\\odot}$), coincident in both position and velocity with the warm, ionized filaments. We also place an upper limit on the H$_2$ mass at the outer edge of the star-forming filament, corresponding to a distance of 60 kpc ($<$0.9 $\\times$10$^9$ M$_{\\odot}$). We measure a strong gradient in the H$\\alpha$/H$_2$ ratio as a function of radius, suggesting different ionization mechanisms in the nucleus and filaments of Abell1795. The total mass of cold molecular gas ($\\sim$7$\\times$10$^{9}$ M$_{\\odot}$) is roughly 30\\% of the classical cooling estimate at the same position, assuming a cooling time of 10$^9$ yr. Combining the cold molecular gas mass with the UV-derived star formation rate and the warm, ionized gas mass, the spectroscopically-derived X-ray cooling rate is fully accounted for and in good agreement with the cooling byproducts over timescales of $\\sim$10$^9$ yr. The overall agreement between the cooling rate of the hot intracluster medium and the mass of the cool gas reservoir suggests that, at least in this system, the cooling flow problem stems from a lack of observable cooling in the more diffuse regions at large radii. ", "introduction": "The search for cooling gas in the cores of galaxy clusters has a rich history. Since it was first discovered that the intracluster medium (ICM) in some galaxy clusters is cooling on timescales shorter than the age of the Universe \\citep[][]{lea73, cowie77, fabian77, mathews78}, considerable time and effort has been spent searching for the byproducts of this cooling. Surveys to detect line emission from cooling gas at $\\sim$10$^5$--10$^6$~K \\citep[OVI; e.g.,][]{bregman01, oegerle01a, bregman06}, $\\sim$10$^4$~K, \\citep[H$\\alpha$; e.g.,][]{hu85, heckman89, crawford99, edwards07, mcdonald10}, $\\sim$10$^3$~K \\citep[warm H$_2$; e.g.,][]{jaffe97, donahue00,edge02,hatch05,jaffe05}, $\\sim$10$^1$--10$^2$~K \\citep[{[}\\ion{C}{2}{]}; e.g.,][]{edge10, mittal11}, and $\\sim$10$^1$~K \\citep[CO; e.g.,][]{edge01,edge03,salome03,salome08} have continued to return the same result: there is considerably less cool gas in cluster cores than predicted by radiative cooling of the hot ICM. It is now assumed that one or a combination of energetic processes balance radiative cooling, allowing only a small fraction of the cooling ICM to reach temperatures below $\\sim$10$^6$~K \\citep[e.g.,][]{peterson03, peterson06}. Leading candidates for this source of heating are cosmic rays \\citep[e.g.,][]{ferland08,mathews09,fujita11} and feedback from the central active galactic nucleus \\citep[AGN; e.g.,][]{mcnamara07, guo08, ma11,randall11}. Assuming some fraction of the ICM cools unimpeded, it will eventually wind up as cold molecular gas. While much work has gone into quantifying the total mass of the cold gas reservoir in cluster cores \\citep[e.g.,][]{edge01,edge03,salome03}, it is considerably more challenging to study the detailed morphology of this gas. Only the Perseus cluster, due to its proximity, has CO detected with the same spatial extent as the soft X-ray- and optically-emitting gas \\citep{hatch05, salome11}. Mapping the distribution of the cold gas is crucial to understanding the complicated relation between cooling and feedback processes, and in estimating the amount of fuel for star formation and black hole growth. In this Letter, we present recent high spatial resolution interferometric observations of CO(1-0) in Abell 1795 (hereafter A1795) from the Combined Array for Research in Millimeter-wave Astronomy (CARMA). This well-studied cluster has a pair of 60 kpc long cooling filaments detected in X-ray \\citep{crawford05}, H$\\alpha$ \\citep[e.g.,][]{cowie83,mcdonald10}, and far-UV \\citep{mcdonald09}. Our observations cover the full extent of the cooling filaments, providing a first estimate of the efficiency of gas cooling from the hot phase to the cold phase removed from the influence of the AGN. In \\S2 we present these data, briefly describing their acquisition and reduction. In \\S3 we present the results of these observations, identifying regions with CO(1-0) line emission and determining the total mass in molecular gas. These results are discussed in the context of the cooling flow model, and compared to previous observations of A1795 and the Perseus cluster in \\S4. Finally, in \\S5 we conclude with a summary of these results and their implications for future work. Throughout this study, we assume a luminosity distance to A1795 of 263 Mpc. ", "conclusions": "We present CARMA observations of the cool core of Abell~1795 at CO(1-0). We find significant amounts of cold molecular gas at clustercentric radii of 0 kpc (3.9 $\\pm$ 0.4 $\\times$10$^9$ M$_{\\odot}$) and 30 kpc (2.9 $\\pm$ 0.4 $\\times$10$^9$ M$_{\\odot}$), and place an upper limit on the H$_2$ mass at 60 kpc ($<$0.9 $\\times$10$^9$ M$_{\\odot}$), assuming a Galactic value of X$_{CO}$. These CO(1-0) peaks are coincident in both position and velocity with the warm, ionized filaments and the brightest star-forming regions. The large amount of cold molecular gas at large radius along the cooling filament implies that the ICM is able to cool efficiently far-removed from the effects of the AGN. In light of these new observations, we compare the ICM cooling estimates to the mass of cooling byproducts and find agreement in the central 10~kpc and along the rapidly-cooling filaments. Furthermore, we find that in these dense regions the local classical cooling rate is in good agreement with both the spectroscopic cooling rate and the cooling rate implied by the cooling byproducts. These results suggest that the cooling flow problem, at least in this system, stems from a lack of observable cooling in the more diffuse (off-filament) regions at radii larger than $\\sim$ 10 kpc. In the near future, ALMA observations will allow extended CO emission to be mapped in a large sample of cool cores, shedding further light on the extent of the cooling flow problem." }, "1207/1207.6455_arXiv.txt": { "abstract": "In this paper we solve the hydrodynamical equations of optically thin, steady state accretion disks around Kerr black holes. Here, fully general relativistic equations are used. We use a new method to calculate the shear tensor in the LNRF (Locally Non-Rotating Frame), BLF (Boyer-Lindquist Frame) and FRF (Fluid Rest Frame). We show that two components of shear tensor in the FRF are nonzero (in previous works only one nonzero component was assumed). We can use these tensors in usual transonic solutions and usual causal viscosity, but we derive solutions analytically by some simplifications. Then we can calculate the four velocity and density in all frames such as the LNRF, BLF and FRF. ", "introduction": "Accretion disks are important in several astrophysical systems. They can be found around Young Stellar Objects(YSO), around compact stellar objects in our galaxy and around several super-massive black holes in Active Galactic Nuclei(AGN). In super massive accretion disks the mass of black hole is $(10^5-10^9)M_{\\bigodot}$. To study such disks we use the general relativity with relativistic hydrodynamics in Kerr metric background geometry. In relativistic Navier-Stokes fluid we have stress-energy tensor which is related to viscosity and is the cause of redistribution of energy and momentum in fluid. This tensor is defined by the four velocity and metric. In the previous studies, the relativistic and stationary solutions of standard black hole disks were solved. Lasota (1994) was the first who wrote down slim-disk equations which include relativistic effects. He also assumed that only r-$\\phi$ component of the stress tensor is nonzero. He used a special form for this component, which was followed by Abramowicz et al(1996 and 1997). Chakrabarti (1996) derived the transonic solutions of thick and thin disks for a weak viscosity. He assumed similar form for viscosity as Lasota (1994). Then Manmoto (2000) derived the global two temperatures structure of advection-dominated accretion flows (ADAFs) numerically by using full relativistic hydrodynamical equations including the energy equations for the ions and electrons. Papaloizou \\& Szuskiewicz (1994) introduced a phenomenological and non-relativistic equation for the evolution of viscous stress tensor (causal viscosity) which has been used by many authors. Gammie \\& Popham (1998) and Popham \\& Gammie (1998) solved ADAFs with relativistic causal viscosity. They used the Boyer-Lindquist coordinates. \\newpage Takahashi (2007b) solved the equations of relativistic disks in the Kerr-Schild coordinates by using the relativistic causal viscosity. In the papers of Gammie \\& Popham and Takahashi they assumed that in the FRF, only the $r-\\phi$ component of shear viscosity is nonzero, then they used the transformation tensors to derive the components of shear stress tensor in the Boyer-Lindquist or Kerr-Schild frames. In present study, we concentrate on the stationary axisymmetric accretion flow in the equatorial plane. We use a new method to calculate the shear tensor and azimuthal velocity of fluids in the locally non-rotating frame (LNRF) by using Keplerian angular velocity. We derive two kinds of shear tensors in LNRF and BLF; in the first one the direction of fluid rotation is the same as that of black hole ($\\Omega^{+}$) and in the second one the direction of fluid rotation is in opposite to that of black hole($\\Omega^{-}$). We calculate the components of the shear tensor for two kinds of fluids in LNRF, BLF and FRF, these calculations show that in FRF there are two nonzero components ($r-t$ and $r-\\phi$ components). But in previous papers, the only nonzero component in FRF was $r-\\phi$ component. The $r-t$ component results from the relativistic calculations of the shear tensor which changes some components of the four velocity. Then, by using these shear tensor components we calculate the four velocity in LNRF, BLF and density in all frames . This paper's agenda is as follows. We introduce the metric and reference frame in $\\S$2. In $\\S$3 basic equations are given. In $\\S$4 the shear tensor is calculated in FRF, LNRF and BLF. We derive four velocity in LNRF and BLF in $\\S$5 and the influence of two important parameters on density and four velocity can be seen in this section. Also, the influence of the $r-t$ component of shear tensor can be seen in this section. Summery and conclusion are given in $\\S$6. ", "conclusions": "" }, "1207/1207.4982_arXiv.txt": { "abstract": "{We show how the mass function of dense cores (CMF) which results from the gravoturbulent fragmentation of a molecular cloud evolves in time under the effect of gas accretion. Accretion onto the cores leads to the formation of larger numbers of massive cores and to a flattening of the CMF. This effect should be visible in the CMF of star forming regions that are massive enough to contain high mass cores and when comparing the CMF of cores in and off dense filaments which have different environmental gas densities.} ", "introduction": " ", "conclusions": "" }, "1207/1207.1569_arXiv.txt": { "abstract": "The accuracy of ground-based astronomical photometry is limited by two factors: photon statistics and stellar scintillation arising when star light passes through Earth's atmosphere. This paper examines the theoretical role of the outer scale $L_0$ of the optical turbulence (OT) which suppresses the low-frequency component of scintillation. It is shown that for typical values of $L_0 \\sim 25 - 50$~m, this effect becomes noticeable for a telescopes of diameter around $4$~m. On extremely large, $30 - 40$~m, telescopes with exposures longer than a few seconds, the inclusion of the outer scale in the calculation reduces the scintillation power by more than a factor of 10 relative to conventional estimates. The details of this phenomenon are discussed for various models of non-Kolmogorov turbulence. Also, a quantitative description of the influence of the telescope central obscuration on the measured scintillation noise is introduced and combined with the effect of the outer scale. Evaluation of the scintillation noise on the future TMT and E-ELT telescopes, predicts an amplitude of approximately $10~\\mu\\mbox{mag}$ for a 60~s exposures. ", "introduction": "Stellar scintillation is the random fluctuation of the radiation flux entering the aperture of the telescope caused by amplitude distortions of light wave passing through the turbulent terrestrial atmosphere \\citep{Tatarsky1967,Roddier1981}. The photometric error due to scintillation, the {\\it scintillation noise}, has been repeatedly studied \\citep[see, e.g.,][]{Young1969,Dravins1998}, as it often determines the fundamental limit of the accuracy of ground-based photometry. The basic dependencies required to calculate scintillation noise have been known by the wide astronomical community for quite a while now \\citep{Heasley1996,Gilliland1993}. However, these relationships were obtained for ideal cases and cannot always be used for accurate prediction of the scintillation noise. This became especially noticeable when accurate measurements of the intensity of the optical turbulence (OT) in the atmosphere at altitudes responsible for the occurrence of scintillation on large telescopes, began to be accessible \\citep{Kenyon2006}. The content of this paper is a theoretical study of two factors affecting the power of stellar scintillation on large and extremely large telescopes: the influence of the OT outer scale and the telescope central obscuration (CO). Section \\ref{sec:theory} recalls the basic description of the phenomenon and its relationship with several parameters of the OT in the atmosphere. In the following section, we assess the influence of the outer scale in the case of different and simplified non-Kolmogorov models. The effect of the CO is considered in Section \\ref{sec:co}, first for the Kolmogorov model, and then for a general case. In the last Section, the conclusions are formulated and a prediction of the scintillation noise for future extremely large telescopes is given. ", "conclusions": "\\label{sec:conclusion} In this paper we considered two factors that modify the scintillation intensity in observations on large ($4 - 10$~m) and extremely large ($20-50$~m) telescopes. The first effect is caused by deviation of the real OT from the Kolmogorov model at the scales of the order of 10~m. In the models which describe the real turbulence with using outer scale $L_0$, the scintillation power at low spatial frequencies is much lower than for pure Kolmogorov spectrum. Effect of the outer scale can be described by an additional function, which depends on the ratio of $D/L_0$. This function is equal to 1 when the telescope aperture is small and decreases with diameter increasing. We considered its general behavior by the example of piecewise power-law spectrum with the salient point at spatial frequency $L_0^{-1}$. Particular features were investigated by numerical integration for four different models: Von Karman and Greenwood-Tarazano models, and two exponential models. For the observed values of $L_0$ and models with a wide intermediate zone, the effect becomes visible on 4~m class telescopes. For measurements with long exposures this effect is more important than in the case of short exposures, and for the extremely large telescopes TMT and E-ELT it can reduce the scintillation power $\\sim 10$ times compared to classical estimates. For the Kolmogorov OT, the effect of amplification of the scintillation power, caused by the CO, inherent in every large telescope, results in the multiplication of the power for circular aperture by the function that depends only on the CO parameter. In the case of models with outer scale, the effect is described by a more complicated way, and it becomes significant for large apertures for both short and long exposures. The significant reduction of scintillation noise, due to the outer scale, enhances the potential of ground-based telescopes to study the variability of astronomical objects at $\\sim 10^{-5}$ level. This accuracy is sufficient to see, e.g., transit of Earth-like planet across the disk of Solar-like star. Drastic increase in the accuracy of fast photometry (up to $\\sim 10^{-4}$) makes it possible to study the micro-variability of many astronomical objects without the accumulation of long time series suitable for temporal spectral analysis." }, "1207/1207.0969_arXiv.txt": { "abstract": "The Chromosphere and Prominence Magnetometer (ChroMag) is conceived with the goal of quantifying the intertwined dynamics and magnetism of the solar chromosphere and in prominences through imaging spectro-polarimetry of the full solar disk. The picture of chromospheric magnetism and dynamics is rapidly developing, and a pressing need exists for breakthrough observations of chromospheric vector magnetic field measurements at the true lower boundary of the heliospheric system. ChroMag will provide measurements that will enable scientists to study and better understand the energetics of the solar atmosphere, how prominences are formed, how energy is stored in the magnetic field structure of the atmosphere and how it is released during space weather events like flares and coronal mass ejections. An integral part of the ChroMag program is a commitment to develop and provide community access to the ``inversion'' tools necessary for the difficult interpretation of the measurements and derive the magneto-hydrodynamic parameters of the plasma. Measurements of an instrument like ChroMag provide critical physical context for the Solar Dynamics Observatory (SDO) and Interface Region Imaging Spectrograph (IRIS) as well as ground-based observatories such as the future Advanced Technology Solar Telescope (ATST). ", "introduction": "\\label{sec:intro} The chromosphere is a deeply complex part of the solar atmosphere. It is named after its red appearance just after the beginning and before the end of a total solar eclipse, caused by emission in the \\Halpha\\ line at $656.3~\\nm$. The same lines that show emission in the flash spectrum allow us to observe the chromosphere on the disk. Early detailed observations \\cite{Secchi1877} already show the intricate structure of the chromosphere. We now know that the magnetized chromosphere is permeated by ``spicules'' at the limb and their on-disk counterparts (``mottles'' and ``fibrils'') \\cite{1968SoPh....3..367B}. Filaments and prominences are ``clouds'' of chromospheric material supported by a complex magnetic field and embedded in the corona that may erupt and produce Coronal Mass Ejections (CMEs). The chromosphere is of integral importance in the mass and energy balance of the outer solar atmosphere, with over 90\\% of the non-radiative energy deposited there \\cite{1983ApJ...267..825W}, and requiring nearly one hundred times the mass and energy flux of the corona for sustenance \\cite{1977ARA&A..15..363W}. The chromosphere remains the most poorly understood region of the outer solar atmosphere. The complex dynamics resulting from the interplay of magnetic field and convectively driven photospheric plasma is incredibly difficult to interpret unambiguously, yet critically important to our understanding of the mass and energy balance of the outer solar atmosphere, including such things as coronal heating, the solar wind, and energetic events collectively known as space weather. The study of the chromosphere is the next frontier in solar physics. The most fundamental unresolved issues in solar physics are tied to the chromosphere: \\begin{compactitem} \\item what are the processes that transport energy and mass in the outer solar atmosphere? \\item how do these processes accelerate the solar wind, and how do they heat the chromosphere and the corona? \\item how are prominences formed, how do they evolve, and how do they relate to CMEs? \\item how is energy stored, and how is it released during flares and CMEs? and \\item what is the nature of the changes in the magnetic structure of the outer solar atmosphere and heliosphere that accompany the 11-year solar cycle? \\end{compactitem} The key to all these questions lies within the magnetic field, and finding the answers depends critically on measuring the chromospheric vector magnetic field and related dynamic processes over the full solar disk. \\figurethree{IBIS_20100804_143736_halpha_core}{IBIS_20100804_143736_halpha_blue}{AIA171_20100804_143736_ibis_aligned}{fig:ibisaia}{IBIS images of \\CaII\\ $854.2~\\nm$ line core (left) and $40$-$60~\\mbox{km}/\\mbox{s}$ in the blue wing (center) in comparison with \\FeIX\\ $17.1~\\nm$ coronal emission observed by SDO Atmospheric Imaging Assembly (AIA) (right). The circular field of view has a diameter of approximately $80\\arcsec$. Note the strong visual correspondence of the chromospheric filter images and the coronal emission. The dark absorbing features in the center panel are known to be associated with a class of spicules that are thought to play a key role in the mass and energy balance of the corona and solar wind.} Magnetism in the solar chromosphere is structured on all spatial scales. At the very finest resolution, we find extremely dynamic features such as spicules, mottles, and fibrils, that can barely be resolved with the best telescopes in the world \\cite{2007PASJ...59S.655D,2009ApJ...705..272R}. An example of chromospheric fine-structure is shown in \\autoref{fig:ibisaia}. New telescopes, such as the 4-m Advanced Technology Solar Telescope (ATST), are being built specifically for high-resolution observations in order to study these kinds of features in detail. High-resolution instruments like the Interferometric Bidimensional Spectrometer (IBIS) \\cite{2006SoPh..236..415C} on the Dunn Solar Telescope (DST) at Sacramento Peak and the CRisp Imaging SpectroPolarimeter (CRISP) \\cite{2006A&A...447.1111S} at the 1-m Swedish Solar Telescope on La Palma are already driving breakthroughs in our understanding of the complex environment of the chromosphere. Those instruments however do not capture the larger scales. EUV imagery from space missions like the Transition Region And Coronal Explorer (TRACE) \\cite{1999SoPh..187..229H}, the Solar and Heliospheric Observatory (SOHO), and now the Solar Dynamics Observatory (SDO) show clearly the complex interconnectivity of the magnetic field over large distances between active regions and the quiet-Sun network. It is clear from those observations that synoptic full-disk measurements of chromospheric magnetic field are needed to investigate the response of the atmosphere at the onset of space weather events like solar flares and CMEs. Magnetic fields have been studied in the solar photosphere for many years by exploiting the Zeeman effect. A degeneracy in atomic energy levels is removed in the presence of a magnetic field, causing a separation in the spectrum. The Zeeman triplet consists of two shifted $\\sigma$ components and one unshifted $\\pi$ component, and the separation of these three components is a measure of the magnetic field strength. In sunspots and active regions, the magnetic field is strong and the Zeeman splitting is comparable to, or even greater than, the line width for many well-known lines such as the \\FeI\\ lines around $630~\\nm$. For weaker fields, such as in the internetwork quiet Sun, the Zeeman splitting is small and the effect is only recognizable as a broadening of the line. Fortunately, the Zeeman components are characterized by distinct polarization properties, and so the weaker magnetic fields can still be investigated by the use of polarimetric instruments. While the first ``magnetographs'' measured only circular polarization, yielding only a diagnostic of line-of-sight component of the magnetic field, modern instruments are routinely used to infer the vector magnetic field in the photosphere through full-Stokes polarimetry, i.e., measurement of both circular and linear polarization properties across the spectral line. In order to derive the magnetic field from the spectro-polarimetric data recorded by the instrument, it is processed with ``inversion'' codes that attempt to fit the line profile simultaneously in the four Stokes parameters with an atmosphere described by a dozen or more magnetohydrodynamic parameters. Both the Hinode and the SDO missions have magnetograph instruments on board that employ this method for routine measurements of photospheric magnetic field. The Spectro-Polarimeter (SP)\\cite{2008SoPh..249..167T} instrument on Hinode is a traditional spectrograph polarimeter that routinely produces measurements with unprecedented spatial resolution and polarimetric accuracy using the \\FeI\\ $630.15$ and $630.25~\\nm$ line pair. The Helioseismic and Magnetic Imager (HMI) on SDO employs full-disk imaging spectro-polarimetry in the \\FeI\\ $617.3~\\nm$ line to produce field measurements at a regular cadence of $45~\\s$. Despite the urgent need for chromospheric vector magnetic field measurements, neither of these instruments nor the upcoming Interface Region Imaging Spectrograph (IRIS) have the capability to make those measurements. The magnetic field in the chromosphere is generally much weaker than the photospheric field. As a result, the two polarized lobes of the Zeeman effect largely cancel. In addition, lines that sample the chromosphere are never formed in local thermal equilibrium (LTE), and the departure from LTE typically includes atomic polarization. A full treatment of the magnetic effects on line polarization is required, including level-crossing interferences and the Hanle effect. While the Zeeman effect is predominantly sensitive to longitudinal magnetic field, the Hanle effect is sensitive to transverse magnetic field. Inversion codes that exploit the complementarity of the Zeeman and Hanle diagnostics, such as the HAZEL code developed at the Instituto de Astrof\\'\\i{}sca de Canarias in Spain are becoming available now. Magnetic field has been successfully measured in regions where the plasma is relatively cool, such as prominences and filaments, or where the field is relatively strong, such as active regions \\cite{2003ApJ...598L..67C,2009A&A...501.1113K}. The Vector Spectromagnetograph (VSM) on the ground-based Synoptic Optical Long-term Investigations of the Sun (SOLIS) \\cite{2003SPIE.4853..194K} telescope makes regular observations in the chromospheric \\CaII\\ $854.2~\\nm$ and \\HeI\\ $1083.0~\\nm$ lines. However, the VSM only measures intensity in both lines and circular polarization in \\CaII\\ $854.2~\\nm$, yielding just a line-of-sight diagnostic. Furthermore, since it is a slit-scanning spectrograph, the cadence of the observations is too slow to study the dynamics of the chromosphere and transient events like flares. Summarizing, there is a considerable need for regular observations of the chromospheric vector magnetic field. Such measurements complement the capabilities of the Hinode, SDO, and IRIS missions and improve our ability to interpret the observations from those missions. Hence, the heliophysics community eagerly awaits the development of an instrument that will provide the community with synoptic global observations of the chromospheric vector field. ", "conclusions": "" }, "1207/1207.5582_arXiv.txt": { "abstract": "In this paper, we study a cosmological model in general relativity within the framework of spatially flat Friedmann-Robertson-Walker space-time filled with ordinary matter (baryonic), radiation, dark matter and dark energy, where the latter two components are described by Chevallier-Polarski-Linder equation of state parameters. We utilize the observational data sets from SNLS3, BAO and {\\it Planck}+WMAP9+WiggleZ measurements of matter power spectrum to constrain the model parameters. We find that the current observational data offer tight constraints on the equation of state parameter of dark matter. We consider the perturbations and study the behavior of dark matter by observing its effects on CMB and matter power spectra. We find that the current observational data favor the cold dark matter scenario with the cosmological constant type dark energy at the present epoch. ", "introduction": "It is not a matter of debate now whether the Universe is accelerating at the present epoch since it is strongly supported by various astronomical probes of complementary nature such as type Ia supernovae data (SN Ia)\\cite{1,2}, galaxy redshift surveys \\cite{3}, cosmic microwave background radiation (CMBR) data \\cite{4,5} and large scale structure \\cite{6}. Observations also suggest that there had been a transition of the Universe from the earlier deceleration phase to the recent acceleration phase \\cite{7}. We do not have a fundamental understanding of the root cause of the accelerating expansion of the Universe. We label our ignorance with the term ``Dark Energy\" (DE), which is assumed to permeate all of space and increase the rate of expansion of the Universe \\cite{8}. On the other hand, the inclusion of DE into the prevailing theory of cosmology has been enormously successful in resolving numerous puzzles that plagued this field for many years. For example, with prior cosmological models, the Universe appeared to be younger than its oldest stars. When DE is included in the model, the problem goes away. The most recent CMB observations indicate that DE accounts for around three fourth of the total mass energy of the universe \\cite{9,Planck13}. However, the nature of DE is still unknown and various cosmological probes on theoretical and experimental fronts are in progress to resolve this problem. The simplest candidate for the DE is the cosmological constant ($\\Lambda$) or vacuum energy since it fits the observational data well. During the cosmological evolution, the cosmological constant has the constant energy density and pressure with the equation of state (EoS) $w_{de}=p_{de}/\\rho_{de}=-1$. However, one has the reason to dislike the cosmological constant since it suffers from the theoretical problems such as the ``fine-tuning\" and ``cosmic coincidence\" puzzles \\cite{10}. Consequently, the dynamic DE models have been studied frequently in the literature. For instance, the Chevallier-Polarski-Linder (CPL) parametrization of the EoS parameter of DE, which was first introduced in \\cite{Chevallier01}, has been frequently constrained with observational data in order to study the nature of dynamic DE (see \\cite{Planck13} for recent constraints from Planck). The $\\Lambda$CDM (cosmological constant + cold dark matter) model, which is the standard model in modern cosmology, has been remarkably successful to describe the Universe on large scales. However, it faces persistent challenges from observations on small scales that probe the innermost regions of dark matter halos and the properties of the Milky Way's dwarf galaxy satellites. See \\cite{vega11,weinberg13} for reviews on the recent observational and theoretical status of these ``small scale controversies\". In this regard, the warm dark matter (WDM) is a plausible dark matter paradigm, which seems to solve many of small scale discrepancies while being indistinguishable from CDM on larger scales. In particle physics, the keV scale sterile neutrinos are believed to account for WDM. On the other hand, the fluid perspective of WDM has been investigated in many studies. In \\cite{muller05}, the bounds on EoS of DM were investigated using CMB, SN Ia and large scale structure data in the cases of no entropy production and vanishing adiabatic sound speed. In \\cite{faber06}, a simple method was suggested for measuring the EoS parameter of DM that combines kinematic and gravitational lensing data to test the widely adopted assumption of pressureless DM. Following the method, the authors in \\cite{serra11} found that the value of the EoS parameter of DM is consistent with pressureless DM within the errors. The authors of \\cite{avelino12,cruz13} investigated the ``warmness\" of the DM fluid constraining cosmological models with constant EoS parameters of DM and DE by considering non-interacting and interacting scenarios of DM and DE. The authors in \\cite{wei13} considered various cosmological models consisting of only DM and DE components by assuming constant and variable EoS parameters of the two components. They found observational constraints on these models using SN Ia, CMB and BAO data, and concluded that WDM models are not favored over the $\\Lambda$CDM model. The authors in \\cite{Calabrese09} investigated the bounds on EoS parameter of DM using WMAP 5 year data and CMB + SDSS + SNLS data in a cosmological model based on spatially flat Friedmann-Robertson-Walker space-time filled with ordinary matter (baryonic) and radiation, DE component acting as a cosmological constant and DM component with constant EoS parameter $w_{dm}$. The Friedmann equation in this model reads as \\begin{equation}\\label{eq1} H=H_{0}\\sqrt{\\Omega_{r}a^{-4}+\\Omega_{b}a^{-3}+\\Omega_{dm}a^{-3(1+w_{dm})}+ \\Omega_{\\Lambda}} \\end{equation} where $a=1/(1+z)$ is scale factor in terms of the redshift $z$; $H_{0}$ is Hubble constant and $\\Omega_i=8\\pi G \\rho_i/(3H_0^2)$ is density parameter for the $i$th component. The authors in \\cite{Calabrese09} stressed that in model \\eqref{eq1}, the background evolution is completely determined by the EoS of DM, and investigated the properties of DM by studying the behavior on its perturbations. In a recent paper \\cite{Lixin13}, the model \\eqref{eq1} is constrained with the currently available observational data. In this work, the tighter constraints are obtained on the EoS parameter of DM due to high quality of observational data. It may be noted that in the studies \\cite{Calabrese09,Lixin13}, the DM is characterized by a constant EoS parameter $w_{dm}$ and the DE candidate is the cosmological constant with constant EoS parameter $w_{de}=-1$. However, the choice of constant EoS parameter for DM is too restrictive \\cite{wei13}. Similarly, candidature of cosmological constant for DE is not satisfactory as discussed earlier. Therefore, in the present work, we consider the naturally motivated CPL parametrizations for the EoS parameters of DM and DE \\cite{wei13}, respectively, given by \\begin{equation}\\label{eq2} w_{de}=w_{dm0}+w_{dma}(1-a), \\end{equation} \\begin{equation}\\label{eq3} w_{de}=w_{de0}+w_{dea}(1-a), \\end{equation} where $w_{dm0}$, $w_{dma}$, $w_{de0}$ and $w_{dea}$ are constants. With these CPL forms of EoS of DM and DE, the Friedmann equation can be written as \\begin{equation}\\label{eq4} H=H_{0}\\sqrt{\\Omega_{r}a^{-4}+\\Omega_{b}a^{-3}+\\Omega_{dm}f(a)+ \\Omega_{de}g(a)} \\end{equation} where \\[f(a)=e^{-3w_{dma}(1-a)}a^{-3(1+w_{dm0}+w_{dma})},\\] \\[g(a)=e^{-3w_{dea}(1-a)}a^{-3(1+w_{de0}+w_{dea})}.\\] It is easy to see that the model \\eqref{eq1} is retrieved from the model \\eqref{eq4} in the particular case $w_{dma}=0$ and $w_{dea}=0$. In the present study, we consider the generalized model \\eqref{eq4} and study the constraints on variable EoS parameters of DM and DE by using the currently available observational data from SNLS3, BAO and {\\it Planck}+WMAP9+WiggleZ measurements of matter power spectrum. One may observe the strong degeneracy in the background evolution of the DM and DE components. We shall deal with this issue later. Next, we assume that DM component interacts with other components only gravitationally. Since DM is believed to be responsible for the gravitational instability and structure formation in the Universe, we consider the perturbations and study the behavior of DM by observing its effects on CMB and matter power spectra. Thus, the main objectives of this study include (i) constraining the CPL EoS parameters of DM and DE with the latest observational data (ii) testing the warmness of DM (iii) testing the behavior of DM by observing its effects on CMB and matter power spectra through perturbations. ", "conclusions": "\\label{sec:dis} Using the observational constraint results displayed in Table \\ref{tab:results}, we now analyze the effects of model parameters $w_{dm0}$ and $w_{dma}$ on the CMB and matter power spectra. We fix the other relevant model parameters to their mean values as given in Table \\ref{tab:results} and vary one of these two model parameters around its mean value. The effects to the CMB TT power spectrum are shown in Figure \\ref{fig:cls}. We see that the positive values of model parameters $w_{dm0}$ and $w_{dma}$ will decrease the equality time of matter and radiation when the other relevant cosmological model parameters are fixed. Consequently, the amplitudes of the peaks of the CMB are depressed and the positions of the peaks are moved to the right side. On the large scale where $l<10$, the curves are increased when the values of $w_{dm0}$ and $w_{dma}$ are positive due to the integrated Sachs-Wolfe effect. \\begin{center} \\begin{figure}[htb!] \\includegraphics[width=9.5cm]{cls.eps} \\caption{The effects of model parameters $w_{dm0}$ and $w_{dma}$ to the CMB TT power spectrum for different values of the model parameters $w_{dm0}$ and $w_{dma}$, where the other relevant model parameters are fixed to their mean values as shown in Table \\ref{tab:results}.}\\label{fig:cls} \\end{figure} \\end{center} Their effects on the matter power spectrum are shown in Figure \\ref{fig:mp}, where the redshift is fixed to $z=0$. The effects of model parameters $w_{dm0}$ and $w_{dma}$ to the matter power spectrum are similar to the ones as analyzed in Ref. \\cite{Lixin13}. The positive values of $w_{dm0}$ and $w_{dma}$ move the matter and radiation equality to earlier times and increase the matter power spectrum. \\begin{center} \\begin{figure}[] \\includegraphics[width=9.5cm]{pks.eps} \\caption{The matter power spectrum at redshift $z=0$ for different values of the model parameters $w_{dm0}$ and $w_{dma}$, where the other relevant model parameters are fixed to their mean values as shown in Table \\ref{tab:results}.}\\label{fig:mp} \\end{figure} \\end{center} From Table \\ref{tab:results}, we see that the 1$\\sigma$ constraints on $w_{dm0}$ are $0.00067_{-0.00067}^{+0.00011}$, which are tighter in comparison to the constraints obtained in the recent studies \\cite{Calabrese09,Lixin13}. Also, the gradual shift of the values of $w_{dm0}$ towards 0 with the advanced and improved observational data shows that the current/future observations are likely to favor the CDM scenario over the WDM. Similarly, the constraints on $w_{de0}$ are $-1.06_{-0.13}^{+0.11}$. Also, the best fit value of $w_{de0}$ is $-1.01$. Thus, current observational data favor cosmological constant type of DE. Finally, we conclude that the $\\Lambda$CDM scenario is favored by the currently available observational data within the framework of the model considered in this study." }, "1207/1207.7314_arXiv.txt": { "abstract": "We show how the 3DVAR data assimilation methodology can be used in the astrophysical context of a two-dimensional convection flow. We study the way this variational approach finds best estimates of the current state of the flow from a weighted average of model states and observations. We use numerical simulations to generate synthetic observations of a vertical two-dimensional slice of the outer part of the solar convection zone for varying noise levels and implement 3DVAR when the covariance matrices are scalar. Our simulation results demonstrate the capability of 3DVAR to produce error estimates of system states between up to tree orders of magnitude below the original noise level present in the observations. This work exemplifies the importance of applying data assimilation techniques in simulations of the stratified convection. ", "introduction": "When using models to describe the temporal evolution of observed complex systems we are confronted with a number of challenges. An immediate difficulty in dealing with this question is that we generally do not know in all detail the current state of the system or the initial condition that is to be used. Lacking such information prevents us from keeping a model-based simulation in step with the behavior of the observed system. \\emph{Data assimilation} techniques offer means to address such challenges for complex systems by keeping a computer simulation (i.e.\\ model) in synchronization with observations of the system it represents. It provides a general framework for simultaneously comparing, combining, and evaluating observations of physical systems and output from computer simulations. The methods used in data assimilation have been developed over several decades, primarily in meteorology and oceanography for the prediction of the future behavior. Data assimilation is used daily in operational weather prediction \\citep{BGK81}, in climate forecast \\citep{PH06} and it was even used to correct the path of the Apollo spacecraft during the first moon landings \\citep{cipra1993engineers}. There is a large and growing body of literature including several monographs \\citep{daley1993,kalnay2003,wunsch2006} and work discussing its theoretical foundations \\citep{L81,lorenc1986,LdT86,ghil1989}. Astrophysical data assimilation has recently been discussed by \\citep{Brun2007}, both in the context of space weather and in solar cycle prediction \\citep{Dikpati2007,choudhuri2007,kitiashvili2008} as well as in the context of dynamo models \\citep{Jouve2011}. Here we focus on the three-dimensional variational (3DVAR) data assimilation technique, also known as sequential approach \\citep{daley1993}, which produces updates of the current state of a model simulation at times when system observations are available. Propagation of model states between times when the system is observed are free simulations of the model initiated at the latest state estimate. An extension of 3DVAR to implicitly incorporate dynamical information is known as four-dimensional variational (4DVAR) data assimilation. 3DVAR dynamically evolves the mean state wheres 4DVAR also evolution other statistical properties of the model dynamics. State estimates produced by 3DVAR are optimal provided that the model is linear and the uncertainties are Gaussian. In other words, 3DVAR states are Best Linear Unbiased Estimates, where \\emph{best} and \\emph{optimal} refer to the lowest possible mean squared error of the estimate \\citep{K60,T97}. For nonlinear models, error statistics may become non-Gaussian even when the initial distribution is Normal and 3DVAR (or 4DVAR) estimates are not longer unbiased. In this case, data assimilation techniques are challenged by the fact that actual applications are typically based on nonlinear processes \\citep{PVT96}. Specifically, that the states exhibited by real systems under observation will diverge from those predicted by a model simulation is clear and this is principally owing to two causes (Palmer \\& Hagedorn 2006): observational error and sensitivity to initial conditions. The first of these is a result of what may be called noise. Since its statistical character may not be known, we may need to make some assumptions about its properties. The second source of error occurs on many complex systems and is referred to as chaotic behavior. This has been known for some time, but only in recent decades serious progress in its understanding has been possible. Sensitivity of the model to initial conditions limits how far into the future predictions can be made \\citep{L93}. Among other challenges present for data assimilation of nonlinear model states like the mismatch between spatial locations of observations and grid positions of the model, and the unpaired model variables to observed quantities make the study towards more effective data assimilation techniques an important and demanding area of research. Despite these challenges and open questions, 3DVAR is widely used in the oceanographic and meteorological communities, and would make a good candidate method to be explored in the context of astrophysical flows. ", "conclusions": "" }, "1207/1207.4910_arXiv.txt": { "abstract": "We give an algorithm for the economical calculation of angles and actions for stars in axisymmetric potentials. We test the algorithm by integrating orbits in a realistic model of the Galactic potential, and find that, even for orbits characteristic of thick-disc stars, the errors in the actions are typically smaller than 2 percent. We describe a scheme for obtaining actions by interpolation on tabulated values that significantly accelerates the process of calculating observables quantities, such as density and velocity moments, from a distribution function. ", "introduction": "When electronic computers first became widely available, it was discovered that orbits in typical axisymmetric galactic potentials usually admit three isolating integrals of motion \\citep{HenonH,Ollongren}. Consequently, by Jeans' theorem, the distribution functions (\\df s) of equilibrium axisymmetric galaxies should be functions of three integrals of motion. Unfortunately, analytic forms of all three integrals are known only for exceptional potentials, so the few three-dimensional galaxy models in the literature that have a known \\df\\ \\citep[e.g.][]{Rowley} employ only the classical energy and angular-momentum integrals $E$ and $L_z$, and therefore lack generality. The action integrals $J_r$, $J_z$ and $L_z$ are particularly useful constants of motion \\citep[e.g.][]{Canary}, and we have previously argued the merits of models in which the distribution function is an analytic function of $J_r$, $J_z$ and $L_z$. To take advantage of these models one should be able to evaluate economically the actions of a star from its conventional phase-space coordinates $(\\vx,\\vv)$. To date we have used two techniques for evaluating actions: (i) torus construction \\citep{KaasalainenB,BinneyM} and (ii) the adiabatic approximation \\citep{B10,BinneyM,SchoenrichB12}. Torus construction is a general and rigorous technique and for some applications it is the technique of choice \\citep[e.g.][]{McMillanB12}. For other applications it is inconvenient because it delivers $(\\vx,\\vv)$ as functions of the actions and angles, rather than the actions and angles as functions of $(\\vx,\\vv)$. The adiabatic approximation delivers actions and angles as functions of $(\\vx,\\vv)$ but it is reasonably accurate only for stars that stay close to the Galaxy's mid-plane. Here we introduce a different approximate way to obtain actions, which, though still approximate, is more accurate than the adiabatic approximation and is valid for stars that move far from the mid plane. ", "conclusions": "We have shown that values of actions and angles accurate to a couple of percent can be obtained for orbits in a realistic axisymmetric model of the Galactic potential by treating the potential as if it were a St\\\"ackel potential. For orbits typical of observed stars belonging to either the thin or thick discs the error in $J_z$ is always less than $\\sim4\\%$ of the average action and is usually significantly smaller. The errors in $J_r$ are always less than 6\\% and usually less than 2\\% of the average action. Even in the era of Gaia it is unlikely that the errors in the measured phase-space coordinates of any star will be small enough that the inaccuracies inherent in our algorithm will dominate the final uncertainties in derived angles and actions. The errors in actions obtained from the adiabatic approximation are larger by a factor $\\sim4$ for thin-disc stars and significantly larger still for thick-disc stars. A possibility that we have not pursued, but which might be important if one needs to model an entire galaxy rather than the extended solar neighbourhood, is to make the inter-focal semi-distance $\\Delta$ a function of $L_z$ and $E$ -- by integrating a few orbits at wide-ranging values of $L_z$ and $E$ it should be possible to choose a suitable functional form for $\\Delta(L_z,E)$. Each action evaluation requires a one-dimensional integral and with the existing code takes $\\sim100\\,\\mu{\\rm s}$ on a laptop. Each angle evaluation takes about twice as long because it requires of order two one-dimensional integrals. Since evaluation of the observables that follow from a \\df\\ requires a great many evaluations of the actions, it is cost-effective to tabulate $(J_r,J_z)$ as functions of the classical integrals $(L_z,E,I_3)$ and we have described an effective scheme for doing this. In a companion paper we illustrate what can be achieved using this scheme by fitting \\df s to observational data for our Galaxy." }, "1207/1207.3192_arXiv.txt": { "abstract": "Magnetic dissipation is frequently invoked as a way of powering the observed emission of relativistic flows in Gamma Ray Bursts and Active Galactic Nuclei. Pulsar Wind Nebulae provide closer to home cosmic laboratories which can be used to test the hypothesis. To this end, we reanalyze the observational data on the spindown power of the Crab pulsar, energetics of the Crab nebula, and its magnetic field. We show that unless the magnetic inclination angle of the Crab pulsar is very close to 90 degrees the overall magnetization of the striped wind after total dissipation of its stripes is significantly higher than that deduced in the Kennel-Coroniti model and recent axisymmetric simulations of Pulsar Wind Nebulae. On the other hand, higher wind magnetization is in conflict with the observed low magnetic field of the Crab nebula, unless it is subject to efficient dissipation inside the nebula as well. For the likely inclination angle of 45 degrees the data require magnetic dissipation on the timescale about 80 years, which is short compared to the life-time of the nebula but long compared to the time scale of Crab's gamma-ray flares. ", "introduction": "\\label{sec:intr} Magnetic fields are often invoked in models of the relativistic jet production by central engines of Active Galactic Nuclei (AGN) and Gamma Ray Bursts (GRB). In these theories the jets are Poynting-dominated at the origin, with the magnetization parameter $\\sigma=B^2/4\\pi\\rho c^2\\gg 1$. This is different from the earlier essentially hydrodynamic, low $\\sigma$, models of relativistic jets in one important aspect. Even strong, high Mach number shocks, in high $\\sigma$ plasma are weakly dissipative compared to their low $\\sigma$ counterparts \\citep[e.g.][]{KC84a,K12}. Moreover, PIC simulations show that the acceleration of nonthermal particles may also be problematic at such shocks \\citep{SS09,SS11a}. This suggests that either the Poynting flux is first converted into the bulk motion kinetic energy via ideal MHD mechanism \\citep[e.g.][]{VK04,KVKB09,L10}, which is then dissipated at shocks, or the magnetic energy is converted directly into the energy of emitting particles via magnetic dissipation, which accompanies magnetic reconnection events \\citep[e.g.][]{DS02,LB03,ZY11,G11,MU12}. In fact, the magnetic dissipation can facilitate bulk acceleration of jets as well \\citep[e.g.][]{DS02}. While AGN and GRBs are very distant sources, which makes their observational studies rather difficult, there exist objects much ``closer to home'' which share similar properties, the Pulsar Wind Nebulae (PWN). They are powered by relativistic winds from neutron stars and these winds are also expected to be Poynting-dominated at their base \\citep[see ][and references therein]{Ar12}. In particular, the Crab nebula is one of the brightest sources of nonthermal emission in the sky throughout the whole observational range of photon energies. Its large angular size (of seven arc minutes), ensures that its spatial structure is well resolved and its relatively small linear size (of several light years) allows direct observations of not only its small-scale structural variability but also its overall dynamics. Because the Crab nebula is such an easy object to observe it has been studied with the level of detail which may never be reached in observations of AGN and GRB jets, and it is rightly considered as a testbed of relativistic astrophysics. The early attempts to built a theoretical model of the Crab nebula using the ideal relativistic MHD approximation resulted in a paradoxical conclusion that the pulsar wind has to have $\\sigma\\sim10^{-3}$ near its termination shock \\citep{RG74,KC84a,EC87,BL92}. A slightly higher magnetization, $\\sigma\\sim10^{-2}$, was later suggested by axisymmetric numerical simulations \\citep{KL03,LDZ04,B05}, although no proper study of this issue has been carried out. The key property of these analytical and numerical solutions is their purely toroidal magnetic field. The strong hoop stress of such field creates excessive axial compression of the nebula in solutions with higher $\\sigma$ and pushes the termination shock too close to the pulsar in the Kennel-Coroniti model, in conflict with the observations. On the other hand, the ideal relativistic MHD acceleration of uncollimated wind-like flows is known to be very inefficient, leaving such flows Poynting-dominated on the astrophysically relevant scales \\citep[e.g.][]{L11,K11}. This striking conflict is known as the $\\sigma$-problem. Attempts have been made to see if $\\sigma$ can be reduced via magnetic dissipation in the so-call striped zone of pulsar winds, where the magnetic field changes it polarity on the length scale $\\lambda_p =cP$, where $P$ is the pulsar period \\citep{C90,LK01}. The dissipation is accompanied by the wind acceleration via conversion of the thermal energy into the bulk kinetic energy of the flow during its adiabatic expansion. Unfortunately, for the wind of the Crab pulsar the dissipation length scale significantly exceeds the radius of the wind termination shock, thus making this mechanism inefficient \\citep{LK01}. \\citet{L03b} has demonstrated that the energy associated with the alternating component of magnetic field of the striped wind can be rapidly dissipated at the termination shock itself, where the characteristic Larmor radius of shock-heated plasma exceeds the wavelength of magnetic stripes. His solution of the shock equations, which accounts for the ``erasing'' of stripes, shows that the post-shock flow is the same as it would be if the dissipation had already been fully completed in the wind. \\citet{SS11b} have used 3D PIC simulations to study the magnetic dissipation and particle acceleration at the termination shock of the striped wind numerically and concluded that efficient magnetic dissipation occurs even when the Larmor radius remains below the stripes wavelength, via rapid development of the tearing mode instability and magnetic reconnection in the post-shock flow. One way or another, this dissipation occurs only in the striped zone and only the alternating component of magnetic field dissipates. Outside of the striped zone, around the poles, the pulsar wind $\\sigma$ remains unaffected by this dissipation and hence very high. As the result, the overall magnetization of plasma injected into the nebula can be much higher than that of the Kennel-Coroniti model, unless the striped zone spreads over almost the entire wind \\citep{C90}. \\citet{L03a} argued that in the polar zone the wind $\\sigma$ can be reduced via the flow acceleration accompanying dissipation of fast magnetosonic waves emitted by the pulsar into the polar zone. However, it seems unlikely that the energy flux associated with these waves can dominate the wind energetics in the polar zone. At least, the 3D numerical simulations of pulsar winds (A.Spitkovsky, private communication) show that their contribution is rather small. Thus, we do not expect $\\sigma$ of the polar zone to be below unity. An alternative solution to the $\\sigma$ problem has been proposed by \\citet{B98}, who argued that the axial compression of the nebula can be reduced via the current-driven kink instability, resulting in more or less uniform total pressure distribution inside the nebula. This would make the overall structure and dynamics of the nebula similar to those in the models with particle-dominated ultra-relativistic pulsar wind. The recent computer simulations of the non-linear development of the kink instability of relativistic z-pinch configurations support this conclusion \\citep{M09,M11}. In this scenario, PWN are supplied with highly magnetized plasma, making magnetic dissipation a potentially important process in their evolution and emission. In this paper, we test whether the magnetic dissipation inside PWN is consistent with the observations of the Crab nebula and its pulsar. The main idea is very simple. First, the timing observations of the Crab pulsar allow us to estimate how much energy has being pumped into the nebula. Second, using the stripe wind model we can calculate how much of this energy is supplied in the magnetic form. Third, a simple dynamical model of the nebula expansion can be used to predict how much magnetic energy is retained by the nebula after adiabatic losses. Finally, the observations of the Crab nebula tell us how much magnetic energy is actually in there and whether the magnetic dissipation is actually required to make the ends meet. ", "conclusions": "\\label{sec:discussion} \\subsection{Magnetic dissipation and the fine structure of the Crab Nebula} The possibility of efficient magnetic dissipation in the Crab nebula raises the question about its observational signatures. What kind of structures if any should we expect to see and where? Unfortunately, our current understanding of magnetic reconnection is not that advanced to make any firm predictions. The most explored and firmly associated with magnetic reconnection phenomena in astrophysics are solar flares. They involve significant restructuring of magnetic fields anchored to the solar surface and distorted by motions in the Sun. It is not clear if such grand events may occur under the conditions of PWN. Smaller scale ``nanoflares'' could be responsible for heating of solar corona and the appearance of bright coronal loops \\citep{P72}, but even this issue has not been settled yet. The so-called ``reconnection exhausts jets'' have been detected in the solar wind via in situ measurements using spacecrafts \\citep[see][and references wherein]{Gos11}. These observations show no evidence of non-thermal particle acceleration or electron heating and do not allow to say how far a spacecraft is from a reconnection site or even if reconnection is still ongoing at the time of observation. Other examples include Earth's magnetosphere and laboratory experiments but these seem to be too specific. For the purpose of identifying the locations of magnetic dissipation in PWN, a non-thermal particle acceleration is its most promising and also likely product. This process could brighten up the interfaces between regions with different orientation of magnetic field. The polarimetric observations of the Crab Nebula by \\citet{BK91} show that in its central region the degree of polarization is about 5 times below that for the synchrotron emission in uniform magnetic field. They explain this result by the presence along the line of sight of several cells with randomly oriented uniform magnetic field. Boundaries of such cells could be the cites of ongoing magnetic reconnection and may appear as arcs or filaments of enhanced nonthermal emission. In fact, the Crab nebula has the most spectacular network of optical filaments but they are made of the line-emitting thermal plasma of the supernova ejecta ionized by the synchrotron radiation of the PWN. In the optical continuum, only the bright cores of these filaments can be seen and only as absorbing features \\citep[e.g.][]{FB90,S98}. The radio emission of the Crab nebula has synchrotron origin and given the results of the optical continuum observations one would not expect to find the filaments in radio images of the nebula. To the contrary, the high resolution VLA radio images do show filamentary structure which is as spectacular as that of the line emission maps \\citep[e.g.][]{BHFB04}. Moreover, the radio filaments seem to coincide with the line-emitting optical filaments. This was noticed already in the early lower resolution study of the nebula with the Cambridge One-Mile radio telescope \\citep{W72}. A localized enhancement of synchrotron emissivity does not have to be related to particle acceleration and may simply reflect a local enhancement of magnetic field. Such enhancement could well arise during interaction between the high speed flow of relativistic plasma inside the PWN with the filaments via the so-called ``magnetic draping'' effect \\citep[e.g.][]{Lt06}. However, in this case one would generally expect the optical non-thermal emissivity to increase as well. Since this is not what is observed, other factors should come into play. The origin of radio emitting electrons (and positrons) of the Crab nebula is a long standing mystery. The most natural assumption is that they come with the wind from the Crab pulsar just like the higher energy electrons, but their number seems to be too high to be accommodated in the current models of pair production in pulsar magnetospheres \\citep{Ar12}. The radio observations of the inner Crab nebula could have settled this issue should they revealed the same features associated with the outflow from the termination shock as in optics and X-rays, or otherwise. Unfortunately, the emerging picture is rather ambiguous. Although radio wisps are observed, they do not coincide with the optical ones and are noticeably slower \\citep{BHFB04}. There is no obvious radio counterpart to the optical and X-ray jet either. Given the strong anisotropy of the pulsar wind, this may indicate that radio electrons and positrons come from different parts of the termination shock. On the other hand, the radio wisps could just be some kind of ripples driven by the unsteady outflow from the termination shock through the PWN. Indeed, the MHD simulations show strong convective motion inside the nebula which brings plasma from outer parts of the nebula quite close to termination shock, where it is pushed out again by the outflow from the termination shock \\citep{C09}. The quantity of radio electrons may be large compared to what is expected in the theory of pulsar magnetospheres, but it is tiny compared to what is available in the line-emitting filaments. It is conceivable that a small fraction of the filament plasma becomes lose and mixes with the relativistic plasma of PWN. In there, its electrons can be accelerated to relativistic energies and produce the observed synchrotron radio emission. \\cite{KC84b} estimate the energy contained in the radio electrons to be of the order of few $\\times10^{48}$erg, which is a sizable fraction of the total internal energy of the nebula. Thus, in situ acceleration of radio electrons requires a substantial source of energy. This could be the energy of the magnetic field injected into the nebula by the pulsar wind, which can indeed be substantial, as we argued in Sec.\\ref{sec:SW}. The magnetic dissipation can be enhanced near the line-emitting filaments when magnetic field lines of different orientation wrap around the same filament. { The magnetic reconnection could also facilitate escape of electrons (and ions) from the filaments into Crab's PWN, otherwise suppressed by the low diffusivity across magnetic field lines.} The other possibility is the second-order Fermi acceleration by the hydromagnetic turbulence driven by various instabilities \\citep{B11}. { The existence of two synchrotron components of different origin is supported by the combined radio and mm-wavelength observations \\citep{BNC02}. The data suggest low energy cutoff around 100~GHz in the emission of the electrons supplied by the termination shock, thus supporting a different origin of radio emitting electrons. However, the relatively smooth matching of radio and optical/infrared components of the integral spectrum may be problematic for this model, particularly when seen in a number of PWN \\citep{B11}.} \\citet{VRV92} state that the VLA images also show filaments which do not have line emitting counterparts. The spectral data do not show any noticeable variation of the radio spectral index across the filaments \\citep{B97}. Since the synchrotron life-time of radio emitting electrons significantly exceeds the age of the nebula, this is not very surprising. Some of the filaments are also seen in the X-ray band \\citep{STF06}. Since the life-time of X-ray emitting electrons is quite short, one would expect to see hardening of the X-ray spectrum of the filaments compared to the diffuse background if these filaments were indeed the cites of particle acceleration. However, the observations give no evidence of such hardening and the photon index is very soft, $\\alpha\\sim3-4$ \\citep{STF06}, creating a problem for any model where these features act as acceleration cites of electrons emitting in X-rays \\citep{STF06}. The continuum optical images of the Crab nebula reveal fine fibrous structure somewhat reminiscent of solar coronal loops \\citep{FB90,H95}. However, it is not clear if this is a product of ``nanoflares'' or simply reflects inhomogeneous structure of magnetic field. \\subsection{Implications for numerical simulations of PWN} The 2D RMHD numerical simulations of the Crab nebula \\citep{KL03,LDZ04,B05,C09} have been very successful in reproducing many key properties of the nebula, such as its jet-torus, the brightness asymmetry, wisps, and even the bright ``inner knot'' \\citep{H95}. In agreement with the observations, the proper motion of jets and wisps produced in the simulations is relatively low, $v=0.2-0.7c$, as expected downstream of an almost purely hydrodynamical shock wave. This success leaves little doubt that the numerical models capture the physics of the nebula quite well. However, the overall low wind magnetization utilized in these models, $<\\!\\sigma\\!>\\simeq 10^{-2}$, is in conflict with what we would expect in the striped wind model without imposing very large magnetic inclination angle of the pulsar. This choice of $\\sigma$ has been influenced by the very low value required in the Kennel-Coroniti model in order to have a termination shock in their 1D solution. However, the flow dynamics of the 2D numerical solutions is already very different, as it involves large scale circulation and mixing. Although \\citet{KL04} did find that, in qualitative agreement with predictions of the Kennel-Coroniti model, the size of the termination shock decreased with $<\\!\\sigma\\!>$, no attempts have been made to study models with $\\sigma\\gg 10^{-2}$. As the result, one cannot claim yet that 2D numerical simulations rule them out. As the shock size is determined by the balance between the wind ram pressure and the total pressure in the nebula, this tendency can be explained by the stronger axial compression of the nebula by the magnetic hoop stress in models with higher $\\sigma$. However, this compression is certainly excessive in 2D models, being enforced by the condition of axial symmetry which does not allow development of the kink instability \\citep{B98}. The 3D numerical study of z-pinch configurations by \\citet{M09,M11} confirms this expectation. Thus, the ultimate answer to the question whether $<\\!\\sigma\\!>\\,\\,\\gg 10^{-2}$ is allowed by the RMHD model will only be found in future 3D simulations of PWN. If at high latitudes the pulsar wind is free from stripes and has high $\\sigma$ then downstream of the termination shock one would expect a very fast flow, with the Lorentz factor $\\gamma\\sim\\sigma^{1/2}$ in the case of perpendicular shock \\citep{KC84a} and even higher in the case of oblique shock \\citep{KL11}. Downstream of a perpendicular shock the flow is subsonic (or sub-fast-magnetosonic to be more precise) and can smoothly decelerate down to $\\gamma\\simeq 1$ inside the nebula. Downstream of an oblique shock it may remain supersonic and a secondary shock will have to appear somewhere on its way. So far the observations of the Crab nebula show no evidence of such a secondary shock or such a fast flow. This may well be related to the low dissipation efficiency of shocks in highly magnetized plasma \\citep[e.g.][]{KC84a,K12}, as well as the inability of such shocks to accelerate non-thermal particles \\citep{SS09,SS11a}. Further investigation is required to clarify this issue. \\subsection{Magnetic dissipation and Crab's gamma-ray flares} The recently discovered strong flares of gamma-ray emission from the Crab nebula at the energies $\\sim 1\\,$GeV with duration about few days \\citep{T11,A11} could be very important for understanding the physics of highly magnetized relativistic plasma. \\citet{KL11} argued that the gamma-rays of these energies could originate from the most compact known bright feature of the Crab nebula, the so-called ``inner knot'', which they explain as a Doppler-boosted emission from the termination shock. However, their model predicts synchronous variability of the knot emission in gamma-rays and optics, which does not seem to be the case \\citep{Ar12}. The only other promising alternative seem to be explosive magnetic reconnection. However, the properties of these flares suggest that they may not be representative of the energetically dominant magnetic dissipation process in the nebula. First, the dissipation time scale given by Eq.\\ref{t-d} is at least three orders of magnitude longer than the typical flare duration. It is possible that the tearing instability produces much smaller structures inside the large scale current sheets, however in this case one would expect a whole spectrum of time scales to be present. Second, the statistical model of flares by \\citep{CL12} gives the total energy release rate which is three orders of magnitude below the spindown power of the Crab pulsar and hence significantly lower than the magnetic dissipation rate given by Eq.\\ref{t-d}. Third, the current reconnection models of these flares involve strong magnetic fields, of order $1000\\,\\mu$G \\citep{UCB11,CUB12}, and/or large bulk Lorentz factors $\\Gamma\\gtrsim\\mbox{few}$ \\citep{KL11,CL12}. Such conditions are not typical for the Crab nebula. Finally, so far the flares have not been identified with any particular kind of events seen at other energies. Given the required conditions for the flares, their most likely location is the polar region near the termination shock, where the freshly supplied plasma can have very high magnetization and streams with ultra-relativistic speeds\\footnote{ A similar conclusion was reached recently by Y.Lyubarsky at a conference presentation (http://www.iasf-roma.inaf.it/Flaring\\_Crab/index.html).}. Large Lorentz factors could also be produced during fast reconnection events inside high-$\\tilde\\sigma$ plasma, which again points out towards the inner polar region of the Crab nebula, where the observations reveal the Crab jet. High magnetization also implies Alfv\\'en speed approaching the speed of light and hence the fastest possible magnetic reconnection speed. \\citet{CUB12} also point out that magnetic field in this region can be much stronger than on average due to the strong axial compression associated with the z-pinch. The region at the base of the Crab jet, the so-called ``anvil'', is in fact the most active region in the nebula \\citep{H02}." }, "1207/1207.2899_arXiv.txt": { "abstract": "One believes there is huge amount of Dark Matter particles in our Galaxy which manifest themselves only gravitationally. There is a big challenge to prove their existence in a laboratory experiment. To this end it is not sufficient to fight only for the best exclusion curve, one has to see an annual recoil spectrum modulation --- the only available positive direct dark matter detection signature. % A necessity to measure the recoil spectra % is stressed. \\vskip 0.3cm \\noindent {PACS:} 95.30.-k, 95.35.+d, 14.80.Ly, 12.60.Jv ", "introduction": " ", "conclusions": "" }, "1207/1207.3910_arXiv.txt": { "abstract": "{} { We present a study of the envelope morphology of the carbon Mira R~For with VLTI/MIDI. This object is one of the few asymptotic giant branch (AGB) stars that underwent a dust-obscuration event. The cause of such events is still a matter of discussion. Several symmetric and asymmetric scenarios have been suggested in the literature.} {Mid-infrared interferometric observations were obtained separated by two years. The observations probe different depths of the atmosphere and cover different pulsation phases. The visibilities and the differential phases were interpreted using GEM-FIND, a tool for fitting spectrally dispersed interferometric observations with the help of wavelength-dependent geometric models.} {We report the detection of an asymmetric structure revealed through the MIDI differential phase. This asymmetry is observed at the same baseline and position angle two years later. % The observations are best simulated with a model that includes a uniform-disc plus a Gaussian envelope plus a point-source. The geometric model can reproduce both the visibilities and the differential phase signatures.} {Our MIDI data favour explanations of the R For obscuration event that are based on an asymmetric geometry. We clearly detect a photocentre shift between the star and the strongly resolved dust component. This might be caused by a dust clump or a substellar companion. However, the available observations do not allow us to distinguish between the two options. The finding has strong implications for future studies of the geometry of the envelope of AGB stars: if this is a binary, are all AGB stars that show an obscuration event binaries as well? Or are we looking at asymmetric mass-loss processes (i.e. dusty clumps) in the inner part of a carbon-rich Mira?} ", "introduction": "The C-star \\object{R For} is classified as a Mira variable in the General Catalogue of Variable Stars \\citep{sam09}. Objects of this variability class show photometric variations with amplitudes larger than 2.5~mag in the $V$ band, and with long periods ($>100$ days). The light curves of Mira variables are usually described as regular with sinusoidal behaviour. Nevertheless, long-term observations of these objects show erratic drops in their brightness \\citep[][and references therein]{bar96}. The variability of R~For was discussed for the first time by \\cite{fea84}, and later by \\cite{leb88}. Both works concluded that the light curve showed a deep minimum in 1983. This decrease was attributed to an increased dust obscuration, the reason for which is (still) unknown. Four scenarios were envisaged by \\cite{fea84} to explain the obscuration event, the first two scenarios were symmetric, and the latter two involved an asymmetric geometry of the envelope. The scenarios involve: \\begin{enumerate} \\item the ejection of a spherical shell by the star; \\item enhanced dust-condensation in an already existing spherical shell; \\item an asymmetric dust ejection in a preferential direction (i.e. a disc, or a more complex structure); \\item an asymmetric dust ejection in a random direction. \\end{enumerate} \\cite{leb88} and \\cite{win94} supported the first two scenarios, which reproduced the photometrical data with dynamic models in spherical symmetry. Advanced 2D and 3D model atmospheres by \\cite{woi05} and \\cite{fre08} predicted the formation of concentrated dust clouds in the inner part of the atmospheres, and a spherical distribution on a large scale. Such a morphology might explain the dust obscuration events. However, up to now those models were not compared with observations. On the other hand, these obscuration events have similarities with those observed for R~Coronae Borealis stars {\\citep[R~CrB;][]{cla12}}. In the latter case, the third and the fourth scenario would be more likely explanations \\citep{bri11, jef12}. \\cite{lea07} used VLTI/MIDI to study the circumstellar environment of the R~CrB star RY~Sgr. This star showed a drop in the $J$ and $H$ light curve similar to the one observed for R~For. The observations of RY~Sgr showed that the star is surrounded by a single dusty cloud located at $\\sim$100 stellar radii ($\\sim$30 AU). \\cite{whi97} suggested a binary-related effect as an alternative explanation. Exceptionally deep minima in the light curves of single O-rich Miras are not observed, but they are fairly common for symbiotic Miras. This letter presents high angular resolution observations in the mid-infrared that allow us to probe the dust-forming region of the nearby AGB star R~For. Observations close in time, with different position angles and baseline lengths, the differential phase information, and the comparison with geometric models can help to distinguish possible deviations from a spherical structure. This allows one to confirm or reject the different scenarios suggested in the past. Observations and data reduction procedures are presented in Sect.~\\ref{obs.sect}. The study of the morphology by means of geometric models is described in Sect.~\\ref{morphstu.sect}, and is discussed in Sect.~\\ref{disc.sect}. \\onlfig{1}{ \\begin{figure}[!thtp] \\includegraphics[angle=0.,width=.5\\textwidth]{u-v.ps} \\caption{\\label{uvcov.fig}\\small{$uv$ coverage dispersed in wavelength of the observations acquired during the programs {\\sc080.D-0231} (grey-scale), and {\\sc084.D-0361} (red-scale). The colour code ranges from 8-12.5~$\\mu$m, with light colours indicating shorter wavelengths. The sample is restricted to the observations used for this work (marked with (b) and (d) in Table~\\ref{obs.tab}).}} \\end{figure} } ", "conclusions": "\\label{disc.sect} As already mentioned, a differential phase signature, similar to that observed for R~For, was detected for two J-type carbon stars. IRAS~1800-3213, presented by \\cite{der07}, shows a signature in the differential phase that is in the same wavelength region (i.e. $8-9\\,\\mu$m) as the one observed for R~For. \\cite{der07} mentioned that long baselines probe the part in the atmosphere where dust emission is strongly (spatially) resolved. The authors speculated that the differential phase signature is caused by an offset of the photocentre between the strongly resolved dust region and the unresolved stellar photosphere. This might be explained by a circumbinary disc. The second object, BM~Gem, presented by \\cite{ohn08}, shows a non-zero differential phase signature between $9-11\\,\\mu$m. The authors interpreted this as due to the presence of a circum-companion disc. The detected asymmetry for R~For excludes the symmetric scenarios 1. and 2. hypothesised by \\cite{fea84}. The plausible explanations are (i) a feature intrinsic to the star, such as a dust blob that moves through the circumstellar environment of R\\,For; or (ii) an unresolved companion that is embedded in the circumstellar envelope. The dust blob might have been ejected in a preferential direction or randomly, as stated in scenarios 3 and 4. The expansion velocity of the outer envelope of R~For is 16.9~km~s$^{-1}$ \\citep{gro99}. Using the determined separation of 42 mas, and assuming that the blob moved with a constant velocity of 5\\,km\\,s$^{-1}$ over the last 30 years, \\citep[a reasonable value for the inner envelope, predicted by dynamic model computations;][]{now10}, it is worthwhile to mention that the event of 1983 would correspond to the moment at which such a structure was formed near the surface of the star. Recent results from investigations of R~CrB stars \\citep{jef12} indicate the presence of dust clouds with different grain-sizes. Nevertheless, \\cite{bri11} did not detect any differential phase signature for these objects, arguing that the signatures of dust clumps are not strong enough to induce significant spectral signatures like the one we observed. This might not be the case for AGB stars because R~CrB stars are much hotter and H-deficient (i.e. no C$_2$H$_2$), therefore carbon nucleation proceeds differently \\citep{goe92}. IRC+10216 could be a very interesting test-case to check if the clumps can produce a phase signature in carbon stars. Indeed, several dust clumps were detected in the atmosphere of this star \\citep{han98, wei98, wei02, lea06} at spatial scales comparable to the one we obtained with the GEM-FIND modelling ($\\sim 30$~AU). The binary option is quite unrealistic for a stellar companion for the following reasons. Archive GALEX observations do not show any evidence of UV excess, which automatically excludes the presence of a warm white dwarf. By translating the flux ratio (see Fig.~\\ref{wave-flux.fig}) into magnitudes, this would imply that the companion has a brightness of $\\sim2.2$\\,mag at 12~$\\mu$m. A faint giant star is very unlikely because such an object would be resolved. A main sequence star also turns out to be too massive and thus would disrupt (because of mass accretion) the spherical symmetry detected at low spatial frequencies for the envelope. We cannot exclude that the companion is a substellar object orbiting the envelope. There are suggestions in literature that such a situation might be common among AGB stars \\citep{sos07}. In this case, only low-mass accretion is expected \\citep{lec07}, leaving a signature only at the high spatial frequencies probed by long baselines. The differential phase signature was detected with the same baselines at the same position angles after two years. Assuming that the flux ratio of the point source over the central star is constant between 2007 and 2009 and that the position of the companion as well as the dust blob would change from 2007 to 2009, the following scenarios are possible: (i) for $1~M_{\\sun}$ central star and a Jupiter mass companion (10$^{-3}\\,M_\\odot$) the orbital period of the companion would be 185\\,yr and would result in a change of the position angle of $\\pm$4\\,$^\\circ$; (ii) assuming a dust blob that is moving outwards at a uniform velocity of 5\\,km\\,s$^{-1}$ would result in a change of the separation of 2.4\\,mas. In both cases the asymmetry would be still observable within the two years with the chosen MIDI configuration. Therefore, the available observations do not allow us to reject any of these two asymmetric scenarios. This letter presents one interpretation for the detected asymmetry in the dust envelope of R~For. More complex morphologies cannot be excluded a priori because of the limited coverage of the available observations. Interferometric imaging with the second-generation instrument VLTI/MATISSE \\citep{lop06} and a detailed polarimetric investigation are needed to clarify the nature of the detected structure. \\begin{figure}[thp] \\includegraphics*[angle=0,width=0.45\\textwidth,bb=62 532 467 795]{wave-fluxratio.ps} \\includegraphics*[angle=0,width=0.45\\textwidth,bb=62 397 467 795]{mol-contributions-Claudia.ps} \\caption{\\label{wave-flux.fig}\\small{Upper panel: Wavelength-dispersed flux ratio values of the best-fitting models. Other panels: Contributions by different sources in typical spectra of C-rich AGB stars for the MIDI wavelength range. The synthetic spectra are plotted with continuum-normalised fluxes $F/F_{\\rm c}$, for detailed descriptions see \\cite{ari09} or Nowotny et al. (2010). The middle panel illustrates the individual molecular absorption due to C$_2$H$_2$ and HCN with the help of a hydrostatic model atmosphere, for details we refer to Fig.\\,4 of \\cite{ari09}. The lower panel shows the characteristic emission due to circumstellar SiC dust on the basis of a dynamic model atmosphere \\citep[e.g.][]{hoe03,now10}, computed following the approach of \\cite{sac11} to artificially impose SiC grains. This dust species was detected in the ISO spectrum of R~For \\citep{cle03}. }} \\end{figure}" }, "1207/1207.5562.txt": { "abstract": "\\noindent The Q/U Imaging ExperimenT (QUIET) is designed to measure polarization in the Cosmic Microwave Background, targeting the imprint of inflationary gravitational waves at large angular scales ( $\\sim$ 1$^\\circ$). Between 2008 October and 2010 December, two independent receiver arrays were deployed sequentially on a 1.4\\,m side-fed Dragonian telescope. The polarimeters which form the focal planes use a highly compact design based on High Electron Mobility Transistors (HEMTs) that provides simultaneous measurements of the Stokes parameters Q, U, and I in a single module. The 17-element Q-band polarimeter array, with a central frequency of 43.1\\,GHz, has the best sensitivity (69\\,$\\mu\\mathrm{Ks}^{1/2}$) and the lowest instrumental systematic errors ever achieved in this band, contributing to the tensor-to-scalar ratio at $r < 0.1$. The 84-element W-band polarimeter array has a sensitivity of 87\\,$\\mu\\mathrm{Ks}^{1/2}$ at a central frequency of 94.5\\,GHz. It has the lowest systematic errors to date, contributing at $r < 0.01$ \\citep{quiet:2012} The two arrays together cover multipoles in the range $\\ell \\approx 25-975$. These are the largest HEMT-based arrays deployed to date. This article describes the design, calibration, performance of, and sources of systematic error for the instrument. ", "introduction": " ", "conclusions": "" }, "1207/1207.3675.txt": { "abstract": "{We present the \\emph{Planck Sky Model} (PSM), a parametric model for the generation of all-sky, few arcminute resolution maps of sky emission at submillimetre to centimetre wavelengths, in both intensity and polarisation. Several options are implemented to model the cosmic microwave background, Galactic diffuse emission (synchrotron, free-free, thermal and spinning dust, CO lines), Galactic \\ion{H}{ii} regions, extragalactic radio sources, dusty galaxies, and thermal and kinetic Sunyaev-Zeldovich signals from clusters of galaxies. Each component is simulated by means of educated interpolations/extrapolations of data sets available at the time of the launch of the \\planck\\ mission, complemented by state-of-the-art models of the emission. Distinctive features of the simulations are: spatially varying spectral properties of synchrotron and dust; different spectral parameters for each point source; modeling of the clustering properties of extragalactic sources and of the power spectrum of fluctuations in the cosmic infrared background. The PSM enables the production of random realizations of the sky emission, constrained to match observational data within their uncertainties, and is implemented in a software package that is regularly updated with incoming information from observations. The model is expected to serve as a useful tool for optimizing planned microwave and sub-millimetre surveys and to test data processing and analysis pipelines. It is, in particular, used for the development and validation of data analysis pipelines within the \\planck\\ collaboration. A version of the software that can be used for simulating the observations for a variety of experiments is made available on a dedicated website. } ", "introduction": "The cosmic microwave background (CMB), relic radiation from the hot Big Bang, carries an image of the Universe at an age of about 380,000 years. CMB anisotropies reflect the inhomogeneities in the early Universe, and are thus an observable of great importance for constraining the parameters of the Big Bang model. For this reason, a large number of experiments dedicated to CMB anisotropy detection and characterisation have observed the sky at wavelengths ranging from the sub-millimetre to a few centimetres (frequencies from a few to a few hundred GHz). Stimulated by CMB science, astronomical observations in this wavelength range have enriched many fields of scientific investigation, ranging from the understanding of emission from the Galactic interstellar medium (ISM), to the study of a large population of extragalactic objects, including radio galaxies, infrared galaxies, and clusters of galaxies. The \\emph{WMAP} satellite, launched by NASA in June 2001, has surveyed the sky in five frequency bands, ranging from 22 to 94\\GHz\\ \\citep{2003ApJS..148....1B}. The data from this mission have provided an all-sky map of CMB temperature, and an unprecedented data set for the study of sky emission at millimetre wavelengths -- both in total intensity and in polarisation \\citep{2011ApJS..192...14J}. %% The \\planck\\ mission\\footnote{\\planck\\ \\emph{(http://www.esa.int/Planck)} is a project of the European Space Agency -- ESA -- with instruments provided by two scientific Consortia funded by ESA member states (in particular the lead countries: France and Italy) with contributions from NASA (USA), and telescope reflectors provided in a collaboration between ESA and a scientific Consortium led and funded by Denmark.}, launched by ESA in May 2009, has started to provide even better observations of the sky in nine frequency bands, ranging from 30 to 850\\GHz\\ \\citep{2010A&A...520A...1T,2011A&A...536A...1P}. These observations, which will be released to the scientific community early 2013, are expected to become the new reference in this frequency range. It has long been recognised that planning future observations and developing methods to analyse CMB data sets both require realistic models and simulations of the sky emission and of its observations \\citep{1998astro.ph.12237G,1999NewA....4..443B,2002A&A...387...82G,2008MNRAS.388..247D,2007ApJ...664..149S,2010ApJ...709..920S}. For the preparation of the analysis of CMB data sets such as those of the \\planck\\ mission, it has proven useful to put together a model of sky emission, the \\emph{Planck Sky Model} (PSM), which can be used to predict or simulate sky emission for various assumptions and various parameter sets. The pre-launch version of this model, based on publicly available data sets existing before \\planck\\ observations (see Section \\ref{sec:psmdata}), is used in particular as a simulation tool in the context of \\planck. An update of the model is planned on the basis of observations obtained by the \\planck\\ mission. This paper summarises the outcome of this preparatory modelling and simulation effort. %% Software and simulations are made available to the scientific community on a dedicated website\\footnote{http://www.apc.univ-paris7.fr/{$\\sim$}delabrou/PSM/psm.html}. The current version of the \\emph{Planck Sky Model} (PSM), described in the present paper, is v1.8. We plan regularly to update models and software tools on the basis of upcoming observations, and when additional sophistication becomes useful for upcoming experiments and scientific investigations. The PSM developers welcome suggestions for improving the models or implementing alternate options. Although nothing prevents running the PSM to simulate sky emission at any frequency, our sky model is actually valid, at present, only for frequencies ranging from about 3\\GHz\\ to about 3\\THz, i.e., wavelengths ranging between 100\\micron\\ and 10\\cm. At lower frequencies, Faraday rotation is important and our polarised maps are not representative of the sky emission, as this effect is presently neglected. Intensity maps, however, should be reasonably accurate down to about 500\\MHz. At frequencies higher than 3\\THz, the complexity of the dust emission is not properly taken into account by the present model, for which the representation of dust emission is based on spectral fits valid below 3\\THz. The paper is organised as follows: in Section~2 we review the basic features of the model as well as the architecture of the package; in Sections~3, 4 and 5 we describe the simulation and implementation of CMB anisotropies, diffuse Galactic and compact object emissions; and finally, in Section~6, we summarise the status of our sky model at the time of this publication and draw some conclusions. %%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%% ", "conclusions": "We have developed a complete, flexible model of multicomponent sky emission, which can be used to predict or simulate astrophysical and cosmological signals at frequencies ranging from about 3\\GHz\\ to 3\\THz. The model, which actually implements several options for each of the components, has been developed for testing component separation methods in preparation for the analysis of \\planck\\ data, but has been used also for simulating data in the context of analyses of \\wmap\\ observations, or for planning future missions such as COrE \\citep{2011arXiv1102.2181T}. Given the high sensitivity of these instruments, the quality of the reconstructed CMB temperature map depends strongly on our ability to remove the contamination by astrophysical foregrounds. It is therefore necessary to build simulations as close as possible to the complexity of the real sky over a large frequency range. We present here what has been achieved through a series of improvements accomplished over several years and the contributions of experts in different fields: CMB, Galactic emission and compact sources, extragalactic radio and far-IR sources, Sunyaev-Zeldovich signals from clusters of galaxies and the intergalactic medium. Note however that the current version of the model cannot be expected to match perfectly data sets that have not been used in the model, in particular those of the \\planck\\ mission and of other upcoming observations. Updates will be proposed as such data sets become available. The maps of each astrophysical component have been repeatedly updated by requiring a close match with the constantly increasing amount of observational data. Fig.~\\ref{fig:sed} shows the spectrum of root mean square (r.m.s.) fluctuations in simulated maps at \\wmap\\ frequencies for $|b|>20^{\\circ}$ and $10^{\\circ}$ resolution. The various diffuse components have different spectral characteristics while the total signal is a good match to the \\wmap\\ 7-year data; the r.m.s. values agree to within better than 5 percent. The sky emission maps observed with WMAP, smoothed to $1^\\circ$ resolution, are compared in Fig. \\ref{fig:compar_psm_wmap} to PSM predicted emission maps at the same frequencies and resolution. The agreement is excellent over most of the sky, except in the galactic ridge and around a few compact regions of strong emission, for which the model does not exactly reproduce the observations. Discrepancies are due to a mixture of uncertainties in the exact resolution of the model (and, in particular, a lack of resolution of the map of synchrotron spectral index), errors in extrapolation of the dust emission, residuals of galactic emission in the predicted CMB map, and insufficient number of parameters used in the current PSM in regions of strong emission (in particular when emission comes from several distinct regions along the line of sight). The model will be refined in future version, in particular using \\planck\\ observations. %%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%% \\begin{figure} \\begin{center} \\includegraphics[width=\\columnwidth]{figures/sed_pol0.png} \\end{center} \\caption{Temperature r.m.s. fluctuations at {\\it WMAP} frequencies for $|b|>20^{\\circ}$ at a resolution of $10^{\\circ}$. The symbols represent the fluctuations in the various diffuse components of the sky model, the total simulated fluctuations, and \\wmap\\ 7-year maps. The total signal is a good match to the \\wmap\\ data.} \\label{fig:sed} \\end{figure} %%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%% %%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%% \\begin{figure*} \\begin{centering} \\includegraphics[width=0.65\\columnwidth,angle=0]{figures/wmap_K.png}\\hspace{0.1in} \\includegraphics[width=0.65\\columnwidth,angle=0]{figures/psm_K.png}\\hspace{0.1in} \\includegraphics[width=0.65\\columnwidth,angle=0]{figures/diff_K.png} \\par\\end{centering} \\begin{centering} \\includegraphics[width=0.65\\columnwidth,angle=0]{figures/wmap_Ka.png}\\hspace{0.1in} \\includegraphics[width=0.65\\columnwidth,angle=0]{figures/psm_Ka.png}\\hspace{0.1in} \\includegraphics[width=0.65\\columnwidth,angle=0]{figures/diff_Ka.png} \\par\\end{centering} \\begin{centering} \\includegraphics[width=0.65\\columnwidth,angle=0]{figures/wmap_Q.png}\\hspace{0.1in} \\includegraphics[width=0.65\\columnwidth,angle=0]{figures/psm_Q.png}\\hspace{0.1in} \\includegraphics[width=0.65\\columnwidth,angle=0]{figures/diff_Q.png} \\par\\end{centering} \\begin{centering} \\includegraphics[width=0.65\\columnwidth,angle=0]{figures/wmap_V.png}\\hspace{0.1in} \\includegraphics[width=0.65\\columnwidth,angle=0]{figures/psm_V.png}\\hspace{0.1in} \\includegraphics[width=0.65\\columnwidth,angle=0]{figures/diff_V.png} \\par\\end{centering} \\begin{centering} \\includegraphics[width=0.65\\columnwidth,angle=0]{figures/wmap_W.png}\\hspace{0.1in} \\includegraphics[width=0.65\\columnwidth,angle=0]{figures/psm_W.png}\\hspace{0.1in} \\includegraphics[width=0.65\\columnwidth,angle=0]{figures/diff_W.png} \\par\\end{centering} \\caption{Comparison of sky emission as observed by WMAP (7-year data) and as predicted by the PSM in the same frequency bands, at a resolution of $1^\\circ$. For each frequency channel, the color scale is the same for WMAP (left column) and PSM prediction (middle column). An histogram equalised color scale is used for the K, Ka, and Q channels, and a linear scale for the V and W channels. Maps are saturated to highlight common features away from the galactic plane. Maps of difference between PSM prediction and WMAP observation are displayed in the right column (note that the color scale is different from that used to display the K, Ka and Q maps), highlighting discrepancies in the galactic plane specifically and at the location of a few regions of compact emission. The agreement is excellent over most of the sky away from the galactic ridge and a few compact regions.} \\label{fig:compar_psm_wmap} \\end{figure*} %%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%% The CMB temperature and polarisation maps, currently based on \\wmap\\ observations (maps and best-fit cosmological model), are complemented with simulations to allow for the weak gravitational lensing by gradients in the large-scale gravitational field. Weak gravitational lensing has been dealt with in the Born approximation. Simulations of CMB temperature maps with primordial non-Gaussianity of local type are also accessible via our sky model. The Galactic diffuse emission model includes five components: synchrotron, free-free, spinning dust and thermal dust radiation, and $^{12}$CO (J=1--0), (J=2--1), and (J=3--2) molecular lines at 115.27, 230.54, and 345.80\\GHz, respectively. For each component, alternative models are available but we propose a combination of models that reproduces best the available data. In this default model the synchrotron emission is based on the 408\\MHz\\ all-sky map by \\citet{1982A&AS...47....1H} extrapolated in frequency exploiting the spectral index map by Miville-Desch\\^enes et al. (2008; model 4). The free-free template is obtained from the \\wmap\\ MEM map, complemented with the H$\\alpha$ all-sky template by Finkbeiner (2003) corrected for extinction as in Dickinson et al. (2003). For thermal dust we adopted model 7 of Finkbeiner et al. (1999), which features spectral variations across the sky. For spinning dust we adopted the template produced by Miville-Desch\\^enes et al. (2008) from an analysis of \\wmap\\ data. The CO emission map is based on the $^{12}$CO(1-0) survey by Dame et al. (2001). Standard intensity ratios with the (J=1--0) line have been used for the (J=2--1) and (J=3--2) transitions. Since all these templates have a resolution insufficient for our purposes, the code allows the possibility to add small-scale fluctuations following Miville-Desch\\^enes et al. (2007). In our model as currently implemented, the only diffuse emissions that are polarised (apart from the CMB) are synchrotron and thermal dust. For synchrotron we rely on the \\wmap\\ 23\\GHz\\ polarisation maps, extrapolated in frequency using the same spectral index map used for intensity. Modeling dust polarisation is made difficult because of the paucity of data. In our model the dust $Q$ and $U$ maps are built assuming a constant intrinsic polarisation fraction, geometrical depolarisation and polarisation angle maps constructed using a Galactic magnetic field model for scales larger than $20^\\circ$ and constraints from 23\\GHz\\ \\wmap\\ polarisation data to reproduce the projected structure at smaller scales due to the turbulent magnetic field of the ISM. While generally compact Galactic sources are part of the diffuse emission map and Galactic point sources are treated in the same way as extragalactic point sources (catalogues do not distinguish between Galactic and extragalactic point sources), ultra-compact \\ion{H}{ii} regions have been included in the model as a separate population. They have been selected from the IRAS catalogue and the extrapolation in frequency of their flux densities is made adopting a grey-body spectrum. Radio counterparts have been searched for in the NVSS and GB6 catalogues. The Sunyaev-Zeldovich effect from galaxy clusters is simulated in two ways. One can start from the epoch-dependent cluster mass function, for the preferred choice of cosmological parameters, plus some recipes to model the density and temperature distribution of hot electrons. A fraction of simulated clusters can be replaced by real clusters drawn from the \\rosat\\ all-sky or the SDSS catalogue. Alternatively, we can resort to $N$-body+hydrodynamical simulations that also contain the SZ emission from the web of intergalactic hot gas. Most radio sources in the model are real, taken from the relatively deep all-sky low-frequency catalogues. Each source has its own spectral properties, either determined directly from multi-frequency data or randomly extracted from the observed distributions of spectral indices. Extrapolations to higher frequencies are made using different spectral indices for different frequency intervals, in order to ensure consistency with multi-frequency source counts. This approach turned out to be remarkably successful in predicting the counts that have been later measured by \\planck. The polarisation degree of each source is randomly drawn from the distributions for flat and steep-spectrum sources determined by \\citet{2004A&A...415..549R} at 20\\GHz, while the polarisation angle is drawn randomly from a uniform distribution. Bright far-IR sources include those taken from the IRAS Point Source Catalog, with flux densities extrapolated in frequency using a grey-body spectrum, plus high-$z$ galaxies strongly gravitationally lensed, based on the model by Perrotta et al. (2003). For fainter galaxies, which make up most of the cosmic infrared background, we adopted the model by Granato et al. (2004). Their clustering properties are modelled following Negrello et al. (2004). The implementation of their spatial distribution is made using the algorithm by Gonzalez-Nuevo et al. (2005). The resulting power spectrum of CIB fluctuations turns out to be quite close to that measured by \\planck\\ Collaboration \\citep{2011A&A...536A..18P}. The mean polarisation degree of individual infrared galaxies can be adjusted as an input parameter, the default value being 1\\%. The polarisation angle is randomly chosen from a uniform distribution. On the whole, these simulations should provide a useful tool for several purposes: to test data analysis pipelines for CMB experiments, identify the most convenient areas of the sky for specific purposes (i.e., areas least contaminated by Galactic emissions), identify the optimal set of frequency channels for CMB temperature and polarisation experiments or for other studies (i.e., for studies of the cosmic infrared background), predict the levels of confusion noise, including the effect of clustering for a given frequency and angular resolution, and much more. Access to documented versions of the package and to reference simulations is available from a dedicated website\\footnote{http://www.apc.univ-paris7.fr/{$\\sim$}delabrou/PSM/psm.html}." }, "1207/1207.1853_arXiv.txt": { "abstract": "Motivated by the recent detection of an enhanced clustering signal along the major axis of haloes in N-body simulations, we derive a formula for the anisotropic density distribution around haloes and voids on large scales. Our model, which assumes linear theory and that the formation and orientation of nonlinear structures are strongly correlated with the Lagrangian shear, is in good agreement with measurements. We also show that the measured amplitude is inconsistent with a model in which the alignment is produced by the initial inertia rather than shear tensor. ", "introduction": "The clustering of matter at late times provides important constraints on cosmological models. Our understanding of the signal is best on large scales, where it can be described by perturbation theory well \\citep[see][]{Peebles}. E.g., the Baryonic Acoustic Oscillations in the power spectrum (BAO), which appear as a spike in the two-point correlation function, lie in this regime \\citep[see][and references to it]{Eisenstein2005}. In addition to the simple correlation function, there are other ways to extract information from the matter distribution. In redshift-space distorted measurements, the two-point correlation function is anisotropic (see \\citealt{Kaiser1987} or the recent work of \\citealt{Schlagenhaufer2012}), and this anisotropy can be used to constrain cosmological parameters. However, certain real-space measures of clustering are also expected to be anisotropic. Galaxy clusters are typically triaxial, and this triaxiality has long been known to align with the surrounding large scale structure \\citep[e.g.,][and references therein]{Smargon2012}. On smaller mass scales, galaxy spins are also known to align with the environment \\citep[e.g.,][]{Lee2001,Zhang2009,Jones2010}. Similar correlations have also been seen in simulations of voids \\citep{Platen2008}. Recently, \\cite{Faltenbacher2012} showed that, in their numerical simulations of hierarchical gravitational clustering, the cross-correlation function between haloes and the surrounding mass was anisotropic: this correlation between halo shapes and large scale structure extended even to the large scales relevant to BAO studies, and affected the zero-crossing of the correlation function. This motivates our work, which attempts to model this anisotropy. Our model, which is described in Section~2, is based on the assumption that halo shapes \\citep{Lee2005, Rossi2011} and orientations \\citep{Lee2000,Lee2001} at late times are correlated with the properties of the initial Lagrangian field from which they formed. This is a fundamental ingredient in models where haloes form from a triaxial collapse \\citep{BM1996, SMT2001}. In such models, the Lagrangian deformation or shear tensor plays a key role, as its eigenvalues can be used to distinguish between haloes, filaments, walls and voids. We illustrate our model for the two extreme cases: haloes and voids. A final section discusses potential applications and extensions of our work. Although our analysis is general, when we illustrate our results, we will assume a $\\Lambda$CDM model with $h = 0.73$, $\\Omega_{\\rm cdm}=0.205$, $\\Omega_{bar}=0.045$, $\\Omega_{\\Lambda} = 0.75$, and $\\sigma_8=0.9$. These values allow a direct comparison with the simulations of \\cite{Faltenbacher2012}, which we provide. ", "conclusions": "\\label{discussion} In this paper, we calculated the anisotropy in the linear density field when conditions are placed on the Lagrangian shear field. If the shear field is strongly correlated with the shapes and orientations of nonlinear haloes, then this calculation should be closely related to the anisotropy of the halo-mass cross-correlation function, which is most easily seen when the mass field around haloes is stacked after aligning along the major axis of the halo \\citep[e.g.][]{Faltenbacher2012}. For haloes, our model (equation~\\ref{main}) captures the main features of the cross-correlation function measurement. The signal along the long axis is stronger than perpendicular to it, but produces a less prominent BAO feature (Figure~\\ref{fig:haloes}). We predict a similar effect for voids (Figure \\ref{fig:voids}). Overall, the signal is slightly weaker on scales below $100 \\mpc$ than in simulations. This may be due to inadequacies in our crude model which relates halo formation to the local shear; or nonlinear evolution may have had a small effect (see Section \\ref{conditions}). Formally, the approximations involved can be summed up as \\begin{equation} P\\big( \\delta (\\mathbf{r})|\\mathbf{\\hat{e}}_h \\big) \\approx P_{GS}\\big( \\delta(\\mathbf{r})|\\mathbf{\\hat{e}}_{sh}=\\mathbf{\\hat{e}}_h, C \\big), \\end{equation} where the lhs is the true distribution of $\\delta$ at $\\mathbf{r}$ relative to a halo oriented in the direction of $\\mathbf{\\hat{e}}_h$, while the rhs is the usual Gaussian conditional probability of $\\delta$ with $ C $ denoting the previously discussed conditions on the shear tensor (Section~\\ref{conditions}) and $\\mathbf{\\hat{e}}_{sh}$ the direction of the eigenvector that belongs to the smallest eigenvalue. Improvements can be devised along the following identity: \\begin{eqnarray} P(\\delta(\\mathbf{r})|\\mathbf{\\hat{e}}_h) & = & \\int P\\big( \\delta(\\mathbf{r})|\\{ \\mathbf{\\hat{e}}_{sh},\\xi_{ii} \\} , \\mathbf{\\hat{e}}_{h} \\big) \\nonumber \\\\ & \\times & P\\big( \\{ \\mathbf{\\hat{e}}_{sh},\\xi_{ii} \\} | \\mathbf{\\hat{e}}_{h} \\big)\\ud \\{ \\mathbf{\\hat{e}}_{sh},\\xi_{ii}\\}, \\label{appr} \\end{eqnarray} where $\\{ \\mathbf{\\hat{e}}_{sh},\\xi_{ii}\\}$ and $\\ud\\{ \\mathbf{\\hat{e}}_{sh},\\xi_{ii}\\}$ are a parametrization of the shear and its volume element respectively. As the final orientation of a halo ($\\mathbf{\\hat{e}}_{h}$) results from the nonlinear evolution of the Lagrangian field, it is a function of the local Lagrangian field and its derivatives. In Section~\\ref{DvsI}, we showed that on large scales in the Gaussian limit higher order derivatives had a small effect on $\\big< \\delta(\\mathbf{r})|\\dots \\big>$ compared to the shear. In this limit, the first term behind the integral in equation~(\\ref{appr}) turns into $ P\\big( \\delta(\\mathbf{r})|\\{ \\mathbf{\\hat{e}}_{sh},\\xi_{ii} \\} , \\mathbf{\\hat{e}}_{h} \\big) \\approx P_{GS}\\big( \\delta(\\mathbf{r})|\\{ \\mathbf{\\hat{e}}_{sh},\\xi_{ii} \\}) $. To improve on this approximation, the nonlinear evolution of haloes has to be understood better. The second term in the integral is equally challenging. A practical approach can be taken by fitting a phenomenological formula to measurements in N-body simulations, similarly to \\cite{Lee2001}, who parametrized the angular momentum of a halo as a function of the shear. E.g., a simple model is given by \\begin{eqnarray} P\\big( \\{ \\mathbf{\\hat{e}}_{sh},\\xi_{ii} \\} | \\mathbf{\\hat{e}}_{h} \\big) & \\approx & \\frac{P(\\{ \\xi_{ii} \\} | C )}{2\\pi} \\times \\nonumber \\\\ & & \\bigg( P_{\\parallel }\\delta_{D}(1-\\mathbf{\\hat{e}}_{h}\\cdot \\mathbf{\\hat{e}}_{sh})+(1-P_{\\parallel}) \\bigg),\\label{ali} \\end{eqnarray} which assumes that the eigenvalues of the shear ($\\xi_{ii}$) are independent of the orientation of the halo ($\\mathbf{\\hat{e}}_{h}$) and that the shear is perfectly aligned with haloes in $P_{\\parallel}\\times 100$ per cent of the time, otherwise its orientation is completely random. With this, the spherical part of equation~(\\ref{main}) remains the same, while the anisotropy gets multiplied by $P_{\\parallel}$. Figure~\\ref{fig:haloes} implies a strong correlation, so $P_{\\parallel}$ must be close to $1$. This is a conjecture that can be verified by a direct measurement of the halo-shear alignment. In general, a more complex model of alignment would introduce higher order Legendre polynomials in the expansion of $\\delta(r,\\mu)$. The same argument holds for voids as well. As the ellipticity of voids is less prominent \\citep{Sheth2004}, their orientation can be measured with a lower accuracy. In a model of alignments, this would increase the randomness. E.g. in equation~(\\ref{ali}), $P_{\\parallel}$ would be smaller thus reducing the measured anisotropy. We also argued that the measured amplitude is inconsistent with a model in which the alignment is produced by the initial inertia rather than shear tensor (Section~\\ref{DvsI}). Absent a model for halo or void formation, our equation~(\\ref{main}) may be treated as a one-parameter family which, given the spherically averaged measurement, describes the anisotropy. This parameter depends only on the local shear, and so may be used to constrain models of halo formation and alignment. Further work can be done to incorporate redshift distortions and nonlinearities into the model. Also, tests on simulations are needed in order to identify systematics that can affect the validity of the model. Finally, we are in the process of checking if this sort of measurement can yield useful constraints on modified gravity models." }, "1207/1207.4192_arXiv.txt": { "abstract": "Most planet pairs in the Kepler data that have measured transit time variations (TTV) are near first-order mean-motion resonances. We derive analytical formulae for their TTV signals. We separate planet eccentricity into free and forced parts, where the forced part is purely due to the planets' proximity to resonance. This separation yields simple analytical formulae. The phase of the TTV depends sensitively on the presence of free eccentricity: if the free eccentricity vanishes, the TTV will be in phase with the longitude of conjunctions. This effect is easily detectable in current TTV data. The amplitude of the TTV depends on planet mass and free eccentricity, and it determines planet mass uniquely only when the free eccentricity is sufficiently small. We proceed to analyze the TTV signals of six short period Kepler pairs. We find that three of these pairs (Kepler-18,24,25) have TTV phase consistent with zero. The other three (Kepler-23,28,32) have small TTV phases, but ones that are distinctly non-zero. We deduce that the free eccentricities of the planets are small, $\\lesssim 0.01$, but not always vanishing. Furthermore, as a consequence of this, we deduce that the true masses of the planets are fairly accurately determined by the TTV amplitudes, within a factor $\\lesssim 2$. The smallness of the free eccentricities suggests that the planets have experienced substantial dissipation. This is consistent with the hypothesis that the observed pile-up of Kepler pairs near mean-motion resonances is caused by resonant repulsion. But the fact that some of the planets have non-vanishing free eccentricity suggests that after resonant repulsion occurred there was a subsequent phase in the planets' evolution when their eccentricities were modestly excited, perhaps by interplanetary interactions. ", "introduction": "\\label{sec:intro} The Kepler mission has detected an abundance of low-mass close-in planets \\citep{Batalhaetal12}. Remarkably, hundreds of them are members of planetary systems \\citep{Lissaueretal11,Fabryckyetal12}. These will likely prove to be a Rosetta stone for deciphering the dynamical history of planetary systems. One of the most intriguing Kepler discoveries is that, while the spacing between planets in a system appears to be roughly random, there is a distinct pile up of planet pairs just wide of certain resonances, and a nearly empty gap just narrow of them \\citep{Lissaueretal11,Fabryckyetal12}. In \\citet{lithwua}, we proposed that dissipation is responsible for this asymmetry, via an effect we termed ``resonant repulsion'' \\citep[also see the independent work by][]{baty}. The eccentricity of planets near resonances can be separated into two parts: a part that is forced by the resonance and is determined by the planets' proximity to resonance ({\\it forced eccentricity}), and a part that is unrelated to the resonance ({\\it free eccentricity})\\footnote{Our forced eccentricity is perhaps more accurately called the forced resonant eccentricity to distinguish it from the more commonly used forced {\\it secular} eccentricity \\citep[e.g.,][]{MD00}. But only the resonant contribution plays a role in this paper.}. If there is dissipation, it damps away the planets' free eccentricities, but the forced eccentricities persist as long as the planets remain close to resonance. As the dissipation continually acts on these forced eccentricities, it extracts energy from the planets' orbits, and in doing so pushes apart any planet pair that happens to lie near a resonance \\citep{lwpluto,paprr}. Hence all such planet pairs end up just wide of resonance, naturally explaining the Kepler result. If it is indeed resonant repulsion that is responsible for the pile up -- and to date no other tenable mechanisms have been proposed -- then there are a number of interesting implications. First, it implies that before resonant repulsion occurred the distribution of spacings was nearly uniform, i.e., that planet pairs were placed with little regard for resonances. Second, it implies that most of the Kepler planets suffered a prolonged bout of eccentricity damping. Third, it implies that the free eccentricities of the planets should be zero today, after the prolonged bout of dissipation. This prediction can be tested using the transit time variations (TTV) recorded by Kepler, as we demonstrate in this paper. A transiting planet that has no companion transits at perfectly periodic times. But one that has a companion deviates slightly from its periodic schedule because of the gravitational tugs from its companion. \\cite{Agol} and \\cite{HolmanMurray} proposed using TTV signals to characterize the companions of transiting extrasolar planets, and this technique has proved to be highly successful both for confirming Kepler candidates and for measuring their masses and eccentricities \\citep[e.g.,][]{CochranKepler18,Ford12,SteffenIII,FabryckyIV}. However, all these studies rely on fitting the observed TTV signals to direct N-body simulations \\citep[e.g.][]{verasford}. Such fits are computationally costly. Moreover, N-body simulations do not provide a dynamically transparent interpretation of the system. Here, we focus on near-resonant pairs because the nearly coherent interactions in such pairs induce particularly large TTV signals. Such pairs account for most of the TTV detections to date in the Kepler database. Motivated by our earlier work on resonant repulsion, we separate the eccentricity into free and forced parts. Interestingly, in so doing, the expression for the TTV near first-order resonance becomes particularly simple. This paper is organized as follows. In Section \\ref{sec:planetparams}, we present new analytical formulae for the TTV from two near-resonant planets, and show that the results agree with N-body simulations. In Section \\ref{sec:data}, we apply the TTV formulae to six planet pairs with published TTV data. In Section \\ref{sec:sum}, we discuss our findings and their implications. ", "conclusions": "\\label{sec:sum} \\subsection{Analytical TTV} We have derived simple analytical expressions for the TTV from two planets near a first order mean motion resonance (Eqs. \\ref{eq:tv}--\\ref{eq:zfreeloc}). These show that the amplitude and phase of the TTV depend on both planet mass and free eccentricity. There is an inherent degeneracy between mass and free eccentricity which in general prevents either from being determined independently of the other. This degeneracy, however, may be (partially) broken under certain circumstances, based on probability arguments. There is a special moment in time when the longitude of conjunction points along the line of sight. If the phase of TTV is zero relative to this time, then it is likely that the free eccentricity in the system is zero. Moreover, if the free eccentricity is zero, the TTV amplitudes can be used to uniquely determine planet masses. When applying this technique to six published systems, we find that three of them are consistent with zero TTV phase, while the other three deviate by less than a radian. This clustering around zero phase can be most naturally explained if all systems have free eccentricity of order a percent or less---which is comparable to the typical distance to resonance in these systems. Furthermore, because the free eccentricities are small, the nominal masses determined by TTV are likely close to the true masses, within a factor $\\lesssim 2$. Without the analytical TTV expressions and relying only on N-body simulations, it would be hard to reach these conclusions. \\subsection{Implications of Small Free Eccentricities} The very small free eccentricities suggest that these planets have experienced damping, as suggested also by the resonant repulsion theory (\\citealp{lithwua}, see also \\citealp{baty}). In that work, we found that if Kepler planets have experienced substantial energy dissipation, but substantially less angular momentum damping, the two planets will be repelled from each other. This would explain the observed pile-up of planets just wide of resonances. A corollary of this theory is that low mass planets in the Kepler sample should have little if any free eccentricity. Although this appears to be confirmed by three of the six systems we analyze, we are puzzled by the small but finite free eccentricities in the other three systems. Assuming that resonant repulsion indeed occurred, it would require that the planets' eccentricities were subsequently excited, perhaps by interplanetary interactions. Such a scenario would argue against tides as the mechanism of dissipation causing resonant repulsion, because tides would have to act over very long times to be effective. Instead, dissipation by a disk of gas or planetesimals are more plausible damping agents. TTV data have also been reported for Kepler 9b/c and Kepler 30b/c \\citep{HolmanKepler9,FabryckyIV}. These pairs contain one or two giant planets. We find preliminary evidence that these systems have large TTV phases. This, if true, will indicate the presence of large free eccentricities in systems of giant planets, in contrast to the lower mass planets discussed here. We are currently analyzing Kepler public lightcurve data to distill more TTV systems. One interesting issue to pursue is whether planet pairs at much shorter or much longer orbital periods have the same characteristics as those analyzed here. Many planet pairs are also near resonance with a third planet. This may bring further complications to our TTV analysis but is not considered here. \\subsection{ Values of Planet Masses} All six systems we analyze have orbital periods between 6 and 15 days, and planet radii ranging from $2$ to $7 R_\\oplus$. We confirm the mass estimates of \\citet{CochranKepler18} for Kepler 18c/d, and the mass estimates of \\citet{FabryckyIV} for Kepler 25b/c, under the assumption that these systems have $|\\E|\\lesssim|\\Delta|$, consistent with their small phase. Densities of these planets range between $0.3$ and $2\\g/\\cm^3$. For Kepler 28b/c, 32b/c, we obtain nominal mass upper limits that lead to density upper limits of $\\sim 3\\g/\\cm^3$. We also argue that these upper limits are likely not too different from the real masses. For Kepler 24b/c and 23b/c, the nominal densities are $\\sim 10\\g/\\cm^3$." }, "1207/1207.3856_arXiv.txt": { "abstract": "We use large cosmological Smoothed-Particle-Hydrodynamics simulations to study the formation and evolution of sub-millimetre galaxies (SMGs). In our previous work, we studied the statistical properties of ultra-violet selected star-forming galaxies at high redshifts. We populate the same cosmological simulations with SMGs by calculating the reprocess of stellar light by dust grains into far-infrared to millimetre wavebands in a self-consistent manner. We generate light-cone outputs to compare directly the statistical properties of the simulated SMGs with available observations. Our model reproduces the submm source number counts and the clustering amplitude. We show that bright SMGs with flux $S > 1$ mJy reside in halos with mass of $\\sim 10^{13} M_{\\odot}$ and have stellar masses greater than $10^{11}~ \\rm M_{\\odot}$. The angular cross-correlation between the SMGs and Lyman-$\\alpha$ emitters is significantly weaker than that between the SMGs and Lyman-break galaxies. The cross-correlation is also weaker than the auto-correlation of the SMGs. The redshift distribution of the SMGs shows a broad peak at $z \\sim 2$, where Bright SMGs contribute significantly to the global cosmic star formation rate density. Our model predicts that there are hundreds of SMGs with $S > 0.1$ mJy at $z > 5$ per 1 square degree field. Such SMGs can be detected by ALMA. ", "introduction": "An array of recent millimetre and sub-millimetre (submm) observations revealed the physical properties of sub-millimetre bright galaxies (SMGs) with luminosities $10^{12} \\sim 10^{13} {\\rm L_{\\odot}}$ \\citep{Swinbank2004, Chapman2005, Capak2008, Coppin2009, Daddi2009, Tamura2009, Knudsen2010, Riechers2010, Hatsukade2011}. The power sources of such luminous SMGs are thought to be very active star formation with $100-1000 \\rm M_{\\odot} yr^{-1}$ and/or active galactic nuclei (AGNs). Ultra-violet photons from massive stars or AGNs are converted to photons in submm to infrared bands by dust grains via their thermal emission. \\citet{Alexander2003} find that, in the Chandra Deep Field North survey, X-rays are detected from more than one-third of bright SMGs, which indicates a significant contribution from AGNs. SMGs provide important information on the star formation rate, the stellar initial mass function, and the chemical evolution of the galaxies at high redshifts. Because SMGs are highly obscured by dust, they often appear dark in UV and optical wavebands. Thus measuring their redshifts by conventional techniques is very difficult. Moreover, identifying the optical counterpart is not easy because the spatial resolution of the currently operating submm telescopes is much worse than that of large optical telescopes. So far, most of the SMGs with measured redshifts are at $z < 3$ \\citep{Swinbank2004, Chapman2005}, but a few SMGs have been found at higher redshifts \\citep{Capak2008, Coppin2009, Daddi2009, Knudsen2010}. SMGs generally show strong clustering \\citep{Webb2003, Blain2004, Scott2006, Weis2009, Cooray2010, Hickox2012}. \\citet{Maddox2010} measured the angular correlation of the $350~\\rm \\mu m$ and $500~\\rm \\mu m$ selected SMGs in the Herschel-ATLAS survey. The SMGs are strongly clustered while the $250~\\rm \\mu m$ selected sample showed weak or no clustering signals. It is thought that SMGs are very massive systems with large gas reservoir ($\\sim 10^{11}~\\rm M_{\\odot}$) and with large stellar masses ($\\sim 10^{11}~\\rm M_{\\odot}$) \\citep{Greve2005, Tacconi2006, Wardlow2011}. Also, SMGs are likely to be ancestors of massive elliptical galaxies in the local universe \\citep{Lilly1996, Smail2004}. \\citet{Chapman2005} report that the redshift distribution of SMGs shows a peak at $z \\sim 2$, similarly to that of AGNs, although the size of the sample is small. Several SMGs have counterparts of star-forming galaxies such as Lyman break galaxies (LBGs) and Lyman $\\alpha$ emitters (LAEs) \\citep{Geach2005, Geach2007, Beelen2008, Daddi2009}. However, many of high-redshift galaxies are not seen as SMGs probably because their submm fluxes are below typical observational flux limits ($> 1 ~\\rm mJy$). Because observations so far detect very bright SMGs, it remains unclear if SMGs are generally associated with LBGs/LAEs. \\citet{Matsuda2007} and \\citet{Tamura2010} failed to detect submm sources at the location of LAEs in a protocluster region even though the target galaxies are very bright in Ly$\\alpha$. \\citet{Dayal2010} use cosmological simulations and found that the submm fluxes of high-$z$ LAEs would be less than $0.1~ \\rm mJy$. They argue, however, that many of such galaxies can be detected by ALMA. \\citet{Yajima2011} study the evolution of submm flux of a simulated galaxy by using detailed three-dimensional radiative transfer calculations. They show that the simulated galaxy in a very early LAE phase can be detected by ALMA. Observing galaxies in the early formation phase is important to understand the early chemical evolution. One of the most important observational quantities is the source number count of SMGs\\citep{Greve2004, Laurent2005, Perera2008, Austermann2009, Austermann2010, Scott2010, Hatsukade2011}. Unfortunately, there remain substantial uncertainties in theoretical models. \\citet{Baugh2005} calculate the source number count of SMGs using a semi-analytic galaxy formation model. Interestingly, in order to reproduce the observed source number count, they make a crucial assumption that the stellar initial mass function (IMF) is 'top-heavy' when starburst occurs (see also \\citet{Lacey2010}). More recently, \\citet{Fontanot2007} reproduced the source number count of SMGs without assuming top-heavy IMF. They argue that the major difference between \\citet{Baugh2005} and their study may be in the gas cooling model. In \\citet{Fontanot2007}, the hot gas in a galaxy is assumed to cool more efficiently than in \\citet{Baugh2005}. However, because of the efficient cooling, the model of \\citet{Fontanot2007} overproduces massive galaxies at lower redshifts. It appears that some efficient star-formation recipe is needed in the semi-analytic models, in order to reproduce the source number count of SMGs. It is interesting to see if the same is true for cosmological hydrodynamic simulations of galaxy formation. \\citet{Dave2010} study the physical properties of SMGs using cosmological simulations. They claim that smooth infalling gas or/and gas-rich satellites are important to produce large star formation rates. Although their finding seems reasonable, they employ a simple model that galaxies with very large star-formation rates are identified as SMGs. They do not calculate dust absorption and the resulting submm flux. It is important and timely to study in detail the statistical properties of SMGs over a wide range of redshift using cosmological hydrodynamic simulations by considering dust absorption and re-emission consistently. In this paper, we study the statistical properties of SMGs such as the stellar mass function, the source number count and the angular correlation function. To this end, we perform large cosmological hydrodynamic simulations based on the standard $\\Lambda$CDM cosmology. Our simulations follow star formation, supernova feedback, and metal enrichment self-consistently. For the galaxies identified in our cosmological simulations, we calculate the spectral evolution and dust extinction to estimate the FIR luminosity. In our earlier work \\citep{Shimizu2011}, we used the same set of simulations to study the statistical properties of LAEs at $z=3.1$. Our model successfully reproduced all the available observational data at $z=3.1$ such as the Lyman-$\\alpha$ luminosity function, the angular correlation functions, and the Lyman-$\\alpha$ equivalent width distribution. Our aim in the present paper is to build a consistent theoretical model of galaxy formation that reproduces the observed properties of both ultra-violet selected galaxies and SMGs. Throughout the present paper, we adopt the $\\Lambda$CDM cosmology with the matter density $\\Omega_{\\rm{M}} = 0.27$, the cosmological constant $\\Omega_{\\Lambda} = 0.73$, the Hubble constant $h = 0.7$ in units of $H_0 = 100 {\\rm ~km ~s^{-1} ~Mpc^{-1}}$ and the baryon density $\\Omega_{\\rm B} = 0.046$. The matter density fluctuations are normalised by setting $\\sigma_8 = 0.81$ \\citep{WMAP}. All magnitudes are expressed in the AB system, and the Ly$\\alpha$ EW$_{\\rm Ly\\alpha}$ values in this paper are in the rest frame. ", "conclusions": "We have studied a variety of statistical properties of SMGs using a large cosmological hydrodynamic simulation. The simulation follows the formation and evolution of star-forming galaxies by employing a new feedback model of \\citet{Okamoto2010}. Earlier in \\citet{Shimizu2011}, the same simulation was used to study the properties of LAEs at $z = 3.1$. We have shown in the present paper that our SMG model is consistent with a number of observational data currently available. The redshift distribution of the SMGs peaks broadly at $z \\sim 2$. Most of the SMGs are at $z > 2$, and the contribution from low-$z$ ($z < 1$) galaxies to the submm source count is small. The contribution from the SMGs to the global star formation rate density increases with decreasing redshift. Bright SMGs contributes significantly to the star formation rate at $z \\sim 2-3$. We calculate the angular two point correlation to quantify the clustering amplitude of SMGs. Our result is consistent with the observations of \\citet{Williams2011}. We also study the angular cross correlation functions between SMGs and LBGs and the correlation between SMGs and LAEs. The former is significantly stronger than the latter. Considering the high star-formation rates of the simulated SMGs, we argue that SMGs are similar population to LBGs. Finally, we explicitly showed that the bright SMGs preferentially reside in massive halos ($> 10^{12}~\\rm M_{\\odot}$) and that their typical stellar mass are greater than $10^{11}~ \\rm M_{\\rm \\odot}$. \\begin{figure*} \\includegraphics[width = 160mm]{CCF_SMG_two_panel.eps} \\caption{ We plot the angular cross-correlation with two other populations of galaxies; the SMG-LAE cross-correlation (squares with error bars) and the SMG-LBG cross-correlation (triangles with error bars). For comparison, we also plot the angular auto-correlation of the SMGs (solid circles with error bars) and the angular cross correlation between the SMGs and dark matter halos with $10^{10}~\\rm M_{\\odot}$ (stars with error bars). We set two detection limits for the submm flux; $0.1~ \\rm mJy$ for the left panel and $1~ \\rm mJy$ for the right panel, respectively.} \\label{CCF} \\end{figure*} In our simple model of a single-sized, supernovae produced dust, the resulting extinction curve is flatter than the standard ones such the SMC extinction curve or Calzetti extinction curve. This mean that the reprocess of UV photons to FIR by dust works more efficiently in our model. However, interestingly, we have found that the characteristic feature of the extinction curve does not significantly affect our main results. For direct comparison, we calculate the FIR luminosities for our galaxy samples using the SMC extinction curve. We then estimate the SMG number count following otherwise the same procedures as in our fiducial model. Fig. \\ref{SMG_NC_COMPARISON} compares the source number count of our model (solid line) and the one calculated using the SMC extinction curve (dashed line). Clearly the difference between our base model and the other is very small. This can be explained by the fact that we normalize the overall luminosity of the galaxies by matching the UV luminosity function at rest-frame $1500~{\\rm \\AA}$ to the observed one (see Fig. \\ref{UV_LF}). Even for a different extinction model, the overall level of extinction in UV is always normalized at $1500~{\\rm \\AA}$. Therefore, the difference in the exact shape of the extinction curves near UV does not affect significantly the SMGs number count in our model, as is explicitly shown in Fig. \\ref{SMG_NC_COMPARISON}. Note, however, that the UV to optical colour of individual galaxies is visibly affected by the extinction curve. Our base model predicts that simulated galaxies appear bluer in UV to optical range than in the case with SMC or Calzetti extinction curve models. Fig. \\ref{SMG_SED} shows, as an illustrative example, the SED of one of our SMG samples from UV to millimetre. The SED for the same galaxy but with SMC extinction curve is also shown. We compare them with the observed SEDs for several SMGs \\citep{Michalowski2010}. Some SMGs have flat (blue) SEDs at UV to optical range whereas others appear redder. The full SEDs are measured only for a limited number of samples. It would be highly interesting to observe SMGs in detail in rest frame optical to UV in order to derive the dust extinction law, and hence the dust size distribution, for star burst galaxies. Next, we discuss the source confusion limit of submm observations. We have seen in previous sections (Fig. \\ref{SMG_STELLAR_MASS} and Fig. \\ref{SMG_MASS}), that some simulated SMGs have large host halo masses with $> 10^{13} \\rm M_{\\odot}$ and/or large stellar masses with $> 10^{12} \\rm M_{\\odot}$. Such massive systems often have multiple substructures (main and satellite galaxies). Basically each such satellite galaxy should be treated as an independent galaxy. Note, however, that the angular resolution of the current submm/mm observation is worse than that of optical observation. For example, the resolution of the AzTEC telescope is $30$ arcseconds which corresponds $260~ \\rm kpc~ (physical~ scale)$ at $z = 2$. The virial radii of the galactic halos in our simulation are only as large as $\\sim 200 ~\\rm kpc~ (physical~ size)$ at $z = 2$. Thus, our treatment that a halo hosts one SMG as a whole is reasonable, even realistic, if we compare with observations with angular resolution of 0.5-1 arcminutes. Finally, we discuss the evolution of the stellar mass function. Our galaxy formation model is defined at high redshift, rather than being calibrated against galaxies at the present epoch. We choose a few model parameters such as the overall normalization of dust extinction to reproduce the observational data of star-forming galaxies at high-$z$ ($z \\sim 3$)\\citep{Shimizu2011}. Contrastingly, previous semi-analytic models (e.g., \\citet{Baugh2005} and \\citet{Lacey2010}) determined the main physical parameters to match the observations of the present-day galaxies. For example, \\citet{Baugh2005} needed to adopt an extreme IMF (top-heavy IMF) for starburst galaxies at high redshift to reproduce the observed SMG number counts. More recently, \\citet{Fontanot2007} reproduced the SMGs number count adopting the Salpeter IMF for all their galaxy samples. However, their model assumes more efficient gas cooling in galactic haloes than in \\citet{Baugh2005} and \\citet{Lacey2010}. Consequently, very massive galaxies with high star-formation rates are formed, which results in disagreement in the stellar mass function at low redshifts ($z < 1$). Clearly it is important to study the stellar mass function of our model. \\begin{figure} \\includegraphics[width = 80mm]{SMG_host_halo_mass.eps} \\caption{The mean submm flux of our simulated SMG samples as a function of their host halo mass. We also plot the halo mass function at $z = 2$ (the solid line). Bright SMGs with $S > 1$ mJy are hosted by dark halos with mass greater than $\\sim 10^{13} M_{\\odot}$. } \\label{SMG_MASS} \\end{figure} \\begin{figure} \\includegraphics[width = 80mm]{SMG_nc_comparison.eps} \\caption{SMG number count at 1.1 mm for the two dust extinction models; our model based on a single-sized dust (solid) and the other using the SMC extinction curve (dashed).} \\label{SMG_NC_COMPARISON} \\end{figure} \\begin{figure} \\includegraphics[width = 80mm]{SMG_SED.eps} \\caption{An example SED of one of our simulated SMGs. The thick solid line and thin dashed line represent the dust-extinct SED for our simple dust model and that for the SMC extinction curve, respectively. We also show the observed SEDs for several SMGs \\citep{Michalowski2010} by open circles, open squares, open triangles and open diamonds. The luminosity of the observed galaxies are arbitrary normalized for this comparison. } \\label{SMG_SED} \\end{figure} \\begin{figure} \\includegraphics[width = 80mm]{smf_z1.eps} \\caption{The stellar mass function at $z = 1$. The solid line is our simulation result. The points with error bars show the observational result from the GOODS NICMOS survey \\citep{Mortlock2011}.} \\label{SMF} \\end{figure} Fig. \\ref{SMF} shows the stellar mass function of our simulated galaxies at $z = 1$. We also plot the observational data (solid points with error bars) of \\citet{Mortlock2011}. Our model is in good agreement with the data from the GOODS NICMOS survey as is clearly seen in Fig. \\ref{SMF}. This provides yet another support for our galaxy formation model that covers a multiple populations, from LAEs, LBGs to SMGs. In order to study the present-day stellar mass function, we have run our cosmological simulation down to $z=0$. Our simulation over-predicts the number density of massive galaxies with stellar mass greater than $10^{11} M_{\\odot}$. Therefore, although we have successfully constructed a consistent model for SMGs and UV-selected galaxies at high redshifts, it appears that our galaxy formation model lacks some physical mechanism(s) that shapes the present-day stellar mass function. As an attempt, we have implemented a model where gas cooling is quenched in large galaxies whose halos' velocity dispersions are greater than 141 km/sec \\footnote{We have done a few test calculations by varying the threshold velocity dispersion and have found that the run with $\\sigma_{\\rm th} = 141$ km/sec reproduces the break of the observed $z=0$ stellar mass function at $M_{*}\\sim 10^{11} M_{\\odot}$.}. Then the resulting stellar mass function closely agrees with the observational data at $z=0$ derived by Li \\& White (2009). We argue that the kind of strong feedback, in terms of star formation efficiency, may be needed for viable galaxy formation models, in order for them to reproduce all the available data from low to high redshifts. Constructing such a perfect model is the ultimate goal of the study of galaxy formation, but it is beyond the scope of the present paper. Further studies on the efficiency of star formation in galaxies are clearly needed using multiple approaches. Together with our previous study on Lyman-$\\alpha$ emitters at $z=3.1$\\citep{Shimizu2011}, we now provide a unified galaxy population model within the standard $\\Lambda$ Cold Dark Matter cosmology, which reproduces simultaneously the statistical properties of UV-selected star-forming galaxies and sub-millimetre galaxies at high redshifts." }, "1207/1207.4009_arXiv.txt": { "abstract": "{The survey of galaxy clusters performed by \\Planck\\ through the Sunyaev-Zeldovich effect has already discovered many interesting objects, thanks to the whole coverage of the sky. One of the SZ candidates detected in the early months of the mission near to the signal to noise threshold, \\name, was later revealed by \\xmm\\ to be a triple system of galaxy clusters. We have further investigated this puzzling system with a multi-wavelength approach and we present here the results from a deep \\xmm\\ re-observation. The characterisation of the physical properties of the three components has allowed us to build a template model to extract the total SZ signal of this system with \\Planck\\ data. We partly reconciled the discrepancy between the expected SZ signal from X-rays and the observed one, which are now consistent at less than $1.2\\,\\sigma$. We measured the redshift of the three components with the iron lines in the X-ray spectrum, and confirmed that the three clumps are likely part of the same supercluster structure. The analysis of the dynamical state of the three components, as well as the absence of detectable excess X-ray emission, suggest that we are witnessing the formation of a massive cluster at an early phase of interaction. } ", "introduction": " ", "conclusions": "The first observations of the multi-wavelength follow-up campaign of \\name, a triple system of galaxy clusters discovered by \\planck, have allowed us to improve our understanding of this object. With the new \\xmm\\ observation we estimated the global properties of each component: the ICM temperatures range from $3.5$ to 5 keV and the total masses within $\\R500$ are in the range $2.2-3\\, 10^{14} \\msol$. We detected the iron K$\\alpha$ lines in the X-ray spectra of each component, and therefore we were able to confirm that components $A$ and $C$ are lying at the same redshift ($z=0.45$). However, given the large angular separation of these two components ($7.5 \\arcmin$, corresponding to $2.6$ Mpc, in the plane of the sky), they have likely not started to interact yet and we did not detect significant excess X-ray emission between these two components. For component $B$, we estimated a larger redshift from X-ray spectroscopy ($z=0.48$), although consistent at two $\\sigma$ with the best fit value for component $A$. A similar indication is supported by the optical data, with the photometric redshifts we retrieved from SDSS DR8. However, given the large uncertainties of our redshift estimates (based both on X-rays and on photometry), a more detailed picture of the three-dimensional structure of \\name\\ will be possible only with the measurement of spectroscopic redshifts for a large sample of member galaxies, that is already foreseen with VLT in our follow-up program.\\\\ Our redshift results are consistent with the three clusters being part of the same supercluster structure, that will eventually lead to the formation of a massive object ($\\simeq 10^{15}\\msol$). This is supported also by our analysis of the galaxy population with SDSS data: the galaxy density maps show the presence of a possible population of inter-cluster galaxies, significant at $3\\sigma$, connecting the whole system (Fig.\\, \\ref{fig:galdens_zcut}). However, the relaxed appearance of component $A$, its large distance ($2.5$ Mpc) in the plane of the sky from component $C$ and along the line of sight from component $B$, as well as the absence of any detectable excess X-ray emission between the components may suggest that we are witnessing a very early phase of interaction.\\\\ Using the X-ray results from the new \\xmm\\ observation, we built a multicomponent model that we used to extract the total SZ signal from \\Planck\\ data. We compared the improved estimate of $Y_{SZ}$ with the prediction from X-rays and we found the latter to be about 68\\% of the measured SZ signal. The discrepancy between these two values is reduced with respect to Paper I and is only marginally significant at $1.2\\sigma$. \\\\ The results from our simulations have shown that an offset as large as $5\\arcmin$ can be expected in the reconstructed $y-$maps for low significance objects, due to noise fluctuations and astrophysical contributions. With this study we have illustrated the expected difficulty of accurately reconstructing the two-dimensional SZ signal for objects with low signal-to-noise ratio. Indeed the instrumental noise and astrophysical contamination compete seriously with the SZ effect at the detection limit threshold. Nonetheless, objects like \\name\\ can be detected with a dedicated optimal filtering detection method, and the SZ signal can be reconstructed assuming priors (such as position, size and relative intensity) from other wavelengths.\\\\ Despite a deep re-observation of this system with \\xmm, the intrinsic limitations of our X-ray data and of the current \\planck\\ SZ maps do not allow us for the time being to assess the presence of possible inter-cluster emission.\\\\ A careful analysis of the galaxy dynamics in the complex potential of this object and of the mass distribution from weak lensing will both be available with our on-going optical follow up program. These observations, combined with the results presented in this paper and with new \\planck\\ data obtained in two other full surveys of the sky, might deliver further clues for the understanding of this peculiar triple system." }, "1207/1207.7185_arXiv.txt": { "abstract": "Helium line observations towards 11 Galactic positions using Westerbork Synthesis Radio Telescope(WSRT) have been reported. These observations were made towards nearby positions where already hydrogen lines were detected at sufficiently high intensity($\\geq$50mK) at 1.4 GHz. This approach gave a fair chance for the detection of helium line as well, keeping in mind the relative abundance(10 \\%) of helium with respect to hydrogen. Care was also taken to avoid the presence of HII regions along the line of sight so that the line emission originates from the extended diffuse low density ionized component, ELDWIM of the Galaxy. The observations have resulted in the detection of helium line towards 5 positions out of 11 with signal to noise ratio(snr) $>$ 4$\\sigma$. An attempt has been made to associate detection/non-detection of helium line to the presence of surrounding HII regions. A weighting scheme that accounts for nearby($<$ 500pc) HII regions, their distances and other factors produces favourable results. It is seen from this weighting scheme that a higher weight favours the detection of helium line while lower weight is associated with non-detection. The idea is to correlate ionization of ELDWIM with the surrounding HII regions. ", "introduction": "The presence of a diffuse Extended Low Density Warm Ionized Medium(ELDWIM) within the inner Galaxy has been established by low frequency($<$2GHz) Radio Recombination Line(RRL or RL for short) observations(Lockman 1976 ; Mezger 1978 ; Ananatharamaiah 1985). This component is understood to have a typical electron density $n_e$ of $1-10 cm^{-3}$ at temperatures of $10^{3}-10^{4}$ K. The ELDWIM was first considered by Mezger(1978) who called it as the \"extended low density fully ionized gas\" which extends from the Galactic center to 13kpc and 100pc above and below the Galactic plane. However continuum radiation from this particular medium was first discovered by Westerhout(1958) in his Dwingeloo survey, who also calculated the corresponding upper limits on density and total mass of ionized hydrogen in the Galaxy. Later investigators have chosen to call this component as Extended Low Density Warm Ionized Medium(Petuchowski \\& Bennett 1993, Heiles 1994). The recent estimated densities of ELDWIM are $1-10 cm^{-3}$(Murray \\& Rahaman 2010) with temperatures of the order of 3000-8000K. The origin, morphology and ionization mechanism of ELDWIM are uncertain. ELDWIM has been thought to be as a collection of evolved HII regions(Shaver 1976) or perhaps it forms the outer envelopes of HII regions(Anantharamaiah 1985). In the later argument it has been suggested that by the sizes of HII region envelopes inferred from observations and due to their large number almost every line of sight intersects these envelopes in the inner Galaxy ($l<40^o$) giving the observed Galactic ridge recombination lines(Anantharamaiah 1986). More recently(Murray \\& Rahaman 2010) ELDWIM has been considered to be a diffuse gas ionized by massive stellar clusters unrelated to HII regions. The present observations aimed at detecting helium lines from ELDWIM to understand its ionization spectrum. The helium ionization potential(24.6 eV) is higher than that of hydrogen(13.6 eV). The detection and amplitude of He line would indicate the ionization condition of ELDWIM. Earlier RL observations have indicated that the ratio of the number of helium to hydrogen ions($N_{He+}/N_{H+}$) in ELDWIM is in the range 0 to 0.054(Heiles et.al 1996b). This ratio is smaller than the generally accepted cosmic abundance of helium, $N_{He}/N_{H} \\sim $ 0.1(Poppi et.al 2007 and references therein). Indicating all helium in ELDWIM to be partially ionized. According to order of magnitude calculations(Heiles et.al 1996b; see also Domgoergen \\& Mathis 1994) this observed ratio in ELDWIM requires the surface temperature of the ionizing star to be $<$35000K, if a standard HII region condition were to be considered. The current knowledge of initial mass function and the total Galactic star formation rate make it difficult to realize such a cool spectrum together with the total Galactic ionization requirement for the ELDWIM and HII regions. This has been called as the \\emph{ionization problem} (Heiles 1996b). However such calculations are not completely exhaustive and alternate ways to explain the ionization spectrum must be found. \\section[]{Observations} Helium RLs are much weaker than hydrogen RLs and hence longer integration time is required to detect them from specific directions. WSRT was used to observe hydrogen and helium RLs from 15 different Galactic positions. In the incoherent addition mode WSRT offers 8 IF bands, each with a bandwidth of 5 MHz. These were used to detect 7 Hn$\\alpha$/Hen$\\alpha$ RLs with n=165 to 171 and 1 Hn$\\beta$ RL with n=208 (Table-1). The primary objective was to detect the Hen$\\alpha$ line by averaging the spectra from the 7 bands. Observations were carried out using dual frequency switching with a shift of 2MHz, keeping the average of the hydrogen and helium rest frequencies at the center of the band. The resolution of the spectra is $\\sim$ 4km/s. The 15 positions were constructed from a list of previously observed(Lockman et.al 1976, 1989; Heiles 1996a) directions which exhibited strong($\\geq$50mK) hydrogen RLs with a single component at 1.4 GHz. A typical line strength of 50mK for hydrogen would imply $T_{He}$ = 0.025$\\times T_{H}$ = 1.25mK for helium, taking into account the mean 0.025 of the earlier mentioned $N_{He+}/N_{H+}$ ratio. With this strength for He lines the proposed integration time aimed at a 3$\\sigma$ detection. Care was also taken to avoid occurence of HII regions along the line of sight. This was to ensure line origin from ELDWIM. A list of relevant transitions and the emerging rest frequencies for hydrogen and helium RLs is given in Table-1. \\\\ \\begin{table} \\begin{center} \\label{transitions} \\begin{tabular}{ccc} \\hline Transition & $\\nu_{H}$ (GHz) & $\\nu_{He}$ (GHz) \\\\ \\hline $165\\alpha$ & 1.450716 & 1.451307 \\\\ $166\\alpha$ & 1.424734 & 1.425314 \\\\ $167\\alpha$ & 1.399368 & 1.399938 \\\\ $168\\alpha$ & 1.374601 & 1.375161 \\\\ $169\\alpha$ & 1.350414 & 1.350965 \\\\ $170\\alpha$ & 1.326792 & 1.327333 \\\\ $171\\alpha$ & 1.303718 & 1.304249 \\\\ $208\\beta$ & 1.44072 & \\\\ \\hline \\end{tabular} \\caption{List of 8 transitions and rest frequencies corresponding to each band.} \\end{center} \\end{table} The 7 spectra from different IF bands naturally had different velocity resolutions. To averge them together they were resampled and shifted using fourier transform interpolation. Which is essentially representing the spectrum by decomposed fourier components. Once the parameters of these components are available the spectrum could be replotted with any resolution. With a common resolution and alignment the spectra were averaged by weighting each spectrum by the inverse of the variance in the spectrum. However it should be noted that the spectra corresponding to H167$\\alpha$ was dropped due to the occurence of H210$\\beta$ line near to the He167$\\alpha$ line. These averaged spectra have been displayed in Figure 1 \\& 2 after a 3-box car smoothing to improve the snr. The parameters obtained from gaussian fits to these spectra are given in Table-2. \\section[]{Data Analysis} The data was acquired in fits format. Auto-correlations were extracted from the fits file using cfitsio library. These auto-correlations were power spectra for 2 polarizations obtained for 1-minute data per frequency setting(LO1 \\& LO2) designated as $T_{on}$ \\& $T_{off}$, with a shift of 2 MHz. The 48 minute observation towards each position gave 48 spectra for each polarization. The frequency switching was done every minute. So each setting had 24 spectra for each antenna, 14 for WSRT. A $T_{on}/T_{off} - 1$ of these spectra removed the background power simultaneously correcting for the gain variation across the band. Since the primary goal was to detect the He line care was taken to avoid any spectrum contaminated with interference. A visual examination of each 1 minute integrated spectrum was done before counting it in the averaging group. An average of all spectra in this group was followed by an appropriate folding to further average the lines appearing in the two parts of the spectrum due to frequency switching. Also the two polarization powers were combined into one at a suitable intermediate stage. \\\\ As mentioned in sec 2 WSRT offers 8 IF bands in the incoherent addition mode which were used to detect the transitions given in Table-1. No helium line could be confidently distinguished or recognized in a single band averaged spectrum. The spectra from the first 7 bands except the one for H167$\\alpha$ were averaged by fourier transform interpolation as explained in sec 2 \\& the previous paragraph. A 3-box car smoothing was further applied to improve the snr of this folded 6 band averaged spectrum. Finally an average system temperature (measured per minute) was used to calibrate the spectrum. These plots have been displayed in Figure 1 \\& 2 for 11 positions out of 15. 4 of the positions had corrupt data and had to be dropped. Gaussian parameter fits to the smoothed spectra have been given in Table-2. \\begin{center} \\begin{figure}[h] \\includegraphics[width=80mm,height=70mm,angle=0]{G10506013.eps} \\includegraphics[width=80mm,height=70mm,angle=0]{G10506014.eps} \\includegraphics[width=80mm,height=70mm,angle=0]{G10506015.eps} \\includegraphics[width=80mm,height=70mm,angle=0]{G10506016.eps} \\includegraphics[width=80mm,height=70mm,angle=0]{G10506017.eps} \\hspace{4.7mm} \\includegraphics[width=80mm,height=70mm,angle=0]{G10506018.eps} \\caption{Positions: G7.2-0.7, G15.8-0.5, G17.4+1.5, G18.6-0.8, G19.2+1.7 G23.4-0.6. Helium line is expected at an offset of -122.2 km/s from the hydrogen line. The obtained gaussian parameters are given in Table-2. The residual after subtracting the fit from data has been shown at an offset of -0.005 along the $T_L$-axis. It should be noted that the final spectra have been corrected for poor baselines by polynomial fitting to portions of the spectrum not containing any astronomical spectral line. Also before polynomial fitting any residual interference was edited out to avoid contribution to the polynomial fit. } \\end{figure} \\end{center} \\begin{center} \\begin{figure}[h] \\includegraphics[width=80mm,height=70mm,angle=0]{G10506019.eps} \\includegraphics[width=80mm,height=70mm,angle=0]{G10506020.eps} \\includegraphics[width=80mm,height=70mm,angle=0]{G10506130.eps} \\includegraphics[width=80mm,height=70mm,angle=0]{G10506131.eps} \\includegraphics[width=80mm,height=70mm,angle=0]{G10506132.eps} \\caption{Positions: G26.5-0.0, G28.0-0.0, G16.3+1.3, G18.0+1.8, G12.6-0.6. Helium line is expected at an offset of -122.2 km/s from the hydrogen line. The obtained gaussian parameters are given in Table-2. The residual after subtracting the fit from data has been shown at an offset of -0.005 along the $T_L$-axis. It should be noted that the final spectra have been corrected for poor baselines by polynomial fitting to portions of the spectrum not containing any astronomical spectral line. Also before polynomial fitting any residual interference was edited out to avoid contribution to the polynomial fit.} \\end{figure} \\end{center} \\clearpage \\begin{table} \\begin{center} \\label{parameters} \\begin{tabular}{cccccccccc} \\hline \\multicolumn{2}{c}{Source} & \\multicolumn{2}{c}{$T_{L}$ (mK)} & \\multicolumn{2}{c}{$V_{LSR}$ (km/s)} & \\multicolumn{2}{c}{$ \\Delta V_{LSR}$ (km/s)} & {$\\sigma$} & \\\\ \\cline{1-8} $l^{o}$ & $b^{o}$ & H & He & H & He & H & He & (mK) & {$N_{He+}/N_{H+}$} \\\\ & & & & & & & & & \\\\ \\hline & & & & & & & & & \\\\ 7.2&-0.7&41.2(0.5)&-&18.0(0.2)&-&31.0(0.4)&-&0.74&$<$0.018 \\\\ & & & & & & & & & \\\\ & & & & & & & & & \\\\ 15.8&-0.5&35.2(0.4)&2.25(0.3)&28.4(0.5)&-102.8(2.0)&36.3(0.8)&25.2(5.0)&&0.064(0.01) \\\\ & & 15.4(0.8) & & 59.0(0.8) & & 27.1(1.2) & & 0.41 & \\\\ & & & & & & & & & \\\\ & & & & & & & & & \\\\ 17.4&+1.5&33.4(2.2)&2.73(0.6)&21.9(0.2)&-102.0(1.2)&21.7(0.9)&11.2(2.8)&& 0.082(0.02) \\\\ & & 18.7(1.6)&&35.7(1.7)&&40.8(1.4)&&0.66& \\\\ & & & & & & & & & \\\\ & & & & & & & & & \\\\ 18.6&-0.8&9.1(1.8)&-&29.3(6.1)&-&34.8(6.7)&-&& \\\\ && 10.3(3.8)&&43.7(0.8)&&17.7(2.8)&&0.69& \\\\ &&30.4(0.1)&&64.5(0.7)&&30.5(1.0&&&$<$0.023 \\\\ & & & & & & & & & \\\\ & & & & & & & & & \\\\ 19.2&+1.7&41.8(6.2)&2.9(0.2)&24.1(1.5)&-98.2(2.6)&34.6(1.1)&35.7(5.4)&&0.069(0.015) \\\\ &&14.8(3.5)&&30.4(0.4)&&18.8(1.8)&&0.41&\\\\ &&4.2(2.1)&&55.4(18.8)&&48.6(18.6)&&& \\\\ & & & & & & & & & \\\\ & & & & & & & & & \\\\ 23.4&-0.6&40.0(5.4)&-&61.4(1.5)&-&32.1(1.7)&-&& \\\\ &&41.4(2.8)&&90.3(2.6)&&42.3(3.1)&&0.7&$<$0.017\\\\ & & & & & & & & & \\\\ & & & & & & & & & \\\\ 26.5&0.0&3.8(0.6)&-&26.1(1.2)&-&14.9(2.9)&-&& \\\\ &&29.3(0.7)&&71.6(0.47)&&25.2(0.85)&&0.78& \\\\ &&75.4(0.5)&&103.4(0.2)&&32.3(0.44)&&&$<$0.01 \\\\ & & & & & & & & & \\\\ & & & & & & & & & \\\\ 28.0&0.0&11.7(0.5)&-&39.4(1.1)&-&29.0(2.3)&-&&\\\\ &&39.0(12.0)&&87.5(7.6)&&34.5(8.2)&&0.55& \\\\ &&24.1(24.4)&&103.7(1.7)&&22.8(5.4)&&& \\\\ &&4.5(0.6)&&138.7(2.7)&&28.0(5.1)&&&$<$0.014 \\\\ & & & & & & & & & \\\\ & & & & & & & & & \\\\ 16.3&+1.3&31.1(0.5)&-&27.1(0.25)&-&32.6(0.6)&-&0.92&$<$0.03 \\\\ & & & & & & & & & \\\\ & & & & & & & & & \\\\ 18.0&+1.8&124.5(0.5)&8.4(1.2)&27.4(0.06)&-92.8(3.5)&28.4(0.1)&26.0(5.2)&0.75&0.067(0.01) \\\\ & & & & & & & & & \\\\ & & & & & & & & & \\\\ 12.6&-0.6&56.1(13.0)&3.2(0.7)&34.2(1.0)&-81.5(4.2)&22.9(2.2)&25.5(6.9)&&0.057(0.026) \\\\ &&22.7(8.0)&&16.9(5.5)&&27.6(4.7)&&0.46& \\\\ &&28.8(2.5&&56.2(1.9)&&27.9(2.1)&&& \\\\ & & & & & & & & & \\\\ \\hline \\end{tabular} \\caption{Gaussian parameters fitted to the spectra in Figure 1 \\& 2. The values in the paranthesis are errors. Regions without helium line detection have upper limits defined as $\\sigma/T_{L}(H)$. All the ratios have been calculated assuming the helium line to be associated with the strongest hydrogen component. It should be noted that the associated carbon line parameters have not be tabulated. Carbon RL is expected at an offset of $\\sim-150 km/s$ w.r.t the hydrogen RL.} \\end{center} \\end{table} \\clearpage \\section[]{Discussion of parameters in Table-2} Table-2 shows that out of 11, 5 positions(G15.8-0.5,G17.4+1.5,G19.2+1.7,G18.0+1.8 \\& G12.6-0.6) exhibit He line detections with snr $>$ 4$\\sigma$. For regions were hydrogen and helium are homogeneously ionized the ratio $N_{He+}/N_{H+}$ is equal to the ratio of observed line strengths of He and H. In case of non-detection of He line an upper limit was obtained by taking the ratio of rms of noise to the strongest H line strength. The highest ratio is towards G17.4+1.5 with a value of 0.082. But this value comes from the assumption that the helium and hydrogen zones overlap. Some spectra exhibit pressure broadening, clearly visible with the H component having extended wings at the edge of the profile. Especially the spectra to be mentioned are G23.4-0.6 and G19.2+1.7. The non-detection or lower intensity of He line than expected implies that it is not ionized or is partially ionized in ELDWIM. Density bound HII regions can act as source of hot ionizing photons(Ferguson et.al 1996; Wood \\& Mathis 2004). The escaping photons will ionize ELDWIM in the vicinity. However the photon spectrum will be drastically modified(Osterbrock 1989; Hoopes \\& Walterbos 2003; Wood \\& Mathis 2004) by the gas surrounding the HII regions. Since any photon that can ionize helium can also ionize hydrogen(Osterbrock, 1989). Only those higher energy photons that escape hydrogen, which is much more abundant than helium, can ionize helium in ELDWIM. This investigation has revealed a high value of $N_{He+}/N_{H+}$ ratio $\\sim$ 0.082 which is close to the primordial cosmic abundance $\\sim$0.1(Poppi et al 2007 and references therein). \\\\ There are also interesting positions like G28.0+0.0 \\& G26.5+0.0 in the Galactic plane which show a lack of He line signature. More observations or a further study of observations towards these and surrounding regions can reveal more about the ionization of ELDWIM in these directions. It is required to have an understanding of distribution of HII regions towards these directions. A non-detection indicates that either helium is under abundant or is not ionized within these regions. The detection of helium line can also be reverted back to the morphology and ionization spectrum originating from the HII regions. Lack of He line indicates that the photons leaking from the surrounding HII regions have a cooler spectrum whereas directions towards which He line is seen have a stronger ionizing spectrum. This investigation has produced 5 positions towards which distinct He line profiles have been observed. Indicating clearly the presence of diffuse ionized helium. RL detections towards nearby HII regions by earlier observers(Heiles et.al 1996a,b) match closely with the $V_{LSR}$ of the current observed lines, but however are stronger. \\\\ The line widths of H and He lines is another interesting aspect. Towards some regions like G17.4+1.5 \\& G15.8-0.5 the line width of He is nearly half of H width, within error bars. This is expected in the case of pure thermal doppler broadening. However turbulence can make He line width more than half of H line width. The width agreement between He and H is a criterion to hold the two lines to be originating from the same region. In the case(G17.4+1.5) where He line width is simply half of H line width it suggests that the lines originate in a low density medium where negligible pressure broadening \\& turbulence is expected. Pressure broadening could contribute to the width of the He line in denser regions and produce extended wings(Smirnov et al 1984). There seems to be significant amount of pressure broadening towards positions like, G23.4-0.6 \\& G19.2+1.7. The non-agreement of He and H line widths indicates that they can originate in different regions. In some cases the width of He line is larger than H line this may be due to blending of different components into one. Further the constraint of observed $V_{LSR}$ difference($\\sim$122.2 km/s) between H and He suggests that the region of origin moves at the same velocity along the line of sight. Based on differential Galactic rotation is the same cloud or a nearby cloud. During gaussian fitting of parameters it has been assumed that the He line is primarily associated with the strongest H line. This seems to be true towards positions G17.4+1.5, G19.2+1.7 \\& G18.0+1.8 where the $V_{LSR}$ for He and H are in good agreement within the error bars. Broad He lines are seen towards positions which have multiple H line components. It was not possible to fit similar multiple components to He line and obtain consistent parameters with H line widths \\& amplitudes. In case of G18.0+1.8 \\& G19.2+1.7 a complete agreement of $V_{LSR}$ between He and H is seen but the widths do not seem to be only thermal. The extra width of He indicates turbulence.\\\\ Another interesting region in the sample is G18.6-0.8 which exhibits multiple H line components. Even though line like features are seen towards $V_{LSR}$ of -112 km/s, the fitted parameters are not consistent with the expected $V_{LSR}$ difference between He and H lines to be originating from the same region. This position seems to cover an HII region at the edge within the beam(Figure 3 \\& 4). It is likely that the prominent H \\& He lines come from different regions for this direction. The set of positions with no He line detection are G7.2-0.7,G23.4-0.6,G26.5+0.0,G28.0+0.0 \\& G16.3+1.3. The position G7.2-0.7 shows a strong H line but has no associated He line which is expected from the fitted parameters to H line at $V_{LSR}$ of -104 km/s. However this position is a good candidate for inspecting ELDWIM and its ionization. G23.4-0.6 is another good candidate for ELDWIM, perhaps with a smaller telescope beam. The H line feature for this position seems to show significant pressure broadening, the most distinct in the sample. This position accepts 2 Voigt profiles instead of 2 Gaussians. This may indicate contribution from denser regions. G16.3+1.3 due to its poor snr may still be a valuable candidate for He line detection. The weighting scheme(sec 5) predicts a He line detection for this position. \\\\ The possible kinematical distances(Sofue et.al 2009) of ELDWIM clouds due to $V_{LSR}$ of the lines for all the 11 positions in Table 2 have been marked in Figure 5 . The identified HII regions(Paladini et.al 2003 ; Anderson et.al 2011) distributed in Galactic longitude and latitude within a distance of 500pc from the line originating region have also been marked on the same plot to show their relative location to the ELDWIM clouds. \\\\ \\section[]{Correlation with HII regions} The obtained values of $N_{He+}/N_{H+}$ ratio for different positions have been tabulated in Table 2 \\& 3. This ratio with the normal accepted primordial abundance of helium to hydrogen is expected to be 0.1(Peimbert et.al 1988; Osterbrock 1989; Baldwin et.al 1991; Reynolds 1995; Heiles et.al 1996b; Madsen et.al 2006; Poppi et.al 2007). Here in most of the cases it has turned out to be smaller than this. When helium optical lines(Reynolds 1995, Madsen et.al 2006) or RL(Heiles et.al 1996b) are observed in HII regions this depression in the ratio is attributed to the smallness of the ionized He zone compared to the H zone which engulfs the former (Osterbrock 1989). However in the case of ELDWIM gas is ionized externally, i.e ionization happens from periphery to inwards. In a diffuse extended gas it is expected that helium and hydrogen are ionized similarly with no morphological difference between them. Assuming the ionizing radiation is strong. Under such a situation the ratio $N_{He+}/N_{H+}$ should reveal the abundance ratio between the two elements assuming complete ionization. However in ELDWIM this is not seen. This ratio is less than the expected abundance ratio. This implies either helium is under abundant or the radiation ionizing the extended diffuse gas is not strong enough to ionize it completely. \\\\ In this paper an attempt has been made to correlate the identified HII regions(Paladini et.al 2003 ; Anderson et.al 2011) around the line originating region with helium RL detection. It is seen from this correlation that the ratio $N_{He+}/N_{H+}$ depends on the distribution of HII regions around them(Figure 5) together with the distance to the cloud from the solar system. This correlation scheme of ionization of ELDWIM with surrounding HII regions has been quantified as a weight W associated with each observed region given by, \\begin{equation} W~=~\\frac{N_{HII} {D_c}^{2}}{{D_c}^{3}{D_c}^{3}}\\left[\\frac{1}{D_{1HII}^2}+\\frac{1}{D_{2HII}^2} + \\cdot \\cdot \\cdot +\\frac{1}{D_{nHII}^2} \\right]~=~\\frac{N_{HII}}{{D_c}^{4}}\\sum_{n}{\\frac{1}{D_{nHII}^2}}. \\end{equation} The weight has been considered to be proportional to the inverse square of the distance to the HII region $D_{nHII}$, as the flux from a source goes as inverse square of the distance. Since every HII region may or may not be a source of ionizing photons the weight has been taken to be proportional to the number of HII regions, $N_{HII}$. The distance to the cloud determines the amount of ELDWIM within the beam. Considering a sphere centered at the line originating region the volume of ELDWIM within the beam is proportional to $D_c^3$ and so is the amount of helium \\& hydrogen. The photons that can ionize helium can also ionize hydrogen. For a given number of photons that can ionize helium, helium ionization is diluted due to the presence of hydrogen. The detection weight has been taken to be inversly proportional to the net amount of helium and hydrogen in view of a weight for the ratio $N_{He+}/N_{H+}$. This contributes $D_c^{6}$ in the denominator. The surface area of this volume ($\\propto $$D_c^2$) measures the photon input into it. This contributes a $D_c^2$ in the numerator. The net contribution from all the above has been accounted by introducing a $4^{th}$ power of distance to the cloud in the denominator on right hand side of Equation (1).The identified HII regions have been marked around the line originating regions in Figure 5 as per their nearest kinematical distances(Sofue et.al 2009). It can be seen from Figure 5 that regions lying beyond 4 kpc from the solar system do not exhibit He line detection. A plot of observed ratio of $N_{He+}/N_{H+}$ Vs corresponding weight $Log_{10}[W_{near}]$for potential candidates representing clean ELDWIM has been given in Figure 6. It can be seen from this plot that higher ratio of $N_{He+}/N_{H+}$ is seen for regions with higher weight, i.e they are grouped towards the right upper corner of the plot. While lower weight regions are grouped towards the left lower corner of the plot. The weights towards different positions have been tabulated in Table-3. \\begin{figure}[h] \\begin{center} \\includegraphics[width=150mm,height=210mm,angle=0]{2.7GHz_pos_all.eps} \\caption{Positions: WSRT beam($\\sim0.5^{o}$) on the first 6 of the 11 positions on 11 cm continuum map(Reich et.al 1990).} \\end{center} \\end{figure} \\begin{center} \\begin{figure}[h] \\includegraphics[width=150mm,height=220mm,angle=0]{2.7GHz_pos_all1.eps} \\caption{Positions: WSRT beam($\\sim0.5^{o}$) on the last 5 of the 11 positions on 11 cm continuum map(Reich et.al 1990).} \\end{figure} \\end{center} \\begin{center} \\begin{figure}[h] \\includegraphics[trim = 1mm 1mm 1mm 120mm, clip, width=63mm,angle=-90]{MarkClouds_013.ps} \\includegraphics[trim = 1mm 1mm 1mm 120mm, clip, width=63mm,angle=-90]{MarkClouds_014.ps} \\includegraphics[trim = 1mm 1mm 1mm 120mm, clip, width=63mm,angle=-90]{MarkClouds_015.ps} \\includegraphics[trim = 1mm 1mm 1mm 120mm, clip, width=63mm,angle=-90]{MarkClouds_016.ps} \\includegraphics[trim = 1mm 1mm 1mm 120mm, clip, width=63mm,angle=-90]{MarkClouds_017.ps} \\includegraphics[trim = 1mm 1mm 1mm 120mm, clip, width=63mm,angle=-90]{MarkClouds_018.ps} \\includegraphics[trim = 1mm 1mm 1mm 120mm, clip, width=63mm,angle=-90]{MarkClouds_019.ps} \\includegraphics[trim = 1mm 1mm 1mm 120mm, clip, width=63mm,angle=-90]{MarkClouds_020.ps} \\includegraphics[trim = 1mm 1mm 1mm 120mm, clip, width=63mm,angle=-90]{MarkClouds_130.ps} \\includegraphics[trim = 1mm 1mm 1mm 120mm, clip, width=63mm,angle=-90]{MarkClouds_131.ps} \\hspace{9mm}\\includegraphics[trim = 1mm 1mm 1mm 120mm, clip, width=65mm,angle=-90]{MarkClouds_132.ps} \\caption{Distribution of HII regions(Paladini et.al 2003 ; Anderson et.al 2011) around the line originating ELDWIM clouds. Kinematical distances are due to Sofue et.al(2009).} \\end{figure} \\end{center} \\clearpage \\begin{table}[h] \\begin{center} \\centering \\label{parameters} \\begin{tabular}{cccccccccc} \\hline \\multicolumn{2}{c}{Source} & & & \\\\ \\cline{1-2} $l^{o}$ & $b^{o}$&$Log_{10}[W_{near}]$& {$N_{He+}/N_{H+}$}& Comments \\\\ & & & & \\\\ \\hline & & & & \\\\ 7.2&-0.7&2.43&$<$0.018&-\\\\ 15.8&-0.5&2.86&0.064&- \\\\ 17.4&+1.5&2.80&0.082&-\\\\ 18.6&-0.8&1.90&$<$0.023&H,He $V_{LSR}$ disagreement \\\\ 19.2&+1.7&2.80&0.069&He width broad, but mulitple H components\\\\ 23.4&-0.6&1.76&$<$0.017&Significant pressure broadening\\\\ 26.5&0.0&1.91&$<$0.01&-\\\\ 28.0&0.0&1.56&$<$0.014&-\\\\ 16.3&+1.3&3.16&$<$0.03&-\\\\ 18.0&+1.8&2.76&0.067&Broad He,but $<$ H width, single H component\\\\ 12.6&-0.6&2.74&0.057&-\\\\ & & & & \\\\ \\hline \\end{tabular} \\caption{Correlation Weights for line originating ELDWIM towards different positions. $W_{near}$ (Eqn (1))is weight for near by ELDWIM clouds. The near and far distances are two possible solutions along the line of sight given by the differential Galactic rotation curve. Here only the near distance has been considered, hence $W \\rightarrow W_{near}$ with $D_c$ being the near cloud distance in Equation(1).} \\end{center} \\end{table} \\vspace{1.8cm} \\begin{center} \\begin{figure}[h] \\vspace{-1.8cm} \\hspace{2.0cm}\\includegraphics[width=80mm,height=110mm,angle=-90]{NHe_NHpVsWnear1.ps} \\caption{Plot of $N_{He+}/N_{H+}$ Vs $Log_{10}\\left[W_{near}\\right]$ for all the 11 positions.} \\end{figure} \\end{center} \\vspace{-1.0cm} \\section[]{Summary} This work presents the results of hydrogen and helium RL observations towards 11 different Galactic positions using WSRT at $\\sim$ 1.4GHz in incoherent addition mode. Out of 11, 5 positions exhibit helium RL detection with snr $>$4$\\sigma$. WSRT provided 8 IF bands each of width 5 MHz which were used to observe 8 distinct RL lines(Hn$\\alpha$/Hen$\\alpha$ with n=165-171 and H208$\\beta$) using frequency switching. Details of observation and data analysis have been given in sec 2 \\& sec 3. These observations aimed at detecting helium RLs from ELDWIM. Since the abundance(~0.1) of helium is smaller than that of hydrogen it is difficult to detect helium RLs compared to the corresponding hydrogen RLs. In view of this a set of 15 positions was constructed by consulting previous RL observations (Lockman et.al 1976, 1989; Heiles et.al 1996a) and the 11 cm continuum map(Reich et al 1990). In this procedure a region near to a previous strong($\\sim$ 50mK) hydrogen RL detection at 1.4 GHz was considered. Next during positioning of the beam on the 11 cm map care was taken not to include any possible HII region within the beam. Line detections from 11 of these positions have been displayed in Figure 1 \\& 2. The fitted gaussian parameters are tabulated in Table-2, along with the obtained ratio $N_{He+}/N_{H+}$. Which is the ratio of line amplitudes of helium and hydrogen. A detailed discussion of line parameters and their implication has been given in sec 4. In general most of the regions seem to be fit to represent ELDWIM. This investigation has produced a high value(0.082) of $N_{He+}/N_{H+}$ ratio which is close to the accepted abundance(0.1) of helium to hydrogen. Indicating the presence of diffuse ionized helium in ELDWIM. Further a weighting scheme(Eqn (1)) has been adopted to show correlation between line orginating ELDWIM and the surrounding HII regions(Paladini et.al 2003 \\& Anderson et.al 2011). According to this scheme a weight is assigned to each region depending on (i)the number of surrounding HII regions within a certain radius from the cloud (ii)their distance from the cloud and (iii)the distance(nearest) of the cloud from the solar system(Figure 5). This correlation has been displayed by plotting $N_{He+}/N_{H+}$ against the obtained weights, $Log_{10}[W_{near}]$. The conclusion from this plot is that regions with high value of $N_{He+}/N_{H+}$ bear a high weight. Indicating a correlation between HII regions and helium line detection. ", "conclusions": "" }, "1207/1207.0888_arXiv.txt": { "abstract": "Near the minor axis of the Galactic bulge, at latitudes $b <-5^\\circ$, the red giant clump stars are split into two components along the line of sight. We investigate this split using the three fields from the ARGOS survey that lie on the minor axis at $(l,b) = (0^\\circ,-5^\\circ), (0^\\circ,-7.5^\\circ), (0^\\circ,-10^\\circ)$. The separation is evident for stars with [Fe/H] $> -0.5$ in the two higher-latitude fields, but not in the field at $b = -5^\\circ$. Stars with [Fe/H] $< -0.5$ do not show the split. We compare the spatial distribution and kinematics of the clump stars with predictions from an evolutionary N-body model of a bulge that grew from a disk via bar-related instabilities. The density distribution of the peanut-shaped model is depressed near its minor axis. This produces a bimodal distribution of stars along the line of sight through the bulge near its minor axis, very much as seen in our observations. The observed and modelled kinematics of the two groups of stars are also similar. We conclude that the split red clump of the bulge is probably a generic feature of boxy/peanut bulges that grew from disks, and that the disk from which the bulge grew had relatively few stars with [Fe/H] $< -0.5$. ", "introduction": "The Milky Way has a boxy/peanut-shaped bulge, as seen in the 2MASS star counts (L{\\'o}pez-Corredoira et al. 2005) and the COBE/DIRBE near-infrared light distribution (Dwek et al. 1995). The long axis of this bulge/bar lies in the Galactic plane and points into the first Galactic quadrant at an angle of about $20^\\circ$ to the Sun-center line (Gerhard 2002). The distribution of stars associated with the boxy/peanut structure shows complexity at higher Galactic latitudes. Near the minor axis, at $b < -5^\\circ$, recent analyses of optical and near-IR stellar photometry by Saito et al. (2011), Nataf et al. (2010) and McWilliam \\& Zoccali (2010) show that the red giant clump stars are split into two components separated by about 0.65 mag along the line of sight. McWilliam \\& Zoccali (2010) argue that the split clump is associated with an underlying X-shaped structure in the bulge. Such X-shaped structure has been noted earlier for other boxy/peanut bulges: (e.g. Whitmore \\& Bell (1988)). We have recently completed the Abundances and Radial velocity Galactic Origins Survey (ARGOS) which is a spectroscopic survey of 28,000 stars, mostly clump giants, in 28 fields of the Galactic bulge and nearby disk. The ARGOS survey provides a clean sample of bulge giants with distances, radial velocities and metallicities. As part of this survey, we have data for about 3000 stars in three two-degree fields along the minor axis, at $(l,b)$ = $(0^\\circ,-5^\\circ), (0^\\circ,-7.5^\\circ)$ and $(0^\\circ,-10^\\circ)$. These fields are hereafter denoted as m0m5, m0m75 and m0m10. We find an obvious split in the magnitude distribution of the red clump stars along the minor axis in the higher-latitude fields, m0m75 and m0m10. The objectives of this paper are (1) to measure the metallicity range of the stars that show the split red clump near the minor axis of the bulge, (2) to determine whether the split red clump has an associated kinematic signature, and (3) to interpret these results in the context of an evolutionary N-body model. Our study complements previous photometric analyses of the stellar density and magnitude distribution and expands upon previous spectroscopic studies: e.g Rangwala \\& Williams (2009); De Propis (2011). ", "conclusions": "The ARGOS survey provides a large sample of bulge stars with accurate abundances and radial velocities, relatively free of contamination from foreground and background stars. In this paper we used ARGOS stars near the minor axis of the bulge to investigate the properties of the split red clump. The split in the red clump magnitude distribution is seen very clearly in the higher latitude fields ($b = -6.5^\\circ$ to $-11^\\circ$) near the minor axis, for stars with [Fe/H] $> -0.5$. The more metal-poor red clump stars ([Fe/H]$ < -0.5$) in the same regions of the bulge do not show the split, and it is almost absent at all metallicities for our m0m5 field ($b = -4^\\circ$ to $-6^\\circ$). We compared the structure and kinematics of the ARGOS stars with an N-body model of a bulge. In the model, the boxy/peanut bulge grew through the bar-forming and bar-buckling instability of a disk of stars. Although the model was not constructed to fit the Galactic bulge, it is a good representation. The orbital structure of the boxy/peanut bulge leads to the central depression in its peanut-shaped density distribution: see Athanassoula (2005). The central depression in turn leads to a split bimodal structure along the line of sight at higher latitudes which closely resembles the split observed in the ARGOS stars and in the Saito et al. (2011) data. At low latitude, the kinematics of the bulge stars in the bright and faint groups of the split red clump are also similar to those for the corresponding stars in the model. In projection, the density distribution would show an X-shaped structure, particularly after unsharp masking: see Whitmore \\& Bell (1988), Athanassoula (2005) and Bureau et al. (2006). It seems likely that the split red clump seen in the Galactic bulge is a generic feature of boxy/peanut bulges which grew from disks. We also conclude that the disk from which the bar/bulge formed would have had relatively few stars with [Fe/H] $< -0.5$." }, "1207/1207.4245_arXiv.txt": { "abstract": "We report the detection of UCF-1.01, a strong exoplanet candidate with a radius 0.66 {\\pm} 0.04 times that of Earth ($R\\sb{\\oplus}$). This sub-Earth-sized planet transits the nearby M-dwarf star GJ 436 with a period of 1.365862 {\\pm} 8$\\times$10$^{-6}$ days. We also report evidence of a 0.65 {\\pm} 0.06 $R\\sb{\\oplus}$ exoplanet candidate (labeled UCF-1.02) orbiting the same star with an undetermined period. Using the {\\em Spitzer Space Telescope}, we measure the dimming of light as the planets pass in front of their parent star to assess their sizes and orbital parameters. If confirmed, UCF-1.01 and UCF-1.02 would be called GJ 436c and GJ 436d, respectively, and would be part of the first multiple-transiting-planet system outside of the Kepler field. Assuming Earth-like densities of 5.515 g/cm$^{3}$, we predict both candidates to have similar masses ($\\sim$0.28 Earth-masses, $M\\sb{\\oplus}$, 2.6 Mars-masses) and surface gravities of $\\sim$0.65 $g$ (where $g$ is the gravity on Earth). UCF-1.01's equilibrium temperature ($T$\\sb{eq}, where emitted and absorbed radiation balance for an equivalent blackbody) is 860 K, making the planet unlikely to harbor life as on Earth. Its weak gravitational field and close proximity to its host star imply that UCF-1.01 is unlikely to have retained its original atmosphere; however, a transient atmosphere is possible if recent impacts or tidal heating were to supply volatiles to the surface. We also present additional observations of GJ 436b during secondary eclipse. The 3.6-{\\micron} light curve shows indications of stellar activity, making a reliable secondary eclipse measurement impossible. A second non-detection at 4.5 {\\microns} supports our previous work in which we find a methane-deficient and carbon monoxide-rich dayside atmosphere. ", "introduction": "\\label{intro} The search for Earth-sized planets around main-sequence stars has progressed expeditiously in the last year. Recent discoveries include two Earth-sized planets (0.868 and 1.03 Earth radii, $R\\sb{\\oplus}$) from the Kepler-20 system \\citep{Fressin2011}, two planet candidates (0.759 and 0.867 $R\\sb{\\oplus}$) from the KIC 05807616 system \\citep{Charpinet2011}, and a three-planet system (0.78, 0.73, and 0.57 $R\\sb{\\oplus}$) orbiting KOI-961 \\citep{Muirhead2012}. The search for a second planet in the GJ 436 system began shortly after the transit detection and confirmed eccentric orbit of GJ 436b \\citep{Gillon2007b,Deming2007}. In 2008, a $\\sim$5-$M\\sb{\\oplus}$ planet on a 5.2-day orbit was proposed by \\citet[later retracted]{Ribas2008} due to three lines of evidence. First, the lack of detectable GJ 436b transits at the time of its 2004 discovery using radial-velocity (RV) measurements \\citep{Butler2004} suggests a change in orbital inclination due to a perturber. Second, given a circularization timescale of $\\sim$30 Myr \\citep{Deming2007} and the estimated 6-Gyr age of the system \\citep{Torres2007}, GJ 436b's non-circular orbit suggests another planet is pumping up its eccentricity. Third, there was evidence of a residual low-amplitude RV signal in a 2:1 mean-motion resonance with GJ 436b \\citep{Ribas2008}. The inferred planet was discredited by orbital-dynamic simulations \\citep{Bean2008,Demory2009} and the absence of transit timing variations (TTVs) with two transit events with the Near Infrared Camera and Multi Object Spectrograph camera on the Hubble Space Telescope \\citep{Pont2009} and over a 254-day span using ground-based H-band observations \\citep{Alonso2008}. \\citet{Ballard2010a}'s analysis of 22 days of nearly continuous observations of GJ 436 during NASA's EPOXI mission ruled out transiting exoplanets $>$2.0 $R\\sb{\\oplus}$ outside GJ 436b's 2.64-day orbit (out to a period of 8.5 days) and $>$1.5 $R\\sb{\\oplus}$ interior to GJ 436b, both with a confidence of 95\\%. Aided by a $\\sim$70-hour {\\em Spitzer} observation of GJ 436 at 8.0 {\\microns}, \\citet{Ballard2010b} postulated the presence of a 0.75-$R\\sb{\\oplus}$ planet with a period of 2.1076 days. However, the predicted transit was not detected in an 18-hour follow-up observation with {\\em Spitzer} at 4.5 {\\microns}. The candidate transit signals in the EPOXI data were likely the result of correlated noise \\citep{Ballard2010b}. In this paper we present {\\em Spitzer} primary-transit observations of UCF-1.01 and UCF-1.02 at 4.5 {\\microns} (including an independent analysis), a phase curve of GJ 436b at 8.0 {\\microns} in which transits of UCF-1.01 are modeled, and a publicly-available EPOXI light curve phased to the period of UCF-1.01. We also include secondary-eclipse observations of GJ 436b at 3.6 and 4.5 {\\microns}. In Section \\ref{sec:obs}, we describe the observations and data analysis. Section \\ref{sec:tide} presents Time-series Image Denoising (TIDe, a wavelet-based technique used to improve image centers) and provides an example analysis using a fake dataset. In Section \\ref{sec:results}, we discuss the specific steps taken with each of the six {\\em Spitzer} datasets, the details of our independent analysis, and transit results from the EPOXI light curve. Section \\ref{sec:disc} describes how we eliminate false positives, our radial-velocity analysis, mass constraints on both sub-Earth-sized exoplanets, and orbital and atmospheric constraints on UCF-1.01. Finally, we give our conclusions in Section \\ref{sec:concl} and supply the full set of best-fit parameters with uncertainties in the Appendix. ", "conclusions": "\\label{sec:concl} In this paper, we announced the detection of UCF-1.01 and UCF-1.02, two sub-Earth-sized transiting exoplanet candidates orbiting the nearby M dwarf GJ 436. Their detections were possible with BLISS mapping and Time-series Image Denoising (TIDe), the latter of which is a novel wavelet-based technique that decreases high-frequency noise in short-cadence, time-series images to improve image centering precision. We presented four transits of UCF-1.01 and two transits of UCF-1.02 at 4.5 {\\microns}, an independent analysis that confirms our best-fit results within 1.5$\\sigma$, an 8.0-{\\micron} phase curve of GJ 436b that includes transits of UCF-1.01, and EPOXI data that are consistent with the presence of a sub-Earth-sized exoplanet. To definitively establish UCF-1.01 as a planet (to be called GJ 436c), we require only a few hours of additional observations, preferably from another telescope or at least at a different wavelength. Establishing UCF-1.02 as a planet (to be called GJ 436d) would likely require an extended observing campaign to constrain its period then successfully predict a transit. Finally, we confirmed the GJ 436b 4.5-{\\micron} results presented by \\citet{Stevenson2010} through an additional non-detection during secondary eclipse; however, we were unable to confirm the strong eclipse depth at 3.6 {\\microns} due to stellar activity. The current data still support a methane-deficient and carbon monoxide-rich dayside atmosphere." }, "1207/1207.3186_arXiv.txt": { "abstract": "We report on \\emph{Chandra} X-ray observations of possible-AGNs which have been correlated with Ultra-high Energy Cosmic Rays (UHECRs) observed by the Pierre Auger Collaboration. Combining our X-ray observations with optical observations, we conclude that one-third of the 21 Veron-Cetty Veron (VCV) galaxies correlating with UHECRs in the first Auger data-release are actually not AGNs. We review existing optical observations of the 20 VCV galaxies correlating with UHECRs in the second Auger data-release and determine that three of them are not AGNs and two are uncertain. Overall, of the 57 published UHECRs with $|b|>10^\\circ$, 22 or 23 correlate with true AGNs using the Auger correlation parameters. We also measured the X-ray luminosity of ESO139-G12 to complete the determination of the bolometric luminosities of AGNs correlating with UHECRs in the first data-set. Apart from two candidate sources which require further observation, we determined bolometric luminosities for the candidate galaxies of the second dataset. We find that only two of the total of 69 published UHECRs correlate with AGNs (IC5135 and IC4329a) which are powerful enough in their steady-state to accelerate protons to the observed energies of their correlated UHECRs. The GZK expectation is that $\\sim 45$\\% of the sources of UHECRs above 60 EeV should be contained within the $z<0.018$ volume defined by the Auger scan analysis, so an observed level of 30-50\\% correlation with weak AGNs is compatible with the suggestion that AGNs experience transient high-luminosity states during which they accelerate UHECRs. ", "introduction": "Identifying the sources of ultra-high energy cosmic rays (UHECRs) is one of the major outstanding goals in astrophysics. Progress has proven difficult due to the rarity of high energy events and because the deflections of the cosmic rays as they travel through Galactic and extragalactic magnetic fields mean that the cosmic ray arrival directions do not point back to their origins. The very most energetic cosmic rays ($E \\gtrsim 6\\times10^{19}$ eV) have attracted attention as a promising path forward. The energy loss due to the GZK mechanism means that CRs with such energies typically have traveled about 100 Mpc or less, significantly limiting the possible sources for such high energy particles. Furthermore, at these energies magnetic deflections of protons may be small enough to allow the identification of the progenitor type based on statistical associations. By 2007, the Pierre Auger Collaboration had compiled enough events to begin drawing statistical correlations between UHECR events and possible sources, and reported a strong correlation between UHECRs and the nearby galaxies listed in the \\citet{VCV} Catalog of Quasars and Active Galactic Nuclei ($12^{th}$ Ed.) \\citep{augerScience07,augerLongAGN}. The scan procedure that was used steps through UHECR energy threshold, maximum angular separation, and maximum VCV galaxy redshift to find the values giving the lowest chance probability compared to an isotropic distribution, then evaluates the likelihood of such a correlation occurring by chance, by performing the same analysis on many isotropic datasets. In the following, the term ``correlated\" referring to a galaxy with respect to a UHECR simply means that the given galaxy falls within the angular and redshift limits Auger proposed based on applying the scan method to the original dataset using the VCV galaxy catalog. Of the 27 cosmic rays with energies above 57 $\\times 10^{18}$ eV that Auger detected before Aug. 31, 2007, twenty correlate within $3.2^{\\circ}$ with VCV galaxies having $z \\leq 0.018$, about 75 Mpc \\citep{augerLongAGN}. There are 21 VCV galaxies correlated with these 20 UHECRs. (More than one galaxy can be correlated with a UHECR and vice versa.) Galaxy catalogs such as VCV are incomplete in the Galactic Plane, so a better comparison is obtained by restricting to $|b|>10^{\\circ}$ where the VCV catalog is more complete. With this restriction there are 22 UHECRs of which 19 UHECRs correlate with 20 galaxies \\citep{zfg09}. This is a much higher correlation than would be expected by chance from an isotropic source distribution, and higher than can be explained by nearby galaxy clustering alone \\citep{zbf10}. More recent Auger data continue to show a significant, albeit less strong, correlation with VCV galaxies \\citep{augerUpdate2010}. An important question is whether the observed correlation with VCV galaxies implies that AGNs are the sources of some or all UHECRs. The VCV catalogue is a list of AGN \\emph{candidates}, so first it must be established whether the correlating galaxies (i.e., the VCV galaxies within 3.2$^\\circ$ of a UHECR) are actually AGNs. \\citet{zfg09} looked at existing observations of the galaxies in VCV that were correlated with the first set of UHECRs and found that only 14 of the 21 galaxies show unambiguous evidence of AGN activity. Three show no signs of AGN activity, while the other four have ambiguous optical spectra. A widely used technique for optical identification of AGNs, so-called BPT line ratios, compares the relative strengths of diagnostic spectral lines \\citep{bpt1981}. The BPT line ratios of the 4 ambiguous galaxies fall within the ranges classified by \\citet{Kauffmann2003} as AGNs but they do not fall within the line-ratio ranges adopted by \\citet{KewleyRatio} as indicative of an active nucleus. AGN activity can often be obscured in the optical bands by dust obscuration, and the Kewley test excludes some known AGNs. Additionally, as many as half of AGNs that are selected based on radio or X-ray properties would not be identified by looking only at their BPT line ratios \\citep{reviglioHelfand06}. Thus, determining whether these 4 ambiguous VCV galaxies have active nuclei requires observations outside of optical wavelengths. The VCV galaxies are not the only population of nearby galaxies found to correlate with the Auger UHECRs. \\citet{bzf10} performed an independent scan analysis between the Auger galaxies and nearby Luminous IR galaxies in the PSCz catalog and found an excess correlation. Restricting to $|b|>10^\\circ$ where PSCz is complete, 13 galaxies of $L_{IR} > 10^{10.5} L_{sun}$ correlate within $2.1^{\\circ}$ with one or more of the 22 Auger UHECR events. Some LIRGs contain an active nucleus with dust absorbing the AGN radiation and re-emitting it thermally, giving rise to large IR luminosities. This raises the question of whether the LIRGs found to correlate may actually be AGNs. \\citet{bzf10} showed that 6 of the 13 correlating IR galaxies are also in VCV and 5 of these do in fact host AGNs\\footnote{The one which is not a confirmed AGN falls within the VCV sample to be tested.}. The other 7 correlating IR galaxies lacked the observations needed to differentiate between star-formation and obscured nuclear activity. A key distinguishing feature of active galaxies is that accretion-driven radiation produces a broad spectrum extending from the IR to the X-ray, known as the broadband continuum. Normal and starburst galaxies, on the other hand, have broadened blackbody spectra due largely to stellar emission, which is peaked sharply in the UV/optical. This means the X-ray to optical flux ratios are considerably larger for AGNs than for normal/starburst galaxies. This has been seen in detailed spectroscopic studies of \\emph{Chandra} Deep Field sources \\citep{Barger2003}. If the active nucleus is obscured, the X-ray to optical ratio becomes even larger since dust absorbs UV and optical photons more readily than X-rays \\citep{Comastri2003}. This makes X-ray observations, particularly when combined with observations in the near-IR, optical or UV, a powerful tool for identifying all classes of AGNs \\citep{Maccacaro1988}. In this paper we use X-ray observations to determine whether the 4 ambiguous VCV galaxies, and the 7 indeterminate PSCz galaxies, have active nuclei. We observed ten of these possible UHECR source galaxies using the \\emph{Chandra} X-ray satellite, while for IC 5169 we used data from a recent XMM-Newton observation. Another interesting question is whether AGNs that correlate with UHECRs have any characteristic features which distinguish them from the AGNs that do not seem to correlate with UHECRs. A particularly relevant property is the luminosity of the AGN, as this helps to constrain the possible cosmic ray acceleration mechanisms \\citep{fg09}. Previously, reliable luminosities have been established for all but one of the correlating AGNs for UHECRs in the first data-release\\citep{zfg09}, \\citep{bzf10}. We observed the remaining AGN, ESO 139-G12, in order to obtain the first robust estimate of its bolometric luminosity. \\begin{table*} \\begin{center} \\caption{Table listing source properties} \\setlength{\\tabcolsep}{0.04in} {\\small \\begin{tabular}{l p{2cm} p{1.8cm} p{1.8cm}p{1.8cm}ccp{1.5cm}} \\hline Source Name & Total Counts \\newline \\scriptsize{ (0.5-10 keV)} & Hard Counts \\newline \\scriptsize{(2-10 keV)} & Flux \\newline (Direct) \\scriptsize{ erg cm$^{-2}$ s$^{-1}$} & Flux \\scriptsize{(webPIMMS)} \\scriptsize {erg cm$^{-2}$ s$^{-1}$} & Rmag & $Log(f_x/f_R) $ & $L_{2-10}$ \\newline \\scriptsize {erg s$^{-1}$} \\\\ \\hline IC 5169& 17 & 6 & 3.3 E-14 & 1.4 E -14 & 12.5$^1$ & -3.0 & 7.4 E 39\\\\ NGC 7591 & 17 & 2 & 6.6 E -15 & 1.4 E - 14 & 12.4$^{3,b}$ & -3.4 & 8.6 E 39 \\\\ NGC 1204 & 13 & 4 & 1.3 E -14 & 1.2 E -14& 13.6$^4$ & -2.9 & 5.8 E 39\\\\ NGC 2989 & 5 & 2 &2.3 E -14 & 4.9 E -15 & 12.5$^{1,2}$ & -3.1 & 8.8 E 39\\\\ \\hline \\hline IC 4523& 6 & 5 & 6.5 E -14 & 6.5 E -15 & 13$^2$ & -2.5 & 3.8 E 40\\\\ ESO 270-G007 & 14 & 5 & 2.2 E -14 & -- & 13$^1$ & -3 & 8.4 E 39\\\\ IC 5186 & 5 & 0 & 5.8 E -15$^a$ & -- & 12.3$^1$ & -3.8 & 3.4 E 39\\\\ ESO 565-G006 & 21 & 8 & 5.5 E -14 & 2.4 E -14 & 12.7$^1$ & -2.7 & 3.2 E 40\\\\ NGC 7648 & 5 & 1 & 3.0 E -15 & -- & 12.2$^{3}$ & -4.4 & 9.7 E 38\\\\ 2MASX J1 754-60 & 3 & 1 & 2.7 E -15 & -- & -- & -- & 1.6 E 39\\\\ IC 5179 & 2685$^c$ & -- & -- & 9.1 E -14 & 11.4$^1$ & -2.98 & 2.5 E 40\\\\ \\hline \\hline ESO 139-G12 & 1070 & 666 & 4.7 E -12$^d$ & -- & 12.9$^{1,2}$ & -0.2 & 2.9 E 43\\\\ \\hline \\end{tabular}} \\label{SourceProp} \\tablecomments{Fluxes are for 2-10 keV. The first group of four rows are the correlating galaxies found in VCV. The following six are those from PSCz. The last row of this group, IC 5179, was observed by XMM-Newton. Only the direct flux determination method is used for diffuse sources. The final row, ESO 139-G12, is a known AGN whose X-ray luminosity had not previously been measured. The 2-10 keV luminosity, $L_{2-10}$, is derived from the direct flux measurement; it is the total X-ray luminosity in the central region and thus an upper limit on X-ray luminosity of a possible AGN, a criterion for the ability to accelerate UHECRs. \\\\ $(a)$ No counts above 2 keV, flux given is for 0.5 - 2 keV band. $(b)$ Johnson magnitude rather than Cousins; this does not affect the result. $(c)$ 25ksec XMM-Newton observation; counts are from 0.5 - 12 keV. $(d)$ Derived from fit, see Figure 1.\\\\ } \\tablerefs{(1) \\cite{ESOCAT1989}, (2) \\cite{HIPASS}, (3) \\cite{Vaucouleurs}, (4) \\cite{6dFsurvey}} \\end{center} \\end{table*} ", "conclusions": "These \\emph{Chandra} observations complete the task of determining which galaxies found to correlate with UHECRs by \\citet{augerScience07} and \\citet{bzf10} have active nuclei, and of determining the bolometric luminosities of the AGNs correlating with UHECRs in the first data-release. The X-ray fluxes of the ten unclassified correlating galaxies studied here all fall within the range typical of normal and starburst galaxies of the same optical magnitude, hence we find no evidence of activity in any of them. Of course, it is impossible to rule out the possibility of very low luminosity nuclear activity which is energetically dominated by the host galaxy, but such weakly active AGNs are not well-motivated candidates for UHECR acceleration anyway \\citep{fg09}. The result, then, is that only 14 out of the original 27 Auger UHECRs (13 out of 22 UHECRs with $|b| \\geq 10^{\\circ}$) correlate with an actual AGN, using the Auger scan parameters to correlate UHECRs with VCV galaxies but discarding candidate sources which are not in fact AGNs. Of the 13 highly luminous IR galaxies ($L_{\\rm IR} \\geq 10^{11.5} \\, L_\\odot$) found by \\citet{bzf10} to correlate within $2.1^\\circ$ with one (or more) of the 22 UHECRs with $|b| \\geq 10^{\\circ}$, we find that only five are also AGNs.\\footnote{We could not run this correlation for the full 27 UHECRs because IRAS does not observe within the Galactic plane.} Thus five of the 27 UHECRs, and one of the 22 UHECRs with $|b| \\geq 10^{\\circ}$, have neither an AGN nor a LIRG within 3.2$^\\circ$ or $2.1^\\circ$ respectively. From the second Auger data set \\citep{augerUpdate2010}, we find 10 UHECRs with $|b|> 10^\\circ$ correlate with confirmed AGNs, one does not, and one remains to be determined. Thus using the full dataset of UHECRs with $|b|>10^\\circ$, where Galactic extinction does not hide source candidates, $(13+10) / (22 + 35) = 0.40^{0.51}_{0.32}$ of the UHECRs correlate with a confirmed AGN using the correlation parameters proposed in the original Auger analysis \\citep{augerScience07}, where the upper and lower values are $1-\\sigma$ limits for Poisson statistics \\citep{gehrels}. It is interesting that when actual AGNs rather than VCV galaxies are used, the agreement between the correlation found in the early and later Auger datasets becomes closer. Using the correct galaxy attributions found here, the first and second data sets individually give correlations to confirmed AGNs of $13/22 = 0.59^{0.80}_{0.43}$ and $10/35 = 0.29^{0.41}_{0.20}$ which differ at the 1.07-$\\sigma$ level, compared to the individual datasets being 1.63-$\\sigma$ from the mean using the uncorrected VCV-attribution. If the last VCV galaxy is confirmed as an AGN, the correlation would rise to $(13+11) / (22 + 35) = 0.42^{0.53}_{0.34} $ of the UHECRs correlating with an AGN, and $11/35 = 0.31 ^{0.44}_{0.22}$ in the second dataset alone. Given the uncertainty in the fraction of VCV galaxies which are actually AGNs, and the fact that the scan parameters were determined prior to removing non-AGNs, it is difficult to assess the significance of the final correlation. Further complicating correlation studies is the incompleteness of source catalogs (especially within the Galactic plane), the composition uncertainty, and magnetic deflection that can obliterate the angular correlation between the CRs and their true source, particularly for UHECRs with charge $Z >1$. Indeed, the new Galactic magnetic field model of \\citet{jf12}, with a more general form for the field constrained by extensive, all sky RM and polarized synchrotron emission data, predicts that Galactic deflections are small in some portions of the sky but are large in others, even for protons. The one established AGN which we observed in order to determine its bolometric luminosity, ESO 139-G12, proves to have $L_{\\rm bol}$ comparable to most of the other 13 correlated AGNs examined in \\citet{zfg09}. Knowledge of $L_{\\rm bol}$ and the energies $E \\equiv E_{20} \\,10^{20}$eV of the correlated UHECRS (89 and 59 EeV) allows us to evaluate $\\lambda_{\\rm bol} \\equiv 10^{-45} L_{\\rm bol} \\, E_{20}^{-2}$, the figure-of-merit introduced by \\citet{zfg09} to quantify the ability of an AGN to accelerate a proton to the energy of the correlated UHECR. A value of $\\lambda_{\\rm bol} \\gtrsim 1$ satisfies the acceleration criterion for protons (c.f., \\citet{fg09}). With $\\lambda_{\\rm bol} \\sim 0.1$, ESO 139-G12 is thus marginal according to standard UHECR acceleration mechanisms for protons. GZK energy losses imply that (taking sources to be uniformly distributed in redshift) about $45$\\% of the sources of protons above 60 EeV should have $z<0.018$. Thus -- taking at face value the 30-50\\% correlation we find between UHECRs and confirmed (albeit weak) AGNs -- all UHECRs could have been produced in galaxies presently hosting (generally weak) AGNs, consistent with the picture of transient production of UHECRs via exceptional, powerful flares flares in very weakly or non-active galaxies \\citep{fg09}. Examples of tidally produced flares in quiescent galaxies have recently been discovered in archival SDSS data \\citep{vf11} and observed (apparently in blazar mode) by the \\emph{Swift} satellite \\citep{Burrows2011}, \\citep{Bloom2011}. The spectral energy distribution of the \\citet{vf11} flares are well-fit by a thin accretion disk model and the resultant bolometric luminosities amply satisfy the minimum luminosity requirement for UHECR acceleration in both cases (GRF, in preparation); depending on how rapidly the evidence of the accretion episode disappears, the host galaxy of a tidal disruption flare may or may not show evidence of weak AGN activity in later observations such that it would appear in a catalog such as VCV. It is also possible, given the uncertainties, that other candidate sources not associated preferentially with AGNs may be responsible for some, most, or all UHECRs. Indeed, other studies have found correlations between the arrival directions of the Auger UHECRs and a variety of extragalactic sources: HI galaxies (Ghisellini et al. 2008), AGNs from the Fermi catalog (Nemmen et al. 2010), Swift-BAT AGNs and 2MRS galaxies \\citep{augerUpdate2010}." }, "1207/1207.4590_arXiv.txt": { "abstract": "We discuss the properties of spiral arms in a N-body simulation of a barred galaxy and present evidence that these are manifold-driven. The strongest evidence comes from following the trajectories of individual particles. Indeed, these move along the arms while spreading out a little. In the neighbourhood of the Lagrangian points they follow a variety of paths, as expected by manifold-driven trajectories. Further evidence comes from the properties of the arms themselves, such as their shape and growth pattern. The shape of the manifold arms changes considerably with time, as expected from the changes in the bar strength and pattern speed. In particular, the radial extent of the arms increases with time, thus bringing about a considerable increase of the disc size, by as much as ~50\\% in about a Gyr. ", "introduction": "\\indent In a series of papers (\\citealt{RomeroGMAG06}, Paper I; \\citealt{RomeroGAMG07}, Paper II; \\citealt{AthaRGM09}, Paper III; \\citealt{AthaRGBM09}, Paper IV; \\citealt{AthaRGBM10}, Paper V), we proposed a theory to explain the formation and properties of spirals and inner and outer rings in barred galaxies. According to it, the backbone of these structures are a bunch of orbits guided and confined by the invariant manifolds associated with the periodic orbits around the saddle points of the potential in the frame of reference co-rotating with the bar. We call our theory manifold theory, or manifold flux-tube theory. Some of the introductory dynamics necessary to follow our work is summarised in \\cite[][Sect. 3.3.2]{Binney.Tremaine.08}. Manifolds and the orbits they guide are described and explained in Paper I, while a relatively lengthy summary, avoiding equations, can be found in Sect. 2 of Paper III. Here we analyse a N-body simulation of an evolving barred galaxy and present evidence that its spiral arms are manifold-driven. In Sect.~\\ref{sec:theory} we give a brief theoretical reminder and describe relevant manifold shapes in simple analytical potentials. In Sect.~\\ref{sec:simul} and \\ref{sec:results} we present the simulation and our results and make comparisons with our theoretical predictions. Further discussion and conclusions are given in Sect.~\\ref{sec:discussion}. ", "conclusions": "\\label{sec:discussion} \\indent In this letter we followed the formation of spiral structure in a N-body simulation of a barred disc galaxy. We witness the formation of a short-lived and recurrent spiral structure, two episodes of which last over roughly 1 Gyr. In between these two episodes the spiral never quite vanishes, but its amplitude decreases considerably. These spirals have the same properties as the manifold-spirals discussed in Papers I to V. We also followed the trajectories of particles in the arms and found that, contrary to what would have been expected for density wave spirals, they do not cross the arms but they move along them, starting off from the vicinity of one of the Lagrangian saddle points ($L_1$ or $L_2$) in the direction of the other one. Thus their trajectories outline the arms. In fact the whole of the trajectory contributes to the spiral structure and these spirals can be called flux-tube manifold spirals. There are several more signs, tell-tale of a manifold origin. For example, particles which outline the spiral have their origin in the outer part of the bar and join the spiral via the vicinity of a point whose location is compatible with that of the saddle Lagrangian points ($L_1$ or $L_2$). Furthermore, they loop in the vicinity of that point before they follow the arm. Another tell-tale sign of the manifold origin of the arm is that in the vicinity of each Lagrangian point the particles split into two groups, the same in simulations and in simple orbital calculations. From all the above it becomes clear that we are witnessing manifold-driven spirals in our simulation. This -- together with the good agreement between manifold-driven and observed rings and spirals found in Papers IV and V and by \\cite{Martinez-Garcia.12} -- argues strongly that manifolds do play a role in spiral and ring formation. Manifolds, however, are not the only possible origin of such structures. Indeed, spirals have been witnessed also in other N-body simulations and interpreted in terms of other theories \\citep{Sellwood.Carlberg.84, DOnghia.VH.12, Grand.KC.12, Sellwood.12}. We already included a note of caution to this avail in Paper V. Our comparison between N-body and theoretical manifolds must necessarily stay qualitative. As any orbital structure work, the manifold theory relies on a few simplifications, the most important of which is that the potential is time-independent in the frame of reference co-rotating with the bar. On the other hand, in simulations and in real galaxies the potential evolves with time, due to redistribution of angular momentum via the resonances \\citep[e.g.][]{Lynden-Bell.Kalnajs.72, Weinberg.85, Athanassoula.02, Athanassoula.03}. However, if this evolution is not too fast, it should not present a problem for our flux-tube manifold theory where the arm is constituted by the whole flow of material guided by the manifolds\\footnote{This is not the case for the theory presented in e.g. \\cite{Voglis.TE.06}, or \\cite{Harsoula.KC.11} which considers the loci of the apsidal manifolds sections. It thus requires a quasi-stationary, non-evolving potential.} and will only lead to a change of the manifold properties with time. Any secular change in the potential is expected to lead to a secular evolution of the manifold properties. Furthermore, our simulations show that, even when the rate of change is considerable, particle trajectories are still guided by manifolds, keep their characteristic signatures and obey the selection rules of permissible paths defined by them. Nevertheless, their shape and extent change as expected, while particles can get trapped or untrapped by these manifolds. Such orbits as described here have also been witnessed as responses to applied analytical potentials \\citep[e.g.][Papers III \\& V]{Danby.65, Patsis.06}, although links to manifolds were not necessarily made. To our knowledge, however, this is the first time that the manifold theory is tested in a realistic self-consistent N-body simulation by following the motion of individual particles\\footnote{ A previous attempt \\citep{Tsoutsis.EV.08} used an unrealistic simulation, with no separate disc and halo components, just a single cylindrical-shaped component, with a vertical to horizontal size ratio of $\\sim$ 0.3 and a velocity dispersion larger than 100 km/sec, i.e. properties very different from those of a spiral galaxy.}. Nevertheless, the simulation discussed here is in no way unique; manifold-driven spirals can be seen in a large number of our simulations, some of which include gas, and will be discussed elsewhere. Comparing the disc extent between times 0.5 and 1.5 Gyrs we see a considerable change, of the order of 50\\%, due to the spirals. This shows that the bar and the associated manifolds can drive the overall evolution of the disc significantly. We will explore this process further in future work." }, "1207/1207.0716_arXiv.txt": { "abstract": "\\textcolor{black}{We consider a binary system made of self-gravitating bodies embedded in a constant and uniform external field $\\bds g$}. We analytically work out several \\textcolor{black}{orbital} effects \\textcolor{black}{induced by a putative} violation of the Strong Equivalence Principle (SEP) \\textcolor{black}{due to $\\bds g$}. \\textcolor{black}{In our calculation,} we do not \\textcolor{black}{assume} $e\\sim 0,$ \\textcolor{black}{where $e$ is the binary's orbit eccentricity}. Moreover, \\textcolor{black}{we do not a priori choose any specific preferred spatial orientation for the fixed direction of $\\bds g$.} Our results do not depend on any particular SEP-violating theoretical scheme. \\textcolor{black}{They} can be applied to general astronomical and astrophysical \\textcolor{black}{binary systems immersed in an external constant and uniform polarizing field}. ", "introduction": "If bodies with non-negligible gravitational binding self-energies, like astronomical and astrophysical objects (planets and their natural satellites, main-sequence stars, white dwarfs, neutron stars), move with different accelerations in a given external field, the so-called Strong Equivalence Principle (SEP) would be violated. While the General Theory of Relativity (GTR) assumes the validity of SEP, it is generally violated in alternative theories of gravity; see \\cite{Dam012,Freire012} and references therein. Potential SEP violations in the Earth-Moon system freely falling in the external gravitational field of the Sun, theoretically predicted by Nordtvedt long ago \\cite{nord1,nord2}, are currently searched for with the Lunar Laser Ranging (LLR) technique \\cite{nord3,Tury07,LLR} in the framework of the Parametrized Post-Newtonian (PPN) formalism. There are projects to test SEP with the Martian moon Phobos and the Planetary Laser Ranging (PLR) technique \\cite{Tury010} as well. As shown by Damour and Sch\\\"{a}fer \\cite{Dam91}, the acceleration due to SEP violation occurring for a two-body system in an external gravitational field $\\bds g$ like, e.g., the Galactic field at the location of a binary pulsar, is, to leading order, \\eqi\\bds A_{\\textcolor{black}{g}}=\\ton{\\Delta_{\\rm p}-\\Delta_{\\rm c}}\\bds g,\\lb{Asep}\\eqf where p and c denote the pulsar and its less compact companion; for a generic body with non-negligible gravitational self-energy\\footnote{For a non-relativistic spherical body of uniform density, it can be cast \\cite{bind1,bind2} $\\mathcal{E}_{\\rm grav}=3M^2 G/5R$, where $G$ is the Newtonian constant of gravitation, $M$ is the body's mass, and $R$ is its radius. For highly relativistic neutron stars, see, e.g., \\cite{latt}.} $\\mathcal{E}_{\\rm grav}$, $\\Delta$ accounts for the SEP-violating difference between the inertial mass $m_i$ and the gravitational mass $m_g$: at first post-Newtonian level, it is \\eqi \\Delta\\doteq\\rp{m_g}{m_i} -1 = \\eta_1\\varepsilon_{\\rm grav},\\ \\varepsilon_{\\rm grav}\\doteq \\rp{\\mathcal{E}_{\\rm grav}}{m_i c^2}, \\eqf where $c$ is the speed of light in vacuum. The external field $\\bds g$ in \\rfr{Asep} is assumed here to be constant and uniform, so that it induces a gravitational analog of the Stark effect \\cite{Dam91}. To maximize the size of potential SEP violations, one should look at compact objects like neutron stars: indeed, it is \\eqi\\left|\\varepsilon_{\\rm grav}^{\\leftmoon}\\right|\\sim 1.9\\times 10^{-11},\\ \\left|\\varepsilon_{\\rm grav}^{\\oplus}\\right|\\sim 4.2\\times 10^{-10},\\ \\left|\\varepsilon_{\\rm grav}^{\\odot}\\right|\\sim 1.3\\times 10^{-6},\\ \\left|\\varepsilon_{\\rm grav}^{\\rm NS}\\right|\\sim 10^{-1}.\\eqf The SEP-violating parameter $\\eta_1$ can be parameterized in various ways depending on the theoretical framework adopted; see \\cite{Dam012,Freire012} for recent overviews. As far as the Earth-Moon system and the PPN framework is concerned, latest published bounds on the Nordtvedt parameter are \\begin{align} \\eta_1 & \\equiv \\eta_{\\rm N}=(4.0\\pm 4.3)\\times 10^{-4}\\ \\cite{Tury07}, \\\\ \\nonumber \\\\ \\eta_1 & \\equiv \\eta_{\\rm N}=(-0.6\\pm 5.2)\\times 10^{-4}\\ \\cite{LLR}. \\end{align} Williams et al. in a recent analysis \\cite{Williams012} released \\eqi \\eta_1\\equiv \\eta_{\\rm N}=(-1.8\\pm 2.9)\\times 10^{-4}. \\eqf In this paper, we deal with the orbital effects due to \\rfr{Asep} in a uniform way by obtaining explicit and transparent analytical expressions valid for the main observable quantities routinely used in empirical studies. In Section \\ref{kepelems} we work out the SEP\\textcolor{black}{-violating} rates of change of the osculating Keplerian orbital elements, which are commonly used in solar system and binaries studies. The SEP\\textcolor{black}{-violating} shift of the projection of the binary's orbit onto the line-of-sight, which is the basic observable in pulsar timing, is treated in Section \\ref{los}. In Section \\ref{radvel} we calculate the SEP\\textcolor{black}{-violating} perturbation of the radial velocity, which is one of the standard observable in spectroscopic studies of binaries: in systems suitable for testing SEP, the pulsar's companion is a white dwarf or a main sequence star for which radial velocity curves may be obtained as well. The primary-to-companion range and range-rate SEP\\textcolor{black}{-violating} perturbations, potentially occurring in systems like the Earth-Moon one, are calculated in Section \\ref{rrate}. Section \\ref{discuss} is devoted to summarizing and discussing our results. \\textcolor{black}{W}e will not resort to a-priori simplifying assumptions about the binary's orbital eccentricity. Moreover, we will not restrict ourselves to any specific spatial orientation for $\\bds g$. Our results are, thus, quite general. They can be used for better understanding and interpreting present and future SEP experiments, hopefully helping in designing new tests as well. Finally, we remark that our calculations are model-independent in the sense that they do not depend on any specific theoretical mechanism yielding SEP violations. Thus, they may be useful when different concurring SEP\\textcolor{black}{-violating} scenarios are considered like, e.g., MOND and its External Field Effect \\cite{Milg09,Blan011}, Galileon-based theories with the Vainshtein mechanism \\cite{Gal012,Ior012}, and variations of fundamental coupling constants as well \\cite{Dam011,Ior011,IorG}. Moreover, there are also other effects, like Lorentz symmetry violations \\cite{BailLV}, which affect the same SEP\\textcolor{black}{-violating} peculiar observables like the eccentricity \\cite{IorLV}. \\textcolor{black}{For a seminal paper on Lorentz-violating orbital effects in binaries, see \\cite{1992PhRvD..46.4128D}.} ", "conclusions": "\\lb{discuss} We worked out \\textcolor{black}{the first time derivatives of some orbital effects} induced by a Stark-like SEP violation in a binary system made of two self-gravitating bodies immersed in an external polarizing field $\\bds g$, assumed constant and uniform with respect to the characteristic temporal and spatial scales of the binary. We provided the reader with full analytical expressions for the SEP\\textcolor{black}{-violating} rates of change of all the six osculating Keplerian orbital elements (Section \\ref{kepelems}), for the projection of the binary's orbit along the line-of-sight (Section \\ref{los}) and its time derivative (Section \\ref{radvel}), and for the range and range-rate (Section \\ref{rrate}). We did not make any a-priori assumption about the orientation of $\\bds g$ in space. We did not make simplifying assumptions about the orbital geometry of the binary. It is important to have at disposal explicit expressions of the different relevant SEP\\textcolor{black}{-violating} effects since each one has a peculiar temporal pattern which may be helpful in separating it from other possible competing signatures. For example, PSR B1620-26 \\cite{psr1,psr2} is a pulsar-white dwarf binary in a relatively wide orbit ($P_{\\rm b}\\textcolor{black}{= 16540653(6)\\ {\\rm s}}\\sim 2\\times 10^2$ d \\textcolor{black}{\\cite{1996ASPC..105..525A}}) orbited by quite distant circumbinary planet-like companion\\footnote{For exoplanets, see \\cite{exop1,exop2,exop3}.} ($P^{'}_{\\rm b}\\sim 3\\times 10^4$ d). For a discussion of the rate of change of the eccentricity in binaries hosting compact objects, see, e.g., \\cite{Freire012}. Our results are important also to accurately asses the overall uncertainty which can be obtained when constraints on $\\Delta$ are inferred by comparing them with the corresponding empirically determined quantities. Indeed, taking into account only the accuracy with which an observable like, say, $de/dt$ can be determined is, in principle, not enough; a propagation of the errors affecting all the parameters \\textcolor{black}{such as $g$ \\cite{gala1,gala2,MW}}, and its orientation $\\hat{\\bds{g}}$, entering the corresponding theoretical prediction must be done as well to correctly infer the total, systematic uncertainty in $\\Delta$. Our expressions are also useful in designing suitable SEP tests and in better interpreting present and future experiments, \\textcolor{black}{not necessarily limited to binary pulsars}, especially when different competing theoretical mechanisms yielding SEP violations are considered. Indeed, we found that the ratio of the node and inclination SEP\\textcolor{black}{-violating} precessions depends only on the inclination itself and on the pericenter." }, "1207/1207.7070_arXiv.txt": { "abstract": "{The majority of all galaxies reside in groups of fewer than 50 member galaxies. These groups are distributed in various large-scale environments from voids to superclusters. } {The evolution of galaxies is affected by the environment in which they reside. Our aim is to study the effects of the local group scale and the supercluster scale environments on galaxy evolution. } { {We use a luminosity-density field to determine the density of the large-scale environment of galaxies in groups of various richnesses. We calculate the fractions of different types of galaxies in groups with richnesses of up to 50 member galaxies and in different large-scale environments from voids to superclusters. } } {The fraction of passive elliptical galaxies rises and the fraction of star-forming spiral galaxies declines when the richness of a group of galaxies rises from two to approximately ten galaxies. On large scales, passive elliptical galaxies become more numerous than star-forming spirals when the environmental density grows to values typical of superclusters. The large-scale environment affects the level of these fractions in groups: galaxies in equally rich groups are more likely to be elliptical in supercluster environments than at lower densities. {The crossing point, where the number of passive and star-forming galaxies is equal, occurs in superclusters in groups that are of lower richness than in voids. Galaxies in low-density environments need to occupy richer groups to evolve from star-forming to passive than galaxies in high-density environments.} Groups in superclusters are on average more luminous than groups in large-scale environments of lower density. These results imply that the large-scale environment affects the properties of galaxies and groups.} { {Our results suggest that the evolution of galaxies is affected by both, the group in which the galaxy resides and its large-scale environment. Galaxies in lower-density regions develop later than galaxies in similar mass groups in high-density environments.}} ", "introduction": "Galaxies form groups and clusters of various sizes. These groups and clusters are distributed in different environments on the large scale. Regions of high galaxy density are superclusters, which are surrounded by filaments and low-density void areas. {The properties of galaxies are affected by their cluster- or group-scale environment. In group environments, galaxies experience baryonic processes, such as ram pressure stripping \\citep{Gunn1972}, viscous stripping, and strangulation \\citep{Larson1980}, and gravitational effects, such as galaxy tidal interactions with other galaxies and the cluster potential \\citep{Moore1996} and galaxy merging. These local processes are believed to modify galaxy morphologies from spiral galaxies to gas poor and spheroidal ones \\citep{Gunn1972, Dressler1980, Postman1984}. } According to the morphology-density relation, elliptical galaxies are more concentrated at the centers of clusters than spiral galaxies {\\citep{Dressler1980,Postman1984}}. Similarly, galaxies in higher density environments are more luminous \\citep{Hamilton1988} and have lower star-formation rates \\citep{Gomez2003}. {The dependence of either the star-formation rate or color on the environment is stronger than the dependency between morphology and the environment \\citep{Kauffmann2004,Blanton2005}. According to \\citet{Baldry2006}, the environmental dependence is at least as important as stellar mass in determining the fraction of red galaxies in a population.} Finally, different types of active galactic nuclei (AGNs) appear in different local environments \\citep{Hickox2009}. In addition to the local group or cluster scale, the properties of galaxies also depend on their environment on large scales. The large-scale morphology-density relation was first found by \\citet{Einasto1987}. According to \\citet{Balogh2004}, the fraction of red galaxies is higher where the surface density of galaxies is higher. \\citet{Porter2008} found that the star-formation rate of galaxies depends on their location in large-scale filaments. \\citet{Skibba2009} found significant color-environment correlations on the 10\\,$h^{-1}$Mpc scale in both spiral and elliptical galaxies. \\citet{Tempel2011} derived luminosity functions of spiral and elliptical galaxies in different large-scale environments. They found that the luminosity function of elliptical galaxies depends strongly on environment, while for spiral galaxies the luminosity function is almost independent of the {large-scale} environment. The large-scale environments of AGN were studied by \\citet{Lietzen2011}. In the same way as on smaller scales, radio galaxies favor high-density environments, while radio-quiet quasars and Seyfert galaxies are mostly located in low-density regions. In high-density large-scale environments, groups and clusters of galaxies tend to be larger and more massive than in low-density environments \\citep{Einasto2005}. {According to \\citet{Einasto2012a}, isolated clusters are also poorer than supercluster members.} Because of this, group-scale effects are also present when studying the large-scale environments. {The effects on galaxy properties on different scales up to tens of Mpc can be distinguished using a luminosity-density method \\citep{Einasto2003}. Group and supercluster scale environments of galaxies were studied together by \\citet{Einasto2008}.} They found that in the outskirt regions of superclusters, rich groups contain more late-type galaxies than in the supercluster cores, where the rich groups are populated mostly by early-type galaxies. \\citet{Einasto2007} found that in the high-density cores of rich superclusters there is an excess of early-type galaxies in groups and clusters, as well as among galaxies that do not belong to any group. {Some studies have used the number-density of galaxies to study the large-scale environments of galaxies. In these studies, the large scale usually refers to scales of a few Mpc. \\citet{Zandivarez2011} found that the Schechter parameters of the luminosity functions of galaxies in groups in high-density environments do not depend on the mass of the group, while in low-density regions some dependence does occur. On the other hand, \\citet{Blanton2007} found that on a few Mpc scales the environment is only weakly related to the colors of galaxies, while the small-scale environment matters more. \\citet{Wilman2010} found no correlation on scales of $\\sim$1\\,Mpc and even an anticorrelation on scales of 2--3\\,Mpc in the fraction of red galaxies. } Of the different ways of determining the environment of a galaxy, some are sensitive to the local scales, which in turn depend on the size of the dark matter halo surrounding the galaxy. Other methods are more sensitive to large scales, which represent the environment in the supercluster-void network {\\citep{Haas2011,Muldrew2012}}. In this paper, we use spectroscopic galaxy and group catalogs based on the Sloan Digital Sky Survey (SDSS) to study the environments of galaxies in groups on both local and large scales. As a measure of the local-scale environment, we use the richness of the group, and for the large-scale environment a luminosity-density field smoothed to scales typical of superclusters. Our goal is to distinguish between the effects that different scales of environment have on galaxies. {The paper is composed as follows: in Sect. \\ref{Data}, we present the data and describe how the group catalog and the large-scale luminosity-density field were constructed. We also describe the galaxy classification criteria. In Sect. \\ref{Results}, we present our results on the environments of galaxies on both the group and large scales. In Sect. \\ref{Discussion}, we compare our results to previous studies and discuss the possible implications of our results for galaxy evolution. } {Throughout this paper we assume a cosmological model with a total matter density $\\Omega_{\\mbox{m}}=0.27$, dark energy density $\\Omega_\\Lambda=0.73$, and Hubble constant $H_0=100h$\\,km\\,s$^{-1}$Mpc$^{-1}$ \\citep{Komatsu:11}. } ", "conclusions": "We have studied the environments of galaxies of the SDSS DR8 at distances of between 120 and 340\\,$h^{-1}$Mpc. We divided the galaxies into different samples based on their spectral properties and morphology. We measured the large-scale environment using a luminosity-density field and the group-scale environment with the group richness. Our main results are the following: \\begin{itemize} \\item{Passive galaxies are located in denser large-scale environments than star-forming galaxies. There is no significant difference between passive elliptical and passive spiral galaxies. } \\item{The fraction of galaxies that are star-forming declines and the fraction of passive galaxies rises as the richness of a group rises from one to approximately ten galaxies. In groups with richnesses of between 20 and 50 galaxies, the fractions of galaxies of different types do not depend on group richness. } \\item{The group richness at which passive galaxies become more numerous than star-forming galaxies depends on the large-scale environment. When the large-scale density grows from levels typical of voids to supercluster regions, the group richness where most galaxies become passive declines. In addition, the fractions of star-forming and passive galaxies in rich groups depend on the large-scale environment: {in voids, the fractions of passive and star-forming galaxies are approximately equal to each other in groups with more than ten galaxies, while in superclusters the fractions of passive galaxies are considerably larger than those of star-forming galaxies.}} \\item{Equally rich groups are more luminous in supercluster regions than in voids.} \\item{The fraction of galaxies with an AGN does not depend strongly on either group richness or the large-scale density.} \\end{itemize} We conclude that galaxy evolution is affected by both the group where the galaxy resides and its large-scale environment. In the future we plan to use numerical simulations to study the plausible baryonic and gravitational processes that determine galaxy evolution in groups of galaxies in different large-scale environments. Another invaluable approach will be to study in more detail the properties that are known to characterise groups e.g. their \\ion{H}{i} content, stellar mass, and X-ray properties." }, "1207/1207.0699_arXiv.txt": { "abstract": "{Cygnus A harbours the nearest powerful radio jet of an Fanaroff-Riley (FR) class II radio galaxy in a galaxy cluster where the interaction of the jet with the intracluster medium (ICM) can be studied in detail. We use a large set of \\emph{Chandra} archival data, \\emph{VLA} and new \\emph{LOFAR} observations to shed new light on the interaction of the jets with the ICM. We identify an X-ray cavity in the distribution of the X-ray emitting plasma in the region south of the Cyg A nucleus which has lower pressure than the surrounding medium. The \\emph{LOFAR} and \\emph{VLA} radio observations show that the cavity is filled with synchrotron emitting plasma. The spectral age and the buoyancy time of the cavity indicates an age at least as large as the current Cyg A jets and not much larger than twice this time. We suggest that this cavity was created in a previous active phase of Cyg A when the energy output of the Active Galactic Nucleus (AGN) was about two orders of magnitude less than today. ", "introduction": "AGN feedback in galaxy clusters has gained a lot of attention in recent studies, e.g.~\\cite{mcnamara}. Most of the AGN in clusters are FR I radio sources. In the local Universe Cygnus A is an exception featuring a powerful FR II radio jet in a galaxy cluster where very energetic interaction effects of the jets with the intracluster medium (ICM) can be studied in detail, e.g.~\\citet{carilli1996,clarke,kaiser,krause}. In this object we can observe jets of relativistic synchrotron emitting plasma plowing into the intracluster medium out to a distance of $\\sim$100~kpc from the nucleus where they end in hotspots clearly seen at radio wavelengths, (e.g.~\\citealt{alexander}) and X-ray images (e.g.~\\citealt{harris}). The internally supersonic jet plasma is thermalised at these hotspots, which advance only at about 2000~km/s due to the extreme density contrast (\\citealt{alexander_book}). The jet plasma flows away from the overpressured hotspots establishing a backflow. Together with the adjacent ICM, this region defines an overpressured bubble which drives a shock into the undisturbed ICM, first seen by~\\cite{carilli1988} in the centre of the Cyg A cluster. Many aspects of Cyg A have been studied in both radio and X-rays, e.g.~\\citet{carilli1994,wilson2000,wilson2006,smith,yaji}. Some progress has been made recently in modeling these jets with hydrodynamical simulations (\\citealt{krause03,krause}, see also~\\citealt{alexander_book}). The observed width of the radio lobes are only reproduced by very light jets, with jet densities of order $10^{-4}$ times the already tenuous surrounding X-ray plasma. This implies a low momentum flux and low hotspot advance speeds, comparable to the general advance speed of the bow shock driven into the surrounding ICM. Consequently, the bow shock assumes an almost spherical shape around the radio source with a wide region of shocked ICM adjacent to the radio lobes. In this paper we use most of the \\emph{Chandra} archival data on Cyg A to study in detail the interaction effects in the region near the AGN inside the radio lobes. In this region we discovered an X-ray cavity very similar to those found in FR I radio sources in clusters and investigated it in detail with the X-ray and radio data to unveil its physical properties and its origin, and the interpretation with a hydrodynamical simulation. For distance related quantities we adopt a flat $\\Lambda$-cosmology with $H_0$=70 km/s/Mpc, $\\Omega_M$=0.3. For the redshift of Cyg A of 0.0561, 1\\arcsec corresponds to 1.09 kpc. ", "conclusions": "We have combined most of the available \\emph{Chandra} X-ray data to construct a very detailed image of the ICM in the central region of Cyg A. The region around the AGN is rich in structure and one of the most striking features is an X-ray cavity south of the nucleus which is reminiscent of similar cavities in other cool cores of galaxy clusters. A comparison of the X-ray image with radio maps from \\emph{VLA} and \\emph{LOFAR} observations shows that the cavity is filled with relativistic plasma with a total enthalpy of $2.2\\times10^{59}$ erg. The most natural explanation in our opinion of the radio plasma filled cavity is an origin from previous activity of the AGN before the present radio lobes and shocked ICM region were created, more than about 30 Myrs ago. The appearance of the cavity and the X-ray morphology in the centre of Cyg A is well reproduced by our simulation of this scenario. In this case the power of the jets at this earlier epoch was about 100 times lower than today also consistent with the simulation results. Thus Cygnus A would have been classified at this epoch as FR I radio source. If our interpretation is correct we have the first evidence that an FR II radio source today was an FR I source at an earlier time." }, "1207/1207.4827_arXiv.txt": { "abstract": "{ The discovery of extensive air showers by Rossi, Schmeiser, Bothe, Kolh\\\"orster and Auger at the end of the 1930s, facilitated by the coincidence technique of Bothe and Rossi, led to fundamental contributions in the field of cosmic ray physics and laid the foundation for high-energy particle physics. Soon after World War II a cosmic ray group at MIT in the USA pioneered detailed investigations of air shower phenomena and their experimental skill laid the foundation for many of the methods and much of the instrumentation used today. Soon interests focussed on the highest energies requiring much larger detectors to be operated. The first detection of air fluorescence light by Japanese and US groups in the early 1970s marked an important experimental breakthrough towards this end as it allowed huge volumes of atmosphere to be monitored by optical telescopes. Radio observations of air showers, pioneered in the 1960s, are presently experiencing a renaissance and may revolutionise the field again. In the last 7 decades the research has seen many ups but also a few downs. However, the example of the Cygnus X-3 story demonstrated that even non-confirmable observations can have a huge impact by boosting new instrumentation to make discoveries and shape an entire scientific community. } ", "introduction": "\\label{sec:gen-overview} Towards the end of the 1930s it was recognised from studies of the effect of the geomagnetic field on cosmic rays that the energy spectrum of the primary particles, not identified as being proton-dominated until 1941, extended to at least 10~GeV. The discovery of extensive air showers in 1938, however, radically changed this situation with the highest energy being pushed up by about 5 orders of magnitude, probably the single largest advance to our knowledge of energy scales ever made. It is now known that the energy spectrum extends to beyond $10^{20}$~eV but it has taken over 60 years to consolidate this picture. In this section we trace the history of the discovery of extensive air showers, show how advances in experimental and theoretical techniques have led to improved understanding of them, and describe how some of the most recent work with contemporary instruments has provided important data on the energy spectrum, the mass composition and the arrival direction distribution of high-energy cosmic rays. These results are of astrophysical importance but additionally some aspects of the shower phenomenon promise to give new insights on hadronic physics at energies beyond that reached by the LHC. The flux of particles falls so rapidly with energy ($\\propto E^{-\\gamma}$ with $\\gamma \\sim 2.7$) that around $10^{14}$~eV it becomes impractical to make measurements of high precision directly: the number of events falling on a detector of a size that can be accommodated on a balloon or a space-craft is simply too small. However at this energy sufficient particles are produced in the atmosphere as secondaries to the incoming primary cosmic rays for some to reach mountain altitudes and, as the energy of the primary increases, even sea level. The transverse momentum acquired by secondary particles at production and the scattering which the shower electrons, in particular, undergo through interactions with the material of the atmosphere are such that the secondaries are spread over significant areas at the observational level. The phenomenon of the nearly-simultaneous arrival of many particles over a large area is called an Extensive Air Shower (EAS): at $10^{15}$~eV around $10^6$ particles cover approximately $10^4$~m$^2$ while at $10^{20}$~eV some $10^{11}$~particles are spread over about 10~km$^2$. It was quickly recognised that the phenomenon of the air shower offered the possibility of answering four major questions: \\begin{enumerate} \\item{{\\bf What particle physics can be learned from understanding air shower evolution?}} A detailed understanding of how an air shower develops is crucial to obtaining an estimate of the primary energy and to learning anything about the mass spectrum of the primary particles. It is worth recalling that when the shower phenomenon was first observed that, in addition to the proton, neutron, electron and positron, only the muon was known, so that a realistic understanding of shower development had to wait until the discovery of the charged pion and its decay chain in 1947 and of the neutral pion in 1950. Indeed, much early thinking was based on the hypothesis that showers were initiated by electrons and/or photons. Once it was recognised that the initiating particle was almost always a proton or a nucleus, the first steps in understanding the nuclear cascade focussed on such matters as whether a proton would lose all or only part of its energy in a nuclear collision and how many pions were radiated in such a collision. A combination of observations in air showers, made using Geiger counters and cloud chambers, of data from studies in nuclear emulsions and of early accelerator information was used to inform the debate. The issues of inelasticity (what fraction of the energy is lost by an incoming nucleon to pion production) and the multiplicity (the number of pions produced) are parameters which are still uncertain at most of the energies of interest. \\item{{\\bf What can be inferred from the arrival direction distributions of the high-energy particles?}} From the earliest years of discovery of cosmic rays there have been searches for directional anisotropies. Hess himself, from a balloon flight made during a solar eclipse in April 1912, i.e.\\ before his discovery flight in August of the same year, deduced that the Sun was not a major source \\citep{Hess:1912ui}. There are a few predictions of the level of anisotropy that might be expected. While there have always been speculations as to the sources, the fact that the primary particles are charged and therefore are deflected in the poorly-known galactic and intergalactic magnetic fields makes it difficult to identify them. One firm prediction was made very early on by Compton and Getting in 1935 \\citep{compton35} that cosmic rays should show an anisotropy because of the motion of the earth within the galaxy. Eventually it was realised that this idea would be testable only with cosmic rays undeflected by the solar wind (discovered much later) so measuring the Compton-Getting effect became a target for air shower experiments. However, as the velocity of the earth is only about 200~km\\,s$^{-1}$, the effect is $\\sim 0.1$\\,\\% and it has taken around 70 years for a convincing demonstration of its discovery. The search for point sources has been largely unsuccessful but one of the motivations for searching for rarer and rarer particles of higher and higher energy has been the expectation that anisotropy would eventually be found. \\item{{\\bf What is the energy spectrum of the primary cosmic rays?}} A power law distribution of cosmic rays was first described by E Fermi in 1949 \\citep{Fermi-49} but until 1966 there were no predictions as to the power law index or to further structures in the energy spectrum. Observations in 1959 had indicated a steepening % at around $3 \\cdot 10^{15}$~eV (the ``knee''), while in 1963 it was claimed from observations made with the first large shower array that the spectrum flattens just above $10^{18}$~eV. However not only were there no predictions of these features, interpretation of them remains controversial. By contrast the discovery of the 2.7~K cosmic background radiation in 1965 led, a year later, to the firm statement that if cosmic rays of energy above $\\sim 4 \\cdot 10^{19}$~eV exist they can come only from nearby sources. It took about 40 years to establish that there is indeed a steepening in the cosmic ray spectrum at about this energy but whether this is a cosmological effect or a consequence of a limit to which sources can accelerate particles is unclear: $4 \\cdot 10^{19}$~eV is within a factor of $\\sim 5$ of the highest energy event ever recorded. \\item{{\\bf What is the mass composition of the primary cosmic rays?}} One of the major tasks of the air shower physicist is to find the mass of the primary particles. This has proved extraordinarily difficult as even if the energy of the primary that produces an event is known, the uncertainties in the hadronic physics make it hard to separate protons from iron. Data from the LHC will surely help but above $10^{17}$~eV one has reached a regime where the centre-of-mass (cms) energies in the collisions are above what is accessible to man-made machines. Indeed it may be that in the coming decades the highest-energy cosmic rays provide a test bed for theories of hadronic interactions, mirroring the fact that cosmic ray physics was the place where particle physics was born in the 1930s. \\end{enumerate} In what follows we have chosen to emphasise the progress made since the 1940s towards answering these four questions through an examination of the development of different techniques, both experimental and analytical, introduced in the last 70 years. While new techniques have enabled air showers to be studied more effectively, it is remarkable how the essentials of what one seeks to measure were recognised by the pioneers in the 1940s and 1950s. Increasingly sophisticated equipment, operated on increasingly larger scales has been developed and has led to some answers to the key questions although many issues remain uncertain. Galbraith \\citep{Galbraith-58} and Cranshaw \\citep{Cranshaw-63} have written books in which details of early work, up to the end of the 1950s, are discussed in more detail than is possible below while in Hillas's classic book on Cosmic Rays \\citep{Hillas-72} there is an excellent discussion of some of the earliest papers in a context which includes fundamental ideas of cosmic rays physics, including shower physics. We now move on by reviewing the history of the discovery of the air shower phenomenon. ", "conclusions": "\\label{sec:Conclusions} In this year, 2012, the centennial of the discovery of cosmic rays will be celebrated all around the globe. The enormous progress that has been made during this period is directly linked to the invention of new experimental tools and instrumentation and could not have been made without the ideas and skills of some ingenious pioneers. Almost no nuclear and particle physics experiment could be done without making use of the coincidence technique but also triggering on rare events, and the construction of calorimeters, as other concepts, have been pioneered in cosmic ray experiments. To remain focussed on cosmic ray and air shower physics, we have omitted in this review the discoveries of new particles made by cosmic ray observations, including the positron, muons, pions, kaons, hyperons, and likely also charmed particles. This part of the history is discussed in \\citep{walter-epjh}. The cosmic energy spectrum has been measured in great detail over more than 32 decades in flux, making this observable unique in Nature. The spectrum initially thought to follow a pure power law distribution has exhibited more and more structure, starting with the discovery of the ``knee'' at about $4 \\cdot 10^{15}$\\,eV by Khristiansen's group at Moscow State University in 1959, followed by the observation of the ``ankle'', first hinted at by Linsley \\citep{Linsley-63a}, at Haverah Park, Akeno, and Fly's Eye in 1991 and the suppression at the GZK threshold in 2008 by the HiRes and Auger observatories. Very recently, a second knee caused by the heavy cosmic ray component has been reported by KASCADE-Grande and it is not unlikely that even more departures from a simple power-law distribution will be exhibited providing important clues about the origin of cosmic rays. Also, great detail about the primary mass could be extracted from the data with remarkable changes seen in the composition coinciding with the the position of the structures in the energy spectrum \\citep{Kampert-12}. The sky in cosmic rays is surprisingly isotropic up to the highest energies and is challenging our understanding of both cosmic ray propagation within the galactic and intergalactic environments and about their sources. Only at the highest energies are departures from isotropy are seen, but data suffer still from statistics. Particles at the upper end of the spectrum have such breath-taking energies, a hundred million times above that provided by the LHC accelerator, that the questions about how cosmic accelerators can boost particles to these energies, and about what is the nature of the particles themselves, are still open and of prime interest. The mystery of cosmic rays is nowadays tackled - and is perhaps going to be solved - by an interplay of sophisticated detectors for high-energy $\\gamma$-rays, charged cosmic rays and neutrinos. Moreover, plans for next generation experiments are being worked out and it is now realized that the true high-energy frontier in Nature provides unique opportunities to test particle and fundamental physics, such as of space-time, at its extreme. Further surprises by future cosmic ray observations are almost guaranteed." }, "1207/1207.3293_arXiv.txt": { "abstract": "We consider the case of a coupling in the dark cosmological sector, where a dark energy scalar field modifies the gravitational attraction between dark matter particles. We find that the strength of the coupling $\\beta$ is constrained using current Cosmic Microwave Background (CMB) data, including WMAP7 and SPT, to be less than 0.063 (0.11) at $68\\%$ ($95\\%$) confidence level. Further, we consider the additional effect of the CMB-lensing amplitude, curvature, effective number of relativistic species and massive neutrinos and show that the bound from current data on $\\beta$ is already strong enough to be rather stable with respect to any of these variables. The strongest effect is obtained when we allow for massive neutrinos, in which case the bound becomes slightly weaker, $\\beta < 0.084 (0.14)$. A larger value of the effective number of relativistic degrees of freedom favors larger couplings between dark matter and dark energy as well as values of the spectral index closer to $1$. Adding the present constraints on the Hubble constant, as well as from baryon acoustic oscillations and supernovae Ia, we find $\\beta< 0.050 (0.074)$. In this case we also find an interesting likelihood peak for $\\beta=0.041$ (still compatible with 0 at 1$\\sigma$). This peak comes mostly from a slight difference between the Hubble parameter HST result and the WMAP7+SPT best fit. Finally, we show that forecasts of Planck+SPT mock data can pin down the coupling to a precision of better than $1\\%$ and detect whether the marginal peak we find at small non zero coupling is a real effect. ", "introduction": "Cosmic Microwave Background (CMB) probes have recently broadened our knowledge of primordial acoustic oscillations to small angular scales, extending previous measurements of temperature power spectrum (Wilkinson Microwave Anisotropy Probe 7, \\cite{Komatsu2011}) up to $l \\sim 3000$ with first compelling evidence of CMB lensing from the South Pole Telescope (SPT, \\cite{k11}) and Atacama Cosmology Telescope (ACT, \\cite{act_2011}). The impact of small-scale CMB measurements and gravitational lensing on cosmology is relevant \\cite{lewis_challinor_2006} and can be used to constrain cosmological parameters and to address one of the major issues of present cosmology, that is to say the nature of dark energy \\citep{verde_spergel_2002, giovi_etal_2003, 2004PhRvD..70b3515A, acquaviva_baccigalupi_2006, hu_etal_2006, Sherwin:2011gv}. The simplest framework for dark energy models considers dark energy as a cosmological constant $\\Lambda$, contributing to about $74\\%$ of the total energy density in the universe and providing late time cosmic acceleration, while a Cold Dark Matter represents about $21\\%$ ($\\Lambda$CDM model). Though theoretically in good agreement with present observations, a cosmological constant is somewhat unpleaseantly affected by the coincidence and fine-tuning problems which seem unavoidable in such a framework. Many alternative models have been proposed, though it is fair to say that so far no one completely avoids these problems. Some encouraging arguments have been put forward in the framework of dynamical dark energy models, where a scalar field (quintessence or cosmon) rolls down a suitable potential \\cite{wetterich_1988, ratra_peebles_1988} possibly interacting with dark matter \\cite{amendola_2000, pettorino_baccigalupi_2008} or gravity \\cite{Matarrese:2004xa, Perrotta:2002sw} and therefore modifying the growth of structure. Usually, one of the features of such dynamical dark energy models is to have a non-negligible amount of dark energy at early times. The amount of early dark energy (\\emph{early} referring to the time of decoupling) influences CMB peaks in various ways and can be strongly constrained when including small scale measurements, as shown for instance in Refs. \\cite{calabrese_etal_2011, reichardt_etal_2011}. In this paper, we consider the case of coupled dark energy models, in which dark matter particles feel an interaction, additional to gravity, mediated by the dark energy scalar field. Such an interaction introduces effectively a coupling between the evolution of the dark energy scalar field and dark matter particles. In this sense, this class of models is both an example in which a non-negligible amount of early dark energy is present as well as a typical scenario of modified gravity theories. When seen in the Jordan frame, a coupling between matter and dark energy can be reformulated in terms of scalar-tensor theories (or $f(R)$ models). This is exactly true when the contribution of baryons is neglected. In the Einstein frame, it is common use to neglect a coupling to baryon within coupled dark energy models, and consider only dark energy - dark matter interactions. Alternatively, in the Jordan frame, scalar-tensor theories ($f(R)$ models) require some sort of screening mechanism (like chameleon \\cite{khoury_etal_2004, Hui:2009kc, 2012arXiv1204.3906U, 2011arXiv1102.5278D} or symmetrons \\cite{2011PhRvD..84j3521H}) that protects the dark energy scalar field and its mass within high density regions, so that local solar system constraints are satisfied. The strength of the coupling affects CMB in several ways, changing the amplitude, the position of the peaks as well as contributing to the Late Integrated Sachs-Wolfe (ISW) effect (manifest at large length scales) and to gravitational lensing (appearing at small length scales in the temperature spectrum). Moreover, the coupling is degenerate with the amount of cold dark matter $\\Omega_{c}$, the spectral index $n$, the Hubble parameter $H(z)$ (see \\cite{amendola_etal_2012} for a review). After recalling the effects of the coupling on CMB, we use a Monte Carlo analysis to constrain the coupling combining WMAP and SPT real data. Furthermore we extend our analysis to forecasting the constraints that Planck data are expected to put on the coupling parameter, combined with mock SPT data. This paper is organized as follows. In Section II we recall the main features of coupled dark energy (CDE) cosmologies. In section III we recall effects of the coupling on the CMB spectrum and describe the methods used, both with regard to the implementation of the numerical code and the data used for this paper. In Section IV we derive the constraints from existing data for several different runs, including effects of the effective relativistic degree of freedom $N_{eff}$, CMB-lensing, curvature and massive neutrinos. Here we also forecast the constraining capability in presence of the forthcoming Planck data, joined with SPT mock data. Finally, in Section V we derive our conclusions. ", "conclusions": "\\label{conclusions} We have considered the possibility that the evolution of dark matter and dark energy might be connected by a constant coupling, of the type illustrated in \\cite{amendola_2000, pettorino_baccigalupi_2008}. We have used current CMB data from WMAP7 and SPT to constrain the coupling parameter $\\beta$. We find that $\\beta$ is constrained to be less than 0.063 (0.11) at $68\\%$ ($95\\%$) C.L. when SPT data are included, with respect to $\\beta < 0.078 (0.14)$ coming from WMAP7 only. We have done a number of tests to check whether this bound depends on the degeneracy with other parameters (lensing, curvature, massive neutrinos, $N_{eff},$ HST/BAO/SNae data). If the effective number of relativistic degrees of freedom $N_{eff}$ is allowed to vary, no much gain is obtained on $\\beta$, which still needs to be $\\beta < 0.074 (0.12)$. We have further considered the effect of CMB-lensing, both with a run which includes no lensing and by marginalizing over $A_L$, a parameter which encodes the rescaling of the lensing power spectrum. $A_L$ is slightly degenerate with $n_s, \\Omega_{DM}h^2$ which in turn are degenerate with $\\beta$, though no direct degeneracy is seen between $\\beta$ and $A_L$. If the assumption of a flat universe is released, constraints on $\\beta$ weaken back almost to the level of constraints given by WMAP only (flat universe), with $\\beta < 0.071 (0.13)$. Degeneracy with massive neutrinos widens the coupling constraints to be $\\beta < 0.084 (0.14)$ when we marginalize over the fraction of massive neutrino species $f_{\\nu}$. We conclude that the bound on $\\beta$ from current data is already strong enough to be quite stable with respect to a better knowledge of other parameters and to all cases considered. When WMAP+SPT are considered (run $cq1$), the best fit value for $\\beta$, though still fully compatible with zero, has a best fit of $\\beta = 0.012^{+0.050}_{-0.012}$. It is interesting to see that when we allow for an effective number of relativistic degrees of freedom, marginalizing over $N_{eff}$ the coupling from WMAP7+SPT data increases to a best fit value of $\\beta \\sim 0.03$. A larger value of $N_{eff}$ favors larger couplings between dark matter and dark energy and values of the spectral index closer to $1$. Including SPT data does not improve significantly constraints on the coupling $\\beta$. Inclusion of additional priors from HST, BAO and SNae moves the best fit to $\\beta=0.041$, again still compatible with zero at 1$\\sigma$. We forecast that the inclusion of Planck data will be able to pin down the coupling to about $1\\%$ and therefore detect whether the small non-zero coupling present in current data is washed away with more data." }, "1207/1207.2295_arXiv.txt": { "abstract": "We present F850LP-F160W color gradients for 11 early-type galaxies (ETGs) at 1.0$<$z$_{spec}<$1.9 selected from the GOODS South field. Significant negative F850LP-F160W color gradients (core redder than the outskirts) have been detected in $\\sim 70\\%$ of our sample within the effective radius R$_{e}$, the remaining 30$\\%$ having a flat color profile consisten with a null gradient. Extending the analysis at R$>$R$_{e}$, enclosing the whole galaxy, we have found that the fraction of high-z ETGs with negative F850LP-F160W color gradients rises up to 100$\\%$. For each galaxy, we investigate the origin of the radial color variation with an innovative technique based on the matching of both the spatially resolved color and the global spectral energy distribution (SED) to predictions of composite stellar population models. In fact, we find that the age of the stellar populations is the only parameter whose radial variation alone can fully account for the observed color gradients and global SEDs for half of the galaxies in our sample (6 ETGs), without the need of radial variation of any other stellar population property. For four out of these six ETGs, a pure metallicity variation can also reproduce the detected color gradients. Nonetheless, a minor contribution to the observed color gradients from radial variation of star-formation time scale, abundance of low-to-high mass stars and dust cannot be completely ruled out. For the remaining half of the sample, our analysis suggests a more complex scenario whereby more properties of the stellar populations need to simultaneously vary, likely with comparable weights, to generate the observed color gradients and global SED. Our results show that, despite the young mean age of our galaxies ($<$3-4 Gyr), they already exhibit significant differences among their stellar content. We have discussed our results within the framework of the widest accepted scenarios of galaxy formation and conclude that none of them can satisfactorily account for the observed distribution of color gradients and for the spatially resolved content of high-z ETGs. Our results suggest that the distribution of color gradients may be due to ETGs forming by different mechanisms. ", "introduction": "A viable way to gather insight on the processes that concur to accrete the stellar mass in early-type galaxies (ellipticals plus S0's, hereafter ETGs) is to analyze the spatial distribution and properties of their stellar content which, in principle, can be directly connected to the events experienced by the galaxies. Indeed, the different scenarios proposed to explain the formation of ETGs give different predictions on their stellar population content. The $revised$ monolithic model predicts that, in a cold dark matter framework, massive ETGs assemble the bulk of their mass at z$>$ 2-3 through the merger of small substructures moving in a common potential well. This initial collapse might be regulated by cold gas streams from the cosmological surroundings, also known as cold accretion~\\citep{dekel09}, the latters becoming less important, for high-mass galaxies, at lower redshift. The subsequent evolution should be mainly characterized by the aging of their stellar populations, with small new episodes of star-formation at z$<$1, related, e.g., to the capture of satellites \\citep{katz91,kawata01,kobayashi04,merlin06}. During the gravitational dissipative collapse, the metal-enriched gas should naturally flows towards the center of the galaxy, leading to an ETG with stellar populations more metal-rich in the center than in the external regions (i.e. negative metallicity gradient). Moreover, because of the deeper potential well, the star-formation is expected to last longer in the central than in the outer regions. This would lead to null or mildly positive age gradients, with stellar populations in the center $\\sim$ 10 $\\%$ younger than those in the outskirts~\\citep{kobayashi04}. Nonetheless the effect of metallicity variation on color profile should be the dominant one, thus in this scenario ETGs are generally expected with negative color gradients. \\\\ The competing formation scheme is the $hierarchical$ scenario. Following the hierarchical assembly of cosmic structures, ETGs are supposed to form through gas-rich (``wet'') mergers of disc galaxies \\citep[e.g.][]{toomre72, delucia06} at high-redshift (z$\\sim$4-5). In this phase, a large fraction of the stellar mass of a galaxy is assembled through central intense bursts of star formation \\citep[e.g.][]{renzini06}. Concurrently, a ``dry'' merger picture has also been advocated, where bright ETGs would form through the merging of quiescent galaxies \\citep[e.g.][]{bell04}. Actually, in the last years a new scheme of mass accretion of massive ETGs, known as \\textit{inside-out growth}, is becoming widely accepted. This scenario is motivated by the observational evidence of ETGs at 1.0$<$z$<$2.5 with effective radii 3-5 times smaller than the mean radius of local ETGs with the same stellar mass \\citep[e.g][]{daddi05,longhetti07}. In this context, supported by a wealth of simulations \\citep[e.g.][]{khochfar06, hopkins09, wuyts10, naab09, bezanson09} ETGs are supposed to be formed at high redshift (z$\\sim$4-5) as compact spheroids result of gas-rich mergers. Then, at lower redshift, compact ETGs would undergo subsequent minor-``dry'' mergers, whose main effect is to add an external low-mass density envelope to the compact core ETGs, enlarging the effective radius while leaving the stellar mass nearly constant. Indeed, this scenario shows some limitations, such as the not plausible number of minor-``dry'' mergers necessary to enlarge the ETGs' size which could produce a scatter in the fundamental plane much larger than the observed one, and its failure to explain the presence of normal ETGs observed at high-z in number similar to the compact ones~\\citep[e.g.][]{nipoti09,saracco11, saracco10}. From a theoretical point of view, the ``wet''-merger scenario predicts ETGs with significant radial age variations. Indeed, the galaxy remnant of a wet merger should be characterized by a central stellar population younger, and more metal-rich, than the outer one, i.e. by a positive (negative) age (metallicity) gradient~\\citep{kobayashi04}. On the contrary, ``dry''-mergers, mixing the pre-existing stellar populations of progenitor galaxies, should dilute any radial variation of age and metallicity, producing flatter distributions (but see also \\citet{dimatteo09}). In the \\textit{inside-out} scenario minor dry-mergers are actually believed to only $add$ an $outer$ low-density envelope on top of a compact core, without mixing the pre-existing stellar-populations but redistributing the star content of the satellites in the outer regions of the compact ETG.\\\\ The picture described so far shows how important can be to spatially resolve the properties of the underlying stellar populations of ETGs in order to pinpoint their formation scenario. The most viable way to carry out information on the radial variation of the stellar content of a galaxy is to investigate its radial color variation, being the color of a stellar population tightly dependent on its age, star-formation time scale, metallicity, and (eventual) presence of dust. \\\\ ETGs at high redshift (hereafter with ``high redshift'' we mean $z>$1) represent the best-suited benchmark to investigate and constrain the mechanisms driving the mass assembly of spheroids due to the short time elapsed since their formation. So far, instrumental limits have allowed only studies of the global (i.e. integrated) properties of high-z spheroids, preventing, indeed, any measure of their spatially resolved information. Recently, the capabilities of the Hubble Space Telescope (HST) have made possible to overwhelm part of these limitations for the first time. Gargiulo et al. (\\citeyear{gargiulo11}, hereafter GSL11) taking advantage of the deep and high-resolution HST Advanced Camera for Surveys (HST/ACS) images of the GOODS South field derived F606W - F850LP color gradients ($\\sim$ (UV -U)$_{\\it{rest frame}}$ at z=1.5, $\\lambda_{eff}$ of F606W and F850LP filter $\\sim$ 5810\\,\\AA\\,and 9010\\,\\AA\\,respectively) for 20 ETGs at 1 $<$ $z_{spec}$ $<$ 2. In their work they ascertained the feasibility of this analysis up to z$\\sim$2 and presented the first spatially resolved information for ETGs at high-z. Despite the short wavelength baseline covered, $\\sim$ 50\\,$\\%$ of the galaxies showed a significant (positive or negative) color gradient. These results clearly showed that, after 3-4 Gyr (at most) from their birth, ETGs do not exhibit a unique spatial distribution of their stellar populations, implying that they followed different mass assembly paths. Unfortunately, the available HST data, mainly optical, prevented the authors to discriminate the drivers of the observed color gradients (e.g. radial variations of age, metallicity, ...) and consequently, to constrain the mechanisms responsible for them. Indeed, at $<$z$>$ $\\sim$ 1.5, the galaxy emission sampled by the F606W and F850LP filters is sensitive to both age and dust variation. Moreover, it is dominated by the youngest ($\\sim$ 1Gyr) stellar populations, missing any information on the distribution of the oldest stellar populations. \\\\ The recent advent of the first HST Wide Field Camera 3 (HST/WFC3) near infrared images for part of the GOODS South area (see Section 2 for details) has opened new possibilities to study color gradients of high-z ETGs and to constrain their origin. In this paper we combine the information provided by the WFC3/F160W-band ($\\sim$ R-band rest-frame at $<$z$>$ $\\sim$ 1.5, $\\lambda_{eff}$ of the filter $\\sim$ 15369\\,\\AA) and F850LP-band emission to derive the F850LP-F160W color gradients ($\\sim$ (U-R)$_{\\it{rest frame}}$ at z=1.5) for a sample of 11 ETGs at $<$\\textit{z}$>$\\,$\\sim$1.5. The bands we selected sample emissions dominated by different stellar populations. In particular, differently from F606W-F850LP color, the F850LP-F160W color is much more sensitive to age variations than dust content. This allows us to extend and to complement the analysis presented by GSL11.\\\\ Actually, \\citet{guo11} presented F850LP-F160W color gradients for 4 massive passively evolving galaxies in the GOODS South area at 1.3 $<$ \\textit{z}$_{spec}$ $<$ 2. They derived the color profiles by measuring optical-NIR colors in concentric annuli, and found that high-z ETGs have cores redder than the outerskirts. The observed radial trend in color gradients is not reproduced by the radial variation of a single stellar population parameter (age, metallicity, extinction), although they found that dust should partially contribute to generate the observed color distribution. In this paper, we examine a sample three times larger than that of \\citet{guo11}, deriving F850LP-F160W color profiles from the 2-dimensional fit of the light profiles in the two bands. Then, we present a new approach to exploit the wealth of available information to constrain the radial variation of the underlying stellar populations. In fact, taking advantage of a unique set of data and an innovative procedure, we are able to constrain the radial variation of stellar population parameters (age, metallicity, dust, star-formation time scale, and initial mass function) and their contribution to produce the color gradients we observe in high-z ETGs. Finally, we compare our findings to the prediction of theoretical formation models.\\\\ Throughout the paper we adopt a standard $\\Lambda$CDM cosmology with H$_{0}$=70km\\,s$^{-1}$\\,Mpc$^{-1}$, $\\Omega$\\,$_{m}$=0.3 and $\\Omega$\\,$_{\\Lambda}$=0.7. All the magnitudes are in the AB system. ", "conclusions": "The results of our analysis establish the heterogeneity of the stellar content in high-z ETGs. For $\\sim$ 50$\\%$ of the galaxies of our sample (6 ETGs) we have found that a pure radial age variation, with the oldest stars located at the center of the galaxy, is able to reproduce the observed radial color profile and global SED. The age gradients we have detected for these galaxies span a range from -0.84 to -0.28 dex per radial decade. For 4 out of these 6 ETGs the color gradients we have observed can be also accounted for by solely metallicity gradients whose strength varies from -0.35 to -1.45 dex per radial decade.\\\\ Due to the lack of similar measurements on other sample of high-z ETGs we have compared the age/metallicity gradient values we have found with those observed in the local ETGs. \\\\ Studies of ETGs at lower redshift affirm that color gradients are mainly due to radial metallicity variations. Assuming the age of the stellar populations constant throughout the galaxy, as we did in the case of metallicity gradient analysis, color gradients in local ETGs turn out to be reproduced by a mean metallicity gradients ranging from -0.16 $\\div$ -0.3 \\citep{peletier90, idiart03, tamura03}. A further confirmation of the metallicity as main driver of local gradients comes out from studies at intermediate redshift which show that the color gradients evolution is better accounted for by the passive evolution of metallicity gradients \\citep{saglia00, hinkley01}. In fact, our results seem to be inconsistent with these findings. In our sample of high-z ETGs we have found that only 4 galaxies out 11 have color gradients well reproduced by pure radial metallicity variation. Moreover, the metallicity gradients we have detected in high-z ETGs ([-0.3$\\div$-1.45]) are systematically steeper than those typically observed in local ETGs ([-0.16$\\div$-0.3]), even if \\citet{ogando05} found that this range become wider ([0.0 $\\div$ -1.0]) for nearby massive (M$_{\\star}>10^{10}$M$_{\\odot}$) ETGs see also \\citep[see also][]{spolaor09}. The steeper metallicity gradients that we have detected derive by a radial metallicity variation from supersolar (2Z$_{\\odot}$) values in the inner regions to subsolar values in the external regions (0.2Z$_{\\odot}$). Although such extreme values of Z have been observed also in few local ETGs \\citep{mehlert03,rickes08}, our results show that metallicity gradients in high-z ETGs of our sample are only marginally comparable with the typical metallicity gradients detected in local ETGs. This result seems to point in favor of a possible evolution of the metallicity gradient in the last 9Gyr. Concurrently, studies on cluster ETGs at low and intermediate redshift show that pure age variations in their stellar populations are not able to account for their color gradients \\citep{saglia00}. In contrast to local results, we have found that for $\\sim$ 50$\\%$ of the galaxies of our sample a radial variation of stellar populations age alone can reproduce the observed color gredients and global SED. In fact, recent studies investigating the simultaneous radial variation of both age and metallicity confirm metallicity gradients ($\\sim$-0.4 dex per radial decade) as the main driver of observed color gradients in local ETGs but found also a small contribution to color variation of positive age gradient ($\\sim$0.1 dex per radial decade) \\citep{labarbera09,wu05}. The age/metallicty degeneracy affecting optical colors does not allow us to consider the simultaneous radial variation of both parameters and hence to detect a possible positive age gradient, whose presence in local ETGs, actually, is still matter of debate. \\\\ To spread light on this issue, high-z ETGs constitute the ideal place to investigate the presence of an age gradient. Indeed, at fixed radial variation of age $\\Delta$age, its effect on color profile is much more enhanced when stellar populations are younger, hence in high-z ETGs. This effect is clearly shown in left panel of Fig. \\ref{evolcol}, where we report the differences observed in the F850LP-F160W color of two stellar populations with age which differ of 2Gyr ($\\tau$ = 0.1Gyr, black solid curve), as a function of the age of the youngest stellar population. The same $\\Delta$age produces a difference in the F850LP-F160W color of the two populations that is $\\sim$ 10 times higher if observed in high-z ETGs (Age $<$ $\\sim$ 4Gyr) respect those observed in local ETGs (Age $\\sim$ 10Gyr). A typical radial variation of 2 Gyr, as the one we measure in the ETGs of our sample, will produce in a local ETG a variation in the F850LP-F160W color of $\\sim$ 0.05 mag, thus at the very limit of the detection. On the contrary, the same age variation will result in a color variation of $\\sim$ [0.3-0.5] mag for high-z ETGs. Red and blue lines report the same of black line, but for pure metallicity variations. In particular, red line show the variation in the F850LP-F160W color observed in two populations with metallicities 2Z$_{\\odot}$ and 0.2Z$_{\\odot}$, while blue line in two populations with metallicities Z$_{\\odot}$ and 0.2Z$_{\\odot}$. In the right panel of Fig. \\ref{evolcol} the color gradient that the above age/metallicity variations would produce when occurring between 0.1R$_{e}$ and 3R$_{e}$. Differently from age variation, the effect of a metallicity variation on color, and hence on gradient, increases with the age of the galaxy of a factor $\\sim$ 2 from high-z ETGs to local ETGs. These plots emphasize how challenging is the detection of an age gradient in local ETGs due to its almost negligible effect on color profile. On the contrary, in high-z ETGs age and metallicity variations produce comparable effect on color profile, thus setting the ideal condition for their detection. \\begin{figure*} \\begin{center} \\includegraphics[width=19.0cm,angle=0]{figure10.ps} \\caption{$Left$ $panel$: Black line shows the difference observed in the F850LP-F160W color of two stellar populations with age differentiating of 2Gyr, as a function of the age of the youngest stellar population. Red line shows the variation in the F850LP-F160W color observed in two populations with metallicities 2Z$_{\\odot}$ and 0.2Z$_{\\odot}$, while blue line for two populations with metallicities Z$_{\\odot}$ and 0.2Z$_{\\odot}$. $Right$ $panel$ Color gradients relative to the color variation in the left panel in the hypothesis that they occur between 0.1R$_{e}$ and 3R$_{e}$.} \\label{evolcol} \\end{center} \\end{figure*} This comparison with local samples is only meant to have an indication on how the results obtained at high-z, both in terms of age/Z gradients and in terms of internal and external age/Z values, relate with the local values. On the other hand, the unknown evolution experienced by ETGs in $\\sim$9 Gyr from z$\\sim$1.5 to z=0 can affect the stellar properties and distribution (e.g. minor mergers triggering secondary burst of star formation) making complex any comparison with local universe. In fact, to properly face on high-z and local values samples of ETGs selected in an homogeneous way, a dataset able to trace the evolution of the same rest-frame color gradient over 9 Gyr, and similar procedures for both the color gradients estimates and the relative analysis should be considered. In a forthcoming paper, taking into account all these factors, we will try to address the origin of the color gradients following their evolution from z$\\sim$2 to z=0. For the remaining five ETGs of our sample, a pure radial variation of a single stellar population parameter is not able to reproduce the observed color gradients and global SEDs. Differently form the previous cases, where color gradients could be reproduced by a pure age or metallicity gradient as well as by a simultaneous radial variation of more than one parameter, these galaxies $need$ a more complex scenario whereby more than one property of the stellar populations have to vary from the center to the periphery to generate the observed color gradients. Thus, our analysis clearly indicates that the properties of the stellar population and their distribution within high-z ETGs do not follow an homogeneous and common scheme. In the following we try to investigate if the theoretical expectations of the widely accepted scenarios of galaxy formation can explain the observed color distributions. In the monolithic revised scenario, color gradients are supposed to be mainly due to metallicity gradient, being the contribution of age null or mild (and positive). Our findings of none correlation between color/metallicity gradients and total mass suggest that the monolithic revised scenario is not the favored mechanisms with which ETGs assembled their mass, although the narrow mass range covered and the assumption of the stellar mass as a proxy of the total stellar mass can affect this conclusion. Theoretical predictions of the \\textit{inside-out-growth} scenario point in favor of compact ETGs with cores \\citep{wuyts10} redder than the outskirt regions. This negative color gradient seems to be due to a combined effect of negative metallicity and positive age gradients, with a non negligible effect of dust \\citep{wuyts10}. To compare our results with the theoretical prediction of this model, we define compact galaxies those ETGs with effective radius more than 1$\\sigma$ smaller than those predicted by the local size-mass relation (SMR) for that mass. In Table 3 we report the compactness C for the galaxies of our sample defined as the ratio R$_{e,z=0}$/R$_{e}$ where R$_{e}$ is the effective radius of the galaxy and R$_{e,z=0}$ is the radius that a galaxy of equal stellar mass would have at z=0 as derived by the SMR of \\citet{shen03}. In our sample, compact galaxies, as we defined them, turn out to have C$\\geq$2. In Fig \\ref{smr} we report the SMR in the F850LP band for our sample (solid symbols) and for local galaxies \\citep[solid line,][]{shen03}. The dashed lines represent the scatter at 1$\\sigma$. It turns out that seven out of 11 galaxies (solid points) of our sample are compact (magenta points) and 4 are normal ETGs (cyan triangles). In Table 3, we report the classification in normal (N) and compact (C) for our galaxies. \\begin{figure} \\includegraphics[width=7.5cm,angle=0]{figure11.ps} \\caption{SM relation for local ETGs (solid line, Shen et al. 2003) and for our sample (solid symbols). The dashed lines are the scatter lines at 1$\\sigma$. We shifted the Shen et al.\u2019s relation by a factor \u223c1.2 towards lower masses to take into account the systematic shift observed in the mass estimations using our models or those adopted by Shen et al. (2003). The circles are compact galaxies, that is, galaxies having the effective radius more than 1$\\sigma$ smaller than those predicted by a local relation for that mass. On the contrary, galaxies having the effective radius comparable at 1$\\sigma$ with those expected by Shen et al.\u2019s relations are classi\ufb01ed as normal (triangle symbols).} \\label{smr} \\end{figure} Compact ETGs in our sample (as well as normal ETGs) show redder cores supporting the \\textit{inside-out-growth} scenario even if we cannot firmly establish the origin and the nature of this gradient. Actually, we cannot test the presence of a positive age gradient, since due to age/metallicity degeneracy we do not treat the case of simultaneous radial variation of age and metallicity. On the other hand, our results show that only two out of seven compact ETGs have color gradients well reproduced by a metallicity gradient.\\\\ Nonetheless, the simulations predicts that the negative color gradients produced by the interplay of age, metallicity and dust should result to correlate with the integrated rest-frame optical color. Following \\citet{wuyts10} in Fig. \\ref{cgvscolor} we show the ratio between the F850LP-band effective radius and F160W-band effective radius as a proxy of the color gradients versus the F850LP-F160W color. The points follow the same convention of Fig. \\ref{smr}. We do not find any correlation neither in the whole sample, nor in the compact selection. \\begin{figure} \\includegraphics[width=5.8cm,angle=-90]{figure12.ps} \\\\ \\caption{The r$_{e,F850LP}$/r$_{e,F160W}$ ratio as proxy of F850LP-F160W color gradient vs the F850LP-F160W global color.} \\label{cgvscolor} \\end{figure} The absence of such correlations cast some doubts on the validity of this scenario as a viable process to assemble the stellar mass in compact ETGs.\\\\ Despite the fact that both the widely accepted formation scenarios do not seem to be able to reproduce the stellar content of high-z ETGs, we have to test the hypothesis whether ETGs can be assembled through a common formation process and the distribution we observed in color gradients can be the final product of subsequent merger events. We have to bear in mind that our galaxies are all younger than 3.5 Gyr. This severely constrains the time each galaxy has at disposal to experience a merger event. Table 1 of \\citet{boylan08} shows the dynamical friction merging time in Gyr, $\\tau_{merger}$, for a host halo with virial mass M$_{host}$ = 10$^{12}$M$_{\\odot}$, measured from numerical simulations. They assumed different initial orbital angular momentum, initial orbital energy, and ratio of initial satellite mass to initial host halo mass and only for two cases they found 3.5 Gyr is enough to complete the merger, while in all the other cases $\\tau_{merger}$ is greater than 4.3 Gyr. Thus, the different stellar content of high-z ETGs does not seem to be due to the effect of subsequent merger events, but primarily to the formation process. Actually, the continuum distribution of ETGs in the size-mass plane, both at high redshift and in the local Universe, together with the systematic direction of the color gradient (negative or null) of high-z ETGs, point toward a common formation process responsible of this continuity. The possible different initial conditions, such as the different time scale of collapsing gas cloud, can be responsible of the observed structural and dynamical differences as we previously suggested \\citep{saracco11, saracco12}." }, "1207/1207.6405_arXiv.txt": { "abstract": "The 21~cm signal produced by non-evaporating primordial black holes (PBHs) is investigated. X-ray photons emitted by accretion of matter onto a PBH ionize and heat the intergalactic medium (IGM) gas near the PBH. Using a simple analytic model, we show that this X-ray heating can produce an observable differential 21~cm brightness temperature. The region of the observable 21~cm brightness temperature can extend to 1--10 Mpc comoving distance from a PBH depending on the PBH mass. The angular power spectrum of 21~cm fluctuations due to PBHs is also calculated. The peak position of the angular spectrum depends on PBH mass, while the amplitude is independent of PBH mass. Comparing this angular power spectrum with the angular power spectrum caused by primordial density fluctuations, it is found that both of them become comparable if $\\Omega_{{\\rm PBH}} = 10^{-11} (M/10^{3}~% M_\\odot)^{-0.2}$ at $z=30$ and $10^{-12} (M/10^{3}~ M_\\odot)^{-0.2}$ at $z=20$ for the PBH mass from $10 ~ M_\\odot $ to $10^8~ M_\\odot $. Finally we find that the Square Kilometer Array can detect the signal due to PBHs up to $\\Omega_{\\rm PBH}=10^{-5} (M/10^{3}~ M_\\odot)^{-0.2}$ at $z=30$ and $10^{-7} (M/10^{3}~ M_\\odot)^{-0.2}$ at $z=20$ for PBHs with mass from $10^2 ~ M_\\odot $ to $10^8~ M_\\odot $. ", "introduction": "primordial black holes (PBHs) could have formed in the early Universe \\citep{1974MNRAS.168..399C,1975ApJ...201....1C}. Although there is no direct evidence of their existence of PBHs, PBHs are attracting attention as a way of constraining physics in the early Universe. In particular, one of the main generation mechanisms of PBHs is the gravitational collapse of an overdense region at the horizon scale when the amplitude of the overdensity exceeds a critical threshold. Therefore the resultant mass function and the abundance of PBHs depend on the amplitude of primordial density fluctuations at the horizon-crossing epoch \\citep{2004PhRvD..70d1502G}. The PBH abundance is expected to be a probe of primordial density fluctuations on small scales, which cannot be accessed by Cosmic Microwave Background or large-scale-structure observations. Constraints on the abundance of PBHs have been extensively studied and continue to be updated \\citep{2010PhRvD..81j4019C}. PBHs with mass less than $10^{15}~$g have evaporated by the present epoch because the evaporation time scale by Hawking radiation is less than the Hubble time scale today \\citep{1974Natur.248...30H}. However evaporation of PBHs generates additional entropy in the Universe after inflation \\citep{% 1976JETPL..24..571Z}, affects big bang nucleosynthesis \\citep{% 1978SvAL....4..185V, % 1978SvA....22..138V,1978PThPh..59.1012M,% 1977SvAL....3..110Z,1980MNRAS.193..593L} and distorts the CMB blackbody spectrum \\citep{2008PhRvD..78b3004T}. PBH evaporation may also produce the observable gamma-ray background \\citep{1976ApJ...206....1P,1991ApJ...371..447M}. According to measurements of these cosmological phenomena, there are strong constraints on the abundance of PBHs with mass less than $10^{15}~$g. PBHs with mass larger than $10^{15}~$g survive in the present Universe. One of the constraints on such PBHs can be set from the fact that the current density parameter of PBHs, $\\Omega_{\\rm PBH}$, cannot exceed the cold dark matter density parameter observed at the present epoch, $\\Omega_{\\rm C}$. Conventionally, the constraint on PBH abundance is given by $\\beta(M)$ which is the fraction of regions of mass $M$ collapsing into PBHs at the formation epoch \\citep{1975ApJ...201....1C}. The constraint on the density parameter of PBHs today, $\\Omega_{{\\rm PBH}}<\\Omega_{\\rm C}$, implies $\\beta < 2 \\times 10^{-18} (M/10^{15}{\\rm g})^{1/2}$ from WMAP 7-year data, i.e., $\\Omega_{\\rm C} =0.22$ \\citep{2011ApJS..192...18K}. Microlensing observations also constrain the abundance of non-evaporating PBHs \\citep{2001ApJ...550L.169A}. \\citet{2008ApJ...680..829R} have obtained the constraint on PBHs with mass larger than $0.1~{\\rm M}_\\odot$, by investigating the effects of such PBHs on cosmic reionization and CMB temperature anisotropies. Future gravitational wave observations are also expected to provide a probe of the massive PBH abundance \\citep{1999PhRvD..60h3512I,2003PhRvL..91b1101I}. In this paper, we evaluate the 21~cm brightness temperature produced by PBHs and study the potential of 21~cm observations to give a constraint on the abundance of PBHs. \\citet{2008arXiv0805.1531M} have investigated the signatures of evaporating PBHs in 21~cm brightness temperature. Accordingly, they have concentrated on PBHs whose mass range is $5 \\times 10^{13}~{\\rm g} \\lesssim M_{\\rm PBH} \\lesssim 10^{17}~{\\rm g} $. On the contrary, here, we focus on non-evaporating PBHs with mass much larger than $10^{15}~$g. After the epoch of matter-radiation equality, gas and matter can accrete onto PBHs. It has been shown that PBHs with large mass could produce X-ray and UV photons through the accretion of matter onto PBHs and these photons heat up and ionize intergalactic medium (IGM)~\\citep{1981MNRAS.194..639C,1995ApJ...438...40G,2001ApJ...561..496M,2008ApJ...680..829R}. Therefore, the heated and ionized IGM gas may produce an observable deviation of the 21~cm brightness temperature from the background even before the birth of the first stars and galaxies~($z>30$). Our aim in this paper is to evaluate this deviation and to discuss the potential of 21~cm observations to constrain the non-evaporating PBH abundance. The paper is organized as follows. In section 2, using a simple model of X-ray photon flux due to the accretion onto a PBH, we evaluate the ionization and temperature profile near a PBH. In section 3, we calculate the spin temperature and the brightness temperature induced by a PBH. In section 4, the angular power spectrum of 21~cm fluctuations due to PBHs are evaluated. Section 5 is devoted to the conclusion. Throughout this paper, we use parameters for a flat $\\Lambda$CDM model: $h=0.7$ $(H_0=h \\times 100 ~ {\\rm km /s /Mpc})$, $\\Omega_{\\rm B}=0.05$ and $\\Omega_{\\rm M}=0.26$. These parameters are consistent with WMAP results \\citep{2011ApJS..192...18K}. ", "conclusions": "We have investigated the 21~cm signal produced by massive PBHs whose masses are larger than $10~M_{\\odot}$. Assuming a power-law spectrum of X-ray photons from an accretion disk, we have studied the ionization and heating of IGM gas near a PBH and evaluated the differential 21~cm brightness temperature. We have shown that a PBH can induce an observable signal of differential 21 ~cm brightness temperature. The size of the region where we can find the differential brightness temperature typically reaches 1--10~Mpc for our interested PBH mass range. The exact size depends on the PBH mass. We have also calculated the angular power spectrum of 21~cm fluctuations due to PBHs. The peak position of the angular spectrum depends on the PBH mass, while the amplitude is independent of the mass. Comparing this spectrum with the angular power spectrum caused by primordial density fluctuations, we have found that both of them become comparable if $\\Omega_{{\\rm PBH}} =10^{-11} (M/10^{3}~ M_\\odot)^{-0.2}$ at $z=30$ and $10^{-12} (M/10^{3}~ M_\\odot)^{-0.2}$ at $z=20$ for PBH's mass from $10 ~ M_\\odot $ to $10^8~ M_\\odot $. If the density parameter is larger than these values, the angular power spectrum due to PBHs exceeds the one from primordial fluctuations and can be measured. In other words, we cannot set constraints on the PBH density parameter below these values from 21~cm observations. If we consider the sensitivity of the SKA-like observation, for example, we can detect the signal of PBHs up to $\\Omega_{\\rm PBH}=10^{-5} (M/10^{3}~ M_\\odot)^{-0.2}$ at $z=30$ and $10^{-7} (M/10^{3}~ M_\\odot)^{-0.2}$ at $z=20$ for PBHs with mass from $10^2 ~ M_\\odot $ to $10^8~ M_\\odot $. The ionization of IGM due to PBHs with such density parameters does not affect the global reionization history of the universe since reionization from each PBH only covers a tiny patch of the universe. Unlike reionization from first stars, therefore, such reionization has little impact on CMB temperature anisotropies. On the other hand, PBHs can heat IGM regions whose scale reaches 1-10Mpc and 21~cm fluctuations are sensitive to the heated IGM regions. Accordingly the PBH density parameter constrained from WMAP data, that is $\\Omega_{{\\rm PBH}} < 10^{-7}$ \\citep{2008ApJ...680..829R}, is several order of magnitude larger than the values at which 21~cm fluctuations from PBHs comparable with those from primordial density fluctuations as we mentioned above. In other words, we can conclude that, if observation instruments or foreground removal are more improved than the current SKA design, 21~cm fluctuation observations have a potential to probe the PBH abundance which is impossible to access by CMB observations. The most theoretical uncertainty in this model is the flux of photons due to the accretion to PBHs. In this paper, we assume that the X-ray photon flux amplitude is a tenth of the Eddington luminosity and has the power law spectrum with $E^{-1}$ for simplicity. We also studied the effect of the power law index by changing to $E^{-1/2}$. The temperature profile shifts to a smaller scale due to the decrease of the ionization efficiency. However this shift is small, and the resultant angular power spectrum does not change much. The amplitude of the luminosity is considered to depend on the matter accretion rate onto a PBH. \\citet{2008ApJ...680..829R} have studied the luminosity for the Bondi-Hoyle accretion in detail. Although the luminosity depends on the PBH mass and feedback effect on the ionization and temperature, they have shown that the luminosity for a PBH with $M=10^3~ M_\\odot$ is roughly a hundredth of the Eddington luminosity at $z>20$. In our model, the brightness temperature profile near a PBH at a certain redshift depends only on the PBH's mass and the amplitude of the angular power spectrum is scaled by $\\Omega_{\\rm PBH}$. Therefore, if we assume that the amplitude of the X-ray photon flux is a hundredth of the Eddington luminosity, the required density parameter for PBHs to dominate the 21~cm fluctuations due to primordial density fluctuations is $10^{-11} (M/10^{4}~ M_\\odot)^{-0.2}$ at $z=30$ and $10^{-12} (M/10^{4}~ M_\\odot)^{-0.2}$ at $z=20$ for PBH's mass from $10 ~ M_\\odot $ to $10^8~ M_\\odot $." }, "1207/1207.6069_arXiv.txt": { "abstract": "This memo describes the software engineering and technical details of the design and implementation of the \\wvrgcal\\ program and associated libraries. This program performs off-line correction of atmospheric phase fluctuations in ALMA observations, using the 183\\,GHz Water Vapour Radiometers (WVRs) installed on the ALMA 12\\,m dishes. The memo can be used as a guide for detailed study of the source code of the program for purposes of further development or maintenance. ", "introduction": "The `{\\tt wvrgcal}' program is an application for off-line correction of atmospheric phase fluctuations in ALMA data based on observations of 183\\,GHz Water Vapour Radiometers (WVRs) that are installed on all ALMA 12\\,m-diameter antennas. The principles of WVR based phase correction, and of the algorithms used in `{\\tt wvrgcal}' are described in previous papers and ALMA memos, e.g., \\cite{2001ApJ...553.1036W,ALMANikolic587,ESONikolic2008}, and forthcoming publications. In this memo we describe the technical details of the software engineering design and implementation of the {\\tt wvrgcal} program. This memo describes version 1.2 of \\wvrgcal\\, which can be downloaded under the GNU Public License at \\url{http://www.mrao.cam.ac.uk/~bn204/soft}. Brief usage instructions for astronomical users of \\wvrgcal\\ can be found in the command line help (\\texttt{wvrgcal --help}), shown in Appendix \\ref{sec:help}. The program is currently shipped with the data reduction environment CASA (from version 3.4), and additional usage instructions can be found in the integrated help within CASA and in the CASA cookbook (\\url{http://casa.nrao.edu/ref_cookbook.shtml}). ", "conclusions": "" }, "1207/1207.5776_arXiv.txt": { "abstract": "We present 279 epochs of optical monitoring data spanning 5.4 years from 2007 January to 2012 June for the largest image separation (22\\farcs6) gravitationally lensed quasar, SDSS J1029+2623. We find that image A leads the images B and C by $\\Delta t_{AB} = (744 \\pm 10)$~days (90\\% confidence); the uncertainty includes both statistical uncertainties and systematic differences due to the choice of models. With only a $\\sim 1\\%$ fractional error, the interpretation of the delay is limited primarily by cosmic variance due to fluctuations in the mean line-of-sight density. We cannot separate the fainter image C from image B, but since image C trails image B by only 2--3 days in all models, the estimate of the time delay between image A and B is little affected by combining the fluxes of images B and C. There is weak evidence for a low level of microlensing, perhaps created by the small galaxy responsible for the flux ratio anomaly in this system. Interpreting the delay depends on better constraining the shape of the gravitational potential using the lensed host galaxy, other lensed arcs and the structure of the X-ray emission. ", "introduction": "SDSS~J1029+2623 \\citep{Inada2006} is the largest image separation lensed quasar, with a maximum separation of 22\\farcs6 that significantly exceeds that of the next largest system (14\\farcs6 for SDSS~J1004+4112; \\citealt{Inada2003}). Although it was first identified with only two images A and B, it actually consists of three images of a $z_s=2.197$ quasar produced by a $z_l=0.58$ galaxy cluster in a rare ``naked cusp'' configuration (\\citealt{Oguri2008}). The fainter image C lies close to image B (1\\farcs8), which would usually mean that B and C should be significantly brighter than A. Instead, the optical flux ratios of the images, A:B:C=0.95:1.00:0.24, show a large anomaly that cannot be reproduced by ellipsoidal models centered near the bright cluster galaxies (\\citealt{Oguri2008}). The quasar is radio loud, and the flux ratio anomaly persists in the radio, albeit with different flux ratios than in the optical (\\citealt{Kratzer2011}). Recent \\textit{Chandra} X-ray observations (\\citealt{Ota2012}) find a cluster mass consistent with lens models and that there is soft X-ray absorption in the spectrum of image C consistent with explaining the optical color differences between the images A, B and C as extinction. After correction for this absorption/extinction, the X-ray, optical and radio flux ratios are broadly consistent and the flux anomalies are due to a small galaxy near image C (\\citealt{Oguri2012b}). As part of a program to better understand this system, including deep \\textit{Hubble Space Telescope} (\\citealt{Oguri2012b}), X-ray (\\citealt{Ota2012}), radio (\\citealt{Kratzer2011}), and weak-lensing (\\citealt{Oguri2012a}) observations, we have been monitoring the lens in the optical since 2007 to measure the time delay. Time delays generally measure a combination of cosmological distances (the Hubble constant to lowest order) and the surface density of the lens at the radius of the images (\\citealt{kochanek2002}). Since the mass distribution of the lens can be independently constrained by the X-ray emission profile (\\citealt{Ota2012}) and additional multiply imaged background galaxies of differing redshifts from the quasar (\\citealt{Oguri2008,Oguri2012b}), cluster lenses have the potential of being excellent cosmological probes if it can be demonstrated that the effects of substructure are controllable. Here we present a light curve for the brighter A and B images of SDSS~J1029+2623 spanning six 8-month observing seasons and measure their time delay. Because image C is faint and very close to B, we cannot independently determine its delay given the quality of our images. Section \\ref{sec:observations} summarizes the available data, Section \\ref{sec:delay} derives the time delay, and Section \\ref{sec:discussion} discusses the results. ", "conclusions": "\\label{sec:discussion} We presented 5.4 years of optical monitoring for the two bright lensed quasar images of the largest image separation gravitationally lensed quasar SDSS J1029+2623. We find that image A leads the images B and C by $\\Delta t_{AB} = (744 \\pm 10)$~days (90\\% confidence). The formal error bar on the time delay, which includes statistical uncertainties and systematic differences, is $\\sim 1.3\\%$ and is in the regime where cosmic variance caused by fluctuations in the mean line-of-sight density is more important than the measurement errors. We find that the effect of microlensing in this system is small. Formally, the detection is statistically significant, but we view the overall evidence for microlensing rather than low level systematic uncertainties as weak. This is the second longest measured time delay after the 822 day delay between images C and A in SDSS~J1004+4112 (\\citealt{Fohlmeister2008}). A detailed interpretation of the measured delay is deferred pending the completion of our analysis of the \\textit{HST} images (\\citealt{Oguri2012b}) and additional spectroscopy of the lensed arcs in this system. However, as an experiment, we fit the lens using a Navarro-Frenk-White model centered near galaxy G2 with a break radius of 78\\farcs0 based on the X-ray data (\\citealt{Ota2012}). We adopt the component positions ($\\pm 0\\farcs05$) from \\cite{Oguri2008} and a time delay of $\\Delta t_{AB} = (744 \\pm 10)$~days that tries to conservatively combine the AIC and BIC results. We used priors of $\\pm 1\\farcs0$ on the position of the model relative to galaxy G2, $\\epsilon = 0.46\\pm0.05$ for the ellipticity of the cluster density, and $\\gamma = 0.05 \\pm 0.05$ for any additional external shear. The value of $\\epsilon$ was determined by the best fit model with $H_0=70$~km~s$^{-1}$~Mpc$^{-1}$. We then fit models as a function of the Hubble constant $H_0$ assuming a flat, $\\Omega_0=0.3$, and $\\Lambda_0=0.7$ cosmological model. We did not include the flux ratios in the fits because of the known flux ratio anomaly. The anomaly appears to be due to a small galaxy near image C (\\citealt{Oguri2012b}) which should have little effect on the overall geometry and the AB time delay. In a normal four-image lens, the constraint on the radial mass profile from the X-ray data would largely eliminate any degeneracies because they are created by uncertainties in the surface mass density at the radius of the images (\\citealt{kochanek2002}). Figure~\\ref{fig:models} shows the resulting goodness of fit as a function of $H_0$ for these simple models. Low values of $H_0$ are disfavored by the astrometry, while high values are weakly disfavored by the priors on the shape of the potential. There is clearly a strong need for additional constraints in order to use this delay for the full characterization of the lens system. There is certainly no difficulty improving the astrometric constraints, since with \\textit{HST} it will be trivial to increase the astrometric precision from the $0\\farcs05$ used here to $0\\farcs01$ or smaller, although the concern here will be whether noise from unmodeled sub-structures in the cluster dominate over the measurement precision. More important, however, is the addition of strong constraints on the shape of the potential, since there is a strong trend requiring flatter density distributions and higher external shears for larger values of $H_0$. We already know from \\cite{Oguri2008} that there is a lensed image of the quasar host galaxy as well as several additional arc systems, and this is confirmed by the new \\textit{HST} observations (\\citealt{Oguri2012b}). The primary difference between these models and the early model of \\cite{Oguri2008}, which predicted a longer delay, is that these models are centered on G2 rather than on G1 based on the X-ray data. Particularly if combined with additional arc redshifts and direct mass models of the X-ray emission, there are no obvious barriers to greatly improving on Figure~\\ref{fig:models}. Where \\cite{Suyu2010} use dynamical measurements of the central lens galaxy to break the \\cite{kochanek2002} or other degeneracies, here we should be able to do so using a combination of the X-ray and arc data. There is, however, significant evidence that the cluster is undergoing a merger, which means that the additional lensing constraints will be more reliable constraints on the mass distribution than further X-ray observations. On the other hand, the ability to combine deeper X-ray observations with the lensing constraints makes this system and excellent laboratory for studying the effects of mergers on the X-ray properties of clusters. In our present data, measuring the B-C time delay is impossible given the data quality. It does shift steadily over the model sequence, decreasing with increasing $H_0$ from 3.4 to 2.1 days. This fully justifies our making the measurements using the combined B/C light curve. It is well worth measuring the BC delay because it is very sensitive to the perturbing galaxy that causes the flux ratio anomaly (see \\citealt{Oguri2007}, \\citealt{Keeton2009}) and hence probes the structure of galaxy halos in cluster environments. Such short delays are extremely difficult to measure accurately in ground based observations because quasars have so little variability power on such short time scales (see \\citealt{Mushotzky2011}) and might require a space-based lens monitoring satellite like the proposed OMEGA (\\citealt{Moustakas2008})." }, "1207/1207.7129_arXiv.txt": { "abstract": "{Models of galaxy evolution assume some connection between the AGN and star formation activity in galaxies. We use the multi-wavelength information of the CDFS to assess this issue. We select the AGNs from the 3\\,Ms {\\it XMM-Newton} survey and measure the star-formation rates of their hosts using data that probe rest-frame wavelengths longward of $\\rm 20\\,\\mu m$, predominantly from deep $\\rm 100\\,\\mu m$ and $\\rm 160\\,\\mu m$ {\\it Herschel} observations, but also from \\emph{Spitzer} MIPS-$\\rm 70\\,\\mu m$. Star-formation rates are obtained from spectral energy distribution fits, identifying and subtracting an AGN component. Our sample consists of sources in the $z\\approx0.5-4$ redshift range, with star-formation rates $\\rm SFR\\approx10^{1}-10^{3}\\,M_{\\sun}\\,yr^{-1}$ and stellar masses $M_{\\star}\\approx10^{10}-10^{11.5}\\,{\\rm M_{\\sun}}$. We divide the star-formation rates by the stellar masses of the hosts to derive specific star-formation rates (sSFR) and find evidence for a positive correlation between the AGN activity (proxied by the X-ray luminosity) and the sSFR for the most active systems with X-ray luminosities exceeding $L_{\\rm x}\\simeq10^{43}{\\rm \\,erg\\,s^{-1}}$ and redshifts $z\\gtrsim1$. We do not find evidence for such a correlation for lower luminosity systems or those at lower redshifts, consistent with previous studies. We do not find any correlation between the SFR (or the sSFR) and the X-ray absorption derived from high-quality {\\it XMM-Newton} spectra either, showing that the absorption is likely to be linked to the nuclear region rather than the host, while the star-formation is not nuclear. Comparing the sSFR of the hosts to the characteristic sSFR of star-forming galaxies at the same redshift (the so-called ``main sequence'') we find that the AGNs reside mostly in main-sequence and starburst hosts, reflecting the AGN - sSFR connection; however the infrared selection might bias this result. Limiting our analysis to the highest X-ray luminosity AGNs (X-ray QSOs with $L_{\\rm x}>10^{44}{\\rm \\,erg\\,s^{-1}}$), we find that the highest-redshift QSOs (with $z\\gtrsim2$) reside predominantly in starburst hosts, with an average sSFR more than double that of the ``main sequence'', and we find a few cases of QSOs at $z\\approx1.5$ with specific star-formation rates compatible with the main-sequence, or even in the ``quiescent'' region. Finally, we test the reliability of the colour-magnitude diagram (plotting the rest-frame optical colours against the stellar mass) in assessing host properties, and find a significant correlation between rest-frame colour (without any correction for AGN contribution or dust extinction) and sSFR excess relative to the ``main sequence'' at a given redshift. This means that the most ``starbursty'' objects have the bluest rest-frame colours.} {}{}{}{} ", "introduction": "One of the most significant observations of modern-day astrophysics is the evidence that the mass of the super-massive black hole (SMBH) in the centre of any galaxy is correlated to the properties of its bulge, parametrised by the spheroid luminosity \\citep[e.g.][]{Magorrian1998}, or the spheroid velocity dispersion \\citep[e.g.][]{Ferrarese2000}. This relation has a small intrinsic dispersion \\citep[e.g.][]{Gultekin2009} which implies an evolutionary connection between the SMBH and the spheroid. The mechanisms that build the super-massive black hole and the bulge of the galaxy are an active galactic nucleus (AGN) and star-formation or possibly merging episodes, respectively. There is additional evidence that the space density of AGNs and cosmic star formation have similar redshift evolution, at least up to redshifts $z\\sim2$ \\citep[e.g.][]{Chapman2005,Merloni2008}. The coeval growth of the SMBH and the host galaxy implies some causal connection between the AGN and star-formation properties \\citep[see][for a review]{Alexander2012}. Theoretical and semi-analytical models of galaxy evolution through mergers assume such a connection, where AGN feedback \\citep[e.g.][]{Hopkins2006,DiMatteo2008} plays a catalytic role. After the SMBH has grown sufficiently massive, the outflows driven by the radiation pressure of the AGN have enough energy to disrupt the cold gas supply which sustains the star formation \\citep[e.g.][]{Springel2005,King2005}, giving rise to the SMBH-bulge relation. The gas supply for both the AGN and the star formation is often thought to come from the galaxy mergers, which are ideal mechanisms for removing angular momentum from the participant galaxies and funnelling gas to the central kpc region \\citep[e.g.][]{DiMatteo2005,Barnes1996}. There is, however, growing evidence that a significant part of galaxy evolution takes place in secularly evolving systems. There is a well-defined relation between the star-formation rate and the stellar mass in local star-forming systems \\citep[see e.g.][]{Brinchmann2004,Salim2007} which defines the so-called ``main sequence'' of star formation. This relation is also found in higher redshift galaxies \\citep[e.g.][]{Elbaz2007,Daddi2007} with a redshift-dependent normalisation. It is also observed that more signs of recent merging activity are found in the morphology of starbursts (defined as star-forming galaxies with star-formation rates higher than the main sequence) than in normal (main-sequence) star-forming galaxies \\citep{Kartaltepe2012}. Mapping the star-formation in high-redshift ($z\\sim1-3$) galaxies using integral field spectroscopy, \\citet{ForsterSchreiber2009} find that about one third of those star-forming galaxies have rotation-dominated kinematics showing no signs of mergers. Moreover, \\citet{Rodighiero2011} have shown that $\\sim90\\%$ of the star-formation density at $z\\sim1-3$ takes place in the main-sequence galaxies. The hosts of AGNs do not seem to significantly deviate from this main sequence \\citep{Mullaney2012,Santini2012}. Similarly, \\citet{Grogin2005} found no apparent connection between mergers and AGN activity at redshifts $0.4\\lesssim z\\lesssim1.3$, a result which is also confirmed by \\citet{Cisternas2011} in a similar redshift range ($0.3\\lesssim z\\lesssim1.0$), and by \\citet{Kocevski2012} at higher redshifts ($1.5\\lesssim z\\lesssim2.5$). In this case, gravitational instabilities of the system may cause the transfer of material to the centre through the formation of bars and pseudo-bulges \\citep{Kormendy2004}. \\citet{Hopkins2010a} and \\citet{DiamondStanic2012} connected the black-hole accretion to the nuclear star-formation. In this study, we expand the search for an AGN-host connection to higher redshifts. Observationally the identification of a connection between the star-formation and accretion rates is challenging, especially at high redshifts. The most efficient way is to isolate the characteristic emission bands of both processes, namely the hard X-ray emission from the hot corona of the AGN and the far-infrared emission from cold dust heated by the UV radiation of massive young stars or radio synchrotron emission from electrons accelerated in supernova explosions. Previous studies using those indicators in deep fields have shown hints of a correlation \\citep[e.g.][]{Trichas2009}, which is more prominent in AGNs with higher luminosities and redshifts \\citep{Mullaney2010,Lutz2010,Shao2010}. These results argue in favour of different mechanisms, secular evolution and evolution through mergers, which take place at lower and higher redshifts (or lower and higher luminosities), respectively. \\citet{Mullaney2012} caution about the effects of both the X-ray (i.e. AGN) and the infrared (i.e. star formation) luminosities increasing with redshift, which could mimic a correlation between those values, especially in samples spanning orders of magnitudes in both $L_{\\rm x}$ and $L_{\\rm IR}$, and find no clear signs of a correlation between $L_{\\rm x}$ and $L_{\\rm IR}$ in moderate luminosity AGNs ($L_{\\rm x}=10^{42}-10^{44}\\,{\\rm erg\\,s^{-1}}$). More recently, \\citet{Mullaney2012b} do find hints of coeval growth of the super-massive black hole and the host galaxy suggesting a causal connection \\citep[see also][]{Rosario2012}. In this paper we use the deepest observations from {\\it XMM-Newton} and {\\it Herschel}, combined with {\\it Chandra} positions and deep multi-wavelength data in the CDFS to investigate the AGN-host connection, expanding to the less well-sampled region of high X-ray luminosities ($L_{\\rm x}>10^{44}\\,{\\rm erg\\,s^{-1}}$) and redshifts ($z>2.5$). We exploit the multi-wavelength information implementing an accurate SED decomposition technique to disentangle the AGN and star-formation signals in the optical and infrared bands, and therefore obtain unbiased star-formation rates for the AGN sample. We also make use of accurate {\\it XMM-Newton} spectra from the deepest 3\\,Ms observation for the first time, to investigate the nature of the AGN - star-formation relation. ", "conclusions": "We select 131 AGNs from the 3\\,Ms {\\it XMM-Newton} survey and measure their star-formation rates using long wavelength far-IR and sub-mm fluxes with rest-frame wavelength above $\\rm 20\\,\\mu m$. For 32 of the 131 sources we are able to derive only an upper limit of the star-formation rate. We take special care in modelling the spectral energy distributions, identifying and removing the AGN contribution, and derive the sSFR and stellar masses of the hosts, comparing them to the AGN properties (X-ray luminosity and absorption). Our results can be summarised as follows: \\begin{enumerate} \\item We find no evidence for a correlation between the sSFR and the X-ray luminosity for sources with $L_{\\rm x}\\lesssim10^{43.5}{\\rm \\,erg\\,s^{-1}}$ and at $z\\lesssim1$. \\item We find a correlation between the sSFR and the X-ray luminosity for sources with $L_{\\rm x}\\gtrsim10^{43}{\\rm \\,erg\\,s^{-1}}$ and at $z\\gtrsim1$. There is no indication that this correlation is a result of a redshift effect, as it is present even when we divide the data into narrow redshift bins. We argue that it is instead a result of the AGN-host co-evolution. which is more prominent for higher luminosity systems, confirming previous results. \\item We do not find any correlation between the star-formation rate (or the specific SFR, or the ``starburstiness'') and the X-ray absorption derived from high-quality {\\it XMM-Newton} spectra, at any redshift or X-ray luminosity. We assume that this is an indication that the X-ray absorption is linked to the nuclear region, and the star-formation to the host. \\item Comparing the sSFR of the hosts to the characteristic sSFR of star-forming galaxies at the same redshift (``main sequence'') we find that the AGNs reside mostly in main-sequence and starburst galaxies, with the mean specific SFR being close the limit between main-sequence and starburst hosts. This reflects the AGN-starburst connection. \\item Higher X-ray luminosity AGNs (X-ray QSOs with $L_{\\rm x}>10^{44}{\\rm \\,erg\\,s^{-1}}$) are found in starburst hosts with average sSFR more than double that of the ``main sequence'' at any redshift above $z\\approx2$. At lower redshifts ($z\\approx1.5$) we find a number of QSOs with low sSFR values, which drive the mean starburstiness of QSOs to a value consistent with that of the overall AGN population. \\item We test the reliability of the colour-magnitude diagram in assessing the host properties, and find a significant anti-correlation between the ``redness'' (deviation of the rest-frame colours from the line dividing red and blue galaxies, without any correction for AGN contribution or dust extinction), and the ``starburstiness'' (the sSFR divided by the ``main sequence'' sSFR at a given redshift). \\end{enumerate}" }, "1207/1207.0015_arXiv.txt": { "abstract": "{The emergence of several unexpected large-scale features in the cosmic microwave background (CMB) has pointed to possible new physics driving the origin of density fluctuations in the early Universe and their evolution into the large-scale structure we see today. } {In this paper, we focus our attention on the possible absence of angular correlation in the CMB anisotropies at angles larger than $\\sim 60^\\circ$, and consider whether this feature may be the signature of fluctuations expected in the $R_{\\rm h}=ct$ Universe. } {We calculate the CMB angular correlation function for a fluctuation spectrum expected from growth in a Universe whose dynamics is constrained by the equation-of-state $p=-\\rho/3$, where $p$ and $\\rho$ are the total pressure and density, respectively. } {We find that, though the disparity between the predictions of $\\Lambda$CDM and the WMAP sky may be due to cosmic variance, it may also be due to an absence of inflation. The classic horizon problem does not exist in the $R_{\\rm h}=ct$ Universe, so a period of exponential growth was not necessary in this cosmology in order to account for the general uniformity of the CMB (save for the aforementioned tiny fluctuations of 1 part in 100,000 in the WMAP relic signal). } {We show that the $R_{\\rm h}=ct$ Universe without inflation can account for the apparent absence in CMB angular correlation at angles $\\theta\\ga 60^\\circ$ without invoking cosmic variance, providing additional motivation for pursuing this cosmology as a viable description of nature. } ", "introduction": "The high signal-to-noise maps of the cosmic microwave background (CMB) anisotropies, particularly those produced by the Wilkinson Microwave Anisotropy Probe (WMAP; Bennett et al. 2003; Spergel et al. 2003) and, more recently, by {\\it Planck} (Planck Collaboration XV 2013), have revolutionized our ability to probe the Universe on its largest scales. In the near future, even higher resolution temperature maps and high-resolution polarization maps, perhaps also tomographic 21-cm observations, will extend our knowledge of the Universe's spacetime and its fluctuations to a deeper level, possibly probing beyond the surface of last scattering. Yet the emergence of greater detail in these all-sky maps has revealed several possible unexpected features on large scales, some of which were first reported by the Cosmic Background Explorer (COBE) Differential Microwave Radiometer (DMR) collaboration (Wright et al. 1996). These include an apparent alignment of the largest modes of CMB anisotropy, as well as unusually low angular correlations on the largest scales. Though viewed as significant anomalies at first (Spergel et al. 2003), these unexpected features may now be explained as possibly being due to cosmic variance within the standard model (Bennett et al. 2013). However, they may also be interesting and important for several reasons. Chief among them is the widely held view that the large-scale structure in the present Universe developed via the process of gravitational instability from tiny primordial fluctuations in energy density. The temperature fluctuations in the CMB, emerging several hundred thousand years after the big bang, are thought to be associated with the high-redshift precursors of the fluctuations that generated the galaxies and clusters we see today. Therefore, an absence of correlations in the CMB anisotropies may hint at required modifications to the standard model ($\\Lambda$CDM), or possibly even new physics, each of which may alter our view of how the Universe came into existence and how it evolved from its earliest moments to the present state. Our focus in this paper will be the possible absence of angular correlation in the CMB at angles larger than $\\sim 60^\\circ$. This feature may be anomalous because the absence of any angular correlation at the largest scales would be at odds with the inflationary paradigm (Guth 1981; Linde 1982). But without inflation, $\\Lambda$CDM simply could not account for the apparent uniformity of the CMB (other than fluctuations at the level of 1 part in 100,000 seen in the WMAP relic signal) across the sky. Thus, if variance is not the cause of the apparent disparity, the standard model of cosmology would be caught between contradictory observational constraints. In this paper, we will therefore explore the possibility that these possible CMB anomalies might be understood within the context of the recently introduced $R_{\\rm h}=ct$ Universe. This cosmology is motivated by a strict adherence to the requirements of the Cosmological Principle and the Weyl postulate, which together suggest that the Universe must be expanding at a constant rate. Additional theoretical support for this conclusion was reached with the recent demonstration that the Friedmann-Robertson-Walker (FRW) metric is apparently only valid for a perfect fluid with zero active mass, i.e., with $\\rho+3p=0$, in terms of the total energy density $\\rho$ and pressure $p$ (Melia 2013b); this is the equation-of-state that gives rise to the $R_{\\rm h}=ct$ condition. We recently showed that the horizon problem, so evident in $\\Lambda$CDM, actually does not exist in the $R_{\\rm h}=ct$ Universe (Melia 2013a), so inflation is not required to bring the CMB into thermal equilibrium following the big bang. The $R_{\\rm h}=ct$ Universe without inflation should therefore provide a meaningful alternative to $\\Lambda$CDM for the purpose of interpreting the CMB angular correlations. ", "conclusions": "It is essential for us to identify the key physical ingredient that guides the behavior of a diagnostic as complicated as $C(\\cos\\theta)$ in Figs.~1 through 5. An episode of inflation early in the Universe's history would have driven all fluctuations to grow, whether in the radiation dominated era, or later during the matter dominated expansion, to very large opening angles, producing a significant angular correlation on all scales. The {\\it Planck} data reproduced in figure~1 (and also the earlier WMAP observations) show that this excessive expansion may not have occurred, if the difference between theory and observations is not due solely to cosmic variance. In the $R_h=ct$ Universe, on the other hand, there was never any inflationary expansion (Melia 2013a), so there was a limit to how large the fluctuations could have grown by the time ($t_e$) the CMB was produced at the surface of last scattering. This limit was attained when fluctuations of size $\\lambda/2\\pi$ had reached the gravitational horizon $R_{\\rm h}(t_e)$. And for a ratio $t_0/t_e\\sim 35,000$--$40,000$, this corresponds to a maximum fluctuation angle $\\theta_{\\rm max}\\sim 30$--$35^\\circ$. This limit is the key ingredient responsible for the shape of the angular correlation function seen in Figs.~2, 3, 4 and 5. Though other influences, such as Doppler shifts, the growth of adiabatic perturbations, and the integrated Sachs-Wolfe effect, have yet to be included in these calculations, they are not expected---on the basis of previous work---to be dominant; they would modify the shape of $C(\\cos\\theta)$ only slightly (perhaps even bringing the theoretical curve closer to the data). The positive comparison between the observed and calculated $C(\\cos\\theta)$ curves seen in these figures offers some support to the viability of the $R_{\\rm h}=ct$ Universe as the correct description of nature. One might also wonder whether the observed lack of angular correlation and the alignment of quadrupole and octopole moments are somehow related. This question was the subject of Raki\\'c \\& Schwarz's analysis (Raki\\'c \\& Schwarz 2007), which concluded that the answer is probably no. More specifically, they inferred that having one does not imply a larger or smaller probability of having the other. However, this analysis was rather simplistic, in the sense that it did not consider whether alternative cosmologies, such as the $R_{\\rm h}=ct$ Universe, could produce the observed alignment as a result of the $\\sim R_{\\rm h}(t_e)$ (non-inflated) fluctuation-size limit, in which case the two anomalies would in fact be related, though only indirectly. In related work, it was shown by Sarkar et al. (2011) that there is no statistically significant correlation in $\\Lambda$CDM between the missing power on large angular scales and the alignment of the $l=2$ and $l=3$ multipoles. If not due to variance, the inconsistency between the standard model and the WMAP data may therefore be greater than each of the anomalies alone, because their combined statistical significance is equal to the product of their individual significances. As pointed out by Sarkar et al. (2011), such an outcome would require a causal explanation. In this paper, we have shown that at least one of these anomalies is not generic to all FRW cosmologies. In fact, the observed angular correlation function is a good match to that predicted in the $R_{\\rm h}=ct$ Universe. This property of the CMB might be pointing to the existence of a maximum angular size $\\theta_{\\rm max}$ for the large-scale fluctuations, imposed by the gravitational horizon $R_{\\rm h}$ at the time $t_e$ of last scattering." }, "1207/1207.2226_arXiv.txt": { "abstract": "In 2003--2012, the INTEGRAL observatory has performed long-term observations of the Large Magellanic Cloud (LMC). At present, this is one of the deepest hard X-ray (20--60 keV) surveys of extragalactic fields in which more than 20 sources of different natures have been detected. We present the results of a statistical analysis of the population of high-mass X-ray binaries in the LMC and active galactic nuclei (AGNs) observed in its direction. The hard X-ray luminosity function of high-mass X-ray binaries is shown to be described by a power law with a slope $\\alpha\\simeq1.8$, that in agreement with the luminosity function measurements both in the LMC itself, but made in the soft X-ray energy band, and in other galaxies. At the same time, the number of detected AGNs toward the LMC turns out to be considerably smaller than the number of AGNs registered in other directions, in particular, toward the source 3C 273. The latter confirms the previously made assumption that the distribution of matter in the local Universe is nonuniform. \\englishkeywords{hard X-ray sources, high-mass X-ray binaries, active galactic nuclei} ", "introduction": "\\label{sec:intro} The all-sky survey that has been performed by the INTEGRAL observatory since 2003 in the hard X-ray ($>20$ keV) energy band has allowed one not only to discover several hundred new X-ray sources (Krivonos et al. 2007, 2010a; Bird et al. 2010) but also for the first time to carry out a fairly comprehensive analysis of the statistical properties of objects of different classes: active galactic nuclei (Sazonov et al. 2007, 2008), high-mass and low-mass X-ray binaries in the inner region of our Galaxy (Lutovinov et al. 2005; Revnivtsev et al. 2008). Over the last several years, the INTEGRAL observatory has performed ultra deep observations of several regions, virtually reaching the limits of possibilities of coded-aperture telescopes. The next improvement in the sensitivity of hard X-ray sky surveys may be expected only with the advent of grazing-incidence orbital telescopes with new-generation multilayered mirrors (NuSTAR, Astro-H, ART-XC/SRG). The field toward the Large Magellanic Cloud (LMC) observed with the INTEGRAL instruments in 2003--2004 and 2010--2012 for more than 7 Ms (the effective exposure time was $\\sim4.8$ Ms) is among the regions with deep coverage. The primary target in these observations was the remnant of Supernova 1987A in order to record the emission lines of the radioactive decay of $^{44}{\\rm Ti}$ synthesized at the time of its explosion (Grebenev et al. 2012a). However, such a long exposure time also made it possible to record more than twenty point hard X-ray sources of different natures (Grebenev et al. 2012b). The goal of this work is a statistical analysis of high-mass X-ray binaries (HMXBs) in the LMC and active galactic nuclei (AGNs) registered in its direction as well as a comparison of our results with those from observations in other energy bands and in other sky regions. A detailed description of the data from the IBIS and JEM-X telescopes of the INTEGRAL observatory (Winkler et al. 2003), a complete list of registered sources, their identification, etc. are presented in the paper of Grebenev et al. (2012b). This work is based on the data from the ISGRI detector of the IBIS telescope obtained in the 20--60 keV energy band; the image processing and reconstruction methods are described in Krivonos et al. (2010b). ", "conclusions": "\\label{sec:concl} Here, we performed a statistical analysis of the population of HMXBs in the LMC and AGNs registered in its direction. We showed that the hard X-ray (20--60 keV) luminosity function of HMXBs could be fitted by a power law with a slope $\\alpha=1.8^{+0.4}_{-0.3}$. This result is in agreement both with the measurements of the luminosity function for HMXBs in the LMC itself, but in the soft X-ray (1--10 keV) energy band, and with the predictions derived from a detailed analysis of a large number of different galaxies. At the same time, the number of detected AGNs toward the LMC turns out to be considerably smaller than the number of such objects registered in other directions, in particular, toward the source 3C 273. This, along with the excess of bright objects at low redshifts in this direction, confirms the previously made assumption that the mass in the local Universe is distributed nonuniformly (see, e.g., Krivonos et al. 2007). \\bigskip ~\\bigskip" }, "1207/1207.2835_arXiv.txt": { "abstract": "Recent numerical simulations suggest that Population III (Pop III) stars were born with masses not larger than $\\sim 100M_\\odot$ but typically $\\sim 40M_{\\odot}$. By self-consistently considering the jet generation and propagation in the envelope of these low mass Pop III stars, we find that a Pop III blue super giant star has the possibility to raise a gamma-ray burst (GRB) even though it keeps a massive hydrogen envelope. We evaluate observational characters of Pop III GRBs and predict that Pop III GRBs have the duration of $\\sim 10^5$ sec in the observer frame and the peak luminosity of $\\sim 5 \\times 10^{50}\\ {\\rm erg}\\ {\\rm sec}^{-1}$. Assuming that the $E_p-L_p$ (or $E_p-E_{\\gamma, \\rm iso}$) correlation holds for Pop III GRBs, we find that the spectrum peak energy falls $\\sim$ a few keV (or $\\sim 100$ keV) in the observer frame. We discuss the detectability of Pop III GRBs by future satellite missions such as {\\it EXIST} and {\\it Lobster}. If the $E_p-E_{\\gamma, \\rm iso}$ correlation holds, we have the possibility to detect Pop III GRBs at $z \\sim 9$ as long duration X-ray rich GRBs by {\\it EXIST}. On the other hand, if the $E_p-L_p$ correlation holds, we have the possibility to detect Pop III GRBs up to $z \\sim 19$ as long duration X-ray flashes by {\\it Lobster}. ", "introduction": "Gamma ray bursts (GRB) are the brightest phenomena in the universe. Long-soft type GRBs are considered to originate from deaths of massive stars such as Wolf-Rayet (WR) stars \\citep{2011arXiv1104.2274H}. The most widely accepted scenario for long GRBs is the collapsar scenario \\citep{1993ApJ...405..273W, 1999ApJ...524..262M}. In this model, after the gravitational collapse of a massive stellar core, a black hole and an accretion disk system is formed and it launches a relativistic jet by magnetic field or neutrino pair annihilation process. If the jet can break out the stellar envelope successfully, a GRB is raised by converting the jet kinetic energy into the radiation energy. Owing to their brightness and detections at high redshift universe, GRBs are expected to be one of the powerful tools to probe the early universe. The development of observational instruments and early follow-up systems enable us to discover some high redshift GRBs. The most distant one ever is GRB 090429B at $z=9.4$ \\citep{2011ApJ...736....7C} and GRB 090423 at $z=8.3$ (e.g. \\citealt{2009Natur.461.1254T, 2009Natur.461.1258S, 2010ApJ...712L..31C}) follows it. If GRBs can be raised by first stars and detected, we will draw informations about the early universe, e.g., the star formation history and the reionization history. First stars in the universe, so called Population III~(Pop III) stars, are considered to be formed from metal free gas in the very early universe. The metal-free primordial gas cools less efficiently compared to the metal-contained present-day gas, which allows the primordial gas to have larger fragmentation masses. Since it was considered that the whole fragmented gas clump collapsed to form a single star, Pop III stars were theoretically predicted to be very massive $\\sim 100 - 1000 M_\\odot$ \\citep{2002Sci...295...93A, 2002ApJ...564...23B}. However, recent studies suggest that this is not always the case and that a massive gas clump can experience further fragmentation to form a binary system \\citep{2009Sci...325..601T, 2010MNRAS.403..45, 2011ApJ...727..110C}. \\cite{2004ApJ...603..383T} and \\cite{2008ApJ...681..771M} suggested that the UV radiation from the central protostar can ionize the surrounding neutral gas and suppress the accretion onto the protostar. They analytically investigated this feed back effect on the protostar evolution and found that Pop III stars finally obtain mass typically $\\sim 140 M_{\\odot}$. More recently, \\cite{2011Sci...334.1250H} performed two-dimensional simulations of the protostar evolution including the above feed back effect and found that Pop III stars finally obtain masses typically $\\sim 40 M_{\\odot}$. They also concluded that the UV radiation from the central star eventually stops the mass accretion and the growth of the star by the evaporation of the surrounding gas. It is considered that metal free Pop III stars do not lose mass, keeping large hydrogen envelopes until the pre-supernova stage, because of the low opacity envelopes \\citep{2002RvMP...74.1015W}. The final fate of a Pop III star depends on the stellar mass \\citep{2003ApJ...591..288H}. After the stellar core collapse, those Pop III stars in the range of $10M_\\odot \\lesssim M \\lesssim 25 M_\\odot$ would explode as supernovae and form neutron stars as remnants. Those in $25M_\\odot \\lesssim M \\lesssim 40 M_\\odot$ form black holes as remnants after the fall back accretion of the envelopes onto the temporally formed neutron stars. More massive stars ($40M_\\odot \\lesssim M \\lesssim 140 M_\\odot$ and $260M_\\odot \\lesssim M$) would fail to blow out their envelopes and promtoly form massive black holes, except those stars in $140M_\\odot \\lesssim M \\lesssim 260 M_\\odot$ who end as pair-instability supernovae due to the explosive nucleosynthesis. These remnant massive black holes are expected to raise various violent phenomena \\citep{2001ApJ...550..372F,2007PASJ...59..771S}. There have been some studies about the productivity of GRBs from massive Pop III stars ($\\gtrsim 100M_{\\odot}$) \\citep{2010ApJ...715..967M,2010MNRAS.402L..25K,2011ApJ...726..107S,2012ApJ...754...85N}. The former two studies assumed a massive black hole surrounded by an accretion disk as an outcome of a massive stellar collapse, and estimated the accretion rate onto the black hole. Then they evaluated the jet luminosity and showed that the burst activity of a Pop III star is observable by current detectors. In \\cite{2011ApJ...726..107S}, they analytically studied the jet propagation in the stellar envelope and showed that massive Pop III stars ($\\sim 900 M_\\odot$) can produce GRBs although they have large hydrogen envelopes, since the long lasting accretion provides enough energy and time for the successful jet breakout. In addition, \\cite{2012ApJ...754...85N} performed two-dimensional relativistic hydrodynamic simulations in which the accretion onto a black hole and the jet production are treated in a self-consistent way for stellar models of massive Pop III stars ($915M_\\odot$), Wolf-Rayet stars (initially $16 M_{\\odot}$), and low mass Pop III stars ($40M_\\odot$). They confirmed the validity of the analytic results in \\cite{2011ApJ...726..107S} and also found that $40 M_{\\odot}$ Pop III stars can be progenitors of GRBs, but did not study their observational characteristics and detectability. The idea of GRBs from blue super giants (BSG) was suggested in \\cite{2001ApJ...556L..37M}. Although they considered the jet dynamics in the stellar envelope, they treated a steady jet and did not reflect the central engine activity caused by the change of the accretion rate. They did not evaluate the possibility of GRBs from BSGs quantitatively. On the other hand, \\cite{2012ApJ...752...32W} discussed gamma-ray transients from Pop III BSG collapsars by investigating the mass accretion of the outermost layers of a star, but did not discuss the jet propagation and the jet break out. Assuming that the conversion efficiency of the accretion energy to the radiation energy $\\sim 10^{-2}$, they found that Pop III BSGs can produce long gamma-ray transients with duration $10^{4-5}$ sec and luminosity $10^{48-49}$ erg ${\\rm sec}^{-1}$. In this paper, we simultaneously investigate both aspects (the jet propagation and the central engine activity) in a self-consistent way by including the following physical processes; the stellar collapse, the non-steady jet injection, and the jet propagation in the stellar envelope. By doing this, we quantitatively discuss the possibility of the jet break out and GRB especially for low mass Pop III stars (around $40 M_{\\odot}$). In \\S2, after introducing the stellar models and the jet propagation models, we investigate the productivity of a GRB focusing on a $40 M_{\\odot}$ Pop III star, which is a Pop III star with the typical mass reported by the state-of-the-art simulation done in \\cite{2011Sci...334.1250H}. In \\S3, we calculate the observational characters, such as the duration $T_{90}$, the peak luminosity $L_p$ and the spectrum peak energy in the observer frame $E_p^{\\rm obs}$, of GRBs from $40 M_{\\odot}$ Pop III progenitors. Then we evaluate the detectability of such Pop III GRBs by future detectors such as {\\it Lobster} and {\\it EXIST} in detail, varying the redshift of a burst. We apply the above discussions to different progenitor models with masses of $30-90 M_{\\odot}$. In the last part of \\S3, we evaluate the light curves of Pop III GRB radio afterglow emissions and their detectability by the Low Frequency Array (LOFAR) and the Expanded Very Large Array (EVLA). \\S4 is devoted to the summary and discussions. ", "conclusions": "GRBs are the brightest phenomena in the universe. Their detections at high $z$ universe ($z \\sim 9$) motivate us to expect GRBs to be one of the powerful tools to probe the early universe. Focusing on the high $z$ universe, we should consider the association of GRBs with Pop III stars. Recent numerical simulations \\citep{2011Sci...334.1250H} suggest that Pop III stars obtain mass typically $\\sim 40M_{\\odot}$ at their birth. Zero metallicity stars are considered not to lose mass during entire life because of the low opacity envelopes \\citep{2002RvMP...74.1015W}. Therefore, they enter into the pre-supernova stage keeping large hydrogen envelopes. According to \\cite{2002RvMP...74.1015W} and \\cite{2010ApJ...724..341H}, Pop III stars end their lives as BSGs or RSGs depending on the amount of primary nitrogen produced in the shell burning. In this paper, we investigate whether such low mass Pop III stars ranging from 30 to 90$M_{\\odot}$ can be progenitors of GRBs. For this purpose, we consider the jet propagation in the stellar envelope and analytically calculate the evolution of the jet-cocoon structure. In BSG envelopes, the jet head velocity is larger than the cocoon velocity all the way except for the very early time and we can expect a successful jet breakout. On the other hand, in RSG envelopes, the cocoon edge reaches the stellar surface as early as or even earlier than the jet head. We confirm that Pop III RSGs have enough largely extended envelopes for jets to be stalled on the way and cannot raise GRBs as shown in \\cite{2003MNRAS.345..575M}. We also confirm that Pop III BSGs are compact enough for the successful jet breakout and have the possibility to raise GRBs as suggested in \\cite{2001ApJ...556L..37M} and \\cite{2012ApJ...752...32W}. It should be noted that the BSG models from Woosley and Heger used above ignored the effect of the rotation on the stellar evolution. \\cite{2008A&A...489..685E} found that when the rotation is included, Pop III stars within our noticed mass range end up as RSGs not BSGs. But \\cite{2008A&A...489..685E} considered the evolution of a star with an extremely high rotation velocity as one half the critical velocity. Recent cosmological simulations (\\citealt{2010MNRAS.403..45} and \\citealt{2011ApJ...727..110C}), however, suggested that Pop III stars are born in binary systems. In this case, the angular momentum which the star forming gas clump initially has is divided into the spin of each star and the orbital angular momentum so that these stars may rotate less rapidly. Therefore, we think that such rapidly rotating stars they considered are rare and the calculations based on BSG models from Woosley and Heger make sense. Using our model, we evaluate observational characters of Pop III GRBs. We predict that although Pop III GRBs radiate as much energy as the most energetic local long GRBs, Pop III GRBs are slightly less luminous than local long GRBs due to their much longer burst duration. Assuming that the $E_p-L_p$ (or $E_p-E_{\\gamma,{\\rm iso}}$) correlation holds for Pop III GRBs, we predict that Pop III GRBs have the much softer (or mildly softer) spectra than local long GRBs in the observer frame. \\cite{2012ApJ...752...32W} predicted that the gamma-ray transients from low metallicity BSGs have duration of $10^{4-5}$ sec and the luminosity of $10^{48-49}$ erg ${\\rm sec}^{-1}$, similar to Pop III GRBs considered here. Their transients are fed by the mass accretion of the outer most layers of stars and the accretion rate onto the BH ($\\sim 10^{-4} M_{\\odot} \\ {\\rm sec}^{-1}$) is much smaller than that in the Pop III GRB case, which imply typically from $\\sim 10^{-3} M_{\\odot} \\ {\\rm sec}^{-1}$ to $\\sim 10^{-2} M_{\\odot} \\ {\\rm sec}^{-1}$. Moreover, they simply assumed a central engine model with the roughly estimated conversion efficiency from the mass accretion to the jet energy as $\\sim 0.1$, while we consider a central engine which is driven by the magnetic process implicitly taking the disk accretion into account with the efficiency of $\\eta \\sim 6.2 \\times 10^{-4}$.\\footnote{Note that \\cite{2012ApJ...752...32W} considered the rotationally supported disk structure and the mass accretion from it so that the meaning of the conversion efficiency is different from ours.} Note that we choose this value so as for Wolf-Rayet stars to reproduce the energetics of local long GRBs. Therefore, although the characters are similar among them, we expect that gamma-ray transients considered in \\cite{2012ApJ...752...32W} are different events from GRBs considered in this paper. We also discuss the detectability of Pop III GRBs by future satellite missions such as {\\it Lobster} and {\\it EXIST} in detail. If the $E_p-E_{\\gamma,{\\rm iso}}$ correlation holds, we have the possibility to detect Pop III GRBs at redshifts $z \\sim 9$ as long duration X-ray rich GRBs by {\\it EXIST}. On the other hand, if the $E_p-L_p$ correlation holds, we have the possibility to detect Pop III GRBs up to $z \\sim 19$ as long duration X-ray flashes by {\\it Lobster}. We briefly comment the expected observable GRB rate per year by {\\it Lobster} using the results of \\cite{2011AA...533A..32D}. We calculate the observed GRB rate per year $dN_{\\rm GRB}^{\\rm obs}/dz$ as \\begin{equation} \\frac{dN_{\\rm GRB}^{\\rm obs}}{dz} = \\frac{\\Omega_{\\rm obs}}{4 \\pi} \\eta_{\\rm beam} \\frac{dN_{\\rm GRB}}{dz}, \\end{equation} where $dN_{\\rm GRB}/dz$, $\\Omega_{\\rm obs}$ and $\\eta_{\\rm beam}$ correspond to the intrinsic GRB rate (the number of on-axis and off-axis GRBs) per year, the detector field of view and the beaming factor of the burst. In Fig. 6 of \\cite{2011AA...533A..32D}, they showed $dN_{\\rm GRB}/dz$ for an {\\it optimistic} case and we use their values. Here, we also adopt the values of $\\eta_{\\rm beam} \\sim 0.01$ and $\\Omega_{\\rm obs} \\sim 0.5$ sr for {\\it Lobster} \\citep{2012IAUS..285...41G}. Optimistically speaking, we predict that {\\it Lobster} detects about 40, 4 and 0.4 Pop III GRBs per year at $z=9, 14$ and 19, respectively. At last, we briefly discuss employed assumptions in this paper. Firstly, we assume that all the stars considered in this paper ($30-90M_{\\odot}$ Pop III stars) form black holes directly after the stellar core collapse (see \\S3.3). It should be noted, however, that this is not always the case especially for less massive stars. Shortly after the onset of the core collapse, a neutron star and a shock wave, which propagates outward or is stalled, are considered to be formed at first. Behind the shock wave, a fall back accretion of the shocked envelope could be present and the continuous accretion onto the neutron star eventually leads to a black hole formation. Although the early activity of the central engine could be affected by the accretion details, i.e., the direct accretion or fall-back accretion, the conclusion of this paper is hardly changed. This is because the mass accretion at the interested time in this paper is coming from the massive envelope so that the central engine already collapsed to a black hole at the corresponding time. In addition, since the energy budget of the shock head is dominated by the envelope accretion, the details of the early phase does not affect the shock evolution in the late phase. For a more massive ($\\gtrsim 40M_{\\odot}$) star, on the other hand, the energy of the shock wave is too low to explode even the portion of the envelope, so a black hole would be formed directly and our assumption is fully justified in this case. Note that the mass threshold between the direct or fall-back induced black hole formation is still under the debate (see e.g. \\citealt{1999ApJ...522..413F}) and beyond the scope of this paper. Secondly, we assume that the whole stellar envelope accretes onto the BH (see \\S3.1). In order to confirm the validity of this assumption, we evaluate the binding energy of each layer of the stellar envelope and compare it with the typical energy of a supernova outgoing shock $\\sim 10^{51}$ erg, which is injected around $\\sim 10$ km from the center. We find that the binding energy becomes larger than $10^{51}$ erg within $r \\lesssim 5 \\times 10^9$ cm. This means that the outgoing shock should stall on the way and we expect little mass ejection. Recently, \\cite{2012MNRAS.423L..92Q} suggested that in the late stage of the stellar evolution, the pre-supernova burning leads to a significant mass ejection from the outer envelope. But they considered only the case of a $40 M_{\\odot}$ star with metallicity $Z = 10^{-4}$ and commented that the amount of mass ejection depends on the metallicity and rotation etc. Thus, the amount of the ejectable mass is uncertain for the progenitors employed here so that we do not consider this effect in this paper." }, "1207/1207.5909_arXiv.txt": { "abstract": "{Direct imaging of circumstellar disks requires high-contrast and high-resolution techniques. The angular differential imaging (hereafter ADI) technique is one of them, initially developed for point-like sources but now increasingly applied to extended objects such as disks. This new field of application raises many questions because the disk images reduced with ADI depend strongly on the amplitude of field rotation and the ADI data reduction strategy. Both of them directly affect the disk observable properties. } {Our aim is to characterize the applicability and biases of some ADI data reduction strategies for different disk morphologies. A particular emphasis is placed on parameters mostly used for disks such as their surface brightness distribution, their width if the disk is a ring, and local features such as gaps or asymmetries. We first present a general method for predicting and quantifying those biases. In a second step we illustrate them for some widely used ADI algorithms applied to typical debris disk morphologies: inclined rings with various inner/outer slopes and width. Last, our aim is also to propose improvements of classical ADI to limit the biases on extended objects. } {Simulated fake disks seen under various observing conditions were used to reduce ADI data and quantify the resulting biases. These conclusions are complemented by previous results from NaCo L' real-disk images of HR\\,4796A. } {As expected, ADI induces flux losses on disks. This makes this technique appropriate only for low- to medium-inclination disks. A theoretical criterion is derived to predict the amount of flux loss for a given disk morphology, and quantitative estimates of the biases are given in some specific configurations. These biases alter the disk observable properties, such as the slopes of the disk surface brightness or the radial/azimuthal extent of the disk. Additionally, this work demonstrates that ADI can very easily create artificial features without involving astrophysical processes. For example, a particularly striking feature appears for a ring when the amplitude of field rotation is too small. The two ring ansae are surrounded by two flux-depleted regions, which makes them appear as bright blobs. This observation does not require any astrophysical process such as dust blown by radiation pressure, as proposed previously in H-band images of HR\\,4796A. } {The ADI techniques behave as spatial filtering algorithms and can bias disk observables. Therefore, the filtering process needs to be properly calibrated when deriving disk parameters from processed images. % } ", "introduction": "The study of circumstellar disks is essential for understanding the formation of planetary systems. Direct imaging of these disks has revealed asymmetries, warps, gaps, troncatures, density waves, or other features that are the results of interactions between the disk and its environment. About 160 disks are now resolved from their visible to thermal emission (http://www.circumstellardisks.org, Stapelfeldt). Debris disks are a particularly interesting class of optically thin disks since planets are already formed, if any, and faint structures in the dust distribution can betray their presence, as proposed for HR\\,4796A in \\citet{Lagrange2012b}. The vicinity of the star and its luminosity however limit the performance of both ground-based \\citep{Racine1999} and space-based \\citep{Schneider2003} high-contrast imaging because of bright quasi-static speckles. Slowly evolving, they add up over time and eventually become the dominant noise source at separations below a few arcseconds, depending on the star brightness \\citep{Macintosh2005} The principle of differential imaging is to subtract a reference frame from the target image to reduce quasi-static speckle noise \\citep{Marois2006}. This reference frame is sometimes also called a reference Point Spread Function \\citep{Lafreniere2007}, hereafter PSF, but we will use here the more generic term reference frame. This frame can be obtained in various ways: either with a reference star or with the target itself observed at a different field of view orientation (ADI), a different wavelength (spectral differential imaging), or polarization (polarimetric differential imaging). This paper focuses on ADI. Angular differential imaging has already achieved significant results on debris disks: \\citet{Buenzli2010} for instance detailed the morphology of the Moth, \\citet{Boccaletti2012} and \\citet{Currie2012} revealed the inner part of HD\\,32297, \\citet{Thalmann2011} and \\citet{Lagrange2012b} studied the ring of HR\\,4796A, \\citet{Lagrange2012} revealed the inner part of the $\\beta$ Pictoris disk and constrained the projected planet position relative to the disk, and \\citet{Thalmann2010} resolved the gap in the transitional disk around LkCa\\,15. However, side effects might alter the apparent morphology of the disk and its observable properties: flux self-subtraction, change in disks slopes, width, etc. The model-matching procedure described by \\citet{Boccaletti2012} to capture the true morphology of HD\\,32297 is a good illustration of how difficult it is to dismiss ADI artifacts. Angular differential imaging uses the differential rotation between the field of view and the pupil occurring on an alt-azimuthal mount to distinguish between the speckle halo and any on-sky source. When the pupil tracking mode is used, the field rotates at the same rate as the parallactic angle while the pupil is stable. The parallactic angle is the angle between the great circle through the object and the zenith, and the hour circle of the object. When extended objects such as disks are imaged, the challenge of ADI is to build reference frames with a speckle pattern highly correlated to the target image but without capturing the flux of the disk. If the disk flux is captured in the reference frame, the reference subtraction will decrease the disk flux in the residual image: this is the problem of self-subtraction. Not only does it lead to a lower signal-to-noise ratio (S/N) for the disk but it also biases the observable parameters of the disk. The objective of this study is twofold: \\begin{enumerate} \\item To guide the observing strategy and data reduction by providing key figures to know how much flux loss is expected for a given disk geometry and field of view rotation. A theoretical criterion is derived to quantify this loss for the specific case of an edge-on disk. \\item To highlight the observable parameters that can be trusted from ADI-based disk images and point out the artifacts potentially created by this technique. \\end{enumerate} We will first describe the simulation procedure, and in Section 3 will build theoretical criteria to analyze the applicability of ADI to disks. A qualitative (Section \\ref{SectionQualitative}) and quantitative (Section \\ref{SectionQuantitative}) description of the biases induced by ADI is then presented before reviewing specific ADI algorithms more adapted for disk reductions in Section \\ref{SectionImprovements}. Systematically investigating the parameter space of these algorithms to derive the most accurate disk parameters is beyond the scope of that section. ", "conclusions": "We discussed the effects of ADI applied to extended circumstellar disks. A theoretical study demonstrated that the field of application of this imaging technique depends on two parameters: the available amplitude of field rotation and the geometrical extension of the object in azimuth. This analysis gives a clear indication of the required amplitude of field rotation to reduce a given disk depending on its morphology, and conversely it gives an indication of the disk minimum inclination for a given field rotation. With reasonable observing time, objects inclined less than $50^\\circ$ are out of the scope of ADI. For all others, self-subtraction is a problem that observers must bear in mind when reducing the data. It affects the whole geometry of the disk by reducing the total disk flux, but also introduces local effects difficult to distinguish from speckle noise patterns. A typical effect already observed for some circumstellar disks reduced with ADI is the enhancement of the ansae for an inclined ring morphology, creating bright blobs along the disk major axis that can be easily misinterpreted. Brightness asymmetries from non-uniform parallactic angle evolution are other striking effects that arise from ADI, especially rADI in which a limited number of images is used to build the reference frames. Measurable astrophysical quantities used to characterize disks such as inner/outer slopes of the SBD or ring FWHM are biased and their estimation must be calibrated. We here provided an insight into how these parameters are affected, depending on the initial disk morphology and the reduction parameters. To evaluate the efficiency of the explored ADI algorithms applied on disks, two key points must be considered: the amount of disk self-subtraction and how well the residual speckles are attenuated. For the first point, masking as used in mcADI is our preferred data reduction strategy, since it replaces the classical separation criterion $N_\\delta$ appropriate for point sources with a binary mask adapted to the expected extension of the object, to efficiently limit self-subtraction. The second point is better addressed with optimization algorithms such as LOCI, but flux losses are much more severe. Our study shows that using larger optimization areas, with an azimuthal extent greater than the radial extent, helps to preserve the disk flux. More importantly constraining the LOCI coefficients to be positive proves to be very effective in this goal." }, "1207/1207.5412_arXiv.txt": { "abstract": "{The smoothed particle hydrodynamics (SPH) technique is a well-known numerical method that has been applied to simulating the evolution of a wide variety of systems. Modern astrophysical applications of the method rely on the Lagrangian formulation of fluid Euler equations, which is fully conservative. A different scheme, based on a matrix approach to the SPH equations is currently being used in computational fluid dynamics (CDF). An original matrix formulation of SPH based on an integral approach to the derivatives, called IAD$_0$, has been recently proposed and is fully conservative and well-suited to simulating astrophysical processes.} {The behavior of the IAD$_0$ scheme is analyzed in connection with several astrophysical scenarios, and compared to the same simulations carried out with the standard SPH technique.} {The proposed hydrodynamic scheme is validated using a variety of numerical tests that cover important topics in astrophysics, such as the evolution of supernova remnants, the stability of self-gravitating bodies, and the coalescence of compact objects.} {The analysis of the hydrodynamical simulations of the above-mentioned astrophysical scenarios suggests that the SPH scheme built with the integral approach to the derivatives improves the results of the standard SPH technique. In particular, there is a better development of hydrodynamic instabilities, a good description of self-gravitating structures in equilibrium and a reasonable description of the process of coalescence of two white dwarfs. We also observed good conservations of energy and both linear and angular momenta that were generally better than those of standard SPH. In addition the new scheme is less susceptible to pairing instability.} {We present a formalism based on a tensor approach to Euler SPH equations that we checked using a variety of three-dimensional tests of astrophysical interest. This new scheme is more accurate because of the re-normalization imposed on the interpolations, which is fully conservative and less prone to undergoing the pairing instability. The analysis of these test cases suggests that the method may improve the simulation of many astrophysical problems with only a moderate computational overload.} ", "introduction": "Multidimensional numerical hydrodynamics is one of the most powerful tools of modern astrophysics to comprehend the cosmos machinery. Among them, the smoothed particle hydrodynamics (SPH) is one of the most widely used techniques because of its ability to describe the evolution of fluids with complicated geometries and a diversity of length scales. Since it was formulated, more than thirty years ago, by \\cite{gm77} and \\cite{lucy77}, it has largely evolved incorporating, little by little, a plethora of methods that makes it competitive compared to grid-based methods of Eulerian type. Details of the modern mathematical formulation, as well as of the main features of the state-of-art of the SPH technique, can be found in the reviews by \\cite{monaghan05}, \\cite{rosswog09}, \\cite{springel10}, and \\cite{price12}. \\cite{gs12} (henceforth paper I) recently suggested that the use of matrix methods, \\cite{dilts99}, in astrophysics could improve the simulations with the SPH technique with an affordable computational cost. In this work, we focus on specific three-dimensional (3D) astrophysical applications of the scheme formulated in paper I. In particular, we choose examples from different fields of astronomy to check the method and highlight its potential advantages over the standard SPH scheme. The suitability of IAD$_0$ for describing hydrodynamic instabilities found in paper I is confirmed by simulating the evolution of a supernova remnant (SNR). As the SNR evolves embedded in an uniform background of particles with negligible gravity, there are no numerical troubles affecting the outer limits of the system. The existence of boundaries becomes relevant to the second test, which is devoted to describing the equilibrium features of polytropes with different indexes and masses. For this problem, the interplay between pressure forces and gravity becomes crucial to ensure that the obtained structures are compatible with the analytical models. We show that the tensor method leads to polytropes where the central density and radius are slightly closer to the theoretical predictions than those obtained with the standard SPH technique. We also explore the ability of IAD$_0$ to describe a very dynamical situation by simulating the coalescence of two white dwarfs. In this case, a catastrophic merging of the stars ensues after a few orbital periods. For this test, the tensor method gives results of, at least, comparable quality to those obtained using the standard SPH scheme, but displaying a more homogeneous mixing of the material of both stars. There are also some differences in the angular velocity distribution of the remnant. For a similar elapsed time, the matrix calculation does not lead to the complete rigid rotation of the core, which is, however, achieved in the simulation using the standard scheme. The text is organized as follows. In Sec.~\\ref{section2}, we review the mathematical formalism linked to IAD$_0$ and discuss its most relevant features. In Sec.~\\ref{section3}, we describe the three astrophysical tests aimed at validating the code and comparing its performance to that of standard SPH. Section~\\ref{section4} is devoted to incorporating thermal conductive transport in the tensor scheme, and to check the resulting algorithm. The benchmarking of the code is done in Sec.~\\ref{section5}. Finally, the main conclusions of our work, as well as some comments about the shortcomings of the developed scheme and future lines of improvement, are outlined in the last section, which is devoted to our conclusions. ", "conclusions": "We have checked the behavior of a novel SPH scheme where gradients are calculated using an integral approach. The main features of the technique, called IAD$_0$, were described in detail in paper I and summarized in Sec.~\\ref{section2} of this paper. The main virtue of the approach relies in that the re-normalization of the derivatives appears naturally, without any degradation of the conservative properties that characterize the SPH technique. Another relevant feature is that the basic mathematical formalism of IAD$_0$ looks very similar to that of a standard SPH technique, making its implementation quite straightforward. As commented in paper I, matrix methods based on a similar, although not identical, formulation have been used in CDF for the past decade (see for instance \\cite{dilts99}), but have never been previously applied to astrophysics. Three test cases of considerable astrophysical interest were selected to validate IAD$_0$, as well as to detect its virtues and weaknesses. The performance of the method in describing the growth of the RT instability in a supernova remnant was analyzed in Sec.~\\ref{remnant}, with the conclusion that IAD$_0$ provides a healthier development of the RT fingers. The method's success in handling hydrodynamical instabilities relies on the re-normalization imposed on the derivatives, thus confirming the results obtained in paper I using toy models. The second test was addressed especially to the applications of the method to describing stellar objects (approached as polytropes with different indexes and known analytical properties). One important question here relies in the treatment of the outer boundary of the object, which is often a controversial point in matrix methods \\cite[]{oger07}. We have entirely explored this issue using several recipes to handle the surface of the polytropes. These include the use of tensor and vector formulations of IAD$_0$, as well as hybrid schemes that use the full matrix expressions in the interior but changes to {\\sl vector-}IAD$_0$ near the surface, according to Eq.~(\\ref{tauijselect}). The best results were obtained when the full tensor IAD$_0$ scheme was used, especially for the most unstable polytrope with index n=5/2. Nevertheless, the amount of numerical noise stored as residual kinetic energy at equilibrium was larger for {\\sl tensor-}IAD$_0$ than for STD. The ability of the tensor method to handle the coalescence and further merging of stellar-like objects was checked in Sec.~\\ref{merger}. Even though for this particular test the simulations did not show any clear advantage of the matrix method, it gave a good depiction of the coalescence process. Despite both schemes being conservative by construction, complete preservation of angular momentum is impeded by the use of the hierarchical cluster method in calculating gravity. Nevertheless, there are indications that the angular momentum-component orthogonal to the orbital plane behaves better in IAD$_0$ than in the standard scheme (Fig.~\\ref{figure12}). The prompt product after the coalescence is a core surrounded by an extended diluted halo of particles moving at Keplerian velocity. The calculations show that the tensor method leads to a more homogeneous mixing of the material of both stars in the core region. As the numerical recipe for implementing AV in both calculations is exactly the same, we conclude that the different behavior is caused by the differences in the algorithm used to compute the kernel derivative. Therefore, it seems that IAD$_0$ is able to provide a more effective hydrodynamical mixing than the standard scheme, but still avoid the penetration of fluids in strong shocks, as suggested in paper I. These differences also affect the distribution of angular velocity in the core of the remnant, which in the STD calculation soon reaches rigid rotation but in the matrix one does not, for a similar elapsed time. A potential weakness of the tensor method in handling these kinds of configurations comes from the computation and further inversion of matrix $\\mathcal T$ calculated with Eq. (\\ref{tauijsph}). In the case of the isolation of a particle, as it could be for those belonging to the most diluted region of the domain, matrix $\\mathcal T$ becomes singular leading to the complete breakdown of the simulation. In this sense, the matrix scheme is less robust than the standard one. Nevertheless, current SPH schemes usually include an algorithm to ensure that the number of neighbors of a given particle remain constant during the simulation. This algorithm is necessary to reliably compute the magnitude $\\Omega$ in Eqs.~(\\ref{momentumLqij}) and (\\ref{energyqij}) and, when working properly, to avoid the singularity problems linked to matrix $\\mathcal T$. Another interesting features of the matrix method are its ability to avoid the pairing instability and the better handling of thermal conduction. The first one leads to less particle clustering, thus improving the quality of the interpolations; while in the second case, the results of Sec.~\\ref{section4} suggest that diffusion-like equations can be handled by IAD$_0$ in a better way than in the standard framework. We also conducted a benchmark test to evaluate the impact of the inclusion of the IAD$_0$ formalism on the wall-clock time for a nominal calculation. According to Table~\\ref{table3} and Fig.~\\ref{figure17}, the computational overload introduced in a serial calculation by the re-normalization of the derivatives is low ($\\simeq 20\\%$), being practically independent of the total number of particles. The scaling of the code with the number of particles remains virtually untouched, and the IAD$_0$-related sections of the code show a good behavior in front of parallelization. Therefore, the analysis of the astrophysical tests discussed in this paper support the main conclusions of paper I, where the basic implementation of the integral approach to the derivatives was discussed and checked. All this suggests that the use of the re-normalized, fully conservative IAD$_0$ approach to the SPH equations may improve, in general, the quality of the simulations in astrophysics with a little computational overload penalty. The smaller amount of viscous dissipation shown by IAD$_0$, compared to STD for the same AV formalism, suggests that it could be of interest not only when handling hydrodynamic instabilities but also when simulating turbulence. Turbulence is at the heart of many astrophysical problems, being of especial relevance to understanding star formation (\\cite{fed10}). In this respect, a full comparison of 3D astrophysical turbulence calculated with a variety of algorithms, both SPH and grid codes, was provided by \\cite{kitsionas09}, \\cite{pricefed10}, and \\cite{kritsuk11}, with the result that both families of codes give similar results for an equivalent number of resolution elements in each direction in space. Nevertheless, these experiments also show that, owing to the higher dissipation, the scaling range of SPH codes is slightly shorter than that of grid-based codes, as demonstrated by \\cite{kitsionas09}. Therefore, it may be of great interest to check the ability of the proposed IAD$_0$ scheme to handle astrophysical turbulence in the near future." }, "1207/1207.2468_arXiv.txt": { "abstract": "We use high resolution cosmological simulations of Milky Way-mass galaxies that include both baryons and dark matter to show that baryonic physics (energetic feedback from supernovae and subsequent tidal stripping) significantly reduces the dark matter mass in the central regions of luminous satellite galaxies. The reduced central masses of the simulated satellites reproduce the observed internal dynamics of Milky Way and M31 satellites as a function of luminosity. We use these realistic satellites to update predictions for the observed velocity and luminosity functions of satellites around Milky Way-mass galaxies when baryonic effects are accounted for. We also predict that field dwarf galaxies in the same luminosity range as the Milky Way classical satellites should not exhibit velocities as low as the satellites, since the field dwarfs do not experience tidal stripping. Additionally, the early formation times of the satellites compared to field galaxies at the same luminosity may be apparent in the star formation histories of the two populations. Including baryonic physics in Cold Dark Matter models naturally explains the observed low dark matter densities in the Milky Way's dwarf spheroidal population. Our simulations therefore resolve the tension between kinematics predicted in Cold Dark Matter theory and observations of satellites, without invoking alternative forms of dark matter. ", "introduction": "There are fewer small satellite galaxies orbiting our Milky Way (MW) galaxy than predicted by the favored Cold Dark Matter (CDM) cosmological model \\citep{Moore1999, Klypin1999, Madau2008, Wadepuhl2011, Brooks2013}. Theories often reconcile the discrepancy between the number of observed satellites and CDM predictions by invoking the suppression of star formation in low mass galaxies, for example by UV heating at reionization \\citep[e.g.,][]{Okamoto2008}. If only the most massive satellites form stars, this can bring the predicted number of luminous satellites down from thousands to tens, in line with observations. Even then a serious problem remains, as the most massive satellites predicted by CDM models are still much too dense compared to what we observe \\citep{Boylan-kolchin2011, Boylan-kolchin2012, Wolf2012, Hayashi2012, Tollerud2012, Martinez2013}. The tension between the observations and the predictions of the CDM model have led some researchers to propose alternative forms of dark matter (e.g., warm or self-interacting) to reduce the central masses of satellites \\citep{Maccio2010, Vogelsberger2012, Lovell2012, Anderhalden2013, Shao2013, Polisensky2013}. However, the highest resolution simulations available to date to study the internal properties of satellites include only the dark matter (DM) component of galaxies, neglecting the effects of baryons \\citep[e.g.,][]{Diemand2007, Springel2008, Boylan-kolchin2012}. \\citet{Zolotov2012} recently examined how baryons impact the DM structure in satellites around a MW-mass galaxy. They demonstrate that supernova (SN) feedback reduces the central DM densities of satellites with $M_{\\star} \\gtrsim 10^7 M_{\\odot}$ before infall \\citep[see also][]{Governato2012}. However, SN feedback alone is not enough to bring the observed densities of satellites in line with observations. After infall, the presence of a baryonic disk in the host galaxy increases the mass loss rate for all satellites via tidal stripping \\citep{Zolotov2012, Arraki2013}. This tidal effect is particularly strong for those satellites that enter with cored DM halos \\citep{Penarrubia2010}, further increasing the discrepancy in the central masses predicted by DM+baryon and DM-only simulations. Previous studies have invoked tidal stripping after infall to reduce the DM masses of satellite galaxies \\citep{Gnedin1999, Taylor2001, Hayashi2003, Kazantzidis2004, Kravtsov2004, Read2006b, Sales2007, Munoz2008, D'Onghia2010, Choi2009, Collins2011}. The additional central mass present in the parent halo of a baryonic run (due to the fact that baryons can cool to the center, unlike DM) lead to an enhanced tidal force not found in DM-only simulations, leading to enhanced stripping. These earlier works examined the evolution in the densities of satellites using idealized models, because cosmological simulations were unable to achieve similar resolutions until recently. The results of these idealized models can be used to parametrize the stripping of mass based on the orbital history of a satellite. However, even after considering the additional tidal stripping that should occur in a cosmologically motivated population of DM-only satellites, studies still could not reproduce a $z=0$ satellite population that matches the observed MW satellite kinematics. Not enough mass is stripped from the most massive satellites to bring them in line with the kinematics observed in the most luminous satellites of the MW \\citep{Read2006b, D'Onghia2010}. \\citet{Read2006b} concluded that the most massive satellites would need to have central density slopes shallower than NFW models to undergo enough stripping to make the theoretical and observational masses consistent. This is because subhalos become more prone to tidal forces as their density slopes become less steep \\citep{Hayashi2003, Kazantzidis2004, Penarrubia2010}. \\citet{Zolotov2012} are the first to use cosmological simulations to probe the combined reduction in mass from both DM core creation and tidal stripping in the satellite population around a MW-mass galaxy. In this paper, we use a complementary sample of satellites to extend the analysis presented in \\citet{Zolotov2012} and further explore the observational consequences of the model. We show that this model can, for the first time in fully cosmological simulations of MW-mass galaxies, reproduce the observed kinematics of dwarf Spheroidal (dSph) satellites in L$^*$ galaxies at $z=0$. We use our results to interpret the observed trend and scatter in the dynamics of the MW and M31 dSphs. We also demonstrate that the inclusion of baryonic physics leads to destruction of a number of luminous satellites that otherwise survive in DM-only simulations, and make new predictions for the surviving subhalo mass and luminosity functions of L$^*$ galaxies. Finally, we compare our simulated dSphs to simulated field dwarfs of the same luminosity, and make predictions for observable differences in the two populations. ", "conclusions": "We have demonstrated that simulations that account for the effects of SN feedback and enhanced tidal stripping on satellites result in a satellite population whose kinematic properties match the observed properties of the Milky Way and M31 satellites. Our findings are in sharp contrast to studies using DM-only simulations, which over-predict the central masses of satellites in comparison to observations. By directly comparing the properties of simulated satellites that include gas hydrodynamics to the same satellites in DM-only simulations, we find: \\begin{enumerate} \\item The majority of satellites simulated with DM-only have $v_{max} > 20$ km/s, grossly inconsistent with the observed kinematics of the Local Group dSph population. We note that this is despite using simulated parent halos that are on the low side of the allowed range for the MW. Although lowering the mass of the MW has been suggested as a solution to the ``too big to fail'' problem \\citep{Vera-Ciro2013, Sawala2012}, our DM-only halos with virial masses of $7-8\\times$10$^{11} M_{\\odot}$ fail to produce a satellite population with kinematics consistent with observations. In contrast, gas-free satellites simulated with baryonic physics have $v_{max}$ values in the range of $6-24$ km/s, matching the observed values of Local Group dSphs. The reduction in $v_{max}$ in these simulated galaxies is due to the combined effect of SN feedback on very luminous satellites (brighter than $M_V = -12$) and enhanced mass stripping for satellites across all luminosities in the presence of the parent galaxy's disk. \\item We find that the velocity dispersions of the simulated satellites are in good agreement with the range of observed dispersions in the MW and M31 dSph satellites. \\item DM-only simulations produce satellites with $2-4$ times more mass in the central 1 kpc than satellites simulated with baryonic physics. \\item Satellites simulated with baryons and DM reproduce the observed scatter $v_c$ for dSphs, while satellites simulated with DM-only do not. A tight $v_{max}$--luminosity relation exists for the satellites prior to infall. After infall, the tidal effects of the baryonic disk in the host galaxy lead to large scatter based on the range of infall times and orbital pericenters of the satellites. \\item Simulations that included baryonic physics have 1/3 fewer satellites that survive to z=0 than DM-only simulations. Six out of the eight DM-only satellites that have no surviving SPH counterpart have orbits that bring them within the central 30 kpc of their host galaxy. It seems therefore likely that the destruction of satellites in the baryonic run is due to tidal heating and shocking at the interface of the disk in the parent halo. These effects alter the predicted luminosity functions for satellites at $z=0$ from the DM-only case. \\item Increased mass loss in tides for the satellites in the simulation with baryons, combined with total destruction of 1/3 of the satellites, shifts the velocity function of satellites at $z=0$ relative to the DM-only predictions. The shift always acts to move satellites toward lower $v_{max}$ values than in the DM-only run. At any given $v_{max}$, we find there are 50\\% fewer satellites expected when baryonic effects are included. \\item Simulated field dwarf galaxies have systematically higher $v_{max}$ and $v_{\\rm{1kpc}}$ values at $z=0$ than satellite galaxies in the same luminosity range. Tidal forces are necessary for satellites to reach the low $v_c$ values observed. Preliminary results from observed field dwarfs suggest our simulated field dwarf velocities are consistent with observations in the local Universe. \\item We find that simulated field dwarfs have similar velocity dispersions to simulated satellite galaxies when compared at the half light radii. These findings reproduce the recent observations of \\citep{Kirby2014}. \\item Satellite galaxies have mass assembly histories that peak at higher redshifts than isolated field dwarf galaxies in the same luminosity range. These early assemblies are likely to be manifest in the cumulative SFHs of the dSphs vs dIrrs. \\end{enumerate}" }, "1207/1207.5886_arXiv.txt": { "abstract": "Evolution of the universe with modified holographic Ricci dark energy model is considered. Dependency of the equation of state parameter and deceleration parameter on the redshift and model parameters are obtained. It is shown that the density evolution of both the non-relativistic matter and dark energy are same until recent times. The evolutionary trajectories of the model for different model parameters are obtained in the statefinder planes, $r-s$ and $r-q$ planes. The present statefinder parameters are obtained for different model parameter values, using that the model is differentiated from other standard models like $\\Lambda$CDM model etc. We have also shown that the evolutionary trajectories are depending on the model parameters, and at past times the dark energy is behaving like cold dark matter, with equation of state equal to zero. \\vspace{0.2in} \\noindent {\\bf Keywords}: Dark energy, Holographic model, Statefinder diagnostic, Cosmological evolution. \\vspace{0.15in} \\noindent {\\bf PACS numbers}: 98.80.Cq, 98.65.Dx ", "introduction": "Observations of distant type Ia supernovae (SNIa) and cosmic microwave background anisotropy have shown that the present universe is accelerating \\cite{Perl98}. This expansion may be driven by a component with negative pressure, called dark energy. The simplest model of dark energy is the cosmological constant $\\Lambda$ which can fit the observations in a fair way \\cite{wein1, sahni1}, whose equation of state is $\\omega_{\\Lambda} = -1.$ during the evolution of the universe. However there are two serious problems with cosmological constant model, namely the fine tuning and the cosmic coincidence \\cite{Cope1}. To solve these problems different dynamic dark energy models have been proposed, with varying equation of state during the expansion of the universe. Holographic dark energy (HDE) is one among them \\cite{Cohen1,Hsu1, Li1}. HDE is constructed based on the holographic principle, that in quantum gravity, the entropy of a system scales not with its volume but with its surface area $L^2$, analogically the cosmological constant in Einstein's theory also is inverse of some length squared. It was shown that \\cite{Cohen1} in effective quantum field theory, the zero point energy of the system with size $L$ should not exceed the mass of a black hole with the same size, thus $L^3 \\rho_{\\Lambda} \\leq L M_P^2,$ where $\\rho_{\\Lambda}$ is the quantum zero-point energy and $M_P = 1/ \\sqrt{8\\pi G}$, is the reduced Plank mass. This inequality relation implies a link between the ultraviolet (UV) cut-off, defined through $\\rho_{\\Lambda}$ and the infrared (IR) cut-off encoded in the scale $L.$ In the context of cosmology one can take the dark energy density of the universe $\\rho_X$ as the same as the vacuum energy, i.e. $\\rho_x = \\rho_{\\Lambda}.$ The largest IR cut-off $L$ is chosen by saturating the inequality, so that the holographic energy density can be written as \\begin{equation} \\rho_x = 3c^2M_P^2L^{-2} \\end{equation} where $c$ is numerical constant. In the current literature, the IR cut-off has been taken as the Hubble horizon \\cite{Hsu1,Li1}, particle horizon and event horizon \\cite{Li1} or some generalized IR cut off \\cite{Gao1,Linsen1,Yang1}. The HDE models with Hubble horizon or particle horizon as the IR cut-off, cannot lead to the current accelerated expansion \\cite{Hsu1} of the universe. When the event horizon is taken as the length scale, the model is suffered from the following disadvantage. Future event horizon is a global concept of space-time. On the other hand density of dark energy is a local quantity. So the relation between them will pose challenges to the concept of causality. These leads to the introduction new HDE, where the length scale is given by the average radius of the Ricci scalar curvature, $R^{-1/2}.$ The holographic Ricci dark energy model introduced by Granda and Oliveros \\cite{Granda1} based on the space-time scalar curvature, is fairly good in fitting with the observational data. This model have the following advantages. First, the fine tuning problem can be avoided in this model. Moreover, the presence of event horizon is not presumed in this model, so that the causality problem can be avoided. The coincidence problem can also be solved effectively in this model. Recently a modified form of Ricci dark energy was studied \\cite{Chimento1} in connection with the dark matter interaction, and analyses the model using \\emph{Om} diagnostic. In this paper we have considered the evolution of the universe in Modified Holographic Ricci Dark Energy (MHRDE) model and obtain the statefinder parameters to discriminate this model with other standard dark energy models. Statefinder parameters is a sensitive and diagnostic tool used to discriminate various dark energy models. The Hubble parameter $H$ and deceleration parameter $q$ alone cannot discriminate various dark energy models because of the degeneracy on these parameters. Hence Sahni et al. \\cite{Sahni2} introduces a set of parameters $\\{r,s\\}$ called statefinder parameters, defined as, \\begin{equation} \\label{statefinder} r = { \\dddot a \\over a H^3}, \\, \\, \\, \\, \\, \\, \\, \\, \\, \\, \\, \\, \\, \\, \\, \\, s = {r - \\Omega_{total} \\over 3 (q - \\Omega_{total} ) /2 }, \\end{equation} where $a$ is the scale factor of the expanding universe and $\\Omega_{total}$ is the total energy density containing dark energy, energy corresponds to curvature and also matter (we are neglecting the radiation part in our analysis). In general statefinder parameter is a geometrical diagnostic such that it depends upon the expansion factor and hence on the metric describing space-time. The $r-s$ plot of dark energy models can help to differentiate and discriminate various models. For the well known $\\Lambda$CDM model, the $r-s$ trajectory is corresponds to fixed point, with $r=1$ and $s=0$ \\cite{Sahni2}. The cosmological behavior of various dark models including holographic dark energy model, were studied and differentiated in the recent literature using statefinder parameters \\cite{Setare1,Malekjani1,Huang1, Mub1}. The paper is organized as follows. In section 2, we have studied the cosmological behavior of the MHRDE model and in section 3 we have considered the statefinder diagnostic analysis followed by the conclusions in section 4. ", "conclusions": "We have studied the modified holographic Ricci dark energy (MHRDE) in flat universe, where the IR cutoff is given by the modified Ricci scalar, and the dark energy become $\\rho_x = 2(\\dot{H}+3\\alpha H^2/2)/(\\alpha - \\beta)$ where $\\alpha$ and $\\beta$ are model parameters. We have calculated the relevant cosmological parameters and their evolution and also analyzed the model form the statefinder view point for discriminating it from other models. The importance of the model is that it depends on the local quantities and thus avoids the causality problem. The density of MHRDE is comparable with the non-relativistic matter at high redshift as shown in figure \\ref{fig:densityevolve} and began to dominate at low redshifts, thus the model is free from the coincidence problem. The evolution of equation of state parameter is studied. The equation of state parameter is nearly zero at high redshift, implies that in the past universe MHRDE behaves like cold dark matter. Further evolution of equation of state is strongly depending on the model parameter $\\beta.$ If the $\\beta$ parameter is positive the equation of state is greater than -1. For negative values of $\\beta$, the equation of state cross the phantom divide $\\omega_x < -1.$ In this model the deceleration parameter starts form around 0.5 at the early times and and starts to become negative when the redshift $z<1.$ . In general we have found that in MHRDE model the universe entering the accelerating phase at times earlier (for allowed range of parameters $\\alpha$ and $\\beta$), than in the $\\Lambda$CDM model. But in particular as the model parameter $\\alpha$ increases, the universe enter the accelerating phase at relatively later times. We have applied the statefinder diagnostic to the MHRDE and plot the trajectories in the $r-s$ and $r-q$ planes. The statefinder diagnostic is a crucial tool for discriminating different dark energy models. The statefinder trajectories are depending on the model parameters $\\alpha$ and $\\beta.$ For positive values of $\\beta$ the $r$ values will decreases from one and for negative $\\beta$ the $r$ will increases form one as the universe evolves. The values of $\\alpha$ and $\\beta$ are constrained using observational data in reference \\cite{Chimento1}, the best fit value is $\\alpha$=1.01, $\\beta$=0.15. The present value of ($r,s$) can be viewed as a discriminator for testing different dark energy models. For the $\\Lambda$CDM model statefinder is a fixed point $r$=1, $s$=0. For positive values $\\beta$ parameter the $r-s$ and $r-q$ plots of MHRDE shows that, the evolutionary trajectories starts form $r=1$ and $q=0.5,$ in the past universe (for the best fit model parameters), which reveals that the MHRDE is behaving like cold dark matter in the past. The further evolution of MHRDE in the $r-s$ plane shows that the present position of MHRDE model in the $r-s$ plane for the best fit parameter is $r_0$=0.59, $s_0$=0.15 and in the $r-q$ plane is $r_0$=0.59, $q_0$=-0.45. The difference between the MHRDE and $\\Lambda$CDM models is in the evolution of the equation of state parameter, which is -1 in the $\\Lambda$CDM model and a time-dependent variable in MHRDE model. A further comparison can be made with the new HDE model \\cite{Setare1}, which gives the present values $r_0 (HDE)=1.357, s_0(HDE)=-0.102$ and $r_0(HDE)=1.357, q_0(HDE)=-0.590.$ So in the $r-s$ plane the distance of the MHRDE model form the $\\Lambda$CDM fixed point is slightly larger compared to the new HDE model for positive values of $\\beta$ parameter. However in the case of MHRDE model the starting point in $r-s$ plane and $r-q$ plane is ($r=1, s=0$ and $r=1, q=0.5$) is same as that in the $\\Lambda$CDM model. For negative values of the $\\beta$ the $r-s$ trajectory we have plotted is different compared to that of positive $\\beta$ values. For negative $\\beta$ values the $r$ value can attains values greater than one as $s$ increases. The present status of the evolution in the $r-s$ plane is $r_0$ =1.325, $s_0$=-0.10 for model parameters $\\alpha$=1.2, $\\beta$=-0.10 and $r_0$=1.321, $s_0$=-0.10. The $r-q$ for $\\alpha$=1.2 and $\\beta$=-0.10 shows that the present state of the MHRDE model is corresponds to $r_0=1.325$ and $q_0$= - 0.63. These values shows that the MHRDE model is different form $\\Lambda$CDM model for the present time when $\\beta$ parameter is negative also. But compared the new HDE model, the present MHRDE model doesn't show much deviation, shows that for negative values the behavior of MHRDE model is almost similar to new HDE model. Irrespective of whether $\\beta$ is positive or negative the MHRDE model is commence to evolve from SCDM model. When $\\beta$ is positive 1 is the maximum value of $r$, on the other hand when $\\beta$ is negative 1 is the minimum value of $r.$ However the exact discrimination of the dark energy models is possible only if we can obtain the present $r-s$ values in a model independent way form the observational data. It is expected that the future high-precision SNAP-type observations can lead to the present statefinder parameters, which could help us to find the right dark energy models. \\newpage" }, "1207/1207.3658_arXiv.txt": { "abstract": "In this work, we present a short review about the high level design methodology (HLDM), that is based on the use of very high level (VHL) programing language as main, and the use of the intermediate level (IL) language only for the critical processing time. The languages used are Python (VHL) and FORTRAN (IL). Moreover, this methodology, making use of the oriented object programing (OOP), permits to produce a readable, portable and reusable code. Also is presented the concept of computational framework, that naturally appears from the OOP paradigm. As an example, we present the framework called PYGRAWC (Python framework for Gravitational Waves from Cosmological origin). Even more, we show that the use of HLDM with Python and FORTRAN produces a powerful tool for solving astrophysical problems. \\bigskip {\\footnotesize {\\bf Keywords}: Computational Physics, Cosmology, Programming Methodology, HLDM.} ", "introduction": " ", "conclusions": "" }, "1207/1207.3528_arXiv.txt": { "abstract": "{Dust attenuation in galaxies is poorly known, especially at high redshift. And yet the amount of dust attenuation is a key parameter to deduce accurate star formation rates from ultraviolet (UV) rest-frame measurements. The wavelength dependence of the dust attenuation is also of fundamental importance to interpret the observed spectral energy distributions (SED) and to derive photometric redshifts or physical properties of galaxies.} {We want to study dust attenuation at UV wavelengths at high redshift, where the UV is redshifted to the observed visible light wavelength range. In particular, we search for a UV bump and related implications for dust attenuation determinations.} {We use photometric data in the $Chandra$ Deep Field South (CDFS), obtained in intermediate and broad band filters by the MUSYC project, to sample the UV rest-frame of 751 galaxies with $0.95 < z < 2.2$. When available, infrared (IR) $Herschel$/PACS\\thanks{$Herschel$ is an ESA space observatory with science instruments provided by European-led Principal Investigator consortia and with important participation from NASA.} data from the GOODS-$Herschel$ project, coupled with {\\it Spitzer}/MIPS measurements, are used to estimate the dust emission and to constrain dust attenuation. The SED of each source is fit using the CIGALE code. The amount of dust attenuation and the characteristics of the dust attenuation curve are obtained as outputs of the SED fitting process, together with other physical parameters linked to the star formation history.} {The global amount of dust attenuation at UV wavelengths is found to increase with stellar mass and to decrease as UV luminosity increases. A UV bump at 2175\\,$\\AA$ is securely detected in 20\\% of the galaxies, and the mean amplitude of the bump for the sample is similar to that observed in the extinction curve of the LMC supershell region. This amplitude is found to be lower in galaxies with very high specific star formation rates, and 90$\\%$ of the galaxies exhibiting a secure bump are at $z < 1.5$. The attenuation curve is confirmed to be steeper than that of local starburst galaxies for 20$\\%$ of the galaxies. The large dispersion found for these two parameters describing the attenuation law is likely to reflect a wide diversity of attenuation laws among galaxies. The relations between dust attenuation, IR-to-UV flux ratio, and the slope of the UV continuum are derived for the mean attenuation curve found for our sample. Deviations from the average trends are found to correlate with the age of the young stellar population and the shape of the attenuation curve.} ", "introduction": "\\label{sec:intro} Although dust is a minor component in galaxies by mass, its effect on the observation of their stellar populations is striking. Dust captures a large fraction of the stellar emission, especially at short wavelengths. This process makes the direct observation of stellar populations from the UV to the near-IR, where they emit their light, insufficient to recover all the emitted photons. Reliable dust corrections are mandatory for measuring the star formation rate in the universe and its evolution with redshift from UV-optical surveys. When dust emission is measured, accurate star formation rates can be derived by combining IR and UV data, but these data are often not available, in particular for deep optical surveys \\citep[e.g.,][]{ilbert10}. As a consequence, it is particularly important to study the dependence of dust attenuation on parameters such as the observed luminosity, the stellar mass, or the slope of the UV continuum, since it could be used to correct large samples for the effect of dust attenuation, at least in a statistical way. Any modelling of stellar populations in galaxies must also include attenuation from interstellar dust. Solving the radiation transfer in model galaxies is the best way to build physical and self-consistent SEDs. These models calculate the effective obscuration and produce attenuation curves as outputs, which are generally very different from the extinction curves affecting the flux of each star \\citep[e.g.,][]{witt00,pierini04,tuffs04,panuzzo07}. However, these sophisticated models rely on numerous free parameters and physical assumptions that are difficult to constrain from the integrated emission from entire galaxies and for very large numbers of objects. Simpler models have been specifically developed to analyse large samples of galaxies, introducing attenuation curves, recipes, and/or templates. The number of free parameters is relatively small. These codes are often developed to measure photometric redshifts and physical parameters such as the star formation rate (SFR) and the stellar mass ($M_{\\rm star}$). With the availability of mid and far-IR data for large samples of galaxies, new codes are emerging that combine stellar and dust emission on the basis of the balance between the stellar luminosity absorbed by dust and the corresponding luminosity re-emitted in the IR \\citep{dacunha08,noll09b}. Attenuation laws are introduced in all these codes, except for those which include a full radiation transfer treatment. The most popular attenuation curve is that of \\citet{calzetti00}, built for local starburst galaxies. Some specific recipes such as a time dependence of dust attenuation are sometimes introduced \\citep{charlot00,panuzzo07}. The attenuation law for local starbursts, based on spectroscopic data, does not exhibit a bump at 2175\\,$\\AA$ such as that observed in the extinction curves of the Milky Way (MW) or the Large Magellanic Cloud (LMC) \\citep{fitz07,gordon03}. The presence of a UV bump in the attenuation curve of galaxies remains an open issue. In the nearby universe, the UV wavelength range has been investigated thanks to GALEX observations, but the results remain controversial. \\citet{wijesinghe11} analyse the consistency of SFR indicators based on GALEX measurements in the far-ultraviolet (FUV) and near-ultraviolet (NUV) bands and fluxes in the H$_\\alpha$ line, and conclude that they must consider an obscuration curve without any 2175\\,$\\AA$ feature. Instead, from a careful analysis of pairs of galaxy SEDs, \\citet{wild11} conclude that the UV slope of the attenuation curve is consistent with the presence of a bump at 2175\\,$\\AA$, a conclusion also reached by \\citet{conroy10b} from an analysis of the GALEX-SDSS colours of galaxies. At higher redshifts, several authors introduce a bump with moderate amplitude to improve photometric redshifts \\citep{ilbert09,kriek11}. Direct evidence of bumps comes from the analysis of the galaxy spectra at $1 < z < 2.5$: Noll and collaborators analyse high quality spectra of $\\sim$~200 galaxies and find significant bumps in at least 30$\\%$ of the sources \\citep{noll05,noll07,noll09a}. In an earlier paper, we analysed SEDs of 30 galaxies in a redshift range between 0.95 and 2.2, observed through intermediate band filters and with IR detections from $Herschel$/PACS, and found evidence for a UV bump in the dust attenuation curve of all these galaxies \\citep[][hereafter paper~I]{buat11b}. Selecting high-redshift galaxies on their observed optical colours is quite common (e.g. Lyman break or BzK galaxies \\citep{reddy08,daddi05}, but most of the time mid- and far-IR data are not available for individual targets, making it difficult to obtain any direct measure of dust attenuation. Studies of dust emission often rely on stacking analyses \\citep{rigopoulou10,burgarella11,reddy12}. The sensitivity of $Herschel$ in the deepest fields allows us to combine stellar and dust emission for individual galaxies at high redshift \\citep{wuyts11}, and to perform SED fitting accounting for both components (paper~I). In this work, we will select galaxies at optical wavelengths with a very good sampling of their UV rest-frame SED: essentially, the sample is UV selected with a redshift range between 1 and 2. When available, IR emission from $Herschel$/PACS will be added to better constrain dust attenuation. SED fitting will be performed to estimate the main characteristics of dust attenuation. The sample and the SED fitting process are described in Sects.~\\ref{sec:sample} and \\ref{sec:SED-fitting}, respectively. The global amount of dust attenuation and its variation with stellar mass and UV luminosity are discussed in Sect.~\\ref{sec:attenuation}. Section~\\ref{sec:attcurve} is devoted to the description of the attenuation law. In Sect.~\\ref{sec:att_slope}, we revisit the relation between dust attenuation and the slope of the UV continuum. Our conclusions are presented in Sect.~\\ref{sec:conclusions}. All magnitudes are given in the AB system. We assume that $\\Omega_{\\rm m} = 0.3$, $\\Omega_{\\Lambda} = 0.7$, and $H_0 = 70\\,{\\rm km\\,s^{-1}\\,Mpc^{-1}}$. ", "conclusions": "\\label{sec:conclusions} We have studied the dust attenuation properties in a sample of 751 galaxies with $0.95 < z <2.2$ observed with intermediate-band filters in order to sample correctly the UV rest-frame continuum. IR data from MIPS (290 galaxies) and PACS (76 galaxies) were used when available. Our results are: \\begin{enumerate} \\item SED fitting is performed with CIGALE. The dust attenuation law is described with two free parameters: its steepness ($\\delta$) and the amplitude of a bump at 2175\\,$\\AA$ ($E_{\\rm b}$). The availability of IR data improves the determination of all the parameters related to dust attenuation. Global parameters such as the SFR, stellar mass, or the amount of dust attenuation are well determined, whereas the detailed star formation history is ill-constrained. Dust attenuations, total IR luminosities, and $\\delta$ are found to be slightly overestimated in the regime of low attenuation and luminosity when IR data are missing. \\item The dust attenuation in the FUV is found to have a large dispersion for low $L_{\\rm FUV}$ ($L_{\\rm FUV} < 10^{10}\\,L_{\\odot}$). The mean value of $A_{\\rm FUV}$ decreases from 2.4 to 1.5\\,mag when $L_{\\rm FUV}$ increases from $10^{9.7}$ to $10^{10.8}\\,L_{\\odot}$. The dust attenuation is found to increase with stellar mass. \\item The parameters describing the dust attenuation curve, $E_{\\rm b}$ and $\\delta$, are found to span a large range of values. This is likely to reflect intrinsic variations of the attenuation curve among galaxies and for different environments within a galaxy. The presence of a bump is confirmed in 20$\\%$ of the sample, at $z < 1.5$ for 90$\\%$ of the confirmed detections. A dust attenuation law steeper than that of \\citet{calzetti00} ($\\delta < 0$) is also confirmed for 20$\\%$ of the sample. These percentages increase to 40$\\%$ for galaxies with IR detections. The mean values of $E_{\\rm b}$ and $\\delta$ derived for the whole sample are similar to those found for the extinction curve of the LMC supershell. Finally, $E_{\\rm b}$ is found to be anti-correlated with the specific SFR of galaxies. \\item The relations between $A_{\\rm FUV}$ and $\\beta$ (the slope of the UV continuum) and IRX ($\\log(L_{\\rm IR}/L_{\\rm FUV})$) and $\\beta$ are derived for our mean attenuation curve and for galaxies of our sample detected in IR. The intrinsic slope of the UV continuum, $\\beta_0$, is found equal to -2.5 and the galaxies follow the IRX-$\\beta$ relation with a low dispersion. Galaxies found above the average IRX-$\\beta$ relation have a high IR luminosity and a dust attenuation curve slightly greyer than our mean law, which can be due to a higher compactness of the star formation regions. The age of the young stellar population (i.e. the intrinsic shape of the UV continuum) is likely to also play a role: for a given $\\beta$, IRX increases when the age of the young stellar population decreases. \\end{enumerate}" }, "1207/1207.4622_arXiv.txt": { "abstract": "{The \\emph{Fermi} Gamma-Ray Burst Monitor (GBM) onboard the Fermi spacecraft currently operates on several trigger algorithms on various time scales and energy ranges. Motivated by the pursuit of faint Gamma-Ray Bursts (e.g. the elusive class of postulated low-luminosity GRBs), here we present the search for untriggered GRBs in the GBM data stream. To this end, I will demonstrate the methods and algorithms which have been developed by the GBM team. As a preliminary result, I am going to highlight the spectral analysis of GRBs which triggered the Swift satellite, but not GBM, and came from positions above the horizon, with a favorable orientation to at least one GBM detector. The properties of these GRBs are then compared to the full sample of GBM GRBs published in the GBM spectral catalogue. We estimate that the lower limit for untriggered GRBs in the GBM data is about 1.6 GRBs per month which corresponds to about 7\\% of the triggered GRBs. } \\FullConference{Gamma-Ray Bursts 2012 Conference -GRB2012,\\\\ May 07-11, 2012\\\\ Munich, Germany} \\begin{document} ", "introduction": "The Gamma-Ray Burst Monitor (GBM) is one of the instruments onboard the \\emph{Fermi} Gamma-Ray Space Telescope \\cite{atwood09} launched on June 11, 2008. Specifically designed for GRB studies, GBM is comprised of a total of 12 sodium iodide (NaI(Tl)) scintillation detectors covering the energy range from 8 keV to 1~MeV and two bismuth germanate scintillation detectors (BGO) sensitive to energies between 150~keV and 40~MeV \\cite{meegan09}. % GBM continuously observes the whole unocculted sky, with its flight software (FSW) constantly monitoring the count rates recorded in the various detectors. For GBM to trigger on a GRB or any other high energy-transient, two or more detectors must have a statistically significant increase in count rate above the background rate. GBM currently operates on 75 (of 119 supported) different trigger algorithms, each defined by its timescale (from 16~ms to 4~s) and energy range ($25-50$~keV, $50-300$~keV, $>100$~keV, and $>300$~keV). In addition, the trigger algorithm is constructed to have two temporally overlapping windows (at half the window length) for all timescales above 16~ms (the so called ``offset'') and each algorithm can be operated on different threshold settings from $0.1\\sigma$ to $25.5\\sigma$. GBM persistently records two different types of science data, called CTIME (fine time resolution, coarse spectral resolution) and CSPEC (coarse time resolution, full spectral resolution). CTIME (CSPEC) data have a nominal time resolution of 0.256~s (4.096~s) (see Fig.\\ref{fig:ctime}) which is increased to 64~ms (1.024~s) when GBM triggers. After 600~s in triggered mode, both data types return to the non-triggered time resolution. For a more detailed review on the data types and trigger properties of GBM we refer to \\cite{meegan92} and \\cite{paciesas12}. Here, I present the preliminary results of an on-ground search in CTIME data for GRBs which were detected by the \\emph{Swift} satellite \\cite{gehrels04} but did not trigger GBM, although having a favorable orientation to the GBM detectors. \\begin{figure}[htbp] \\begin{center} \\includegraphics[width=0.7\\textwidth]{ctime.eps} \\caption{Typical CTIME light curve of one NaI detector over a full day in the energy interval from $10-1000$~keV. One can clearly identify the overall background variations and SAA passages during which the detectors are switched-off.} \\label{fig:ctime} \\end{center} \\end{figure} ", "conclusions": "" }, "1207/1207.7237_arXiv.txt": { "abstract": "The total number of true, likely and possible planetary nebulae (PN) now known in the Milky Way is about 3000, approximately twice the number known a decade ago. The new discoveries are a legacy of the recent availability of wide-field, narrowband imaging surveys, primarily in the light of H$\\alpha$. The two most important are the AAO/UKST SuperCOSMOS H$\\alpha$ survey \u2013 SHS and the Isaac Newton photometric H$\\alpha$ survey - IPHAS, which are responsible for most of the new discoveries. A serious problem with previous PN catalogues is that several different kinds of astrophysical objects are able to mimic PN in some of their observed properties leading to significant contamination. These objects include H~II regions and Str\\\"{o}mgren zones around young O/B stars, reflection nebulae, Wolf-Rayet ejecta, supernova remnants, Herbig-Haro objects, young stellar objects, B[e] stars, symbiotic stars and outflows, late-type stars, cataclysmic variables, low redshift emission-line galaxies, and even image/detector flaws. PN catalogues such as the Macquarie/AAO/Strasbourg H$\\alpha$ Planetary Nebula catalogue (MASH) have been carefully vetted to remove these mimics using the wealth of new wide-field multi-wavelength data and our 100\\% follow-up spectroscopy to produce a compilation of new PN discoveries of high purity. During this process significant numbers of PN mimics have been identified. The aim of this project is to compile these MASH rejects into a catalogue of Miscellaneous Emission Nebulae (MEN) and to highlight the most unusual and interesting examples. A new global analysis of these MEN objects is underway before publishing the MEN catalogue online categorizing objects by type together with their spectra and multi-wavelength images. ", "introduction": "The MASH Miscellaneous Emission Nebulae catalogue (MASH-MEN) represents the compilation of about 450 new emission line sources identified as PN candidates but now removed as contaminants mostly prior to publication of MASH-I (\\cite[Parker \\etal ~2006]{ParkerAl2006}) and MASH-II (\\cite[Miszalski \\etal ~2008]{MiszalskiAl2008}). This was achieved by careful application of spectral and multi-wavelength image diagnostic criteria as described in \\cite[Frew \\& Parker (2010)]{FrewParker2010}. ", "conclusions": "The MASH-MEN Project was instigated first to house the many interesting but non-PN mimics uncovered during the MASH survey and subsequently to classify the mimics according to type into a new catalogue of miscellaneous emission nebulae (MEN). Application of our careful diagnostic processes have resulted in the MASH PN catalogues being of high purity. Some of these contaminants are very interesting in their own right and deserve follow-up and further study. We are now applying the same processes to other extant PN catalogues to remove their many mimics." }, "1207/1207.2138_arXiv.txt": { "abstract": "We describe the construction of a suite of galaxy cluster mock catalogues from N-body simulations, based on the properties of the new ROSAT-ESO Flux-Limited X-Ray (REFLEX II) galaxy cluster catalogue. Our procedure is based on the measurements of the cluster abundance, and involves the calibration of the underlying scaling relation linking the mass of dark matter haloes to the cluster X-ray luminosity determined in the \\emph{ROSAT} energy band $0.1-2.4$ keV. In order to reproduce the observed abundance in the luminosity range probed by the REFLEX II X-ray luminosity function ($0.01 33 \\rm \\, eV$ and $EW > 40 \\rm \\, eV$. Similarly, \\citet{Ballo2011} used {\\it XMM-Newton} data to fit the 0.4--10 keV spectrum and found an $EW = 38^{+11}_{-9} \\rm \\, eV$. \\\\ 3C~111 has been detected in $\\gamma$-rays by {\\it CGRO}/EGRET \\citep{Hartman1999, Sguera2005,Hartman2008}, and was included in the first {\\it Fermi}/LAT catalogue \\citep{Abdo2010firstsourcecatalog}. In the second {\\it Fermi}/LAT catalogue, the source was omitted since it had no longer been significantly detected \\citep{Fermisecondsourcecat}. 3C~111 is likely variable in the $\\gamma$-ray regime \\citep{Ackermann2011}. Using 24 months of {\\it Fermi}/LAT data, \\citet{Grandi2012} found 3C~111 to be detectable during a short time-period ($\\Delta t \\sim 30-60$ days), limiting the radius of the emission region to be $R<0.1$ pc (assuming a Doppler factor of $\\delta =3$) based on to causality arguments. The detectability of the source in the GeV energy range coincided with an increase in the flux in the millimeter, optical and X-rays regimes, indicating the emission is likely to emerge from the same region. Since the outburst in millimeter, optical, and X-rays is associated with the ejection of a bright radio knot, this indicates that the GeV emission originates from the radio core within 0.3 pc of the central supermassive black hole. \\\\ To understand the physical processes causing the emission in 3C~111, we first study the X-ray to $\\gamma$-ray spectrum to find whether the emission is the product of thermal (Seyfert-like) or non-thermal processes (blazar-like) or a combination of both. We then model the spectral energy distribution (SED) to understand the processes dominating the broad-band radiation of 3C~111 from radio to high energies. ", "conclusions": "The origin of the high-energy emission from non-blazar AGNs remains unclear. \\citet{Marscher2002} suggested that with the acceleration of the inner regions of the accretion disc a shock front will stream along the jet and the expulsion of a superluminal bright knot will follow. \\citet{Grandi2012} show, using {\\it Fermi}/LAT data, that the GeV emission of 3C~111 appears to originate from a compact knot confined to within 0.1 pc. This knot is clearly separate from the core and placed 0.3 pc from the central engine.\\\\ By analysing the X-ray spectrum and the broad-band SED, we have studied the nature of the high-energy emission of the radio galaxy 3C~111. We have presented an X-ray spectrum between 0.4 keV and 200 keV using data of 3C~111 acquired by several instruments and showed that the best-fit model is an absorbed cut-off power-law with both a reflection component and a Gaussian component to account for the iron line at 6.4 keV. The values we found for the reflection and high-energy cut-off are similar to those found in Seyfert galaxies, which would indicate that there is a thermal core visible. The cut-off can also originate from non-thermal processes and the EW of the iron line is variable and smaller than expected for Seyfert galaxies. We therefore conclude that the X-ray spectrum is mainly of thermal origin, but there may be a small non-thermal contribution. \\\\ Using the X-ray spectrum, together with $\\gamma$-ray data from {\\it Fermi}/LAT and archival deabsorbed radio and infrared data, we modelled the broad-band SED of 3C~111 using a single-zone synchrotron self-Compton model. This model is non-thermal and we also did not need to include an additional thermal component to model the SED. Since the X-ray emission is likely to have a combined thermal and non-thermal origin, the SSC model we used may overestimate the non-thermal contribution in the X-ray band and should therefore be considered an upper limit. \\\\ In conclusion, it seems that the high-energy emission from 3C~111 consists of both thermal and non-thermal components. In the X-ray spectrum, the thermal components manifest themselves in terms of an iron line and reflection. The non-thermal component is visible through the variability in the EW of the iron line. The high-energy cut-off can be the result of either thermal or non-thermal inverse Compton scattering, but our present spectrum does not allow us to distinguish which process is occurring. The broadband SED can be modelled with a non-thermal model, but it is possible there is a thermal component that we are unable to discern with the current data set." }, "1207/1207.5575_arXiv.txt": { "abstract": "We study C/O white dwarfs with masses of 1.0 to 1.4 $M_\\odot$ accreting solar-composition material at very high accretion rates. We address the secular changes in the WDs, and in particular, the question whether accretion and the thermonuclear runaways result is net accretion or erosion. The present calculation is unique in that it follows a large number of cycles, thus revealing the secular evolution of the WD system. We find that counter to previous studies, accretion does not give rise to steady state burning. Instead, it produces cyclic thermonuclear runaways of two types. During most of the evolution, many small cycles of hydrogen ignition and burning build a helium layer over the surface of the white dwarf. This He layer gradually thickens and progressively becomes more degenerate. Once a sufficient amount of He has accumulated, several very large helium burning flashes take place and expel the accreted envelope, leaving no net mass accumulation. The results imply that such a system will not undergo an accretion induced collapse, nor will it lead to a SN Type Ia, unless a major new physical process is found. ", "introduction": "The prevailing scenarios leading to Type Ia SNe can be divided into two classes - those of singly (SD) and those of doubly degenerate (DD) systems. In the first, a degenerate WD accretes from its binary companion, and accumulates sufficient mass to approach the Chandrasekhar limit, where CO detonation or the collapse to a neutron star takes place \\citep[e.g.,][]{Hillebrandt2000,Pod2008}. In the second class, Ia's occur from the merger of a binary WD system \\citep[e.g.,][]{Webbink1984}. Since no model within the above classes of progenitors comes without significant caveats, there is still no consensus model for Type Ia SNe. When considering the merger of WDs, one major caveat is the merger rate. The theoretical and observed rates of super-Chandrasekhar mergers is typically an order of magnitude smaller than the observed rate of Ia's, unless carbon detonation can take place in sub-Chandrasekhar systems as well \\citep{vanKerkwijk2010,Badenes2012}. Another problem is that merger scenarios would tend to produce explosions which are more heterogeneous than observed Ia's, both in terms of light curves and ejected mix of elements. Systems where a WD accretes from a non-degenerate companion, on the other hand, are subject to other theoretical and observational limitations. Perhaps the most important one is that of the accretion rate. In their review, \\cite{Hillebrandt2000} pointed out that although SD is the favoured progenitor model, the major problem has always been that nearly all possible accretion rates can be ruled out by strong arguments. When hydrogen is accreted at sufficiently low rates, it can cool and become degenerate. When a sufficient amount of gas accumulates, the hydrogen ignites to produce an unstable thermal flash \\citep{Schwarzschild1965,Rakavy1968}. Since the nuclear luminosity may reach extremely high values, of the order of $10^{12}L_{\\odot}$ for a relatively long period of time, the flash can dynamically eject the accreted mass. In fact, this mechanism naturally explains classical novae eruptions \\citep{Iben1984}. Because lower accretion rates allow more mass to accumulate, lower accretion rates produce stronger eruptions. Moreover, the observations indicate that the ejecta contain as a rule, more heavy elements than in the accreted matter. It is obvious that if He burning and beyond does not take place, then these heavy element originate form the underlying WD, which is eroded in this process. \\cite{Prialnik1995} and \\cite{Yaron2005} found that all WD's accreting at rates ${\\dot m}\\lesssim 10^{-7}M_{\\odot}/yr$ erode in mass. At somewhat higher accretion rates, there is net accumulation, but with a low efficiency since most of the mass is still ejected. One way to avoid the hydrogen flashes associated with the low accretion rate, and the consequent mass loss, is to accrete pure helium, which may avoid ignition \\citep{Iben1991, Iben1994}. However, population syntheses suggest that such progenitor systems, which appear as AM CVn stars can contribute at most, about $10^{-2}$ of the SN Ia rate \\citep{HeAccretionPopulation}. At the opposite limit, of very high accretion rates, \\cite{Iben1984} argued that the WD does not experience hydrogen shell flashes, but instead, its envelope expands to a radius $R\\approx 1050(M/M_{\\odot}-0.5)^{0.68}R_{\\odot}$, which for a massive WD is much larger than the orbital separation of typical cataclysmic variables. Once a common envelope is formed, heavy mass loss should then prevent the mass of the C/O WD from reaching the Chandrasekhar limit. In between the above ranges, are accretion rates for which the released gravitational binding energy is close to the Eddington limit, around $10^{-7}$ to $10^{-6}M_{\\odot}/yr$ (depending on the WD mass). It was suggested that the high accretion should lead to a quiet hydrostatic burning of H and He \\citep{Nomoto1982}. Various attempts to calculate the effect of high accretion rate were carried out under different approximations (e.g., \\citealt{Nomoto1982}, using the technique developed by \\citealt{Nomoto1977}). In particular, the assumption of steady state ignores the time dependent evolution of the accretion system. However, this assumption simplifies the calculation because time dependent models have two major obstacles. First, the latter require very fine mass shells, and second, the full evolution requires following a large number of small flashes. \\cite{Fujimoto1982} investigated the thermal properties of hydrogen shell burning on accreting white dwarfs. Assuming hydrostatic equilibrium (no dynamic effects were allowed), Fujimoto found that the hydrogen burns steadily on the surface of the WD. The gravitational energy release was simulated by $${\\partial \\over \\partial q}\\left({L_r \\over M}\\right)=\\epsilon_g,$$ where $L_r$ is the radiative flux and $q=M_r/M$. The burning shell exhibited periodic flashes, all treated under hydrostatic equilibrium. As Fujimoto assumed steady state, the rate at which the mass of the WD grew was equal to the assumed accretion rate. \\cite{Prialnik1995} alleviated the assumption of steady state. The have shown that one must calculate at least several dozens of cycles before the effects of the arbitrary chosen initial conditions decays away. In some cases even 1001 cycles were followed \\citep{Epel2007}. The calculations were extended by \\cite{Yaron2005} with basically the same result. In particular, it was found that WDs accreting at rates ${\\dot m}\\gtrsim 10^{-7}M_{\\odot}/yr$ grow in mass. In addition to the above theoretical arguments, interesting observational constraints should be considered as well. Some argue in favour of SD scenarios and some argue against. The SD scenarios imply that the former companion stars should remain after the explosion, and be observable in association with nearby remnants. And indeed, a G-star was claimed to be the former companion of Tycho Brahe's 1572 supernova \\citep{ruiz2004binary}. On the other hand, the opposite claim was also made. No former companion was found to be associated with SNR 0509-67.5, down to $M_v = + 8.4$ \\citep{schaefer2012absence}. In this respect, one should also mention the lack of any radio detection in SN 2011fe, which took place only 8 Mpc away. This was used, under several theoretical assumptions, to place an upper limit of order 10$^{-8}$M${}_{\\odot}$/yr, on the companion wind prior to the supernova explosion \\citep{Horesh}. Using {\\sc Chandra} and {\\sc HST} archival data, SN 2011fe could also be used to place a limit on the accretion spectrum. If the WD was emitting at the Eddington luminosity, then its temperature should have been within $60$ eV $\\gtrsim kT \\gtrsim 10$ eV \\citep{Liu2012}. In contrast, \\cite{Voss2008} reported the disappearance of an X-ray source from the location of a Type Ia supernova in NGC 1404. They estimated the source to have been emitting $2-3 \\times 10^{37}$ erg/sec in X-rays. In terms of statistics, high rate accretion implies that SD progenitor systems should be sufficiently bright as to be easily detected. It is presently not clear at all, whether SD progenitors can be identified with any of the known cataclysmic binaries, the known symbiotic systems, or whether perhaps with the very bright supersoft X-ray sources \\citep[][]{Iben1994,Gilfanov2010}. Given the present state of uncertainty, further study of the different progenitor scenarios is necessary. Here we shall address one particularly pressing question, which is whether systems accreting hydrogen at high rates can actually accumulate mass. We begin in \\S\\ref{sec:Details} by describing our numerical code, and the delicate points requiring attention when carrying out a multi-cycle evolution. In \\S\\ref{sec:results} we describe our results. Here we will see that there are two types of thermonuclear runaways (TNRs), those which burn H and expel no matter, and larger ones burning He, which do do expel. We will therefore discuss in \\S\\ref{sec:energetics} the energetics of how this is possible, even though the specific energy release in helium burning is small. We end with conclusions in \\S\\ref{sec:conclusions}. ", "conclusions": "\\label{sec:conclusions} High rate accretion of solar composition material onto WDs prevents degeneracy. According to general wisdom, this should give rise to quiet steady state burning in which the helium ash from the hydrogen burning in the envelope should accrete onto the core, and secularly increase its mass. Since eruptions are expected to disrupt the accretion process, our first conclusion is that instead of a steady state process, the accretion and burning take place in eruptive cycles. However, this cyclic behaviour is different from classical nova eruptions. We find that the accreted material is only moderately degenerate before it ignites. And once it does, it burns completely. Therefore, any ejecta will contain no detectable hydrogen. Clearly then, the non-existence of hydrogen in observed ejecta does not immediately imply that the accreted matter did not contain any hydrogen as well. Since the lack of observable hydrogen in ejecta is generally considered as one of the strongest evidence suggesting the accretion of pure helium, such claims should be considered more cautiously. The last, but most important conclusion we have reached is that high accretion rates, of $10^{-6}M_{\\odot}$, do not lead to the a secular increase in the mass of the white dwarf, because giant helium eruptions take place after a sufficient amount of helium accumulates. These eruptions expel all the accreted layer and even a small amount of the underlying WD. We find that helium flash explosions take place for WDs with masses up to 1.35$M_\\odot$. From eq.~\\ref{eq:expul}, however, there should be no upper limit above which helium cannot expel the material---although more compact systems would require more energy to expel material out of them, they also release more energy during the accretion. The above accretion rate is very high and can be maintained by only few systems. The results for a lower accretion rate, of $10^{-7}M_{\\odot}$ is qualitatively different. There are no giant helium flashes, and there is a net increase of the WD mass. However, because the mass accretion efficiency is small, it requires the donor to fine tune the mass transfer over a large integrated total mass. It therefore seems unlikely that large increases in the WD mass are possible. The above conclusions depend on the energy release by the hydrogen burning. Thus, systems in which helium is accreted will not have their envelope bloated, and any helium eruption will not lead to large mass loss. This possibility is presently under investigation. Another important point that is the subject of further investigation is the effects that will arise once super-Eddington conditions are allowed to develop \\citep{ShavivNovae}. This should modify the above conclusion for two main reasons. On one hand, the super-Eddington luminosities drive mass loss through continuum driven winds. On the other hand, a higher luminosity will inhibit the atmosphere from bloating, such that the helium flashes will not necessarily be able to drive the envelope away. Thus, it is not a priori clear how the above picture will be modified." }, "1207/1207.0399_arXiv.txt": { "abstract": "~ We have performed an extensive numerical study of coalescing black-hole binaries to understand the gravitational-wave spectrum of quasi-normal modes excited in the merged black hole. Remarkably, we find that the masses and spins of the progenitor are clearly encoded in the mode spectrum of the ringdown signal. Some of the mode amplitudes carry the signature of the binary's mass ratio, while others depend critically on the spins. Simulations of precessing binaries suggest that our results carry over to \\emph{generic} systems. Using Bayesian inference, we demonstrate that it is possible to accurately measure the mass ratio and a proper combination of spins even when the binary is itself invisible to a detector. Using a mapping of the binary masses and spins to the final black hole spin, allows us to further extract the spin components of the progenitor. Our results could have tremendous implications for gravitational astronomy by facilitating novel tests of general relativity using merging black holes. ", "introduction": " ", "conclusions": "" }, "1207/1207.1160_arXiv.txt": { "abstract": "{ We review a molecular dynamics method for nucleon many-body systems called the quantum molecular dynamics (QMD) and our studies using this method. These studies address the structure and the dynamics of nuclear matter relevant to the neutron star crusts, supernova cores, and heavy-ion collisions. A key advantage of QMD is that we can study dynamical processes of nucleon many-body systems without any assumptions on the nuclear structure. First we focus on the inhomogeneous structures of low-density nuclear matter consisting not only of spherical nuclei but also of nuclear ``pasta'', i.e., rod-like and slab-like nuclei. We show that the pasta phases can appear in the ground and equilibrium states of nuclear matter without assuming nuclear shape. Next we show our simulation of compression of nuclear matter which corresponds to the collapsing stage of supernovae. With increase of density, a crystalline solid of spherical nuclei change to a triangular lattice of rods by connecting neighboring nuclei. Finally, we discuss the fragment formation in expanding nuclear matter. Our results suggest that a generally accepted scenario based on the liquid-gas phase transition is not plausible at lower temperatures. } ", "introduction": "Due to the progress of computers, numerical simulations became increasingly capable in tackling complicated problems in nuclear physics. Generally, numerical simulations can be classified into two types: macroscopic and microscopic simulations. The former, macroscopic simulations, deal directly with the macroscopic quantities which we are interested in. We need to introduce physics models which describe how the quantities are connected with each other. On the other hand, microscopic simulations are based on the degrees of freedom of the constituent elements. The necessary inputs are equations of motion and the interactions among the elements. The properties of the total system are obtained later by analyzing the resultant information of these constituent elements. Microscopic simulations have several advantages: 1) We need only a few assumptions on the model. 2) We may obtain unexpected results. 3) We may find a physical principle (the law governing the elements) if we obtain suitable observables. This review article is about the molecular dynamics (MD) simulations of nuclear matter. First, we give an overview of the history of the MD models in nuclear physics. The simulation study of nuclear dynamics originated from the formulation of the time-dependent Hartree-Fock (TDHF) theory in 1930 \\cite{Dirac}. TDHF deals with the time evolution of many-fermion systems and is an approximation of the time-dependent Schr\\\"odinger equation with the use of a single Slater determinant. However, it was in the 1970's that the TDHF was first solved numerically \\cite{Bonche}. Due to the limitation of computer power, only low-energy phenomena were studied in the early stage. As the computational power drastically increased after the 1980's, applications to higher-energy phenomena with larger numbers of degrees of freedom and also improvements of the framework to include correlations, etc., have been made. However, TDHF cannot describe the heavy-ion collision process at higher energies where the degrees of freedom drastically increase. The reaction mechanism depends on the incident energy and the impact parameter, as shown in Fig.\\ \\ref{reaction-mechanism}. Particularly, the Fermi energy is the key quantity to characterize the mechanism. It is because the reaction mechanism is determined by the competition between Fermi motion inside the nuclei and the relative motion of colliding nuclei. Above the Fermi energy (medium -- high energy), the region where the colliding nuclei overlap with each other will break into fragments, i.e., fragmentation occurs. This fragmentation process is driven by the energy of the relative motion between the two nuclei converted into thermal excitations, which are generated by two-body collisions. Each collision is regarded as a transition from a Slater determinant into another. Such a change of the wavefunction cannot be described by TDHF with a single Slater determinant. \\begin{figure} \\begin{center} \\includegraphics[width=0.5\\textwidth]{reaction-mech.ps} \\caption{\\label{reaction-mechanism} Schematic diagram of heavy-ion reaction mechanism. } \\end{center} \\end{figure} At higher energies, the above-mentioned two-particle collision becomes an important process which determines the reaction mechanism in the heavy-ion collisions. The time evolution of the phase-space distribution function in heavy-ion collisions is described by a Boltzmann-type equation of motion (EOM), with a smooth change by a Newtonian equation and dissipation by the two-body collision process, which is called Boltzmann-Uehling-Uhlenbeck (BUU), Vlasov-Uehling-Uhlenbeck (VUU), or Boltzmann-Nordheim-Vlasov (BNV) equation. If one omits the collision term, this framework can be regarded as a classical limit ($\\hbar\\rightarrow 0$) of TDHF \\cite{CYWong}, i.e., the Vlasov equation. This Boltzmann-type equation can be numerically solved by a test-particle method: The fluid elements in a 6-dimensional phase space are replaced by a classical particle and the phase-space distribution function is obtained by counting the number of those particles in the 6-dimensional mesh (three dimensions for coordinate space and the other three dimensions for momentum space). For sufficiently large numbers of test particles, the 6-dimensional particle density is conserved throughout the time evolution in a mean-field potential described by the Vlasov equation. The two-body collision process of test particles violates the conservation of 6-dimensional phase-space distribution. In the 1980's many works on the heavy-ion collisions in the medium -- high energy region have been made via BUU simulations with the test-particle method. \\begin{figure} \\begin{center} \\includegraphics[width=0.5\\textwidth]{mean-field.ps} \\caption{Schematic explanation on the time evolution of two-particle correlation. The curves indicate the distribution functions and the circles are their representative test particles. Refer to the text for details. \\label{figCorrelation}} \\end{center} \\end{figure} Molecular dynamics simulation for nuclear systems has been developed in the late 1980's. Aichelin and St\\\"ocker have proposed quantum molecular dynamics (QMD) model to simulate heavy-ion collisions from medium to high energies \\cite{QMD}. This framework is obtained by reducing the number of test particles in BUU simulations so that each particle represents one nucleon. By this reduction of test particles, it became possible to describe many-body correlation of the system. Let us take the two-body correlation as an example. If two nucleons stay within a distance so that their distribution functions (solid curves in Fig.\\ \\ref{figCorrelation}) overlap with each other, the representative two test particles (circles in Fig.\\ \\ref{figCorrelation}) can be very close to each other (corresponding to the left panel of Fig.\\ \\ref{figCorrelation}) and also can be far (right panel). In the former case, those two nucleons may be bound to form a cluster while, in the latter case, they may diverge from each other as time passes. However, with a huge number of test particles, we obtain only one time evolution of the distribution function which corresponds to the average of many events. This example of the two-body correlation effect on the fragment formation is related to the variation between time evolutions of different events. It is also natural that the two-body correlation is important for the description of spatial fluctuation in a event. In the mean-field calculation, on the other hand, both fluctuations in space and the variation between events will be washed out. The QMD model assumes a direct product of nucleon single-particle wavefunctions as a total wavefunction and a Lagrangian with a non-relativistic kinetic energy and a potential energy from effective interactions among nucleons. The single-particle wavefunction is assumed to be a Gaussian wavepacket with a fixed width. The EOM of the wavefunction is derived from a variational principle with the above Lagrangian, and results in a classical EOM with a Hamiltonian as a function of coordinates of those Gaussian wavepackets. In the QMD model, the stochastic two-body collision process is added to the time evolution by the Hamilton EOM. The final state of the two-body collision process is checked so that it obeys the Pauli principle, i.e., the condition on the phase-space density. QMD is named ``quantum'' due to (1) the many-body correlation or fluctuation in density caused by the EOM and the collision term, (2) the stochasticity in the collision process, (3) the Pauli blocking in the final state of collision, and (4) the use of Gaussian wavepackets for single-particle wavefunctions. However, the actual feature of QMD simulation is rather classical. First, the time evolution of the system concerns only the centroids of wavepackets. Their width in the coordinate space, which is a fixed parameter, appears only in the interaction among the particles by means of the double folding. The width in the momentum space gives rise to a part of the kinetic energy. However, this energy is spurious, i.e., it will never be effective, since it is constant during the time evolution. Second, the Pauli principle, which yields the fermionic momentum distribution, is not basically taken into account. When the number of the test particles per nucleon is large enough in the Boltzmann type simulation, the phase-space density is conserved in the moving frame of a fluid element. In spite of the many-body nature obtained by the reduction of the number of test particles, it sacrifices the fermionic nature of the system. One of the most serious problems of QMD is in the description of the ground state. Due to the lack of fermionic characteristics, the energy minimum states of QMD model violate the Pauli principle, and all the particles degenerate into zero in the momentum space so that they overestimate the binding energy. We cannot use the energy-minimum states as the initial conditions of collision simulations. If we prepare initial conditions with appropriate binding energies, the constituent nucleons would be moving. This motion makes the initial condition unstable against the emission of nucleons. To take the fermionic characteristics into account, we need to introduce explicitly the antisymmetrization of the wavefunction \\cite{refFMD,refAMD,FMDreview}. Fermionic molecular dynamics (FMD) \\cite{refFMD} and antisymmetrized molecular dynamics (AMD) \\cite{refAMD} have been proposed in 1990 and 1992, respectively. They assume a Slater determinant of Gaussian wavepackets as the wavefunction of the system. In FMD, the widths of nucleons are time-dependent variables and the kinetic energies of wavepackets are not spurious. In AMD, the widths of wavepackets are constant in time but the zero-point center-of-mass kinetic energies of fragments are removed in a phenomenological way. They have succeeded in describing the ground state properties of light nuclei as well as the dynamical processes of low-energy heavy-ion collisions. The problem is, however, a huge amount of computing cost to solve the equations of motion of FMD and AMD, which is proportional to the fourth power of the particle number $N$ (cf. $\\propto N^2$ for QMD). Thus the use of FMD and AMD has been limited to small systems with the total number of particles up to a few hundreds. In this situation, a new phenomenological way to mimic the Pauli principle was introduced in QMD \\cite{Pei91}. Wilets \\etal \\cite{Wilets} and then Dorso \\etal \\cite{Dorso} developed a repulsive two-body potential so-called the Pauli potential. It is a function of not only the distance in the coordinate space, but also of the distance in the momentum space. This repulsive potential acts between nucleons with the same spin and isospin so that it prevents those particles from coming close in the phase space. Note that, in this framework, simulated ideal Fermi gases contain the potential energy which comes from the Pauli potential. It is counted as a part of the nuclear potential energy when one determines the parameters of effective potential. Due to the momentum dependence of the Pauli potential, constituent nucleons have non-zero values of the momentum in the ground state keeping their velocities at zero; thus the above-mentioned spurious emission of nucleons is avoided. Since the appearance of the QMD model with the Pauli potential, it became possible to carry out simulations of systems with a large number of nucleons. One interesting target was low-density (below the saturation density) nuclear matter in compact stars such as crusts of neutron stars and cores of supernovae. In low-density nuclear matter, exotic structures called ``nuclear pasta'' have been predicted by Ravenhall \\etal \\cite{Ravenhall83} and Hashimoto \\etal \\cite{Hashimoto84}. There, nuclear matter cannot be uniform due to a negative partial pressure of nucleons and should be clusterized. With increasing density, the shape of the cluster changes from droplet, rod, slab, tube, bubble, and then uniform. The name of ``pasta'' comes from the similarity of the rods to ``spaghetti'' and the slabs to ``lasagna'', etc. Since Ravenhall \\etal and Hashimoto \\etal have proposed, many works have been done on nuclear matter with pasta structures. Most of them are based on the Wigner-Seitz (WS) approximation, in which a unit cell with the dimensionality 1, 2, and 3 is replaced by the same volume of the plate, cylinder, and sphere, respectively. The WS approximation is useful and saves much CPU time. However, the use of WS cell should be a strong constraint on the structure and only simple structures are allowed. On the other hand, MD simulation is a microscopic framework which does not need any assumption on the structure and the reaction mechanism. Since a QMD model is capable of simulating systems with a huge number of particles, we have applied it to low-density nuclear matter and its inhomogeneous structures. In Sec.\\ \\ref{sec-matter}, we present the results of our QMD simulations. Apart from adapting QMD to low-energy phenomena, other efforts have been made to the opposite direction, i.e., an attempt to describe high-energy phenomena. Sorge \\etal have proposed relativistic quantum molecular dynamics (RQMD) \\cite{refRQMD} in 1989. The main improvements of RQMD from QMD are (1) the Lorentz covariance in the interaction, kinematics, and the two-body collisions and (2) the inclusion of baryon resonances, strange particles, and the string excitations in the two-body collision process. Simulations of high-energy heavy-ion collisions have been carried out to analyze experiments with $E/A\\approx 1$ -- $200$ GeV at the SIS, AGS, and SPS facilities. At further high energies, interaction between particles becomes less important and the production of mesons and excitations of baryon resonances are essentially important. A set of computational codes called ``ultrarelativistic QMD'' (UrQMD) is developed with a collision term highly tuned-up to include various kinds of baryons, mesons, and their excited states \\cite{refUrQMD}. It is distributed on the internet and is often used by many people for simulations of heavy-ion collisions at RHIC experiments. ", "conclusions": "In this article, we have reviewed the molecular dynamics approaches to nucleon many-body systems, focusing on the quantum molecular dynamics (QMD) model. This method was originally developed for the study of heavy-ion collisions to describe multifragmentation which the time-dependent Hartree-Fock (TDHF) theory failed to explain. With the help of the Pauli potential to take account of the Pauli principle, QMD has been applied to dense nuclear matter in addition to heavy-ion collisions. A great advantage of MD approaches is that we can study the dynamical process of nucleon many-body systems without any assumptions on the nuclear structure. QMD, in particular, is very suitable to study inhomogeneous nuclear matter in the mesoscopic to macroscopic scales because QMD makes it possible to simulate large systems with many nucleons. To show the power of this method, we have presented some of our works using QMD. First, we have explained our QMD model in Sec.~\\ref{sec-formulation}. We have seen that this model is designed to give the correct saturation properties and reasonable equations of state (EOS) of nuclear matter, and give a good agreement of the binding energies of light nuclei with $A\\lesssim 8$ including alpha particle and also heavy nuclei with $A\\gtrsim 40$. These points are important for the reliability of the predictions by this model, which we have presented in the remaining part of this article. From Sec.~\\ref{QMDmatter-maru} to \\ref{sect_qmdsn}, we have shown a series of our studies about the structure of nuclear matter at subsaturation densities. It has been predicted that nuclear ``pasta'' phases, states of matter with nuclei of rod-like and slab-like shapes, can be the ground state of matter in this density region. Using QMD, we have shown that the pasta phases can actually be formed by cooling hot uniform nuclear matter and by compressing a bcc lattice of spherical nuclei. These results strongly suggest that the pasta phases are formed in the cooling process of hot neutron star crusts and by the compression of matter in the collapse of supernova cores. These results have important implications on the mechanical strength of the neutron star crust, the cooling process of hot protoneutron stars, the mechanism of glitches, etc. Making an EOS table for core collapse simulations, taking into account the existence of the pasta phases including their effects on the neutrino opacity, is an important direction. In Sec.~\\ref{qmd-expand}, we have discussed fragment formation in expanding nuclear matter. We have developed a method to describe isotropically expanding matter using the periodic boundary conditions. Using this method together with our QMD model, we have calculated the fragment mass distribution and have found that it shows an exponential decay for rapid expansion and a power-law decay for slow expansion. Our analysis suggests that multifragmentation at lower temperatures occurs not by the liquid-gas phase transition but by some other mechanisms, e.g., the formation of cracks in the solid-like expanding region. Molecular dynamics simulations have been playing important roles in nuclear physics: both in the study of the nuclear structure and reaction. They can successfully describe the structure of nuclei and statistical properties of heavy-ion collisions taking account of many-body correlations and fluctuations. In addition to nucleon many-body systems, MD simulations are used also in QCD studies: UrQMD is now commonly used to analyze heavy-ion collisions with a quark-gluon plasma \\cite{refUrQMD}, and some dynamical properties of quark matter have been studied by a MD approach\\cite{qmmd,akimura}. One of the most important and challenging directions is to incorporate the wave nature of the quantum mechanics in the MD approach, which is based on the particle picture. A new framework beyond QMD and FMD/AMD in this direction is highly awaited \\cite{AMDV}. By such a breakthrough, MD should be a promising approach also for studying the dynamics of fission and fusion, which is a long-standing mportant problem in nuclear physics." }, "1207/1207.4329_arXiv.txt": { "abstract": "{The radial velocity (RV) technique is a powerful tool for detecting extrasolar planets and deriving mass detection limits that are useful for constraining planet pulsations and formation models.} {Detection limit methods must take into account the temporal distribution of power of various origins in the stellar signal. These methods must also be able to be applied to large samples of stellar RV time series} {We describe new methods for providing detection limits. We compute the detection limits for a sample of ten main sequence stars, which are of G-F-A type, in general active, and/or with detected planets, and various properties. We use them to compare the performances of these methods with those of two other methods used in the litterature. } {We obtained detection limits in the 2-1000 day period range for ten stars. Two of the proposed methods, based on the correlation between periodograms and the power in the periodogram of the RV time series in specific period ranges, are robust and represent a significant improvement compared to a method based on the root mean square of the RV signal. } {We conclude that two of the new methods (correlation-based method and local power analysis, i.e. LPA, method) provide robust detection limits, which are better than those provided by methods that do not take into account the temporal sampling.} ", "introduction": "The radial velocity (hereafter RV) technique is a powerful tool for detecting planets, but also for deriving detection limits (i.e. the upper limit to possible planet masses for different periods). Detection limits indeed represent invaluable information for either the study of specific objects, or and above all the derivation of quantitative constraints on the formation processes of planets. This is true for detection limits obtained with RV as well as direct imaging. Different criteria have been used to compute these detection limits. One is based on the root mean square (rms) of the RV data compared to the rms of the planetary RV \\cite[][for the principle]{galland05} and has been used by \\cite{lagrange09}. This rms-based method is very fast, but can significantly overestimate the detection limit in some cases. The signal of the planet (which has a certain period) is compared to the rms of the whole signal, which may contain strong power at periods very different from the planet period. This was the case for \\object{$\\beta$ Pic} \\cite[][hereafter paper I]{lagrange12}: \\object{$\\beta$ Pic} is a pulsating star, with a strong power in the domain 20-30 minutes, which dominates the rms computed over the RV signal, but this power is much weaker in the frequency domain in which we search for planets. Detection limits based on this rms are then overestimated. We therefore derived other methods that allow us to take into account the temporal behavior of the stellar noise. In paper I, we presented the first results for \\object{$\\beta$ Pic} and showed that we could improve the detection limits significantly for periods in the range from a few days to a few hundreds days. Here, we present these methods in detail and use them on a sample of ten stars (including \\object{$\\beta$ Pic} for comparison purposes) with various characteristics. We test their robustness, as the detection limit depends on the available data (temporal sampling) and the temporal structure of the stellar noise. These tests are made on a limited number of stars, that are representative of our large (250 stars) sample covering either early-type (A, F) stars or young solar-type stars. We also added for comparison purposes a slowly rotating solar-type main-sequence (MS) star. Their v.$\\sin$i ranges from a few km/s up to 17 km/s. Our long-term goal is to use these methods to obtain detection limits for the $\\sim$250 stars of our complete set of MS A-F stars, some of which are young stars sample that have been surveyed in the northern and southern hemispheres during searches for planets. Given the characteristics of our stellar sample, we focus mainly on Jupiter-mass planets. The star sample is described in Sect.~2. We compute the detection limits using four different methods: rms, correlation, peak, and {\\it local power amplitude} LPA, respectively described in Sect.~3 to 6. In Sect.~3 to 6, we also present the results and test the robustness of each of these methods for each of the relevant parameters. In Sect.~7, the detection limits obtained with the different methods are discussed and compared to each other, and we discuss specific stellar cases. We present our conclusions in Sect.~8. ", "conclusions": "We have determined and discussed robust detection limits in the range 4--800 days for ten stars, including \\object{$\\beta$ Pic}, using several new methods (Fig.~\\ref{res_a}). We have compared the obtained detection limits, as well as the robustness, of three new methods (two of which were introduced in paper I). These three methods have the advantage of taking into account the temporal distribution of the power in the observed RV, owing to both the temporal sampling and the presence of power of various origins in the stellar signal, and not only the temporal sampling as is done in bootstrap methods such as either the one described in \\cite{zechmeister09} and \\cite{wittenmyer06} or the rms method. The three methods tested in this paper are sensitive to different aspects of the temporal distribution of the power and temporal sampling. The correlation-based detection limit is related to the pattern introduced into the periodogram by the presence of planets (with power being introduced at a given period and interacting with the temporal sampling), but also to the global power in the periodograms. The peak detection limit results from the comparison between the amplitude of the peak that corresponds to the planet closest to the planet period, and the amplitudes of other peaks. Finally, the LPA method is related to how the total power due to the planet in a given period range compares to the same power in the stellar periodogram. We compared these methods to both the rms-based method, which is computed in the SAFIR program \\cite[][]{galland05,lagrange12}, and the bootstrap method described in \\cite{zechmeister09}, as this method is widely used. The correlation-based and LPA methods give the lowest detection limits in all cases (the latter often being better), while the rms method usually gives the largest, although the ratio varies significantly from one star to another. The peak method is not as efficient as the other two. For this small sample, the improvement with respect to the rms method for both methods is clearly illustrated by the ratio of the power at the origin of the rms RV to the power in the periodogram at the period we consider. The correlation-based and LPA methods are also the two most robust of the three. The correlation-based method uses a small number of parameters and is very robust, although it is not a practical method for actually detecting planets in an observed time series. It is indeed difficult to determine the threshold in the correlation between the periodograms and the impact of the presence of several planets on the pattern observed in the periodogram. The LPA method is also very robust, although the period window has a significant impact on the result. The peak method however is not as robust, as in some cases it is impossible to identify a planet peak in the periodogram owing to the temporal sampling of the observations, even for a very high planet mass. It also relies on many parameters. The main limitation of the peak method, despite it being the closest to the method actually used to detect a planet on a RV time series, is that one needs to identify automatically the peak corresponding to the planet period: in a significant number of cases, the planet peak is indeed not the largest one, even for a very massive planet, owing to the temporal sampling. In addition, the study of the robustness provides an estimation of the uncertainty in the derived detection limits. We conclude that the rms method is ideal for achieving a quick look. Both this method and the bootstrap method provides an efficient determination of the detection limits. However, we point out that, at least for the stars we have studied (F-G stars, with stellar activity and / or planets, MS stars), these methods may not give the best results, as they do not take into account the temporal response of the RV signal (due to either, for example, stellar activity or the presence of planets). Therefore, to obtain a more robust estimate, we recommand using both the correlation-based method and the LPA method, especially for times series corresponding to many observations that have well-defined peaks in the temporal periodogram. The use of more than one method is also useful for estimating the uncertainty in the detection limits. This study is in principle limited by the size of the sample and selection effects. However, even for a small number of stars that cover a large range of parameters, it has allowed us to derive some clear indications of to the expected improvement. An improvement exists even for stars with a very low RV jitter, so that the whole sample exhibits a coherent behavior." }, "1207/1207.3023_arXiv.txt": { "abstract": "We present the results from a search for HI emission from a sample of newly discovered dwarf galaxies in the M81 group. HI is detected in three galaxies, all of which are classified as BCDs. The HI masses of these galaxies are $\\sim 10^{6}$~\\ms, making these some of the lowest mass BCDs known. For these three galaxies FUV images (from GALEX) and \\ha\\ images (from the Russian 6m BTA telescope) are available.The \\ha\\ emission is very faint, and, in principle could be produced by a single O star. Further, in all cases we find offsets between the peak of the FUV emission and that of the \\ha\\ emission. Offsets between the most recent sites of star formation (i.e. those traced by \\ha) and the older sites (i.e. those traced by FUV) would be natural if the star formation is stochastic. In spite of the expectation that the effects of mechanical feedback from star formation would be most directly seen in the smallest galaxies with low gravitational potentials, we only see tentative evidence of outflowing HI gas associated with the star forming region in one of the galaxies. ", "introduction": "\\label{sec:int} Accurate determination of the faint end slope of the Luminosity Function in groups of galaxies is very important to resolve the discrepancy between the number of dwarf galaxies observed and that predicted by dark matter simulations of hierarchical structure formation \\citep{kly99,moo99}, especially in low-density environments \\citep{bar07,rob07}. An excellent candidate low-density environment group for carrying out a deep census of dwarfs is the M81 group of galaxies, for which it is possible to detect galaxies as faint as M$_R \\sim -9$, and for which accurate TRGB distances can be determined for individual galaxies. \\citet{chi09} present a census of the dwarfs in this group, determined from a large area survey using the Megacam instrument on the CFHT. A total of 22 candidate dwarf galaxies were found in this survey. Here we present follow up HI observations of some of these newly discovered dwarf galaxies. The observations were done at the Giant Metrewave Radio Telescope (GMRT) and targeted a sub-sample of galaxies that showed evidence for active star formation, and as such was expected to have a significant HI content. At the redshift of the M81 group, HI emission from group members is expected to overlap in velocity with emission from HI in our Galaxy. An interferometric survey is hence the only reliable method for robustly detecting gas associated with the dwarfs in this group. Deeper follow-up observations were obtained using the GMRT and the Westerbork Synthesis Radio telescope (WSRT) for the galaxies in which HI was detected. The target galaxies were selected based on the original Megacam CFHT survey data. However, subsequently FUV data from the GALEX satellite as well as \\ha\\ data from the Russian 6m BTA telescope also became available. We present and analyse this data in conjunction with the HI data from the GMRT and the WSRT. The galaxies for which we do have GMRT detections have HI masses $\\sim 10^{6}$ \\ms, and star formation rates $\\sim 10^{-4}$~\\ms/yr. They represent some of the smallest mass star forming galaxies known. Feedback into the ISM from star formation is expected to be most pronounced for small galaxies, and we use our data to explore this issue too. Below, in Sec.~\\ref{sec:obsr} we present the sample, briefly describe the observations and results. The results are discussed in detail in Sec.~\\ref{sec:res}, and the summary and conclusions presented in Sec.~\\ref{sec:summary}. ", "conclusions": "\\label{sec:summary} Our search for HI in the newly discovered dwarf galaxies in the M81 group has resulted in three new detections. The detected galaxies are some of the lowest mass BCDs known, with star formation rates $\\sim 10^{-4}$ \\ms/yr. The star formation can be traced by both \\ha\\ and FUV emission, and as is the case for most other dwarfs, the observed \\ha\\ flux is less than what would be predicted from the star formation rate deduced from the FUV. This is likely to be a result of stochastic sampling of the upper end of the mass function. The \\ha\\ emission is also offset from the peak of the FUV emission. We find tentative evidence of HI gas outflow at the edge of the stellar disc of one of the galaxies. Even though this result agrees with the expectation that the effects of mechanical feedback from star formation would be most directly seen in the smallest galaxies, surprisingly no clear evidence of feedback associated with star formation is seen in the HI discs of the other two galaxies." }, "1207/1207.5814_arXiv.txt": { "abstract": "We present our X-ray imaging spectroscopic analysis of data from deep {\\it Suzaku} and {\\it XMM-Newton} Observatory exposures of the Virgo Cluster elliptical galaxy NGC 4649 (M60), focusing on the abundance pattern in the hot interstellar medium (ISM). All measured elements show a radial decline in abundance, with the possible exception of Oxygen. We construct steady state solutions to the chemical evolution equations that include infall in addition to stellar mass return and SNIa enrichment, and consider recently published SNIa yields. By adjusting a single model parameter to obtain a match to the global abundance pattern in NGC 4649 we infer that introduction of subsolar metallicity external gas has reduced the overall ISM metallicity and diluted the effectiveness of SNIa to skew the pattern towards low $\\alpha/Fe$ ratios, and estimate the combination of SNIa rate and level of dilution. Evidently, newly-introduced gas is heated as it is integrated into, and interacts with, the hot gas that is already present. These results indicate a complex flow and enrichment history for NGC 4649, reflecting the continual evolution of elliptical galaxies beyond the formation epoch. The heating and circulation of accreted gas may help reconcile this dynamic history with the mostly passive evolution of elliptical stellar populations. In an appendix we examine the effects of the recent updated atomic database {\\em AtomDB} in spectral fitting of thermal plasmas with hot ISM temperatures in the elliptical galaxy range. ", "introduction": "There are two basic approaches to studying the formation and evolution of elliptical galaxies. One may directly examine the assembly history of the baryonic component by observing how the space density and morphological demographics of the population of elliptical galaxies, and their progenitors, develop over time. Alternatively one may conduct detailed investigation of those spectro-photometric properties that reflect their histories, and the scaling relations among these properties. While surveys at a range of redshifts indicate significant growth in both the number and sizes of ellipticals, ``archeological'' investigation finds that elliptical galaxy stellar populations are mostly in place at high redshifts ($z>2$ for massive systems) and passively evolve thereafter. Resolution of this apparent paradox is crucial for validation and elucidation of the prevailing hierarchical assembly paradigm of galaxy formation, where dynamical evolution and the triggering of star formation are presumed to be interconnected. The archeological approach traditionally relies on optical photometry and spectroscopy of the stellar population; however, the hot interstellar medium (ISM) provides a rich complementary site of diagnostic data that is accessible by means of X-ray observation (Loewenstein \\& Davis 2010, hereafter Paper I, and references therein; Pipino \\& Mattecucci 2011, hereafter PM11). The physical properties of the ISM in giant elliptical galaxies reflect the distinctive history and nature of these systems and, as such, markedly differ from those in the ISM of spiral galaxies such as the Milky Way. The most striking contrast is that, while mass loss from evolved stars is a primary source of gas in both spirals and ellipticals, most of the mass return in the latter is promptly heated to the high temperatures corresponding to the stellar velocity dispersion and does not currently participate in an ongoing star-gas cycle \\citep{mb03}. This distinct ecology provides a repository of information about processes in the distant and recent past of elliptical galaxies. The ISM is more responsive to energetic events than the stellar population which is, to first order, passively evolving. Evidence of feedback processes subsequent to the establishment of the stars and the nuclear supermassive black hole (SMBH) can only be found in the ISM. The ISM mass within elliptical galaxies is a small fraction of the stellar mass return integrated over the post-star-formation history of an elliptical galaxy, implying the existence of some steady and/or episodic means of gas removal into the intergalactic, or some circumgalactic, medium. A successful gasdynamical model for ellipticals must explain ISM that generally are, nevertheless, sufficiently massive to rule out persistent supersonic or transonic galactic winds, and also display a large scatter in gas-to-stellar mass ratio ($M_{\\rm gas}/M_{\\rm stars}$; Mathews et al. 2006 and references therein). While ram pressure stripping may be important at times for some ellipticals in rich clusters, large scale flows driven by Type Ia supernovae (SNIa) and active galactic nuclei (AGN) likely dominate the redistribution of the ISM and its removal from the galaxy. The amount of energy associated with SNIa exploding at the estimated rate in ellipticals is sufficient to drive a galactic wind. Although effects such as depletion through the formation of dust (PM11) may reduce the discrepancy, the high expected ISM Fe abundance is at odds with X-ray observations (e.g., Paper I, PM11), and a certain amount of fine-tuning in the SNIa rate as a function of time is required for SNIa-driven winds to explain the $M_{\\rm gas}/M_{\\rm stars}$ scatter without driving out virtually all of the gas and rendering ellipticals undetectable as diffuse soft X-ray sources. Feedback from active galactic nuclei (AGN), implicated in quenching star formation in ellipticals \\citep{sch07,cat09} and establishing the scaling relations between SMBH mass and stellar mass or velocity dispersion (e.g., DeBuhr, Quataert, \\& Ma 2011 and references therein), may also drive galactic flows at later times. AGN interaction with hot ISM is evident in {\\it Chandra} X-ray Observatory observations of several ellipticals \\citep{mn07,nul09}. AGN feedback is fundamentally self-regulating and intermittent: the SMBH is fueled by an initial inflow that is subsequently reversed as the AGN powers up, thus cutting off the supply of gas and enabling the cycle to restart as gas once again flows inward towards the now-dormant SMBH. There has been great progress in implementing and applying hydrodynamical simulations that include various feedback prescriptions \\citep{mbb04,bmhb09,cop10}, however it remains to be seen whether the full range of ISM observables -- X-ray luminosities, temperatures, metallicities and abundance patterns -- and their galaxy-to-galaxy variations can be accurately and self-consistently modeled without additional, and perhaps fundamental, adjustment \\citep{mb03}. In addition, many ellipticals are embedded in an extended hot intergalactic or circumgalactic medium, and may interact with their environment in a variety of ways \\citep{mj10}. In particular ellipticals may accrete some combination of primordial gas and gas ejected from winds at earlier epochs \\citep{pm04,dfo11}. Such an external medium may also confine outflows. Recent X-ray measurements of elemental abundances in the hot ISM of ellipticals beg consideration in determining the future direction of these models, and constraining more general theories of the chemical evolution of these systems. The ISM abundance pattern may be a particularly sensitive diagnostic of dynamical processes in ellipticals. If the sole source of interstellar gas is stellar mass loss, the ISM abundance pattern in ellipticals will reflect that in evolved stars found to have $[\\alpha/Fe]_{\\rm stars}>0$ ($[\\alpha/Fe]_{\\rm stars}$ is defined as the log of the abundance ratio, with respect to Fe, of elements primarily produced through $\\alpha$ capture to Fe -- relative to the solar ratio). Because the internal injection and flow of energy, and the various mechanisms of mass exchange with the external environment, each imprint distinctive departures from the baseline ISM abundance pattern determined by local stellar mass loss, this pattern may be analyzed to probe for signatures of these processes. Measurement of elliptical galaxy abundance patterns encompassing a broad range of elements, and extending to large radii may be made with the {\\it Suzaku} Observatory, facilitated by the low internal background and relatively sharp energy resolution of the XIS CCD detectors. In Paper I we derived the abundance pattern in the elliptical galaxy NGC 4472 from analysis of {\\it Suzaku} spectra, supported by analysis of co-spatial {\\it XMM-Newton} Observatory spectra. Application of simple chemical evolution models to these data, led us to conclude that the abundances may be explained by a combination of $\\alpha$-element enhanced stellar mass loss and direct injection of ejecta from SNIa exploding at a rate $\\sim 4-6$ times lower than the standard value. In addition, we discovered abundance anomalies in the sense that no published set of SNIa yields could simultaneously reproduce the inferred Ca and Ar, and Ni abundances; and (confirming prior results summarized in Paper I) in the sense that standard core collapse nucleosynthesis models evidently overproduce O by $\\sim 2$. In this paper we adopt a broadly similar approach in our investigation of NGC 4649 (M60). As is the case for NGC 4472, NGC 4649 is a giant elliptical galaxy in the Virgo cluster with an old stellar population enhanced in $\\alpha$-elements -- but with several notable differences both optically and in X-rays. In particular NGC 4649 has a substantial major axis rotation (Brighenti et al. 2009, and references therein), and is considerably more compact in X-rays. We have adjusted our spectral analysis procedures in response to the distinctive X-ray characteristics of NGC 4649, and have revised and expanded our models to consider possible episodes of inflow, and a recently published set of SNIa yields. We detail our data reduction and spectral analysis procedures in Sections 2 and 3, where we present our derived NGC 4649 hot ISM thermal and chemical properties and their radial variation. In Section 4 we focus on interpreting the global abundance pattern in the context of the relative contributions of metal enrichment from stellar mass return, SNIa, and inflow of extragalactic material using steady state solutions to the equations of chemical evolution. Section 5 includes a summary of our conclusions, and discusses possible implications of our results. In an appendix we examine the effects of the recent updated atomic database {\\em AtomDB} in spectral fitting of elliptical galaxy hot ISM such as NGC 4649. ", "conclusions": "The distinctive history of giant elliptical galaxies results in the creation and maintenance of an extensive ISM, dominated by hot gas, that contains a fossil record of that very history: evolution determines ecology, ecology enables archeology, archeology illuminates evolution. Based on estimates of supersolar stellar $\\alpha$-to-Fe ratios and SNIa rates in excess of $>0.1$ SNU, one naively expects an ISM abundance pattern in elliptical galaxies that is supersolar across the board, and increasingly so for elements more proficiently synthesized in SNIa. In actuality, measured ISM abundances are solar or subsolar for all elements and do not deviate strongly from solar ratios, possibly excepting Ni. A simple reduction in the SNIa rate (or by assuming that SNIa ejecta fail to mix into the hot ISM; Brighenti \\& Mathews 2005) does not resolve this puzzle -- the effective rate would need be reduced to very low values to explain the modest Fe abundance {\\it level}, begging explanation for why the {\\it ratios} do not then mirror the $\\alpha$-element enhanced pattern in the mass-losing stars. We bring these issues into sharper focus, and attempt to make progress on their resolution, through measurement of the abundance pattern in the elliptical galaxy NGC 4649 derived from deep {\\it Suzaku} and {\\it XMM-Newton} observations. The measurement of abundances in the hot ISM of elliptical galaxies remains problematic, despite the high quality of the spectra extracted from these data due to effects of angular resolution in the case of {\\it Suzaku}, background systematics in the case of {\\it XMM-Newton}, and lingering atomic physics uncertainties (see Appendix A). Nevertheless, by demanding a degree of cross-mission consistency, and focusing most on conclusions that rely on relatively model-independent abundance ratios integrated over the galaxy, we find that robust constraints on elliptical galaxy evolution can be inferred. Towards this end, we compare the galaxy-wide average ISM abundance pattern with one-zone steady state solutions to the equations of chemical evolution that depend on a single parameter that characterizes the relative contributions of SNIa, stellar mass loss, and inflow to the ISM metal inventory. We vary this parameter to fit the data, utilizing the most recent solar standard abundance scale \\citep{agss09} and SNIa yields \\citep{mae10}, and reproduce the observed abundance pattern in NGC 4649, thus lending some insights into the contributions of various sources of ISM metal enrichment. Our approach has its drawbacks and limitations, and should be considered provisional as we expand the data analysis sample and refine our formalism. By adopting the steady-state limit of equations (1) and (2), we focus on the relative contributions of stellar mass loss, SNIa, and inflow/outflow (in the broad sense in which they are defined here), but gloss over the fact that these are not strictly co-eval and that the dilution factor refered to as ``inflow'' may in part correspond to a pre-exisiting group environment that continues to be reflected in gas phase abundances at large galactic radii. Encapsulating the relative impact of stellar mass return, dilution, and SNIa on the pattern of abundance ratios in a single parameter has the virtue of facilitating accurate model/data comparison, and emphasizes their interplay and degeneracy -- e.g., the challenge in distinguishing the combination of low SNIa rate and low dilution from that of high SNIa rate and high dilution. However, this complicates the interpretational framework somewhat while, at the same time, limiting it by leaving these unbroken degeneracies in place. We are confident that these degeneracies may be broken as the sample, and our confidence in absolute abundance measurements, grow. We find that a key ingredient in models that successfully reproduce the observed pattern is dilution via inflow of subsolar metallicity extragalactic gas with an (assumed) abundance pattern similar to that in the stars at an average rate comparable to, or greater than, that of stellar mass loss. As a result the overall ISM metallicity is reduced while, at the same time, the effectiveness of Fe group enhancement from direct SNIa injection is diluted. Although SNIa rates of 0.1 SNU or more are not required, we now find that they may be accommodated (and, in fact, that specific rates $<0.02$ SNU are excluded by limits on the ISM $[\\alpha/Fe]$ ratio relative to that in the stars) provided that a significant fraction of the ISM is of external origin. Astrophysically reasonable magnitudes of SNIa enrichment and external dilution may play off against one another, as shown in Figure 15, to produce subsolar abundances in roughly solar proportions as observed. In a similar vein, \\cite{bmhb09} modeled the two-dimensional gasdynamics of NGC 4649 and found that inflow of circumgalactic gas with metallicity comparable to the mean stellar value could dilute the effects of SNIa enrichment (in their models, $\\theta_{\\rm SNIa}/\\theta_{\\rm MR}=0.84$ using our scale) on the Fe abundance by a factor of $\\sim 3$. We cannot precisely identify the origin and means of delivery of the extragalactic material. Inflow may be quasi-steady, or intermittent following the episodes of outflow that must occur to prevent the gas from accumulating to levels beyond what is observed \\citep{cpm09}. The reservoir from which gas accretes may be an intracluster, intragroup, or filamentary intergalactic medium; and may have been previously ejected from the same galaxy (or other galaxies). Alternatively, the extragalactic material may originate in discrete instances of galaxy merging with gas-rich systems. Some combination of these mechanisms would seem most likely, with their relative importance depending on galaxy history and environment \\citep{lab11}. The level and pattern of ISM metal enrichment will reflect this and, with detailed modeling and additional observational analysis, may prove an important diagnostic of the nature of post-formation elliptical galaxy interactions. Thus, the X-ray abundance pattern confirms the complex flow and enrichment history of NGC 4649 proposed in \\cite{bmhb09}, and reflects the dynamic, continually evolving nature of elliptical galaxies implied by optical observations \\citep{fab07}. The earliest epochs of elliptical galaxy formation are characterized by the interplay of major mergers, inflow of primordial gas, vigorous star formation, and powerful outflows \\citep{cpm09,man10,arr10}. Even after the quenching of star formation has halted the main formation epoch, ellipticals do not evolve purely in a passive manner in isolation \\citep{tho10}. Evidence of mergers and other interactions, indicate that ellipticals continue to grow up to the present day \\citep{tra00b,vd08,tal09}, as required to explain the observed increase in the size of ellipticals and the continual build up of the red sequence \\citep{njo09,vdw09,sha10,rob10,tfd11,dam11,cas11}. In addition, there are signs of the presence of modest amounts of more recently formed stars -- often in ellipticals with morphological signs of interaction \\citep{tra00b,kav08,kav09,san09,kav10,veg10}. However, evolution subsequent to the primary star forming epoch must proceed in such a way as to be consistent with the old and passively evolving nature of elliptical galaxy stellar populations inferred from their individual properties and scaling relations \\citep{pm06,pm08}. Attempts to resolve this apparent paradox, alluded to in the introduction, often appeal to the process of ``dry'' merging \\citep{tal09,kav10,coo11} that involves very little gas -- and hence new star formation (as contrasted with the $z>2$ ``wet'' mergers where stars are efficiently formed). The observed scaling relations and predominance of old stellar populations are thus preserved. However it is not clear that mergers of this type occur in sufficient numbers \\citep{hop10}; and, most galaxies interacting with ellipticals are not, in fact, gas-free \\citep{dep10,so10}. Massive elliptical galaxies are themselves not gas-poor, and the observed properties of the hot ISM are often neglected in these considerations. We have shown that the chemical structure of the hot ISM in ellipticals not only confirms a continual environmental impact, but demonstrates that the accreted gas is hot or is heated as it is introduced and interacts with pre-existing hot gas. This implies a reduction of the star formation efficiency of any cold gas that may be newly introduced either in the form of mergers with small gas-rich galaxies or via a smoother accretion process \\citep{ker09}. \\lastpagefootnote The entire one-dimensional ``humidity-based'' classification of mergers\\footnote{in addition to ``wet'' and ``dry'', mergers have also been referred to as ``moist'' \\citep{san09} and ``damp'' \\citep{for07}} implicitly assumes a one-to-one correspondence between gas content, dissipation, and star formation that breaks down in giant ellipticals known to be filled with hot gas. An encounter involving such a galaxy is never gas-poor, and may involve substantial dissipation if the companion galaxy is also gas-rich. However, these mergers may prove to be a hostile setting for new star formation, and are effectively ``dry.'' Alternatively, the late-time accretion may be dominated by hot mode accretion \\citep{vdV11}. Understanding elliptical galaxy evolution requires tracking the three ``phases'' of baryonic matter -- stars, (cold) star-forming gas, and (hot) inert gas -- and how their mutual exchange of mass, metals, and energy as a function of age, environment, and type of merger proceeds as the galaxy evolves subject to internal and external influences \\citep{del10,lu11}. The rapid build-up of the bulk of the stars that compose present-day giant elliptical galaxies was quenched by some process that maintained a galaxy ecology free of significant amounts of cold gas. Suggesed mechanisms, that also may account for the observe dichotomy in the characteristics of galaxies, include quenching via gravitational heating \\citep{db08,bd11,nb07,jno09,cat09}, and via AGN feedback (Bell et al. 2012 and references therein). We see that ISM abudances are a potential diagnostic of the continuing operation of star formation quenching. Our conclusions are driven by the robust analysis results indicating that ISM abundances are solar or less with ratios that lie between the $\\alpha$-element enriched pattern expected from pure enrichment from stellar mass loss and the O/Ne/Mg-poor pattern expected from the combination of stellar mass loss and direct SNIa injection at $\\theta_{\\rm SNIa}\\sim 1$. As NGC 4649 is a member of the Virgo Cluster and displays a compact X-ray morphology, the evolution of its hot ISM (e.g., stripping of the outer regions) may be affected by the interaction of the galaxy (or a subgroup to which it belongs) with the ICM. Nevertheless, given the universality of the departure of observed elliptical galaxy hot ISM abundances from the naive expectations described above, we suggest that the following main conclusions we draw for NGC 4649 may be generalized: \\begin{itemize} \\item ISM abundance patterns indicate that metal enrichment from stellar mass loss and from direct injection of SNIa, and dilution from infall are all significant when averaged over the residence time of the ISM. SNIa injection prevents the ISM from reflecting the supersolar stellar $[\\alpha/Fe]$ ratio, and infall dilutes the overall ISM metallicity. \\item An amount of low metallicity gas exceeding that originating in stellar mass loss rate must be introduced into the galaxy if one is to reconcile the ISM Fe abundance with the standard specific SNIa rate of 0.16 SNU. \\item We find further support to the picture of ellipticals as ``open ecosystems'' that continually grow, and exchange mass and metals with the intergalactic environment and/or other galaxies. \\item The gas that encroaches into or is introduced into the ISM must already be hot, or must be heated by mixing with the hot gas and/or some other internal mechanism in order to explain why these processes are accompanied by little or no star formation. \\end{itemize}" }, "1207/1207.0093_arXiv.txt": { "abstract": "We show how the existence of a relation between the star formation rate and the gas density, i.e.\\ the Kennicutt-Schmidt law, implies a continuous accretion of fresh gas from the environment into the discs of spiral galaxies. We present a method to derive the gas infall rate in a galaxy disc as a function of time and radius, and we apply it to the disc of the Milky Way and 21 galaxies from the THINGS sample. For the Milky Way, we found that the ratio between the past and current star formation rates is about $2-3$, averaged over the disc, but it varies substantially with radius. In the other disc galaxies there is a clear dependency of this ratio with galaxy stellar mass and Hubble type, with more constant star formation histories for small galaxies of later type. The gas accretion rate follows very closely the SFR for every galaxy and it dominates the evolution of these systems. The Milky Way has formed two thirds of its stars after $z=1$, whilst the mass of cold gas in the disc has remained fairly constant with time. In general, all discs have accreted a significant fraction of their gas after $z=1$. Accretion moves from the inner regions of the disc to the outer parts, and as a consequence star formation moves inside-out as well. At $z=0$ the peak of gas accretion in the Galaxy is at about $6-7 \\kpc$ from the centre.% ", "introduction": "Star formation is the fundamental process that shapes galaxies into different classes. Although the majority of stars in the local Universe are found in spheroidal systems, most of the star formation is contributed by disc galaxies of the later types (beyond Sb). The key ingredient for star formation is cold ({\\it star-forming}) gas, which is present almost exclusively in disc galaxies. However, the amount of cold gas currently available in galaxy discs appears rather scant. \\citet{Kennicutt98a} estimated that disc galaxies have current star formation rates ranging from a few to about $\\simeq 10 \\moyr$. Thus, considering a typical gaseous mass of a few $10^9 \\mo$, the gas consumption time scale (i.e.\\ the time needed to exhaust the gas fuel with a constant star formation rate) is always of the order of a few Gigayears. This result, known as the gas-consumption dilemma \\citep{Kennicutt83}, suggests the need for continuous accretion of cold gas onto galaxy discs \\citep[e.g.][]{Sancisi+08}. Several pieces of evidence show that disc galaxies should collect fresh gas from the environment in order to sustain their star formation. In the Milky Way, the star formation rate (SFR) in the solar neighbourhood appears to have remained rather constant in the last $\\sim10 \\Gyr$ \\citep[e.g.][]{Twarog80, Rocha-Pinto+00, Binney+00}, suggesting a continuous replenishment of the gas supply. Simple (closed-box) models of chemical evolution for our Galaxy predict too few metal deficient \\textit{G-dwarf} stars than observed \\citep{Searle&Sargent72, Pagel&Patchett75, Haywood01}. These observations are easily explained by accounting for infall of fresh unpolluted gas \\citep[e.g.][]{Chiappini+97, Chiappini+01}. Observations of Damped Lyman Alpha systems show almost no evolution in the neutral gas content of structures in the Universe \\citep[][]{Zwaan+05,Lah+07,Prochaska+09}. \\citet{Hopkins+08} pointed out that this constancy of gas density can be explained assuming a rate of gas replenishment proportional to the universal SFR density. \\citet{Bauermeister+10} converged to a similar result when comparing the evolution of the molecular gas depletion rate with that of the cosmic star formation history (SFH). Finally, the derivations of SFHs for galaxies with different stellar masses consistently show that late type systems do have a rather constant SFR throughout the whole Hubble time \\citep[e.g.][]{Panter+07}. These findings, other than being the signature of {\\it downsizing} in cosmic structures, point at a continuous infall of gas onto galaxies of late Hubble types. The way gas accretion into galaxies takes place is still a matter of debate. The classical picture states that thermal instabilities should cause the cooling of the hot coronae that surround disc galaxies and be the source of cold gas infall \\citep[e.g.][]{Maller&Bullock04, Kaufmann+06}. However, recent studies have shown that, due to a combination of buoyancy and thermal conduction, hot coronae turn out to be remarkably stable and thermal instability does not appear to be a viable mechanism for gas accretion \\citep{Binney+09, Joung+11}. On the other hand, accretion may take place in the form of cold flows but the importance of this process is expected to drop significantly for redshift $z<2$ \\citep{Dekel&Birnboim06, vdVoort+11}. Observations of local gas accretion at 21-cm emission seem to show too little gas around galaxies in the form of \\hi\\ (high-velocity) clouds to justify an efficient feeding of the disc star formation \\citep{Sancisi+08, Thom+08, Fraternali09}. However, UV absorption towards quasars and halo stars point to the possibility that ionised gas at temperatures between a few 10$^4$ and a few $10^5 \\K$ could fill the gap between expectations and data \\citep{Bland-Hawthorn09, Collins+09, Lehner&Howk11}. In this paper, we estimate gas accretion in galactic discs indirectly by comparing basic physical properties of galaxies today. Our model relies on the existence of a law relating the star formation rate density (SFRD) and the gas surface density ($\\Sigma_{\\rm gas}$) holding at every redshift, i.e.\\ the Kennicutt-Schmidt (K-S) law \\citep{Schmidt59, Kennicutt98a}. This approach has similarities with the model of \\citet{Naab&Ostriker06} that we describe in detail in Section \\ref{sec:discussion}. The paper is organized as follows. Section \\ref{sec:model} describes our method. In Section \\ref{sec:data} we apply our model to the Milky Way disc and to a sample of external discs. In Section \\ref{sec:discussion} we discuss our results. Section \\ref{sec:conclusions} sums up. ", "conclusions": "\\label{sec:conclusions} In this paper we have proposed a simple method to derive the amount of gas needed for the star formation to proceed in a galactic disc, using the Kennicutt-Schmidt law. We found that in typical disc galaxies (Sb or later types) most star-forming gas is not in place in the disc at the time of formation but needs to be slowly acquired from the surrounding environment. We derived the gas accretion rate as a function of time for the Milky Way and other 21 disc galaxies from the THINGS sample. We parametrized the SFH with a simple polynomial function with a parameter ($\\gamma$) that defines the slope of the SFH, $\\gamma=1$ stands for a flat SFH. We summarize our results as follows: \\begin{enumerate} \\item the disc of the Milky Way as a whole (not only the Solar neighborhood) formed stars in the past a rate that was $\\sim2-3$ times larger than now; \\item the steepness of the SFH is a function of galaxy mass and Hubble type, with late type galaxies having nearly flat or inverted SFH; \\item gas accretion is tightly linked to star formation, a constant ratio between the two is maintained for a large fraction of the life of a galaxy, pointing at a reciprocal interplay; \\item galaxy discs must have experienced accretion of large fraction of their mass (in gaseous form) after $z=1$, unless the K-S law had a strong evolution with $z$; \\item gas accretion progressively moves from the inner to the outer regions of galaxy discs, as a consequence star formation also proceeds inside-out; \\item the efficiency of gas accretion is less than 1 and decreases with time in all galaxies expect some late-type dwarfs, as a consequence the gas consumed by star formation is not completely replenished and the galaxy eventually stops forming stars; \\item the present-time accretion rates we derive are very uncertain, but even so there are indications that some large spirals may not need ongoing gas accretion. \\end{enumerate} An obvious development of this study would be to incorporate metallicity. Our results on the amount of gas infall needed for the Galaxy are broadly in agreement with the amount of infalling material needed by chemical evolution models. Our independent determination of the infall rate as a function of time and radius could be taken as an input in chemical evolution models to obtain a better understanding of the accretion mechanisms that feed the star formation in galaxy discs." }, "1207/1207.3090_arXiv.txt": { "abstract": "We examine the agreement between the observed and theoretical low-mass ($< 0.8\\,M_{\\odot}$) stellar main sequence mass-radius relationship by comparing detached eclipsing binary (DEB) data with a new, large grid of stellar evolution models. The new grid allows for a realistic variation in the age and metallicity of the DEB population, characteristic of the local galactic neighborhood. Overall, our models do a reasonable job of reproducing the observational data. A large majority of the models match the observed stellar radii to within 4\\%, with a mean absolute error of 2.3\\%. These results represent a factor of two improvement compared to previous examinations of the low-mass mass-radius relationship. The improved agreement between models and observations brings the radius deviations within the limits imposed by potential starspot-related uncertainties for 92\\% of the stars in our DEB sample. ", "introduction": "The disagreement between the theoretical and observational low-mass, main sequence mass-radius (henceforth MR) relationship has been recognized for nearly four decades \\citep{Hoxie1970,Hoxie1973}. Although, only in the past two decades has the disagreement become overwhelmingly apparent with the reduction of observational uncertainties \\citep[for an excellent review, see][]{Torres2010} and the development of sophisticated low-mass stellar models \\citep{BCAH98}. The primary line of evidence stems from the study of detached double-lined eclipsing binaries (hereafter DEBs) with additional support garnered by direct measurements of stellar radii via interferometry \\citep[e.g.,][]{Berger2006,vBraun2012}. These observations routinely suggest that stellar evolution models systematically under predict stellar radii by 5~--~15\\% and over predict effective temperatures at the 3~--~5\\% level. However, it is presently not clear whether the routinely quoted 5~--~15\\% disagreement is representative of true radius discrepancies or whether there are other factors contributing to the derivation of such large radius errors. One such factor derives from the fact that previous studies focusing on the comparison between models and observations have generally applied a limited sample of isochrones to their data. Largely, these sets are comprised of 1~Gyr and 5~Gyr, solar metallicity isochrones. This is predominantly a consequence of the limited age and metallicity range of currently available low-mass stellar models. Age and metallicity effects are less important in the low-mass regime, but the stringent uncertainties quoted by observational efforts preclude the use of such a limited set of isochrones. For example, \\citet{Burrows2011} discovered non-negligible radius variations in brown dwarfs and very-low-mass stars ($<0.1\\,M_{\\odot}$) when allowing for a more comprehensive set of metallicities. Isochrones with metallicities spanning a range characteristic of the local galactic neighborhood are therefore essential to accurately assess the validity of stellar evolution models. Furthermore, one must also consider that the population of well-characterized DEBs has, until recently, consisted of eight systems. While unlikely, it is not unimaginable that those eight systems were more the exception than the rule in terms of their lack of consistency with stellar evolution models. Since publication of the \\citet{Torres2010} review, the population of well-characterized, low-mass DEBs has more than doubled. The availability of this new data allows for a more accurate statistical characterization of the agreement (or lack of) between the MR relationship defined by models and observations. When discrepancies are observed, they are typically attributed to the effects of a large scale magnetic field \\citep[e.g.,][]{Ribas2006, Lopezm2007, Morales2008, Morales2009a, Torres2010, Kraus2011} as DEBs are often found in tight, short-period orbits with periods under three days. Tidal interactions and angular momentum conservation act to synchronize the orbital and rotational periods of the components, increasing the rotational velocity of each star in the process. The dynamo mechanism, thought to be responsible for generating and sustaining stellar magnetic fields, is amplified as a result of the rotational spin-up and enhances the efficiency of magnetic field generation within the star. Each component in the binary system is then more able to produce and maintain a strong, large-scale magnetic field than a comparable single field star. The effects of a large-scale magnetic field are thought to be two-fold: convective motions within the star are suppressed and the total surface coverage of starspots is increased. In both cases, a reduction in the total energy flux across a given surface within the star occurs, forcing the stellar radius to inflate in order to conserve flux \\citep{Gough1966}. Recent attempts at modeling these effects have indicated that an enhanced magnetic field is a plausible explanation, although the primary physical mechanism affecting the structure of the star is still debated \\citep{MM01,Chabrier2007,MM11}. Regardless of the precise physical mechanism, magnetic fields should betray their presence through the generation of magnetic activity in the stellar atmosphere. If magnetism is responsible for the observed inflated stellar radii, then correlations should be expected between individual stellar radius deviations and magnetic activity indicators (i.e., chromospheric H$\\alpha$ and CaII H \\& K emission, coronal x-ray emission, etc.). Tantalizing evidence of such correlations has been reported previously by \\citet{Lopezm2007} and \\citet{Morales2008}. However, recent evidence appears to stand in contrast with the current theory. Two systems, LSPM J1112+7626 \\citep{Irwin2011} and Kepler-16 \\citep{Doyle2011}, were discovered that have wide orbits with approximately forty-one day periods. Despite this, both appear to display discrepancies with stellar evolution models. In these systems, the component stars should be evolving individually with mutual tidal interactions playing a negligible role in the overall angular momentum evolution. The stars should be spinning down over time due to magnetic breaking processes \\citep{Skumanich1972}, meaning the stars should not be as magnetically active compared with short period binary systems. The contrast is particularly evident for LSPM J1112+7626, where a rotation period of sixty-five days was detected via starspot modulation in the out-of-eclipse light curve. Gyrochronology suggests that the system has an age of approximately 9 Gyr \\citep{Barnes2010} and implies further that the secondary is likely slowly rotating and should, therefore, not shown signs of strong magnetic activity or an inflated radius. A third system also appears to defy the current hypothesis. KOI-126 \\citep{Carter2011} is a hierarchical triple system with two low-mass, fully convective stars in orbit around a 1.35 $M_{\\odot}$ primary. The two low-mass stars are orbiting each other with a period of 1.77 days. Therefore, they should show signs of inflated radii due to enhanced magnetic activity. However, it has been shown that the two low-mass, fully convective stars were in agreement with model predictions when considering their super-solar metallicity and the age of the higher mass primary \\citep{Feiden2011}. This agreement was further confirmed by \\citet{Spada2012}. Metallicity has been proposed previously as a solution to the observed MR discrepancies, but for the case of single field stars \\citep{Berger2006}. This was contradicted shortly thereafter by \\citet{Lopezm2007}, most notably for DEBs. Although, we must consider that the radius discrepancy-metallicity correlation is severely complicated by the fact that metallicities of M dwarfs are notoriously difficult to determine observationally. Finally, developments in light curve modeling of spotted stars has generated interesting results. The presence of large polar spots may alter the light curve analysis of DEBs by modifying the eclipse profile. These modifications lead to 2~--~4\\% uncertainties in the derived stellar radii \\citep{Morales2010, Windmiller2010, Kraus2011}. Thus far, only two DEBs (GU~Boo and CM~Dra) have been thoroughly tested for their sensitivity to spots. Systematic uncertainties may therefore dominate the error budget, casting a shadow of doubt on the observed radius discrepancies, which are often made apparent due to the minuscule random uncertainties. The uncertainties and developments outlined above have motivated us to reevaluate the current state of the low-mass MR relationship. In what is to follow, we use a large grid of theoretical stellar evolution isochrones in an effort to compare the low-mass models of the Dartmouth Stellar Evolution Program (DSEP) with DEB systems that have well constrained masses and radii. We then explore how potentially unaccounted for systematic uncertainties have the ability to create the appearance of discrepancies when neglected and mask real ones when considered. Section \\ref{sec:data} will present the DEB sample followed by a description of the stellar models in Section \\ref{sec:models}. The isochrone grid and fitting procedures will be explained in Section \\ref{sec:proc}. Results will be presented in Section \\ref{sec:result} followed by a discussion of the implications of our findings in Section \\ref{sec:disc}. We conclude with a brief summary of the entire study in Section \\ref{sec:summary}. ", "conclusions": "\\label{sec:disc} Data presented in this study hint at two competing explanations for the occurrence of the differing model and observational MR relations. It is entirely plausible that stellar evolution models are not incorporating key physical processes that can account for the observed discrepancies. This view is not new and has always accompanied discussions of low-mass models \\citep{Hoxie1970,Hoxie1973,CB97,BCAH97,BCAH98}. Typically cited is our incomplete knowledge dealing with the complex array of molecules present in M-dwarf atmospheres as well as the lack of structural changes induced by a large-scale magnetic field. Included in the latter are both the effects on convective energy transport and the emergence of spots on the stellar photosphere. The other scenario is one in which the neglect of systematic uncertainties is driving the apparent discrepancies. We have shown that by allowing for realistic variation in age and metallicity of the models, the radius residuals are of the same magnitude as the potential systematic uncertainties. Note, this is without any modification to the solar calibrated mixing-length or the inclusion of any non-standard physics. Previous studies have indicated that systematics may help alleviate the size of the radius residuals, but have required additional modifications to the models in order to fully reconcile models with observations. It is now clear that we must work to constrain and minimize systematic uncertainties in observations of DEBs to allow for them to provide an accurate test of stellar evolution models. \\subsection{Radius Deviations} The MAE between our models and the observations was 2.3\\%, a factor of two improvement over the canonical \\emph{minimum} of 5\\%. We found deviations of no more than 4\\%, with the exception of a few stars, instead of ubiquitous 5~--~15\\% errors, as is often quoted. Despite improving the situation, panel~A of Figure~\\ref{fig:drad} illustrates that the models are still unable to fully reproduce the observed stellar radii. Accepting the factor of two improvement presented in Section~\\ref{sec:result1}, the paradigm of broad disagreement between models and observations is shifted to one where agreement is broad, and large discrepancies are an exception. With the radius deviations typically less than 4\\%, an evaluation of the systematic errors becomes imperative. Formerly, systematic uncertainties of about 4\\% were incapable of relieving radius deviations greater than 5\\%. Stellar evolution models still appeared to disagree with DEB observations even after the inclusion of systematic errors. The reason for the factor of two improvement is twofold. First, we have calculated models with a finer grid of metallicities. DSEP utilizes an EOS that enables models to be more reliably calculated for super-solar metallicities, allowing for a greater range of stellar compositions to be considered. Low-mass stellar models with super-solar metallicity have previously been unavailable for comparison with low-mass DEB data. Thus, the ability to extend our model set to super-solar metallicities allows for more flexibility in attempting to match the observed properties of DEB components. Second, a far larger number of low-mass DEBs with precisely measured radii were available to us as opposed to previous studies. Before the publication of TAG10, there were only eight systems that met the criteria necessary to accurately constrain stellar evolution models. Following the TAG10 review, the total number of systems that met the necessary criteria more than doubled with the addition of ten newly characterized DEBs. These additional systems appear to be more in line with the results of standard stellar evolution theory. However, well-known discrepant systems remain noticeably discrepant and still require further explanation. We must now ask, ``what belies the current discrepancies between models and observations?'' Figure~\\ref{fig:drad} favors the hypothesis that non-standard physics are absent from current stellar evolution models. Our larger data set allows us to notice that stars of similar masses from different DEB systems appear to be discrepant at varying levels, an effect a single set of standard models cannot correct. However, this is contingent upon the accuracy our age and metallicity predictions. Until we have better empirical age and metallicity estimates for the various systems, it is too difficult to ascertain the true level of discrepancy for any individual star. The efficiency of convective energy transport is of greatest interest. It is possible that convection is naturally inefficient. Although, we gather from Figure~\\ref{fig:drad} that suppression of convective energy transport must be tied to a stellar property that is largely independent of mass. Simple parametrization of the suppression of convection is too uniform over a given mass regime to fully account for the observed differences in stellar radii for stars with similar masses (see also Appendix~\\ref{sec:vmixl}). Any effort, either theoretical or observational, to constrain the physics of convection in low-mass stars will lend crucial insight. The most favored option, is that convection is not intrinsically inefficient, but that a large-scale magnetic field acts to suppress convective motions \\citep{MM01,Chabrier2007,MM11}. Stellar evolution models self-consistently incorporating the effects of large-scale magnetic fields will help on this front. Observations of cool-star magnetic field strengths and topologies will then provide a means of judging the validity of any new models. One final hypothesis is that starspots may affect the structure of stars and generate the inflated radii we observe. \\citet{Chabrier2007} investigated such a possibility by artificially reducing the total stellar bolometric luminosity in an effort to mimic the effects of spots. They found that radius discrepancies were relieved with their parametrization. However, it is still not apparent whether spots reduce the total bolometric flux or if they locally shift flux to longer wavelengths \\citep{Jackson2009}, preserving the total luminosity. Ultimately, if starspots are required in stellar evolution models, their inclusion is necessitated in the analysis of DEB light curves. Quantifying the effect starspots may have on observed DEB light curves is extremely difficult. Obtaining accurate knowledge of the total surface coverage of spots, the total number of spots, their individual sizes, temperature contrasts, and their overall distribution on the stellar surface is nearly an impossible task given only a light curve. From a theoretical perspective, finding a proper parametrization to mimic the effect of spots on a three dimensional volume within the framework of a one dimensional model provides its own complications. Currently, the only feasible method to include spots, is to include their potential effects on the radius measurement uncertainties. Unfortunately, while the inclusion of fixed 3\\% radius uncertainties in our analysis was able to alleviate many of the noted radius discrepancies, it also created a situation where the measurement uncertainty was on the order of the typical radius deviation. We are presented with a case where the observations are no longer effective at testing the models and the manifestation of most radius discrepancies can be attributed to under estimated error bars. At this point, we require observations that have been rigorously vetted for potential systematic uncertainties and are able to still provide mass and radius measurements to better than 2\\%. \\subsection{Radius-Rotation-Activity Correlations} Direct measurements of low-mass magnetic field strengths are rare, especially among fast rotating stars \\citep{Reiners2012}. Without a direct measure of the magnetic field strength, we are forced to rely on indirect measures to probe correlations between stellar magnetism and the appearance of inflated stellar radii. Ideally, these indirect measures are intimately connected with the dynamo mechanism, thought to generate and maintain stellar magnetic fields, or are the product of magnetic processes in the stellar atmosphere. The preferred indirect measures are typically stellar rotation or the observation of magnetically driven emission (H$\\alpha$, CaII H and K, X-ray). \\subsubsection{Rotation} \\label{sec:rot} Typically, low-mass DEBs are found in tight, short period orbits ($<$~3d; see Figure~\\ref{fig:porb}). Tidal interactions spin-up the individual components and allow them to remain rapidly rotating throughout their life cycle. Stellar dynamo theory dictates that the large-scale magnetic field strength is tied to the rotational properties of a star \\citep[i.e., a rapidly rotating star should be more magnetically active than a comparable star that is slowly rotating;][]{Parker1979,Reiners2012b}, providing a natural starting point for our investigation. We performed two independent statistical tests on the distribution of radius residuals as a function of the orbital period ($P_{\\textrm{orb}}$). Our primary objective was to determine whether rapidly rotating systems produce, on average, larger radius deviations than systems perceived to be slow rotators. Rotational periods were assumed to be synchronous with the orbital period unless a separate value for the rotational period was cited in the literature. The statistical tests performed were a Kirmogorov-Smirnov (K-S) test and another whereby we tested the probability of obtaining a given distribution of residuals via a Monte Carlo method. Both tests were performed on the observed difference in mean absolute error (MAE)\\footnote{We selected the MAE over the RMSD as a measure of the mean radius deviation of a given ensemble in order to reduce the weight of any individual outlier in the final mean.} between two data bins. The data bins were divided at preselected values of $P_{\\textrm{orb}}$, identified visually as vertical dashed lines in Figure~\\ref{fig:porb}. Comparing the radius deviations with the rotational periods, we find no evidence of any dominant correlation. Figure~\\ref{fig:porb} displays the residual data as a function of the orbital period, with frame A showing the full range of observed periods and frame B highlighting the ``short period'' regime. The correlation of radius deviations with orbital period has been studied previously \\citep[e.g.,][]{Kraus2011} where a significant difference between the two bins was observed around $P_{\\textrm{orb}}$ = 1.5 days. We performed the statistical tests using three values for the orbital period (1.0d, 1.5d, and 2.0d) that defined the two period bins. \\begin{figure} \\plotone{f3.eps} \\caption{The theoretical Rossby number, Ro = $P_{\\textrm{rot}}/\\tau_{\\textrm{conv}}$, versus the relative radius error. Ro is tied directly to the theoretical stellar dynamo mechanism and is empirically related to the ratio of a star's x-ray to bolometric luminosity (a magnetic activity indicator). Asterisks in maroon are stars with known x-ray flux measurements.} \\label{fig:rossby} \\end{figure} \\begin{figure*}[!ht] \\epsscale{0.95} \\plotone{f4.eps} \\caption{Relative error between observational and model radii for stars with detected x-ray emission. X-ray data is drawn from the \\emph{ROSAT} All-Sky Survey and combined with the radius residuals derived from this study. Data are shown as maroon filled circles. Illustrated are two least-square regressions performed on the data. The light-blue, dashed line demonstrates a non-negligible slope of $\\sim25\\pm17$ across all of the data, while the indigo, dash-dotted line excludes the two most discrepant points, UV Psc B and YY Gem.} \\label{fig:xray} \\end{figure*} We confirm the results of \\citet{Kraus2011} and find a $3.1\\sigma$ difference in the distribution of radius deviations around 1.5d. Systems with $P_{\\textrm{orb}} < 1.5$d had a MAE of 3.4\\% while the longer period systems had a MAE of 1.0\\%. While this is tantalizing, we can not necessarily attribute any physical significance to this particular division. We should expect this difference to be present for any two subsamples. However, we fail to find any evidence for a statistically meaningful difference when we divide the subsamples at 1.0d and 2.0d. Therefore, the significant difference noted at 1.5d is likely a spurious statistical result\\footnote{\\citet{Kraus2011} posit the difference may actually be a by-product of the DEB light curve analysis methods. Providing a further examination is outside the scope of this study.}. Inclusion of more long-period DEBs will be instrumental in providing a robust conclusion. Until those systems are discovered, there does not appear to be a physically meaningful explanation for why the divide should be made at 1.5d, but then also not hold for a division at 1.0d. The rotational (or orbital) period is not necessarily the most appropriate proxy for the magnetic field strength or potential magnetic activity level that we could select. It would be ideal for the rotational parameter to have some connection to intrinsic stellar properties. For instance, rotational velocity normalizes the rotational period to the stellar radius, providing a distinction between two main sequence stars of different masses that may have similar rotation periods. Optimally, the rotational variable in question would also provide a direct link to either observable magnetic activity or a theoretical description of stellar magnetism. Accordingly, we advocate the use of the Rossby number (hereafter Ro). Ro is defined as the ratio of the rotational period of the star to its convective overturn time, Ro~=~$P_{\\textrm{rot}}/\\tau_{\\textrm{conv}}$, and measures the strength of the Coriolis force acting on the vertical motion of convection cells. The dimensionless quantity Ro appears directly in standard mean-field dynamo theory \\citep[$\\alpha$-$\\omega$ dynamo;][]{Parker1979}\\footnote{For fully convective stars, an $\\alpha$-$\\omega$ dynamo cannot operate due to the lack of a tachocline. Instead, it is thought that an $\\alpha^2$ dynamo can efficiently generate a magnetic field. Since it is the $\\alpha$ mechanism that is related to the strength of the Coriolis force, the Rossby number should be just as applicable to fully convective stars \\citep{Chabrier2006,Browning2008}.} and is intimately related to the ratio of the stellar x-ray luminosity to bolometric luminosity \\citep{wright2011}. The latter quantity has been shown to be a strong indicator of a stellar corona heated to over $10^6$~K by magnetic activity \\citep{Vaiana1981,Pallavicini1981,Noyes1984}. One cut in Ro was performed at Ro~=~0.1, illustrtated in Figure~\\ref{fig:rossby}. The selection of Ro~=~0.1 approximately corresponds to Ro$_{\\textrm{sat}}$, or the value of Ro associated with an observed saturation of the stellar dynamo apparent in the ratio of the stellar x-ray luminosity to the bolometric luminosity \\citep[i.e., coronal saturation;][]{wright2011}. Intuitively, this suggests that all points with Ro values less than 0.1 are, presumably, sufficiently active so as to display inflated radii. Given our current understanding of coronal saturation, it is difficult to ascertain how strong of correlation is expected to exist. With that said, we assume that stars with very low Ro values should show at least a marginal degree of inflation compared to stars with higher Ro values, thereby indicating we should observe at least some evidence of a correlation. There appears to be no emergent correlation between Ro and radius deviations, at least for the present sample of data. We find no significant difference in the distribution of data points observed as a function of Ro. However, inspection of Figure~\\ref{fig:rossby} highlights the need for more data in order to draw a definitive conclusion. Selection of Ro~=~0.1 has the unfortunate effect of creating a bin with a small sample population, potentially affecting the statistics. Field stars in wide binary systems are a good starting point for studies interested in populating the high Ro region of Figure~\\ref{fig:rossby}. Caution must further be taken as secondary stars in the longer period ($>$~15d) systems in our sample do not have independently measured rotation periods, meaning they could potentially have different Ro values than are presented here. Visual inspection of Figure~\\ref{fig:rossby} leads us to the same conclusion, if we ignore the two most discrepant points. Since we were comparing MAE values, the outliers do not significantly affect the results of the statistical analysis. Arguably, the MAE is not an effective measure of the degree of \\emph{inflation} of each sample as it treats deflated radii the same as inflated radii. Instead, the actual direct mean may provide a more compelling statistical measure. Thus, we ran our statistical analysis on the direct mean error. Overall, the typical degree of inflation among the radii of low-mass stars was found to be about 1.6\\%. No statistically significant correlations with either $P_{\\textrm{orb}}$ or Ro were uncovered. The most significant result was found for the period cut at 1.5d, where the K-S test indicated a significant difference. However, the MC method produced a difference in the mean error of the two populations at 2.2$\\sigma$, below our significance threshold of 3.0$\\sigma$. All other bin divisions yielded results significant at only about the 1$\\sigma$ level. Curiously, if we consider only data points with measured x-ray fluxes (see below), there is evidence that systems with lower Ro values may have larger radius discrepancies. However, we are then prompted to explain why the trend does not continue to higher Ro values, a question we are currently not equipped to answer. There is still some ambiguity with the presence of the points near Ro = 0.01, which appear to contradict the presence of any definite correlation. Deeper x-ray observations of x-ray faint low-mass DEBs will clarify this ambiguity. \\subsubsection{X-ray Activity} \\label{sec:xray} An interesting extension of our discussion in the previous subsection, is to compare our derived radius deviations with the observed x-ray to bolometric luminosity ratio (hereafter $R_{x} = L_{x}/L_{bol}$). Since no correlation was noted with Ro, we expect that no correlation will be observed with $R_{x}$, as it has been shown to be intimately connected with Ro \\citep{wright2011}. \\citet{Lopezm2007} previously performed such a comparison and found a clear correlation between $R_{x}$ and radius deviations. Her comparison was performed under different modeling assumptions, which may have influenced the results. Specifically, she compared all observations to a 1~Gyr, solar metallicity isochrone from \\citet{BCAH98}. Without variation in age and metallicity, the observed radius discrepancies may have been over estimated (or under estimated in the case of YY Gem). A subsample of sixteen DEBs from this study have identified x-ray counterparts in the \\emph{ROSAT} All-Sky Survey Bright and Faint Source Catalogues \\citep{Voges1999,Voges2000}. Our analysis follows that of \\citet{Lopezm2007} whose previous analysis contained a fraction of our current x-ray detected sample. X-ray count rates were converted to x-ray fluxes according to the formula derived by \\citet{Schmitt1995}, \\begin{equation} F_x = \\left( 5.30 \\textrm{HR} + 8.31\\right)\\times 10^{-12} X_{\\textrm{cr}} \\end{equation} where HR is the x-ray hardness ratio\\footnote{There are typically two hardness ratios listed in the \\emph{ROSAT} catalog, HR1 and HR2. The \\citet{Schmitt1995} formula requires the use of HR1.}, $X_{\\textrm{cr}}$ is the x-ray count rate, and $F_x$ is the x-ray flux. Luminosities in the x-ray spectrum were computed with distances determined from either \\emph{Hipparcos} parallaxes \\citep[preferred when available;][]{hip07} or near-infrared photometry. Photometric distances were estimated using the luminosity calculated from the Stefan-Boltzmann relation assuming the observationally determined radius and effective temperature. Absolute magnitudes were then derived using the PHOENIX AMES-COND model atmospheres, adopting the best fit isochrone metallicities. In an effort to reduce errors introduced by the theoretical atmospheres, distances were computed using the average distance modulus derived from \\emph{2MASS} J and K band photometry \\citep{2mass}. We derived $R_x$ values for all sixteen stars in our x-ray sample to ensure internal consistency. Slight discrepancies between values presented here and those of \\citet{Lopezm2007} are due entirely to differences in the adopted distances. Attributing the x-ray flux contribution from each star in an DEB system is a difficult task. As such, L\\'{o}pez-Morales performed her analysis using three reported empirical scaling relations \\citep{Pallavicini1981,Fleming1989}: \\begin{itemize} \\item Case 1 -- each component contributes equal weight. \\item Case 2 -- proportional to the respective rotational velocity, $v\\sin i$, of each component. \\item Case 3 -- proportional to the square of the rotational velocity, $(v\\sin i)^2$, of each component. % \\end{itemize} We present results from all three cases in Figure~\\ref{fig:xray}. Immediately we notice that the size of the stellar radius deviations appears to correlate with $R_x$. A linear least-square regression performed on each data set (light-blue dashed line in Figure~\\ref{fig:xray}) suggests the slope for each case (1 -- 3) is 26, 24, and 22 ($\\pm$ 17), respectively, all with reduced-$\\chi^2$ values of $\\sim$8. Pearson $r$ coefficients were 0.72, 0.69, and 0.63, respectively. The statistics suggest that the likelihood of uncorrelated sets of data producing these particular correlations are 0.2\\%, 0.3\\%, and 0.9\\%, for case 1, 2, and 3, respectively. We can therefore rule out the null hypothesis with greater than 99\\% confidence. Visually, however, we notice that the correlation is largely driven by the presence of YY~Gem and UV~Psc~B, located in the upper-right region of each panel in Figure~\\ref{fig:xray}. If we were to remove those three points (YY~Gem appears as a single point), the correlation vanishes and we only observe an offset from the zero point (Indigo dash-dotted line in Figure~\\ref{fig:xray}). Furthermore, the distance to YY~Gem is highly uncertain. Assuming it is associated with Castor AB provides a distance of about 15~pc (see Appendix~\\ref{sec:appendix}), but our photometric analysis, as described above, places YY~Gem at a distance of approximately 11~pc, 4~pc closer to the Sun than the Castor AB system. Instead of selecting a single distance estimate, we averaged the two estimates and adopted $d=13\\pm 2$~pc, which also happens to be the distance assumed by \\citet{Delfosse2003}. The fact that our statistical correlation critically hinges upon three points, two of which are strongly distance dependent, is a cause for concern and implies that the statistics should be interpreted with care. If we believe the strong statistical correlation, then we are presented with a situation where the rotational data and the x-ray data disagree. This may be due to the physics underlying the stellar dynamo or those underlying x-ray saturation. However, there are two further interpretations that are contingent upon the role systematics play. First, we have that the correlation is entirely real and the presence of non-inflated stars with $R_x$~$\\sim$~0.0007 are outliers in the relation. Second, systematics play a large role, as proposed in Section~\\ref{sec:result}. Here, the truly deviant stars exhibit very strong x-ray emission ($R_x > 0.001$), while non-inflated stars show lower, varying levels of x-ray emission. For this view to hold, the relation between the level of radius inflation and magnetic field strength can not be linear. Strong magnetic fields would induce significant radius inflation while moderate and weak fields would produce little or no inflation. Accepting, on the other hand, that the statistical correlation is spurious, we are left with a picture that is coherent with our rotation analysis. Namely, magnetic activity may not be the leading cause for all of the observed inflated radii. Here, systematics may still play a role in producing stars that appear inflated, but that are consistent with the models, leaving a few discrepant stars. Naturally inefficient convection, potentially dependent on particular stellar properties, may be operating. However, magnetic activity cannot be fully ruled out, as we do not yet have a fully self-consistent description of the interaction of magnetic fields with the stellar interior and atmosphere. It may be that magnetic fields acting within YY~Gem and UV~Psc have a more noticeable affect on stellar structure, as higher mass stars are affected more by changes to the properties of convection (see Appendix~\\ref{sec:vmixl}). We tend to favor a hybrid interpretation. Here, most of the observed ``inflation'' is an artifact of unaccounted-for systematics, but significantly discrepant stars (YY~Gem, UV~Psc) are associated with very strong magnetic activity. We believe that CM~Dra probably fits into the latter category due to a push in the literature towards subsolar metallicity. Effects of a large-scale magnetic field are presumably mass dependent, with higher mass stars showing a greater propensity to become inflated owing to their sensitivity to changes in convective properties. Whether there is a characteristic magnetic field strength that induces substantial radius deviations is unclear. Dynamo saturation and the saturation of magnetic activity, as evidenced by the flattening of the Ro-$R_x$ relation in \\citet{wright2011}, is not yet fully understood, but may provide further insight into the apparent disagreement between our rotation and x-ray analyses. Finally, clarity will be obtained with a better distance measurement to YY~Gem, either from ground-based parallax programs or with eventual results from \\emph{Gaia}. Of all the points in Figure~\\ref{fig:xray}, the existence of a positive correlation is most dependent upon the distance to YY~Gem. An accurate distance, as well as a reanalysis of its association with the Castor system, has the ability to not only relieve the ambiguity present in the x-ray data, but also to provide insight into the necessary constraints for the system (i.e., is YY~Gem really about 400 Myr old?). As further low-mass DEB systems are discovered, x-ray observations are strongly encouraged in order to develop a coherent picture of how radius deviations correlate with magnetic activity. This study focused on reevaluating the current state of agreement between the theoretical and observational low-mass, main sequence MR relationship. The DEB systems used in the analysis were required to have quoted random uncertainties in the mass and radius below 3\\% in order to provide an effective test of stellar evolution models. A large grid of DSEP models were computed with variation in age and metallicity characteristic of the local galactic neighborhood. Best fit isochrones were derived by allowing the age and metallicity to be optimized while maintaining the constraint that the system be coeval with a single composition. DEBs with reliably determined ages or metallicities were compared with a restricted set of isochrones in the range allowed by the observational priors. Overall, we find 92\\% of stars in our sample are less than 4\\% discrepant with the models, largely representing a factor of two improvement over the canonical 5 -- 15\\% deviations. Our results suggest that low-mass stars with radii that deviate significantly from model predictions are exceptions to general agreement. Discrepancies may also be the result of unaccounted for systematic uncertainties (i.e., starspots) that may as large as 4\\%. With uncertainties as large as the typical radius deviations, we find it difficult to draw the firm conclusion suggesting that models are in broad disagreement with observations. Instead, we are left with a situation where the observational uncertainties may be too large to provide an adequate test of stellar models. The combination of random and systematic uncertainties for the sample of low-mass DEBs must be constrained and minimized below the 2\\% level before accurate model comparisons may be made. Radius correlations with orbital (rotational) period, Ro, and $R_{x}$ were also considered. No distinct trends were identified with either orbital period or Ro. However, we find evidence for a strong correlation between radius deviations and $R_{x}$ (previously noted by \\citet{Lopezm2007}) in contradiction with our Ro analysis. The trend is not as tight as that derived by \\citet{Lopezm2007}, owing to the age and metallicity variations allowed by our analysis. This correlation is also largely contingent upon the veracity of the distance estimate to YY~Gem. Accurately determining the distance to YY~Gem and evaluating its association with the Castor~AB quadruple would alleviate much of the uncertainty. Finally, we must not leave the theoreticians out of the spotlight. The degree to which a magnetic field can alter the interior structure of low-mass stars is still only partially known and further investigations are required. Development of models with self-consistent magnetic field perturbations will begin to shed light on this unknown. Comparisons between predicted magnetic field strengths from self-consistent models and observational data (either direct or indirect) will provide a measure of the validity of the ability of magnetic fields to inflate stellar radii. Whatever the final solution may be to this long-standing problem, it is now apparent that the level of inflation required by theory is not a ubiquitous 5~--~15\\%, but only so in extreme cases." }, "1207/1207.1606_arXiv.txt": { "abstract": "{The broad 30 \\mic feature in carbon stars is commonly attributed to MgS dust particles. However, reproducing the 30 \\mic feature with homogeneous MgS grains would require much more sulfur relative to the solar abundance. Direct gas-phase condensation of MgS occurs at a low efficiency. Precipitation of MgS on SiC precursor grains provides a more efficient formation mechanism, such that the assumption of homogeneous MgS grains may not be correct. Using a Monte Carlo-based radiative transfer code, we aim to model the 30 \\mic feature of the extreme carbon star LL Peg with MgS dust particles. We find that for LL Peg this modeling is insensitive to the unknown MgS optical properties at $\\lambda <10 $ \\mic. When MgS is allowed to be in thermal contact with amorphous carbon and SiC, the amount of MgS required to reproduce the strength of 30 \\mic feature agrees with the solar abundance of sulfur, thereby resolving the reported \\emph{MgS mass problem}. We conclude that MgS is a valid candidate to be the carrier of the 30 \\mic feature when it is part of a composite grain population that has optical properties representative of an ensemble of particle shapes.} ", "introduction": "\\label{sect:intro} \\vspace{-0.1cm} The broad 30 \\mic feature in the thermal continuum emission of carbon stars was identified in 1985 by \\citet{goe1985} as due to magnesium sulfide (MgS). \\citet{beg1994} presented optical data for MgS and compared them to the thermal continuum emission of CW Leo to confirm MgS as the likely carrier of the 30 \\mic feature. The \\emph{Infrared Space Observatory} (ISO; \\citeauthor{kes1996}~\\citeyear{kes1996}) observed a large sample of carbon stars (e.g.~\\citeauthor{yam1998}~\\citeyear{yam1998}; \\citeauthor{jia1999}~\\citeyear{jia1999}; \\citeauthor{szc1999}~\\citeyear{szc1999}; \\citeauthor{hri2000}~\\citeyear{hri2000}; \\citeauthor{vol2002}~\\citeyear{vol2002}) displaying a diversity in strength, shape and width of the feature. In an analysis of this data set, \\citet{hon2002} concluded that the shape of the 30 \\mic feature can be best reproduced when the grains are not perfect spheres. \\citet{zha2009a} modeled the 30 \\mic feature in the proto-planetary nebula HD 56126 using pure MgS grains. Assuming the grains to be irradiated by unattenuated stellar light, their analysis needed the optical properties of MgS at wavelengths $\\lambda < 10$ \\mic. Unfortunately, such measurements are lacking, as yet. Assuming relatively high absorption efficiencies in this regime, these authors required an amount of MgS up to ten times the amount of available atomic sulfur to explain the strength of the 30 \\mic feature. Owing to this \\emph{mass problem}, Zhang and collaborators argued against MgS as the carrier of the feature. In this letter, we report on the observational evidence for composite grains in the outflow of the high mass-loss rate asymptotic giant branch (AGB) star LL Peg, also known as AFGL 3068. We show that MgS is a viable candidate to explain the 30 \\mic feature independent of the absorbing efficiencies of these grains at $\\lambda < 10$ \\mic. Moreover, we show that if one assumes thermal contact between the dust species in the outflow, the \\emph{mass problem} in HD 56126 reported by \\citet{zha2009a} does not occur in LL Peg. Finally, we discuss the formation of composite grains in AGB outflows. ", "conclusions": "\\label{sect:conclusion} We have modeled the ISO spectrum of the high-density carbon star LL Peg with a dust composition consisting of a-C, Fe, SiC, and MgS. The (high) density and temperature structure in the envelope of this source allow one to model the 30 \\mic feature with MgS dust particles, independent of the unknown MgS optical properties at $\\lambda < 10$ \\mic. We showed that MgS is a viable candidate to be the carrier of the 30 \\mic feature. An ensemble of particle shapes works significantly better than spherical grains to explain the shape of the feature. Thermal contact between the dust species is required to ensure that the amount of MgS dust in the envelope of LL Peg does not exceed the solar abundance of sulfur, thereby avoiding the \\emph{mass problem} as reported by \\citet{zha2009a} for the post-AGB star HD 56126. Achieving thermal contact between all dust species is possible if these species form in some kind of heterogeneous composite grain structure." }, "1207/1207.6210_arXiv.txt": { "abstract": "Since there are several ways planets can survive the giant phase of the host star, we examine the habitability and detection of planets orbiting white dwarfs. As a white dwarf cools from 6000\\,K to 4000\\,K, a planet orbiting at 0.01\\,AU would remain in the Continuous Habitable Zone (CHZ) for $\\sim$8\\,Gyr. We show that photosynthetic processes can be sustained on such planets. The DNA-weighted UV radiation dose for an Earth-like planet in the CHZ is less than the maxima encountered on Earth, hence non-magnetic white dwarfs are compatible with the persistence of complex life. Polarisation due to a terrestrial planet in the CHZ of a cool white dwarf is 10$^2$ (10$^4$) times larger than it would be in the habitable zone of a typical M-dwarf (Sun-like star). Polarimetry is thus a viable way to detect close-in rocky planets around white dwarfs. Multi-band polarimetry would also allow reveal the presence of a planet atmosphere, providing a first characterisation. Planets in the CHZ of a 0.6\\,\\M\\ white dwarf will be distorted by Roche geometry, and a Kepler-11d analogue would overfill its Roche lobe. With current facilities a Super-Earth-sized atmosphereless planet is detectable with polarimetry around the brightest known cool white dwarf. Planned future facilities render smaller planets detectable, in particular by increasing the instrumental sensitivity in the blue. ", "introduction": "\\label{sec:intro} The search for habitable Earth-like planets is a major contemporary goal of astronomy. As the detection of exoplanets is biased towards systems with small differences in mass, radius and luminosity between star and planet \\citep[e.g.,][]{Haswell10}, M-type main sequence stars have become prime targets in the search of Earth-like planets in the habitable zone. M dwarfs evolve slowly: their planets might remain within a continuously habitable zone (CHZ)\\footnote{Range of planet orbital distances at which the planet is habitable for a minimum of 3\\,Gyr.}, i.e. harbouring surface liquid water, for several Gyr, providing ample time for the advent of life on a rocky planet. With an effective temperature (\\Teff) $\\leq$6000\\,K, cool white dwarfs (CWD) are also promising hosts of rocky planets in the habitable zone. White dwarfs initially cool down rapidly, with temperature decreasing by thousands of degrees in $\\sim$3\\,Gyr \\citep{salaris10}. At \\Teff$\\sim$6000\\,K, crystallisation slows the cooling process. This produces a habitable zone which endures for up to 8\\,Gyr \\citep{agol}, well in excess of the time required for life to arise on Earth. White dwarfs provide a stable luminosity source without the potentially damaging radiation produced by stellar activity in M dwarfs. A planet orbiting close to a white dwarf would synchronize within 1000\\,yr and would have a stable orbit, as the planet would not raise tides on the star \\citep{agol}. The major issue for the presence of a planet close to a white dwarf is survival during the host star's giant phase. \\citet{faedi11} review several mechanisms which would result in a planet orbiting a white dwarf. \\citet{charpinet11} found two Earth-sized bodies in a very close orbit ($\\sim$0.007\\,AU) around a post-red-giant star proving that planet-sized objects can survive the post-main sequence evolution phases of their host star. Further evidence for the existence of rocky bodies close to white dwarfs comes from the presence of metallic lines (e.g. Mg and Fe) in the spectra of DZ white dwarfs \\citep{zuckerman}. Heavy metals in the atmospheres of these stars can only be explained by atmospheric ``pollution\" caused by the accretion of terrestrial-like planets or planetesimals \\citep[see e.g.,][]{farihi10,melis11,klein11,gensicke}. The low luminosity of CWDs creates a habitable zone at only $\\sim$0.01\\,AU, ten times closer than for M dwarfs. This facilitates the detection of small bodies orbiting white dwarfs: \\citet{agol} showed that transits of Mars-sized planets in the white dwarf CHZ would be 1\\% deep, easily detectable with present-day ground-based facilities, even for rather faint stars, though searches for planetary companions to white dwarfs have so far been unsuccessful \\citep{friedrich,faedi11}. Polarimetric techniques can also be used to discover and characterize exoplanets. As shown by \\citet{seager00} and \\citet{stam08}, as a planet orbits around the host star, the amount of polarisation varies regularly, showing maxima when the planet is near quadrature. The amplitude of the variation depends mainly on the orbital inclination, $i$, with no polarimetric variability for a face-on orbit. The detection and measurement of regular polarisation variations permit discovery of exoplanets, with an efficiency dependent on $i$, similar to that of the radial velocity planet detections. Spectropolarimetry of planet-hosting stars could characterise the atmosphere of an exoplanet \\citep{stam08}, something now only possible for transiting exoplanets. ", "conclusions": "" }, "1207/1207.4215_arXiv.txt": { "abstract": "A dense ionized cloud of gas has been recently discovered to be moving directly toward the supermassive black hole, \\sgr, at the Galactic Center. In June 2013, at the pericenter of its highly eccentric orbit, the cloud will be approximately 3100 Schwarzschild radii from the black hole and will move supersonically through the ambient hot gas with a velocity of $v_p \\approx 5400$~km~s$^{-1}$. A bow shock is likely to form in front of the cloud and could accelerate electrons to relativistic energies. We estimate via particle-in-cell simulations the energy distribution of the accelerated electrons and show that the non-thermal synchrotron emission from these electrons might exceed the quiescent radio emission from \\sgr\\ by a factor of several. The enhanced radio emission should be detectable at GHz and higher frequencies around the time of pericentric passage and in the following months. The bow shock emission is expected to be displaced from the quiescent radio emission of \\sgr\\ by $\\sim 33$~mas. Interferometric observations could resolve potential changes in the radio image of \\sgr\\ at wavelengths $\\lesssim 6$~cm. ", "introduction": "Recent sub-millimeter observations revealed a dense ionized cloud of gas known as G2, rapidly approaching \\sgr, the black hole at the Galactic Center (Gillessen et al.\\ 2012). The cloud is on a highly eccentric trajectory, with a 2011 distance from the black hole of $1.8\\times10^{16}$\\,cm. The pericentric passage, which is expected to occur in mid 2013, will bring the cloud within $4\\times10^{15}$\\,cm from the supermassive black hole. Given the mass of the black hole, which is determined through observations of nearby stellar orbits to be $M=4.3\\times10^6\\;\\msun$ (Ghez et al.\\ 2008; Gillessen et al.\\ 2009), this pericentric distance is only $R_p=3100 \\,R_S$, where the Schwarzschild radius $R_S=1.27\\times10^{12}$\\,cm. The accretion flow around the black hole extends to the Bondi radius $\\sim10^5R_S$ (e.g., Yuan, Quataert \\& Narayan 2003) and powers the multiwavelength emission observed from it. The flux at 1~GHz is $\\simeq 0.5$~Jy, rising to $\\approx 4$~Jy at 500~GHz, before rapidly declining at higher frequencies. The radio emission has been successfully modeled as synchrotron radiation from relativistic electrons, either in a radiatively inefficient accretion flow (ADAF) (Narayan, Yi, \\& Mahadevan 1995; \\\"Ozel, Psaltis \\& Narayan 2000) or in a jet (Falcke \\& Markoff 2000). At the lowest end of the spectrum --- $\\nu \\sim 1-10$\\,GHz --- the radio emission shows flux variability on a time scale of months to years with a root mean square amplitude $\\sim10\\%$, (Zhao et al. 1989; Falcke 1999; Macquart \\& Bower 2006). At its pericentric passage, the gas cloud will interact with the accretion flow, and this may significantly change its dynamics. Here, we show that a bow shock is likely to develop as the cloud plows through the hot, tenuous plasma at $R_p$. In \\S2 we calculate the properties of the bow shock and in \\S3 we estimate the energy distribution of electrons accelerated at this shock. In \\S4 we calculate the extra radio emission that will result from these accelerated electrons. For likely electron energy distributions, the additional emission is $\\sim10$\\,Jy at frequencies $\\sim1-10$~GHz. This is well above the quiescent emission from \\sgr\\ and should be easily detectable. We also show that the flux increase will be accompanied by significant changes in the spectral index. In \\S5 we summarize our findings and argue that interferometric observations could resolve potential changes in the radio image of \\sgr\\ caused by the interaction of the cloud with the accretion flow. ", "conclusions": "The passage of the recently discovered cloud of gas G2 near \\sgr\\ presents a unique opportunity to study the dynamics and properties of hot gas in the vicinity of the black hole at the Galactic Center. In this Letter, we showed that a bow shock may form during the pericentric passage of the cloud, and we investigated the flux enhancement in the $1-100$~GHz frequency range that will arise as a result of particle acceleration in the shock front. We ran first-principles PIC simulations for shock parameters appropriate to the bow shock and thereby obtained realistic estimates of the energy distribution of accelerated electrons. Using these results, we calculated the likely synchrotron emission from the bow shock and found that the additional flux might exceed the quiescent emission from \\sgr\\ by up to an order of magnitude at GHz frequencies. This suggests that there is a good chance of detecting enhanced radio emission as G2 plows through the ambient hot medium around the time of pericentric passage. Since the cooling time of the accelerated electrons is estimated to be long, the enhanced emission should continue well after the encounter. There are order unity uncertainties in the parameters we have assumed for the bow shock, and hence the predictions made in this Letter are not likely to be quantitatively accurate. We have allowed for some of these uncertainties while computing the hatched regions shown in the two panels in Figure~\\ref{fig:sgrA_radio}. An additional uncertainty is whether or not a bow shock will form in the first place. Gillessen et al. (2012) discuss a compression shock moving into the cloud, which will inevitably be accompanied by an external shock moving into the ambient medium, the bow shock in our model. In most models of G2 (Burkert et al. 2012; Miralda-Escude 2012; Schartmann et al. 2012; Murray-Clay \\& Loeb 2012), the cloud retains some level of integrity during its pericentric passage and thus is likely to develop an external shock. However, if the cloud is completely shredded by Kelvin-Helmholtz or other instabilities before a bow shock forms, then our synchrotron emission estimates will no longer be valid. Given the pericentric distance of 3100~$R_S$, which corresponds to a projected angle of $\\sim 33$~mas, the bow shock emission should be displaced from the quiescent radio emission of \\sgr\\ by the same amount. At wavelengths $\\lesssim 6$~cm, this angular distance is larger than the size of the scattering ellipse of \\sgr\\ (Bower et al.\\ 2006) and ought to be resolved by interferometric observations. The scatter-broadening of the radio image of \\sgr\\ is believed to be caused by a compact foreground interstellar cloud that is at least $\\sim100$\\,pc from the black hole (Frail et al. 1994). Thus, the broadening is unlikely to be affected by any gas stripped from G2 during its pericentric encounter. In addition to the bow shock and the associated prompt radio synchrotron emission considered in this Letter, it is expected that the cloud G2 will also shed mass as it interacts with the ambient hot gas. A likely early signature of the increase in the gas density at a few thousand $R_S$ is a change in the observed Faraday rotation above a GHz, which may provide the first estimates of the increase in the mass accretion rate. As it moves inwards, this gas will cause the mass accretion rate on to the central black hole to be enhanced over a period of many years. Such an increase could cause a secular change in the radio flux of \\sgr\\ on a time scale of ten years to several decades, accompanied by changes in the ``silhouette'' of the black hole that could be monitored by future interferometers (Moscibrodzka et al. 2012)." }, "1207/1207.3329_arXiv.txt": { "abstract": "The Advanced CCD Imaging Spectrometer is an instrument on the Chandra X-ray Observatory. CCDs are vulnerable to radiation damage, particularly by soft protons in the radiation belts and solar storms. The Chandra team has implemented procedures to protect ACIS during high-radiation events including autonomous protection triggered by an on-board radiation monitor. Elevated temperatures have reduced the effectiveness of the on-board monitor. The ACIS team has developed an algorithm which uses data from the CCDs themselves to detect periods of high radiation and a flight software patch to apply this algorithm is currently active on-board the instrument. In this paper, we explore the ACIS response to particle radiation through comparisons to a number of external measures of the radiation environment. We hope to better understand the efficiency of the algorithm as a function of the flux and spectrum of the particles and the time-profile of the radiation event. ", "introduction": "\\label{sec:intro} The Chandra X-ray Observatory, the third of NASA's great observatories in space, was launched just past midnight on July 23, 1999, aboard the space shuttle {\\it Columbia}\\cite{cha2}. After a series of orbital maneuvers, Chandra reached its operational orbit, with initial 10,000-km perigee altitude, 140,000-km apogee altitude, and 28.5$^\\circ$ inclination. In this evolving high elliptical orbit, Chandra transits a wide range of particle environments, from the radiation belts at closest approach through the magnetosphere and magnetopause and past the bow shock into the solar wind. The Advanced CCD Imaging Spectrometer (ACIS), one of two focal plane science instruments on Chandra, utilizes frame-transfer charge-coupled devices (CCDs) of two types, front- and back-illuminated (FI and BI)\\cite{acis}. Soon after launch it was discovered that the FI CCDs had suffered radiation damage from exposure to soft protons scattered off the Observatory's grazing-incidence optics during passages through the Earth's radiation belts\\cite{gyp00}. Since mid-September 1999, ACIS has been protected during radiation belt passages and there is an ongoing effort to prevent further damage and to develop hardware and software strategies to mitigate the effects of radiation damage on data analysis\\cite{odell}. Our primary measure of radiation damage on the CCDs is charge transfer inefficiency (CTI). The eight front-illuminated CCDs had essentially no CTI before launch, but are strongly sensitive to radiation damage from low energy protons ($\\sim$100~keV) which preferentially create traps in the buried transfer channel. The framestore covers are thick enough to stop this radiation, so the initial damage was limited to the imaging area of the FI CCDs. Radiation damage from low-energy protons is now minimized by moving the ACIS detector away from the aimpoint of the observatory during passages through the Earth's particle belts and during solar storms. Continuing exposure to both low and high energy particles over the lifetime of the mission slowly degrades the CTI further.\\cite{odell,ctitrend} The two back-illuminated CCDs (ACIS-S1,S3) suffered damage during the manufacturing process and exhibit CTI in both the imaging and framestore areas and the serial transfer array. However, owing to their much deeper charge-transfer channel, BI CCDs are insensitive to damage by the low-energy ions that damage FI devices. Since early in the Chandra mission, procedures have been implemented that protect the focal plane instruments during times of high radiation. \\cite{odell} ACIS is translated out of the focal plane, providing protection against soft protons, and is powered off. Three types of procedures are in place; planned protection during radiation-belt transits, autonomous protection triggered by the on-board radiation monitor, and manual intervention based upon assessment of space-weather conditions. The Chandra weekly command load includes automatic scheduled safing of the focal-plane instruments during radiation belt passages. The timing of radiation belt ingress and egress are determined using the standard AP8/AE8 environment with a small additional pad time to protect against temporal variations. Solar storms are detected either by the on-board radiation monitor or by ground operations monitoring of various space weather measures, such as from NASA's Advanced Composition Explorer (ACE)\\cite{ace}, the NOAA Geostationary Operational Environmental Satellite system (GOES), and the planetary index Kp. The on-board radiation monitor cannot detect protons at hundreds of keV, which are the most damaging to ACIS, so on-board protection is supplemented by other measures of the radiation environment. The Electron, Proton, Helium INstrument (EPHIN) is a particle detector on-board the Chandra spacecraft used to monitor the local particle radiation environment. It is sensitive to electrons in the energy range 150~keV--50~MeV and protons from 5--49~MeV. Chandra-EPHIN is very similar to the EPHIN detector onboard SOHO.\\cite{ephin} Until December 2008, EPHIN rates in three channels were monitored by the spacecraft computer which can command radiation shutdowns during solar storms. The monitored channels were P4, sensitive to protons with 5.0--8.3~MeV, P41, sensitive to protons with 41--53~MeV and E1300, sensitive to electrons with 2.64--6.18~MeV. As the spacecraft insulation has aged and degraded, thermal control of some subsystems, including EPHIN, has become more difficult. Elevated EPHIN temperatures cause anomalous noise, which in some EPHIN channels can be significant and occasionally dominate the signal.\\cite{odell} To prevent against false triggers due to EPHIN noise, EPHIN was reconfigured in December 2008 into a mode that does not distinguish particle species. At that point, the spacecraft computer began monitoring only two EPHIN channels, one of which was subsequently deleted due to increasing noise. The remaining monitored EPHIN channel, E1300$^\\prime$, is basically the union of the original E1300 and P41 channels; hence, it is sensitive to high-energy (41--53 MeV) protons. Concern about the effectiveness of EPHIN going into the future motivates looking for other on-board measures of the radiation environment. Consequently, the spacecraft computer now monitors the anti-coincidence shield rate of the other Chandra focal-plane instrument, the High Resolution Camera (HRC)\\cite{odell}. As the HRC shield responds to only penetrating charged particles (protons $>$ 30~MeV), it monitors approximately the same proton energies as does the E1300$^\\prime$ channel. A second measure, examined here, is using ACIS itself. We first addressed the potential use of ACIS as its own particle monitor in Ref.~\\citenum{radmon1} (Paper I). As the primary purpose of ACIS is always to collect astrophysical data, the particle monitoring processes need to be secondary in using the available on-board resources and need to be flexible in dealing with the numerous observing modes and X-ray source types. Ideally the scientific user should be completely unaware of the simultaneous radiation monitoring during his/her observation. The monitor needs to detect as many instances of enhanced radiation environment as possible, without unnecessarily interrupting observations due to bright X-ray sources or hardware anomalies. Paper I was an initial exploration into the efficacy of a particular proxy for the radiation environment---the threshold crossing rate. ACIS performs X-ray event finding and recognition tasks on-board in real time. Each CCD frame is examined to find pixels with pulseheights above a pre-determined threshold value. These pixels are considered candidates to be X-ray events and are subjected to further processing. The number of pixels above the event detection threshold in a CCD frame, or threshold crossing rate, is calculated and saved as part of the standard event processing on-board and is sensitive to both the X-ray and particle intensity. We have developed an ACIS flight software patch that keeps track of the mean threshold crossing rate (pixels/second/row) for both FI and BI CCDs and signals the Chandra On-Board Computer (OBC) when this rate exceeds a predetermined value and is increasing. Ref.~\\citenum{pgfradmon} from this conference describes in more detail the algorithm, the software patch, and determination of the optimal parameters for the patch. The flight software patch was installed on ACIS in November 2011 and updated with more optimal parameters in April 2012. The patch has operated as expected with no impact on regular science operations and has correctly indicated an enhanced radiation environment on two occasions. The Chandra OBC was patched in May 2012 and will now respond to any ACIS radiation triggers with the standard radiation protection procedures. To better match the changing quiescent particle background which is anti-correlated with the solar cycle, the parameters of the patch will be re-evaluated a few times a year and updated as necessary. In this paper, we explore the response of the ACIS radiation monitor to the Chandra particle environment. In particular, we examine the historical data to better understand the nature of the particle events that would and would not have triggered the ACIS radiation monitor. ", "conclusions": "" }, "1207/1207.3435_arXiv.txt": { "abstract": "{Observations of the atomic and molecular line emission associated with jets and outflows emitted by young stellar objects provide sensitive diagnostics of the excitation conditions, and can be used to trace the various evolutionary stages they pass through as they evolve to become main sequence stars.} {To understand the relevance of atomic and molecular cooling in shocks, and how accretion and ejection efficiency evolves with the evolutionary state of the sources, we will study the far-infrared counterparts of bright optical jets associated with Class I and II sources in Taurus (T Tau, DG Tau A, DG Tau B, FS Tau A+B, and RW Aur).} {We have analysed {\\it Herschel}/PACS observations of a number of atomic (\\oi63\\um, 145\\um, \\cii158\\um) and molecular (high-J CO, \\ho, OH) lines, collected within the Open Time Key project GASPS (PI: W.~R.~F. Dent). To constrain the origin of the detected lines we have compared the obtained FIR emission maps with the emission from optical-jets and millimetre-outflows, and the measured line fluxes and ratios with predictions from shock and disk models.} {All of the targets are associated with extended emission in the atomic lines; in particular, the strong \\oi~63~\\um\\, emission is correlated with the direction of the optical jet/mm-outflow. The line ratios suggest that the atomic lines can be excited in fast dissociative J-shocks occurring along the jet. The molecular emission, on the contrary, originates from a compact region, that is spatially and spectrally unresolved, and lines from highly excited levels are detected (e.g., the o-\\ho\\,8$_{18}$ - 7$_{07}$ line, and the CO J=36-35 line). Disk models are unable to explain the brightness of the observed lines (CO and \\ho\\, line fluxes up to 10$^{-15}$-6~10$^{-16}$ \\wm). Slow C- or J- shocks with high pre-shock densities reproduce the observed \\ho\\, and high-J CO lines; however, the disk and/or UV-heated outflow cavities may contribute to the observed emission.} {Similarly to Class 0 sources, the FIR emission associated with Class I and II jet-sources is likely to be shock-excited. While the cooling is dominated by CO and \\ho\\, lines in Class 0 sources, \\oi\\, becomes an important coolant as the source evolves and the environment is cleared. The cooling and mass loss rates estimated for Class II and I sources are one to four orders of magnitude lower than for Class 0 sources. This provides strong evidence to indicate that the outflow activity decreases as the source evolves. } ", "introduction": "Theoretical models \\citep[e.g., ][]{shu94,konigl00} predict a tight correlation between the accretion of matter onto a young star and the ejection in winds and/or jets. Measurements of stellar accretion and mass loss \\citep[e.g., ][]{hartigan95} support the general picture presented by these models, but the uncertainties of these measurements are too large to provide a quantitative test of the predictions of the ratio of the mass accretion rate to the mass loss rate. Sources in the earliest stages in their evolution (Class 0) are not visible, and are often indirectly identified by means of their strong ejection activity, which is manifested in the form of bipolar parsec-scale molecular outflows often observed at millimetre wavelengths \\citep[e.g., ][]{bachiller96}. The ejection associated with evolved, optically visible T Tauri stars (i.e., Class II) is instead usually traced by bright blue- and red- shifted forbidden emission lines present at optical and near-infrared (NIR) wavelengths \\citep[e.g., ][]{hartigan95}. For Class I sources and some Class II sources both the molecular outflow and the optical jet have been observed \\citep{gueth99,pety06}. These observations show that the two components are connected. The optical/NIR forbidden lines trace hot ($\\sim$10$^4$ K) atomic gas, which is believed to have been extracted from the disk, and accelerated in the observed fast and collimated {\\it jets} (velocities up to hundreds of \\kms\\, and jet widths smaller than 200 AU). The millimetre observations, instead, trace cold ($\\sim$10-100 K) and slow (tens of \\kms) gas, which is thought to be ambient gas that has been set into motion by the jet propagation (i.e., {\\it jet-driven molecular outflow}, e.g., \\citealt{raga93,cabrit97}). However, collimated high velocity molecular gas (velocity up to $\\sim$60 \\kms, \\citealt{lefloch07}) has also previously been detected at millimetre wavelengths, questioning this simple picture, and suggesting that molecules can also be extracted from the disk and accelerated in the jet \\citep{pontoppidan11,panoglou12}.\\\\ In this context, observations at far-infrared wavelengths allow us to trace the intermediate {\\it warm gas component} in the jet/outflow system. Previous observations from the Infrared Space Observatory (ISO, \\citealt{kessler96}) targeting outflow sources have shown emission in a large number of atomic (\\oi, \\cii) and molecular (\\ho, CO, OH) lines. Despite the very low spectral and spatial resolution offered by ISO, analysis of the line fluxes and ratios indicates that the bulk of the detected \\oi\\, and molecular emission is most likely to be excited in the shocks occurring along the jet/outflow \\citep{nisini96,nisini99,giannini01}. The line fluxes have previously been used to estimate the cooling in the atomic and molecular lines, and to quantify the outflow efficiency as the ratio between the total luminosity radiated away in the far-infrared lines, L~(FIR), and the source bolometric luminosity, L$_{bol}$ \\citep{giannini01, nisini02}. However, because of its limited sensitivity, ISO observations have been restricted to studies of bright and extended outflows from Class 0 and I objects. The ESA {\\it Herschel} Space Observatory ({\\it Herschel}) has allowed, for the first time, observations of the FIR counterparts of optical jets associated with Class I and Class II sources whose environment has been largely cleared. As the source evolves, the accretion/ejection activity is expected to decrease, with the surrounding cloud material being either accreted onto the star, or dispersed by the jet. As a consequence, the optical jet will become visible while the emission at far-infrared wavelengths should be expected to be fainter and less extended than in Class 0 sources. However, FIR observations of the ejection activity associated with more evolved Class I and II sources is interesting for the following reasons: \\begin{itemize} \\item[-] The source, disk, and accretion properties of T Tauri stars are well-known, and so we can study the correlation between the detected ejection phenomena and these properties. Thus, we can estimate the mass ejection to mass accretion ratio; \\item[-] It is usually the case that Class 0 sources are observable only at millimetre wavelengths, whilst the ejection activity from T Tauri stars is detected only in the optical. FIR lines can however be detected from Class 0 to Class II sources, therefore facilitating a way to form an evolutionary picture of jet activity; \\item[-] The detection of molecular emission in sources whose environment has been cleared may support the idea of a disk-wind molecular component providing strong constraints to existing models of jet launching \\citep{panoglou12}. \\end{itemize} Thus we have analysed the FIR emission from five Class I and II sources in Taurus (d$\\sim$140 pc). These sources were observed as part of the {\\it Herschel} Open Time Key project GASPS ({\\it GAS in Protoplanetary Systems}, PI: W. R. F. Dent) using the PACS integral field spectrometer \\citep{poglitsch10} on board {\\it Herschel} \\citep{pilbratt10}. These PACS observations provide maps of the emission in a number of atomic and molecular lines (see, \\citealt{mathews10} and Dent et al., {\\it in preparation}). The five sources presented in this paper form a subset of the Taurus sample analysed in Howard et al., {\\it in preparation}. These sources were selected on the basis of their association with bright and extended stellar jets detected in the typical \\sii, \\azii, and \\oi\\, optical forbidden lines \\citep[e.g., ][]{hartigan95,hirth97}. {The details of the observations, and the applied data reduction processing, are described in Sect.~\\ref{sect:obs}. In Sect.~\\ref{sect:results} we show the obtained spectra and maps and compare the spatial distribution of the atomic and molecular emission with that of the associated optical jets. In Sect.~\\ref{sect:discussion} we will compare the observed line fluxes and ratios with predictions from both disk and shock models. % Hence, we will use the observed line fluxes and the results from shock modelling to estimate the far-infrared cooling and the mass loss rate. The comparison with the values estimated for Class 0 and I sources will allow us to place Class II sources into an evolutionary picture. Finally, in Sect.~\\ref{sect:conclusions} we summarise our conclusions. ", "conclusions": "\\label{sect:conclusions} In this paper we have analysed {\\it Herschel}/PACS integral-field spectroscopic observations of Class I and II sources in Taurus which are known to drive bright optical jets. Thanks to the {\\it Herschel} sensitivity (100-1000 times larger than ISO) we are able to detect the FIR counterpart of optical jets from the selected Class I and II sources for the first time. An exception is T Tau which is a bright multiple system unresolved with PACS, consisting of a Class II source and a Class I binary system, and associated with at least two jets, which has been observed with ISO \\citep{spinoglio00}. We investigate the origin of the detected atomic and molecular lines by carefully evaluating the spatial distribution of the emission on the PACS detector and by comparing line fluxes and ratios with predictions from disk and shock models. The results of our analysis are summarised below: \\begin{itemize} \\item[-] {\\it the emission in the atomic \\oi\\, and \\cii\\, lines is extended and spatially correlated with the optical jet emission}. In two cases (DG Tau B and RW Aur) we also detect a consistent offset in velocity in all the spaxels where the \\oi~63~\\um\\, line is detected which indicates a gas velocity in agreement with the values measured for the associated optical jets. \\\\ % \\item[-] {\\it the emission in the molecular \\ho, CO, and OH lines is spatially and spectrally unresolved}. % However, by using the DENT grid of models we show that for typical low mass YSO and T Tauri star parameters the irradiated disk surface is unlikely to produce the observed large \\ho, CO fluxes (up to 10$^{-16}$ and 10$^{-15}$ \\wm, respectively) even when the source is associated with a strong X-ray field. Slow C- and J- shocks (V$_{shock} \\le$40 \\kms\\, and V$_{shock} \\le$30 \\kms, respectively), on the other hand, can reproduce the observed line fluxes for an emitting area of diameter of a few tens to a few hundreds of AU. Thus, a shock origin is favoured.\\\\ \\item[-] {\\it high-J CO lines (up to CO J=36-35) and \\ho\\, lines from high excitation levels (up to E$_{up}\\sim$1070 K) are detected} similarly to what was observed by \\citet{vankempen10} and \\citet{herczeg12} for Class 0 and Class I outflow sources (NGC 1333 IRAS 4B and HH 46/47, respectively). This suggests that lines from high excitation levels can be shock excited if the density is high enough.\\\\ \\item[-] {\\it the extended atomic emission may be produced by fast J-shocks}. Shocks with velocities higher than 30 \\kms\\, with a radiative precursor \\citep{hollenbach89} strongly dissociate and ionize the gas giving rise to high \\oi\\, and \\cii\\, line fluxes, in agreement with the observed line ratios (\\oi\\, 63/145$\\sim$15-30, \\cii / \\oi$\\le$0.17). Excess \\cii\\, emission may be due to UV-heated gas in the outflow cavity walls.\\\\ \\item[-] {\\it molecular emission may originate instead in slow C- or J- shocks}, which preserve molecules \\citep{flower10}. High pre-shock densities are required to populate the high excitation \\ho\\, levels and reproduce the observed line ratios (i.e. J-shocks with n$ \\sim$10$^4$-10$^5$ \\cmc\\, or C-shocks with n$ \\sim$10$^6$-10$^7$ \\cmc). We cannot exclude, however, that the disk and the warm gas in the outflow cavity walls are contributing to the observed emission. \\\\ \\item[-] {\\it the cooling in the FIR lines is decreasing as the source evolves:} for the Class II sources in our sample the cooling is from one to four orders of magnitude lower than for Class I and 0 sources (L~\\oi\\,$\\sim$10$^{-4}$--10$^{-3}$ \\lsol, L~\\ho\\, and L~CO $\\sim$10$^{-4}$ \\lsol). The molecular cooling is decreasing more abruptly as the source evolves indicating that for Class 0 sources the main coolants are water and CO, while in Class I and II \\oi\\, becomes an important coolant. \\\\ \\item[-] {\\it the mass loss rate for the Class II sources in our sample is up to three orders of magnitude lower than for Class 0 and I sources}, i.e. \\mjet (\\oi~63~\\um) $\\sim 10^{-8} - 10^{-7}$ \\msolyr. \\\\ \\item[-] {\\it the mass loss rates inferred from the \\oi~63~um\\, line are larger than or comparable to values obtained from optical and NIR forbidden lines,} implying higher mass ejection to mass accretion ratios, up to 0.35. This may have important implication for jet launching models. \\\\ \\end{itemize} The above summary places the FIR emission from Class II and I jet sources within an evolutionary picture. The Taurus optical-jet-sources studied in this work show FIR atomic and molecular emission similar to that previously observed with ISO for Class 0 and Class I sources, including a highly excited molecular component. However, the emission associated with Class II sources is fainter and more compact (in particular the molecular component), and the FIR line cooling and mass loss rates are one to three orders of magnitude lower than those estimated for Class 0 and I sources." }, "1207/1207.6032_arXiv.txt": { "abstract": "The WR binary CV Serpentis (= WR113, WC8d + O8-9IV) has been a source of mystery since it was shown that its atmospheric eclipses change with time over decades, in addition to its sporadic dust production. The first high-precision time-dependent photometric observations obtained with the MOST space telescope in 2009 show two \\textit{consecutive} eclipses over the 29d orbit, with varying depths. A subsequent MOST run in 2010 showed a seemingly asymmetric eclipse profile. In order to help make sense of these observations, parallel optical spectroscopy was obtained from the Mont Megantic Observatory (2009, 2010) and from the Dominion Astrophysical Observatory (2009). Assuming these depth variations are entirely due to electron scattering in a $\\beta$-law wind, an unprecedented 62\\% increase in $\\dot{M}$ is observed over one orbital period. Alternatively, no change in mass-loss rate would be required if a relatively small fraction of the carbon ions in the wind globally recombined and coaggulated to form carbon dust grains. However, it remains a mystery as to how this could occur. There also seems to be evidence for the presence of corotating interaction regions (CIR) in the WR wind: a CIR-like signature is found in the light curves, implying a potential rotation period for the WR star of 1.6 d. Finally, a new circular orbit is derived, along with constraints for the wind collision. ", "introduction": "Ever since \\citet{1972A&A....20..333A} showed an IR excess in some WC9 stars, it has been known that certain late-type WCs produce dust. This dust is composed of amorphous carbon grains \\citep{1987A&A...182...91W} and its formation is also favored in suitable WC + O binaries (i.e. binaries with large enough separations with respect to each star's luminosity, so that the ionizing flux from the stars cannot prevent dust formation), presumably because of the high densities attained in the shocked region between the colliding winds. CV Ser (= HD 168206 = WR113, $\\alpha$ (J2000.0) = 18:19:07.36, $\\delta$ (J2000.0) = -11:37:59.2, v = 9.2) is a long-studied WC8d+O8-9IV spectroscopic binary \\citep{1945ApJ...101..356H} with atmospheric eclipses and a 29.704d period \\citep{1996RMxAC...5..100N}. Following the first published light-curve by \\citet{1949PZ......7...36G}, various other light curves have shown different eclipse depths or even no eclipse whatsoever (e.g. \\citealt{1963ApJ...137.1080H, 1970AcA....20...13S, 1970ApJ...160L.185K, 1977Obs...97....76W, 1985Ap&SS.109...57L}). Different explanations were given, including the possibility that the authors had used the wrong orbital period \\citep{1971A&A....11..407C}. However, even when no eclipses were found in the optical continuum, \\citet{1972SvA....15..955C} showed that the system was still eclipsing in the $\\lambda 4653$ emission line (confirmed by \\citealt{1972PASP...84..635M}). Since then, two MOST (Microvariability and Oscillations of STars; see below) runs conducted in 2009 and 2010 also show varying depths, possibly implying a varying mass-loss rate. CV Ser was also shown to produce dust \\citep{1975A&A....40..291C}. This phenomenon was seen as a plausible explanation for the variation in its optical eclipses. It has since been classified as a persistent dust producer \\citep{1995IAUS..163..335W}. Several studies have been carried out to refine the orbital solution of CV Ser. Most of them found a quasi circular orbit (e.g. \\citealt{1981ApJ...245..195M}) but more recently, \\citet{1996RMxAC...5..100N} have cast some doubt on that result, finding an eccentricity of 0.19. Initially, the goal of this project was to monitor stochastic short-term absorption features in CV Ser's light curve involving light from the orbiting O star in order to try and link them to clumps in the WR wind. Because of its high sensitivity, the MOST space telescope seemed like the perfect instrument to carry out this research. The system was observed in 2009 for more than a complete orbital cycle, since the anticipated absorption associated with the clumps should vary with the orbital phase, depending on the part of the WR wind illuminated by the O star along the line of sight (see Fig.~1). \\begin{figure*} \\includegraphics[width=5.0in]{sketch} \\caption{Intuitive cartoon to show how it is expected that the absorption by the clumps along the line of sight should depend on the orbital phase.} \\label{fig:1} \\end{figure*} If the clumps produce observable absorption throughout the orbit, one would expect to find shorter and deeper variations near $\\phi = 0$ (WR at inferior conjunction) and almost no random variations (unless one of the stars undergoes intrinsic variability) around $\\phi = 0.5$. Therefore, the scatter of the residuals from a given light curve fit should vary with phase, being larger around $\\phi = 0$. These variations would be a few mmag deep, as expected from the typical relative density enhancement in clumps. Section~\\ref{sec:obs} will briefly summarize the observations (both photometric and spectroscopic) of CV Ser taken for this study. The specific results for each observing run are presented in section~\\ref{sec:anal}, and then are briefly discussed in section~\\ref{sec:disc}. ", "conclusions": "\\label{sec:disc} After having eliminated the other possible causes for the change in eclipse depth observed in the 2009 MOST photometry, the deduced 62\\% increase in $\\dot{M}$ over one 29.704d orbital period is a truly remarkable finding, unprecedented for WR stars. Even more intriguing is the fact that the following year, the mass-loss rate had gone back to a lower value, suggesting that it may vary considerably over short to long timescales. The derived values of $\\dot{M}$ are also fairly low for a WR star. According to \\citet{2002A&A...392..653C}, variations in wind density should cause the spectral type to change, but this effect is not observed here, adding to the mystery, although these results are not necessarily incompatible since it is noted that the dependence on the wind density is weaker for late-type WC stars (WC8 and WC9). In any case, this definitely constitutes a challenge to both theorists and observers, since there is a dire need for a theoretical explanation as to what could drive such an impressive variation of the mass-loss rate, while constant monitoring would be required in order to get a better idea of the long-term behaviour of this parameter. One alternate explanation could conceivably be the presence of thin-shell instabilities in the wind collision zone. However, this will only affect the light curve dips stochastically, not systematically. In any case, it is not so much the details of the wind collision which determine the eclipse's shape and depth; rather it is the global wind of the WR component in which the O star orbits. As for the primary goal of this study, it was not possible to systematically link the stochastic photometric variability to the clumping phenomenon. There was no clear phase dependence of the scatter of the light curves. The variability is then probably due to noise (the combination of the unquantifiable absorption due to clumping and other stochastic phenomena), but also to CIRs in the WR wind, as the analysis of the light curve residuals suggests. The repeated signature (both in 2009 and 2010) offers a rather robust clue towards that hypothesis. This finding is very interesting in the actual context of CIRs in WR stars, since more and more CIR candidates are showing up (e.g. \\citealt{2009ApJ...698.1951S}, \\citealt{2011ApJ...735...34C}, \\citealt{2011ApJ...736..140C}). Even though they have mainly been considered as an exception up until now, CIRs in WR winds might just prove to be the rule after all. Further studies will be necessary in order to find cyclical spectroscopic evidence supporting this possibility. However, their origin remains mysterious, especially since no magnetic fields have been detected in the extended atmospheres of WR stars yet. Nevertheless, this scenario should not be ruled out since at the base of the winds, where the CIRs originate, the magnetic field might have much higher values than our current detection limits but still is not detected simply because this region is obscured by the dense wind (or perhaps due to cancellation in small dipole loops as seen on sunspots). Non radial pulsations (NRP) or starspots have also been suggested as possible causes for this phenomenon \\citep{1996ApJ...462..469C}. Finally, what was initially thought to be a possible dust event during the 2010 observations cannot conclusively be determined as statistically significant. However, the perceived asymmetry in the light curve is reminiscent of the model presented by \\citet{1998A&A...329..199V}. Although this model was used to characterize the condensation of dust clouds and its effect on a light curve, and as such is very unlikely linked to these observations (since such a dust event should create a much deeper dip in the light curve and the probability of it occurring exactly at the same time as the eclipse is rather low), it still gives way to a very important question : what role can dust play in CV Ser's bizarre behaviour? Could a variation of the quantity of dust persistently produced by the WR star affect the light curve in the same way it has been inferred that a variation of the total mass-loss rate might? In order for such an explanation to arise, one must first postulate that the dust distribution in the wind is the same as the free electron and hence overall density distribution, in order to preserve the system's geometrical properties and to reproduce the same eclipse profile found in the light curve. As to whether this assumption makes physical sense, it is hard to determine whether the dust could be produced (or maintained) following such a spatial distribution, since the production of dust in late WC stars is not yet well understood. However, it is fairly straightforward to calculate what quantity of dust would be needed to produce a similar effect to that of an increased mass-loss rate on the light curve. Assuming Mie scattering, we get a cross-section, for one grain of average radius $a$, of $\\sigma_{d} = Q \\pi a^{2}$, where $Q$ is the cross-section efficiency. Choosing $a$ to be $\\sim 0.1$ $\\mu$m (as in WR140: \\citealt{2003ApJ...596.1295M}), we find that $Q \\sim 2$ for wavelengths of the order of 5000 \\AA , giving a value of $\\sigma_{d} \\sim 6.3 \\times 10^{-10}$ cm$^{2}$. However, using a density inside the dust grains of $\\rho_{g} = 2$ g/cm$^{3}$, we find that each grain is composed of roughly $4 \\times 10^{8}$ carbon atoms. Since the most common carbon ion in the WR wind is C{\\sc iii}, for each atom of carbon we should expect to find 2 free electrons. Then, the Thomson cross-section of the number of free electrons needed to combine with $4 \\times 10^{8}$ carbon atoms to produce one grain is $\\sigma_{e, tot} = 8 \\times 10^{8} \\times \\sigma_{e} = 5.3 \\times 10^{-16}$ cm$^{2}$. Therefore, scattering by the equivalent quantity of recombined carbon in the form of a dust grain is about $10^{6}$ times more efficient, therefore an appreciable increase in the depth of the eclipse could be caused by recombination of a negligibly small quantity of carbon and its condensation in the form of dust. Possible support for this is seen in the UBV light curve of CV Ser obtained from 1984 to 1994 by Dzhapiashvili (Anthokhin, priv. com.) in which the eclipse depth varies strongly with wavelength. Therefore, if dust plays any role in these varying eclipse depths, the 62\\% value for the increase of $\\dot{M}$ over an orbital period in the 2009 data obtained in this study should probably be considered as an upper limit. Unfortunately, not much more can be deduced about the production of dust in CV Ser, except that it remains a very interesting phenomenon and should be further studied. Indeed, with its troubled history and intriguing behaviour, CV Ser might prove to be a key system for understanding the production of dust in WC+O binaries. It does not appear clear whether this process is necessarily due to the wind collision zone or if it originates in wind shocks and is then intrinsic to the WR component. In conclusion, this work possibly raises more questions than it answers, but we conclude without a doubt that CV Ser is a very important system that might hold the answer to old problems. Hopefully, our findings will motivate the community to take a deeper look into this remarkable object in the years to come." }, "1207/1207.6518_arXiv.txt": { "abstract": "{Stars are born deeply embedded in molecular clouds. In the earliest embedded phases, protostars emit the bulk of their radiation in the far-infrared wavelength range, where Herschel is perfectly suited to probe at high angular resolution and dynamic range. In the high-mass regime, the birthplaces of protostars are thought to be in the high-density structures known as infrared-dark clouds (IRDCs). While massive IRDCs are believed to have the right conditions to give rise to massive stars and clusters, the evolutionary sequence of this process is not well-characterized.} {As part of the Earliest Phases of Star formation (EPoS) Herschel Guaranteed Time Key Program, we isolate the embedded structures within IRDCs and other cold, massive molecular clouds. We present the full sample of 45 high-mass regions which were mapped at PACS 70, 100, and 160\\,$\\mu$m and SPIRE 250, 350, and 500\\,$\\mu$m. In the present paper, we characterize a population of cores which appear in the PACS bands and place them into context with their host molecular cloud and investigate their evolutionary stage.} {We construct spectral energy distributions (SEDs) of 496 cores which appear in all PACS bands, 34\\% of which lack counterparts at 24\\,$\\mu$m. From single-temperature modified blackbody fits of the SEDs, we derive the temperature, luminosity, and mass of each core. These properties predominantly reflect the conditions in the cold, outer regions. Taking into account optical depth effects and performing simple radiative transfer models, we explore the origin of emission at PACS wavelengths.} {The core population has a median temperature of 20\\,K and has masses and luminosities that span four to five orders of magnitude. Cores with a counterpart at 24\\,$\\mu$m are warmer and bluer on average than cores without a 24\\,$\\mu$m counterpart. We conclude that cores bright at 24\\,$\\mu$m are on average more advanced in their evolution, where a central protostar(s) have heated the outer bulk of the core, than 24\\,$\\mu$m-dark cores. The 24\\,$\\mu$m emission itself can arise in instances where our line of sight aligns with an exposed part of the warm inner core. About 10\\% of the total cloud mass is found in a given cloud's core population. We uncover over 300 further candidate cores which are dark until 100\\,$\\mu$m. These are possibly starless objects, and further observations will help us determine the nature of these very cold cores.} {} ", "introduction": "\\label{s:intro} \\begin{figure*} \\includegraphics[width=\\textwidth]{galaxyface_zoom2} \\caption{Face-on schematic view of the distribution of IRDCs (green triangles) and ISOSS sources (red squares) and the HMSC 07029 (blue triangle) in the Milky Way. The background image is an artist's impression of the Milky Way based on the GLIMPSE survey, credit R. Hurt [SSC-Caltech], adapted by MPIA graphics department. The kinematic distances to each object are derived using the \\citet{Reid2009} model. \\label{fig:gal_distrib2}} \\end{figure*} Star formation is a critical ingredient in a broad range of astrophysical phenomena, yet there are fundamental components of the process -- particularly in the early stages -- that remain poorly understood. Over the past decades, a basic framework for the formation of low-mass stars has developed beginning with gravitationally bound pre-stellar cores \\citep{WardThompson2002}, evolving into Class 0 and Class I protostars then Class II and Class III pre-main sequence stars \\citep{ShuAdamsLizano1987,Andre1993}. Such a sequence for the formation of high-mass stars has not yet been established. Several good candidates for massive young cores have been identified using the 170\\,$\\mu$m ISOPHOT Serendipity Survey (ISOSS) \\citep{Lemke1996,Krause2003,Krause2004,Birkmann2006} or sensitive millimeter surveys \\citep[e.g.][]{Klein2005,Sridharan:2005}. However, upon further investigation, most have been found to already host (deeply embedded) low- to intermediate-mass protostars \\citep[e.g.][]{Motte_cygX,Hennemann2008,BeutherHenning2009}. It is the stage previous to the onset of massive protostar formation that continues to elude observers. Many gaps remain in our understanding of how massive stars and clusters form, beginning with the elusive initial conditions. Do massive stars result from the gravitational collapse of cold, very massive cores \\citep{Evans2002,Beuther_ppv} like their low-mass siblings? What role does the environment play in determining the ultimate fate of such cores? As a (massive) protostar evolves, how drastically do its properties change, and what impact does it have on its surroundings? In order to study these very early, embedded phases of (massive) star formation, access to the far-infrared wavelength regime, where the peak of the cold dust radiation ($T_{dust} \\sim$ 10-20\\,K) is located, is critical. In addition, high angular resolution is needed to study massive star-forming regions which usually reside several kiloparsecs from the Sun. The {\\em Herschel} far-infrared satellite \\citep{A&ASpecialIssue-Herschel} drastically improves our ability to peer deep into the dense regions where such young cores are embedded. The Earliest Phases of Star formation (EPoS) Guaranteed Time Key Program \\citep[P.I. O. Krause; ][]{A&ASpecialIssue-Henning,A&ASpecialIssue-Linz,A&ASpecialIssue-Beuther,A&ASpecialIssue-Stutz} is a PACS and SPIRE photometric mapping survey which targets objects known to be in the cold early phases of star formation. There are two main components to the EPoS sample: 15 isolated, low-mass globules at various evolutionary stages, from starless to Class I, and 45 high-mass regions, which are mostly larger, high density molecular cloud complexes containing a range of objects within their boundaries. In \\citet{A&ASpecialIssue-Stutz}, Nielbock et al. (2012, submitted), and a forthcoming comprehensive study by Launhardt et al. (in prep.), the low-mass part of the EPoS sample is investigated. In this overview, we focus on the high-mass star-forming regions. Our sample is comprised of objects known as infrared-dark clouds (IRDCs), which were first discovered in silhouette against the bright Galactic background in the mid-infrared with the ISOCAM instrument \\citep{Perault1996} and the MSX satellite \\citep{egan_msx} at 15 and 8\\,$\\mu$m, respectively. Several surveys in millimeter and sub-millimeter continuum and spectral lines have followed \\citep[e.g.][]{carey_msx, carey_submmIRDC, Teyssier2002, Pillai_ammonia, rathborne2006, ragan_msxsurv, Vasyunina2009} and have established that massive IRDCs harbor the precursors and early phases of massive star and cluster formation. We select 29 IRDCs, most of which are in the inner quadrants of the Galaxy (see Figure~\\ref{fig:gal_distrib2}). Assuming the IRDCs lie on the near side of the Milky Way (see Section~\\ref{ss:distance}), they coincide with the Scutum-Centraurus spiral arm, which (in the first quadrant) overlaps with the Molecular Ring \\citep{Jackson_galdistr_IRDCs}. Also part of our sample are sources discovered in the ISO Serendipity Survey (ISOSS) at 170\\,$\\mu$m, which are seen to harbor cold, massive clumps at large Galactocentric distances. We present the first comprehensive results of the {\\em Herschel} PACS and SPIRE imaging survey of 45 massive targets as part of the EPoS survey. The goals of this study are as follows: (1) give a general characterization of the sample based on {\\em Herschel} data in concert with existing complementary datasets; (2) characterize point sources embedded within the targeted clouds; and (3) connect the point source properties to the overall cloud structure and environment. % ", "conclusions": "\\label{s:conclusion} We present an overview of the first results of the Earliest Phases of Star Formation survey with {\\em Herschel}, focusing on the sample of 45 high-mass regions. The goal of the work presented here is to present the EPoS sample of IRDCs and ISOSS sources as a whole and profile the population of unresolved point sources, which we call cores, within them. We use PACS point source flux densities to construct and fit modified blackbodies to the spectral energy distributions of each core and use the fit to estimate its temperature, luminosity, and mass. The main results of this work are as follows: \\begin{itemize} \\item We extract 496 point sources in the 45 IRDC structures in our sample. Their sizes range from 0.05 to 0.3~pc, which are consistent with ``cores'' in the global context of star formation \\citep{BerginTafalla_ARAA2007}. We model the SED of the cores based on the 70, 100, and 160\\,$\\mu$m point source fluxes. We find a wide range in core luminosities (0.1 to 10$^4 \\lsun$, median 16\\,$\\lsun$) and masses (0.1 to a few 10$^3 \\msun$, median 4\\,$\\msun$). The dust temperatures range from 13 to 30\\,K (median ~20\\,K). \\item The fluxes at 70, 100, and 160\\,$\\mu$m are good predictors of the core luminosity, with the tightest correlation at 160\\,$\\mu$m. We perform simple radiative transfer models which show that in cores housing protostars, emission at these wavelengths are determined mainly by the internal source properties. For starless cores, our models show that an amplified external radiation field can cause emission at these wavelengths to reach levels found in protostellar cores. Further work is needed to determine the effects of protostars with various parameters and that of anisotropic heating from neighboring sources. % \\item Most (66\\%) of the cores have a counterpart at 24\\,$\\mu$m. Cores with 24\\,$\\mu$m counterparts tend to be marginally more massive and more luminous on average than their 24\\,$\\mu$m-dark brethren. To the extent which other surveys (e.g. {\\em Spitzer} and molecular line observations) overlap with our sample, we find that the pre-existing evidence for star formation activity (e.g. YSO colors, outflow activity) almost always coincides with the presence of a 24\\,$\\mu$m counterpart, leading us to conclude that such cores contain protostars. Cores without a 24\\,$\\mu$m counterpart may also harbor protostars, but have not yet been probed for supporting evidence for embedded sources, or they may be starless cores with some level of external heating. Our radiative transfer models show that external heating is unlikely to account for the 24\\,$\\mu$m emission. In order to detect a 24\\,$\\mu$m counterpart, the inner core region containing the warm dust heated by an internal protostar must be exposed via a protostellar outflow clearing a cavity in the outer pare. This leads us to conclude that when 24\\,$\\mu$m counterparts are detected, it is because a core is in a more evolved state. \\item Cores that have a counterpart at 24\\,$\\mu$m are warmer on average than cores without a 24\\,$\\mu$m counterpart. We conclude that while the mass and luminosities of these two populations are similar, the warmer cores have protostars that have heated a larger volume of dust within the core, which moderately increases the average core temperature traced by the PACS observations. We find that cores with larger fluxes at 24\\,$\\mu$m tend to have higher outer-core temperatures. Warmer cores also correspond to ``bluer'' colors on the PACS color-color plot ([70] - [100] vs. [100] - [160]). \\item In addition to the 496 cores cataloged in this paper, we find an additional 312 candidate cores which appear only at 100 and 160\\,$\\mu$m. We can not model these cores because the SED is too poorly constrained, but we find that their [100] - [160] color is consistent with yet colder dust temperatures, thus these may represent an evolutionary stage prior to protostellar formation. \\item The cores are warmer than typical temperatures of large-scale cloud structures in which they are embedded. We conclude that at PACS wavelengths, the core fluxes are due to the effects of internal heating sources and heating from the external radiation field does not play a significant role. We emphasize here that since we probe on {\\em core scales} where the cores are embedded in high column density structures which shield them from the ambient radiation field. Further modelling of the dust temperature and column density structure on the ``clump'' and ``cloud'' scales will enable us to determine more precisely on what scales temperature gradients caused by external heating play an important role. \\item Cores within ISOSS sources tend to be 24\\,$\\mu$m-bright and less massive than cores in IRDCs. They serve as our probes of cores in the outer Galaxy, where we are biased against finding ``dark clouds'' in the same sense as IRDCs in the inner-Galaxy. As molecular gas is known to be more concentrated in the inner Molecular Ring of the Galaxy, these sources offer unique insight into the early phases of isolated ``massive'' star formation. \\end{itemize} The legacy of the EPoS data will reach far beyond the overview presented here. In this paper, we focus on the core population, much of which was until now completely obscured by heavy extinction. These embedded cores will play a role in the further interpretation of the EPoS data, which will include a full treatment of the radiative transfer from the small to large scales. Future work will also address the detailed large-scale structure in dust temperature and column density, which together can better characterize the environments of the embedded cores." }, "1207/1207.4806_arXiv.txt": { "abstract": "{In a five hour $\\rm H\\alpha$ exposure of the northwest region of the Coma cluster with the 2.1m telescope at San Pedro Martir (Mx), we discovered a 65 kpc cometary emission of ionized gas trailing behind the SBab galaxy NGC 4848. The tail points in the opposite direction of the cluster center, in the same direction where stripped HI had been detected in previous observations. The galaxy shows bright HII regions in an inner ring-like pattern, where the star formation takes place at the prodigious rate of $\\rm \\sim 8.9~ M_{\\odot}~ yr^{-1}$. From the morphologies of the galaxy and the trailing material, we infer that the galaxy is suffering from ram pressure due to its high velocity motion through the intergalactic medium. We estimate that $\\sim 4 \\times 10^9$ $\\rm M_{\\odot}$ of gas is swept out from the galaxy forming the tail. Given the ambient conditions in the Coma cluster ($\\rho_0=6.3\\times 10^{-27} ~\\rm g~ cm^{-3}$; $\\sigma_{vel}=940 ~\\rm km~s^{-1}$), simulations predict that the ram pressure mechanism is able to remove such an amount of gas in less than 200 Myr. This, combined with the geometry of the interaction, is indicative of radial infall into the cluster, leading to the conclusion that NGC 4848 has been caught during its first passage through the dense cluster environment. } ", "introduction": "Yagi et al. (2010) reported a deep (4.5 hour integration) H$\\alpha$ survey covering the central $0.5\\times 0.5 \\rm ~deg^2$ of the Coma cluster (A1656) with the Suprime-Cam mounted on the Subaru telescope. Unexpectedly for an evolved cluster such as Coma, these authors found that almost every star-forming member has its own spectacular complex of diffuse, ionized, gaseous trails extending dozens of kpc behind the optical extent of the galaxies, and sometimes harboring star-forming compact knots. They revealed 14 such systems, including those previously reported by Yagi et al. (2007) and Yoshida et al. (2008) in the Coma cluster. Similar examples can also be found in A1367 (Gavazzi et al. 2001), in Virgo (Yoshida et al. 2002, 2004; Kenney et al. 2008), and in A3627 (Sun et al. 2007). These features suggest that the galaxies were recently captured by the cluster gravitational potential and are now infalling toward the cluster center (Yagi et al. 2010). Unfortunately, the field of view of the Suprime-Cam missed by less than 5 arcmin the position of NGC 4848 which is another obvious candidate for possible extended $\\rm H\\alpha$ emission. Bothum \\& Dressler (1986) had indeed listed NGC 4848 as one of the dozen unusually active galaxies found in the Coma cluster. \\object{NGC 4848} (CGCG 160-055; Zwicky et al. 1961-68) is a bright ($M_B$=-20.5) SBab:edge-on (RC3, de Vaucouleurs et al. 1991) galaxy that lies at the northwest (N-W) periphery of the X-ray emitting region in the Coma cluster. It has a vigorous star-formation rate of $\\rm \\sim 9 ~M_{\\odot} ~yr^{-1}$ as derived from the $\\rm H\\alpha$, ultraviolet (UV), far-infrared (FIR), and radio-continuum emission (see \\S \\ref{galaxy}). Observations in the 21 cm line of HI (Gavazzi 1989; Bravo Alfaro et al. 2001) revealed a moderately deficient HI content ($Def_{\\rm HI}=0.46$), displaced in the N-W direction, as opposed to its H$_2$ content, which appears normal and centrally concentrated (Vollmer et al. 2001). The asymmetry in the HI distribution suggests that the galaxy is experiencing ram pressure (Gunn \\& Gott 1972) owing to its high velocity motion through the intergalactic medium (IGM). This discrepancy is expected since unless ram pressure stripping is severe, only the atomic phase of the gas distributed at the galaxy periphery is removed, while the $\\rm H_2$, bound deep within the galaxy potential well, is mostly unaffected by ram pressure (Combes et al. 1988; Kenney \\& Young 1989; Boselli et al. 2002; Fumagalli \\& Gavazzi 2008). Numerical hydrodynamical simulations of galaxies subject to ram pressure stripping in rich clusters (e.g. Kapferer et al. 2009; Tonnesen \\& Bryan 2009, 2010, 2012; Ruszkowski et al. 2012) reveal that in much less than 1 Gyr these galaxies lose all of their gas when the density of the IGM and the transit velocity are as high as in the Coma cluster. Consistent results are found both with and without magnetic fields. Extended gaseous tails form and the gas is shocked and heated by turbulence (Yoshida et al. 2004, 2012; Kenney et al. 2008), producing compact knots where radiative cooling takes place favoring the star formation. \\begin{table*}[!t] \\begin{center} \\caption{Log book of the imaging observations. } \\label{Tab1} \\begin{tabular}{cccccccc} \\hline \\hline \\noalign{\\smallskip} Telescope & Date & CCD & Pix & Filter & Tint & Nexp &seeing\\\\ & & & (arcsec) & (\\AA) & (sec) & & (arcsec)\\\\ \\noalign{\\smallskip} \\hline \\noalign{\\smallskip} INT & 20 Mar 1999 & $4 \\times 2048 \\times4100$ EEV & 0.33 & B & 300 & 2 & 1.0\\\\ INT & 28 Apr 2000 & $4 \\times 2048 \\times4100$ EEV & 0.33 & 6725 (80) & 1200 & 3 & 1.3\\\\ INT & 28 Apr 2000 & $4 \\times 2048 \\times4100$ EEV & 0.33 & (Gunn) $r'$ & 300 & 3 & 1.3\\\\ TNG & 09 Feb 2001 & $1024 \\times 1024$ NICS & 0.25 & H & 60 & 9 & 0.8\\\\ SPM & Apr 2012 & $1024 \\times 1024$ EEV & 0.35 & 6723 (80) & 600 & 30 & 1.4\\\\ SPM & Apr 2012 & $1024 \\times 1024$ EEV & 0.35 & (Gunn) $r$ & 60 & 26 & 1.4\\\\ \\noalign{\\smallskip} \\hline \\hline \\end{tabular} \\end{center} \\end{table*} In 2000, we serendipitously discovered a low surface-brightness $\\rm H\\alpha$ emission trailing behind NGC 4848 (not reported by Iglesias-P{\\'a}ramo et al. 2002) in a one-hour exposure of the central $1\\times 1 \\rm ~deg^2$ of the Coma cluster with the Wide Field Camera (WFC) at the Isaac Newton Telescope (INT, La Palma). However, this extended emission was only marginally detected and follow-up observations were required. Similar extended features were detected in deep GALEX images by Smith et al. (2010) showing several knots of recent star-formation along the tail. In 2012, we acquired additional five-hour observations with the San Pedro Martir (SPM) telescope using narrow-band $\\rm H\\alpha$ filters. The resulting stacked six-hour exposure, which we present in this work, is sufficiently deep to allow a robust determination of the flux in the tail. Throughout this paper, we assume $\\rm H_0=73~km~s^{-1}~Mpc^{-1}$, thus NGC 4848 is at the distance of 95.5 Mpc, that of the Coma cluster. ", "conclusions": "There is unanimous consensus (Yagi et al. 2010, Yoshida et al. 2012) that the morphological disturbances suffered by several late-type galaxies in the Coma cluster are caused by the dynamical interaction with the IGM, namely by the ram pressure mechanism (Gunn \\& Gott 1972), including NGC 4848 (Gavazzi 1989, Vollmer et al. 2001). The question is how long this process has been acting for and whether the galaxy is infalling into the cluster for the first time or, as maintained by Vollmer et al. (2001), it entered the cluster environment about 1 Gyr ago and already crossed the cluster center 400 Myr ago. From the length of the tail, we inferred that the galaxy motion takes place primarily in the plane of the sky, probably at $V \\rm \\sim 1330~km~s^{-1}$ ($\\sqrt 2 \\sigma_{vel}$). Unfortunately, we cannot directly measure this velocity since from the redshift we can infer only the component along the line of sight ($V_{//}$). Subtracting from the heliocentric velocity of NGC 4848 (Zabludoff et al. 1993), the mean velocity of Coma cluster galaxies (Gavazzi et al. 2010), we obtained $V_{//}=257~\\rm km~s^{-1}$. There are two possibilities: that the orbit is circular, with its angular momentum in the plane of the sky, or that it is radial. In the first case, the ratio of the velocity along the line of sight to the tangential velocity $V_{//}/V = 257/1330 = 0.19$ gives the angle between $V_\\perp$ and $V$. Assuming that the present projected radial distance of NGC 4848 from the cluster center is 0.75 Mpc, we conclude that the radius of the circular orbit is 4 Mpc, i.e. exceeds the virial radius of Coma by almost a factor of two. We are much more likely to conclude, however, that the orbit is radial or at least highly eccentric, and that the present motion of NGC 4848 is toward the center of the cluster at $V \\rm \\sim 1330~km~s^{-1}$, as suggested by the direction of the H$\\alpha$ tail pointing in the opposite direction. More evidence of infall into the Coma cluster is provided by the velocity distribution of late-type galaxies (LTG) which in general hardly follow a Gaussian distribution but one that is skewed toward either lower or higher velocities than the overall mean velocity (Biviano et al. 1997, Boselli \\& Gavazzi 2006). The velocity dispersion profile of the LTGs in clusters is found to be consistent with orbits more radial than those of early-type galaxies (ETG), providing a picture in which possibly all spirals have not yet crossed the virialized cluster core, and may even be on a first (infall) approach toward the central, high-density region. In the ram pressure scenario, the amount of stripped gas can be computed from the equilibrium between the gravitational binding energy and the dynamical pressure. The calculation requires the density profile of the gas in the IGM. Following Cavaliere \\& Fusco-Femiano (1978), the density radial profile is well-approximated by an isothermal sphere \\begin{equation} \\rho_{\\rm{IGM}}=\\rho_0\\left[1+\\left(\\frac{r}{r_c}\\right)^2\\right]^{-3\\beta/2}, \\end{equation} where $\\rho_0$ represents the central density and $r_c$ the scale-length. Using the parameters estimates of Mohr et al. (1999) for the Coma cluster (all quantities having been recomputed for $\\rm H_0=73~km~s^{-1}~Mpc^{-1}$) namely $\\rho_0=6.3\\times 10^{-27} ~\\rm g~ cm^{-3}$, $r_c=0.26 \\rm ~Mpc$, and $\\beta=0.7$, we computed that at the present position of NGC 4848 (0.75 Mpc from the center) the IGM density is $\\rho_{0.75}=6\\times 10^{-28} ~\\rm g~ cm^{-3}$. The radius at which ram pressure becomes efficient can be estimated as (Domainko et al. 2006) \\begin{equation} R_{\\rm{strip}} = 0.5 R_{\\rm{HI}} ln \\left (\\frac{GM_{\\rm{star}}M_{\\rm{HI}}}{v^2 \\rho_{\\rm{IGM}} 2 \\pi R_{\\rm{star}}^2 R_{\\rm{HI}}^2} \\right), \\end{equation} while the stripped mass is \\begin{equation} M_{\\rm{strip}} = M_{\\rm{HI}} \\left(\\frac{R_{\\rm{strip}}}{R_{\\rm{HI}}}+1\\right) exp\\left(-\\frac{R_{\\rm{strip}}}{R_{\\rm{HI}}}\\right). \\end{equation} In this calculation, we assumed exponential profiles for both the stars and the interstellar gas, with $R_{\\rm star}=5.0$ kpc (computed on the H band image) and $R_{\\rm HI}=1.8 \\times R_{\\rm star}=9.0$ kpc for the HI disk (Boselli \\& Gavazzi 2006). This yielded $R_{\\rm{strip}} = 8.5$ kpc and $M_{\\rm{strip}} = 4.7 \\times 10^9 \\rm ~M_\\odot ~\\sim 0.75 ~M_{HI~orig}$. This is in remarkably good agreement with the mass of stripped gas ($3.6\\times 10^9 \\rm ~M_\\odot$) computed in \\S \\ref{trail}, which in turns is consistent with the missing mass of atomic hydrogen ($4.1\\times 10^9 \\rm ~M_\\odot$) computed from the HI deficiency parameter. Using the ram pressure simulation by Kapferer et al. (2009) (the case with $V_{rel}=1000 \\rm ~km~s^{-1}$) and assuming the IGM density computed above ($\\rho_{0.75}=6\\times 10^{-28} ~\\rm g~ cm^{-3}$), we derived that 65 \\% of the original HI content is stripped in about 200 Myr. In spite of the different conditions found in the Virgo cluster with respect to the Coma Cluster, a similarly short timescale ($\\sim 100 \\rm ~Myr$) is reported by Boselli et al. (2006) for an almost complete ablation of the atomic gas from NGC 4569. This time is significantly shorter than the crossing time in the Coma cluster ($1.6\\times 10^9$ yr, Boselli \\& Gavazzi 2006), supporting the conclusion that NGC 4848 is on its first passage through the cluster core. This time would be sufficient to remove most of its gas, in contradiction to the 1 Gyr timescale proposed by Vollmer et al. (2001)." }, "1207/1207.1562_arXiv.txt": { "abstract": "{The latest observations of line and continuum spectra emitted from the extended narrow line region (ENLR) of the Seyfert 2 galaxy NGC 7212 are analysed using models accounting for photoionization from the active nucleus and shocks. The results show that relatively high (500--800 \\kms) shock velocities appear on the edge of the cone and outside of it. The model-inferred AGN flux, which is lower than $10^{-11}$ photons cm$^{-2}$ s$^{-1}$ eV$^{-1}$ at the Lyman limit, is more typical of low-luminosity AGN, and less so for Seyfert 2 galaxies. The preshock densities are characteristic of the ENLR and range between 80--150 cm$^{-3}$. Nitrogen and sulphur are found depleted by a factor lower than 2, particularly at the eastern edge. Oxygen is depleted at several locations. The Fe/H ratio is approximately solar, whereas the Ne/H relative abundance is unusually high, 1.5--2 times the solar value. Modelling the continuum spectral energy distribution (SED), we have found radio synchrotron radiation generated by the Fermi mechanism at the shock front, whereas the X-rays are produced by the bremsstrahlung from a relatively high temperature plasma.} ", "introduction": "Morphological and spectroscopic studies of Seyfert galaxies have shed light on the nature of active galactic nuclei (AGN). Specifcally, the spectra of type 2 Seyferts are generally dominated by strong forbidden and permitted lines revealing the physical properties of the extended narrow line region (ENLR). The spectra indicate that radiation from the AGN dominates the radiation field; however, line ratios corresponding to low ionization levels and neutral lines could be explained by underlying shock wave hydrodynamics. The full width at half maximum (FWHM) of the lines reveal velocities up to 1000 \\kms. Modelling of line and continuum spectra observed in many single regions throughout the ENLR provided a new dimension to our knowledge. For instance, the complex nature of \\object{NGC 7130}, in terms of the intermingled activity of starbursts, shocks, and of an active nucleus and their mutual location, could be traced by Contini et al. (\\cite{Con02a}). In such objects, e.g. \\object{NGC 4388} (Ciroi et al. \\cite{C03}) or \\object{Mrk 298} (Radovich et al. \\cite{Rad05}), the relative importance of the radiation flux from the AGN, of shocks, and of radiation from starbursts could be determined in the different regions of the galaxy. This method, which is useful for isolated galaxies, becomes fundamental for the analysis of Seyfert nuclei in merging systems originating from collisions, even in case of multiple (generally double) nuclei that are not yet clearly identified as such by the observations. For instance, the biconical structure of the high excitation region which emerges from the toroidal obscuring material, can be fragmented by collision, and patches of highly ionized material can appear at the edges of the ENLR, as e.g. in \\object{NGC 3393} (Cooke et al. \\cite{coo00}). The interaction of the ENLR galactic matter with fast shocks, star formation, dust grain destruction etc. are amply analysed and summarised by Ramos Almeida et al. (\\cite{R09}) for \\object{NGC 7212} and other Seyfert 2 galaxies. \\object{NGC 7212} (z=0.0266) belongs to a compact group of interacting galaxies, with two galaxies in a clear ongoing merger (Wasilewski 1981). Spectropolarimetry studies show the presence of a hidden broad line region (BLR) (Tran \\& Kay 1992), clumpy nuclear morphology, and irregular diffuse emission. Mu{\\~n}oz Mar{\\'{\\i}}n et al. (\\cite{M07}) also claim that these galaxies manifest some cases in which in addition to the extended emission, the UV light stems from the knots or clumps likely produced by star clusters. They identify the bright clumps to the south as star-forming regions with a certain contribution from the AGN. Exploring whether photoionization or shocks dominate in NGC 7212, Bianchi et al. (\\cite{bi06}) found a striking resemblance of the [\\ion{O}{III}] structure with soft X-ray emission, a clear indication that shocks are at work. In an accompanying paper, Cracco et al. (\\cite{Cr11}) presented new observations of NGC 7212 and suggested that its ENLR gas has an external origin, likely due to gravitational interaction effects in act in this triple system. The optical spectra detailed by Cracco et al. are rich in number of lines: e.g. oxygen from three ionization levels ([\\ion{O}{III}]5007, [\\ion{O}{II}]3727 and [\\ion{O}{I}]6300), \\ion{He}{II} 4686 and \\ion{He}{I} 5876, [\\ion{Ne}{III}]3869, [\\ion{N}{II}]6548,6583, the two [\\ion{S}{II}] lines 6716 and 6731, lines from higher ionization levels, e.g. [\\ion{Ar}{IV}]4713 and [\\ion{Fe}{VII}]6087, in addition to H$\\beta$ and H$\\alpha$. The line ratios constrain the physical conditions of the emitting gas and the relative abundances. These authors pointed out for the first time the presence of an extended ionization cone in NGC 7212, with high values of [\\ion{O}{III}]/H$\\beta$, up to 3.6 kpc. In fact, the Veilleux \\& Osterbrock (\\cite{VO87}) diagnostic diagrams indicate that the main source of ionization, at least in the entire field of view of the integral field spectra, is the AGN. However, high values of [\\ion{N}{II}]/H$\\alpha$ and [\\ion{S}{II}]/H$\\alpha$ are observed towards the region of gravitational interaction between the interacting galaxies, suggesting a possible combination of ionization by the active nucleus and shocks. On the other hand, no contribution from stellar ionization is required to explain the observed emission line ratios. An archival broad-band Hubble Space Telescope (HST) image of NGC 7212 obtained with the F606W filter revelas a structure composed of clouds or filaments (see fig.\\,19 in Cracco et al. \\cite{Cr11}). The evidence of dust located in the ENLR also supports the idea that the gas in the cone has a filamentary structure. In addition, comparing the more significant line ratios (e.g. [\\ion{O}{III}]/H$\\beta$, [\\ion{O}{II}]/H$\\beta$, [\\ion{N}{II}]/H$\\alpha$, etc) with diagnostic diagrams, Cracco et al. found depletion of heavy elements such as N and O relative to solar. The emission line profiles of [\\ion{O}{III}] inside the ionization cone display blue wings in the northern side of the cone, and red wings in the southern side, while in the nucleus the profiles are symmetric. This suggests the presence of gas in radial motions, which is confirmed also by the analysis of high resolution echelle spectra characterized by multiple kinematical components at different velocities. If radial motions are in act in the ENLR of NGC 7212, shock ionization is expected to play a non negligible role. Moreover, it is well known that the higher the flux from the photoionizing source, the stronger the [\\ion{O}{III}] lines are, while emission lines such as [\\ion{O}{II}], [\\ion{N}{II}], etc., namely low ionization level lines, are significant when the gas is affected by strong shock. Consequently, we decided to reproduce the observational data presented by Cracco et al. (\\cite{Cr11}) by means of models accounting for the coupled effect of photoionization and shock, improving their interpretation, and adding information concerning the physical and chemical conditions throughout the galaxy. The structure of the paper is as follows: observations of the optical spectra are described in Sect.\\,\\ref{obs}. In Sect.\\,\\ref{model} the models and the modelling method are explained. Modelling results of the line and continuum spectra are presented in Sect.\\,\\ref{model_result} and in Sect.\\,\\ref{sed}, respectively. Discussion and concluding remarks follow in Sect.\\, \\ref{disc} and Sect.\\,\\ref{result}. \\begin{figure} \\centering \\includegraphics[width=8.5cm]{continuo_num.ps} \\caption{2D map of the stellar continuum at 5500 \\AA. The numbers identify the spectra in Fig.\\,\\ref{fig4b}. The image size is $16\\times16$ square arcsec. North is up and East is to the left.} \\label{fig1} \\end{figure} ", "conclusions": "\\label{result} In the previous sections we presented an analysis of the line spectra observed by Cracco et al. (\\cite{Cr11}) in the different positions throughout the ENLR of NGC 7212, adopting calculation models which account for the coupled effect of shocks and photoionization. The results yield new reliable estimates about the intensity and distribution of the flux from the AGN, the velocity field, the preshock density, the relative abundances, etc. In Seyfert galaxies the radiation from the active nucleus is the main photoionization mechanism. The AGN in NGC 7212 is characterised by a flux intensity lower than 10$^{11}$ photons cm$^{-2}$ s$^{-1}$ eV$^{-1}$, at the Lyman limit, similar to that of low luminosity AGN (e.g. Contini \\cite{Con04}), rather than to Seyfert 2 (e.g. Contini et al. \\cite{Con02a}, Ciroi et al. 2003). The AGN flux peaks in correspondence with the maximum starlight flux. In the ENLR of NGC 7212, as in most Seyfert galaxies, the single lines show complex profiles. This indicates that the hydrodynamical field is complex. Shock waves created by winds and jets from the stars and/or by collision of matter are generally present. Characteristic of shocked clouds is the profile of the density downstream, where the gas cools due to recombination processes. Compression speeds up recombination, enhancing the intensity of lines corresponding to low ionization levels. The spectra emitted by shocked gas are therefore different from those emitted from clouds heated and ionized by radiation from an external source. The relative importance of photoionization and shocks in NGC 7212 throughout most of the 256 spatial elements simultaneously observed by MPFS, results from the consistent modelling of line and continuum spectra. Each spectrum stems from a complex grid of models which account for the parameter in various ranges. The parameters are constrained by the \\ion{He}{II}/H$\\beta$ line ratios ($F$), by [\\ion{S}{II}]6716/6731 (\\n0), by [\\ion{O}{III}]:[\\ion{O}{II}]:[\\ion{O}{I}] (\\Vs and $F$), etc. The shock velocity is also constrained by the continuum SED. In conclusion, the interaction of galaxy-galaxy or of a galaxy with matter from the interstellar or intracluster medium can reveal shocks by line ratios which cannot be readily explained by pure photoionization models (e.g. Villar-Mart{\\'{\\i}}n et al. 1999), by high electron temperatures leading to X-ray emission and by radio synchrotron radiation. Radiation and collisional processes are strongly intermingled in ionizing the gas (e.g. Contini 1997). The degree to which shocks contribute to the gas ionization in the different regions of NGC 7212 can be found only by consistent modelling of line and continuum spectra. Our results agree with Tadhunter et al. (2000), who claim that shocks contribute significantly to the ionization of the ENLR of Seyfert galaxies, even if photoionization contributes substantially in most of AGNs (Robinson et al. 1987, Villar-Mart{\\'{\\i}}n et al. 1997)." }, "1207/1207.3084_arXiv.txt": { "abstract": "We present a study of the resolved star-forming properties of a sample of distant massive ($M_{\\*} > 10^{11}M_{\\odot}$) galaxies in the GOODS NICMOS Survey (GNS), based on deep Hubble Space Telescope imaging from the GOODS North and South fields. We derive dust corrected UV star formation rates (SFRs) as a function of radius for 45 massive galaxies within the redshift range $1.5 < z < 3$ in order to measure the spatial location of ongoing star formation in massive galaxies. We find that the star formation rates present in different regions of a galaxy reflect the already existent stellar mass density, i.e. high density regions have higher star formation rates than lower density regions, on average. This observed star formation is extrapolated in several ways to the present day, and we measure the amount of new stellar mass that is created in individual portions of each galaxy to determine how the stellar mass added via star formation changes the observed stellar mass profile, the S\\'{e}rsic index and effective radius over time. We find that these massive galaxies fall into three broad classifications of star formation distribution: 1) Total stellar mass added via star formation is insignificant compared to the stellar mass that is already in place at high redshift. 2) Stellar mass added via star formation is only significant in the outer regions ($R > 1\\rm{kpc})$ of the galaxy. 3) Stellar mass added via star formation is significant in both the inner $(R < 1\\rm{kpc})$ and outer regions of the galaxy. These different star formation distributions increase the effective radii over time, which are on average a factor of $\\sim16\\pm5\\%$ larger, with little change in the S\\'{e}rsic index (average $\\Delta n = -0.9\\pm0.9$) after evolution. We also implement a range of simple stellar migration models into the simulated evolutionary path of these galaxies in order to gauge its effect on the properties of our sample. This yields a larger increase in the evolved effective radii than the pure static star formation model, with a maximum average increase of $\\Delta R_{e}\\sim 54\\pm 19\\%$, but with little change in the S\\'{e}rsic index, $\\Delta n \\sim-1.1\\pm1.3$. These results are not in agreement with the observed change in the effective radius and S\\'{e}rsic index between $z\\sim2.5$ and $z\\sim0$ obtained via various observational studies. We conclude that star formation and stellar migration alone cannot account for the observed change in structural parameters for this galaxy population, implying that other mechanisms must additionally be at work to produce the evolution, such as merging. ", "introduction": "\\label{sec:Intro} One of the least understood aspects of galaxy evolution is the star formation rates in galaxies, how these vary across individual galaxies, and influence galaxy properties. A key way to address galaxy evolution directly is to understand how the nearby galaxy population was put into place and evolved from higher redshift galaxies, which we can now observe in near complete mass-selected samples up to $z=3$ (e.g. Daddi et al. 2007; Conselice et al. 2011). One major finding of high redshift studies is that massive galaxies $(M_{*} > 10^{11}M_{\\odot})$ have significantly smaller effective radii than low redshift galaxies of similar mass (e.g. Daddi et al 2005; Trujillo et al. 2006a, 2006b, 2007, 2011; Buitrago et al. 2008, 2011; Cimatti et al. 2008; van Dokkum et al. 2008, 2010; Franx et al. 2008 ; van der Wel et al. 2008; Damjanov et al. 2009; Carrasco et al. 2010; Newman et al. 2010; Szomoru et al. 2011; Weinzirl et al. 2011). Several physical processes have been proposed to explain this strong size evolution within the massive galaxy population at $z < 2$. These can be divided into two distinct categories, external processes such as gas poor (\u201cdry\u201d) mergers (e.g. Khochfar \\& Silk 2006; Naab et al. 2009) and cold gas flows along cosmic web filaments (e.g. Dekel et al. 2009; Conselice et al. 2012 submitted) as a means for puffing up the stellar components of these massive galaxies, or internal processes such as adiabatic expansion resulting from stellar mass loss and strong AGN-fuelled feedback (e.g. Fan et al. 2008, 2010; Hopkins et al. 2010A; Bluck et al. 2011). One process that has not been looked at in detail is the internal star formation distribution present within massive galaxies at high redshift, and whether this can account for the observed structural evolution. This can now be examined due to high resolution data from the GOODS NICMOS Survey taken with the ACS and NICMOS-3 instruments on the Hubble Space Telescope (Conselice at al. 2011). We know that galaxies evolve significantly in stellar mass from observational studies showing that half of the stellar mass of present day galaxies is already in place by $z \\sim1$ (e.g. Brinchmann \\& Ellis 2000; Drory et al. 2004; Bundy et al. 2006; P\\'{e}rez-Gonz\\'{a}lez et al. 2008; Mortlock et al. 2011). The most massive galaxies $(M_{*} > 10^{11}M_{\\odot})$ appear on average to have red rest-frame colours which we expect to see for galaxies dominated by old stellar populations (Saracco et al. 2005; Labb\\'{e} et al. 2006; Conselice et al. 2007; Gr\\\"{u}tzbauch et al. 2011). However, Bauer et al. (2011b) show that $\\sim80\\%$ of these massive red galaxies likely harbour dusty star formation. This star formation over cosmic time could contribute large amounts of stellar mass to massive galaxies, and depending on where this mass is created could affect their observable structural properties as they evolve. In the merger scenario, estimates for the total number of major mergers experienced by massive galaxies since $z=3$ is $N_{m} = 1.7 \\pm 0.5$ (Bluck et al. 2009). This would imply an average stellar mass increase of, at best, a factor of two due to major mergers. However over the same period of time the effective radius of massive galaxies has increased on average by a factor of three for disk\\--like galaxies, and a factor of five for spheroid\\--like galaxies, effectively building up stellar mass in the outer regions of galaxies (see e.g. Buitrago et al. 2008, 2011; Trujillo et al. 2007; Carrasco et al. 2010; van Dokkum et al. 2010). This additional stellar mass could arise from star formation already present at high redshift within these outer regions. To date studies have only looked at the total star formation rates of these galaxies as a whole (e.g. P\\'{e}rez-Gonz\\'{a}lez et al. 2008; Cava et al. 2010; van Dokkum et al. 2010; Bauer et al. 2011; Gr\\\"{u}zbauch et al. 2011; Viero et al. 2011; Hilton et al. 2012), but have not examined the locations of the star formation within these galaxies. Thus, we combine the observed stellar mass profiles with the observed star formation profiles of high redshift massive galaxies in order to measure the effect stellar mass added via star formation over $\\sim 10$ Gyr has on different spatial regions, and to the total stellar mass profile. We also ascertain whether this star formation can account for the observed size evolution. Along with size evolution within the massive galaxy population there is also a change in overall morphology. The present day universe is populated by massive galaxies with early\\--type morphologies (e.g. Baldry et al 2004, Conselice et al. 2006). At earlier epochs, $z>1.5$, observational studies have found that the massive galaxy population is dominated by galaxies with late\\--type morphologies (e.g. Buitrago et al. 2011; Cameron et al. 2011; van der Wel et al. 2011; Weinzirl et al. 2011). This morphological shift can be seen via a change in S\\'{e}rsic index from low values, $n < \\sim2.5$ denoting a possible late\\--type morphology, to high values, $n > \\sim2.5$ denoting a possible early\\--type morphology. In the hierarchical model of galaxy evolution there are many methods that can drive morphological evolution. These methods include in situ star formation producing disk\\--like systems (e.g. Dekel et al. 2009; Oser et al. 2010; Ricciardelli et al. 2010; Wuyts et al. 2010; Bournaud et al. 2011), and/or mergers with satellite galaxies producing a more spheroid\\--like system (e.g. Khochfar \\& Silk 2006; Hopkins et al. 2009; Feldmann et al. 2010; Oser et al. 2010). We therefore also investigate how in situ star formation over cosmic time changes the S\\'{e}rsic index of the massive galaxies, and ascertain whether this process can account for the observed morphological changes. This paper is set out as follows: Section 2 discusses the GOODS NICMOS Survey, the galaxy sample, and how the data used in this paper was obtained. Section 3.1 examines the stellar mass radial density distributions of the massive galaxies. Section 3.2 describes how the stellar mass density added via star formation is calculated. In Section 3.3 we examine the evolved galaxy profiles. Section 4.1 presents the findings of how the structure and size of the massive galaxies is altered by star formation. In Section 4.2 we introduce a simple stellar migration model to the stellar mass added by star formation in order to gauge the effect this has on structures and sizes. Section 5 and 6 contain the discussion and summary of our findings, respectively. Throughout this paper we assume $\\Omega_{M}=0.3$, $\\Omega_{\\Lambda}=0.7$ and $H_{0}=70$ km s$^{-1}$ Mpc$^{-1}$. AB magnitudes and a Salpeter IMF are used throughout. ", "conclusions": "\\subsection{Size Evolution} As stated previously, recent studies over the last few years have found evidence for a dramatic size evolution of massive galaxies over the past 10 billion years (e.g Daddi et al. 2005; Truijllo et al. 2007; van Dokkum et al. 2010; Buitrago et al. 2008, 2011). Current estimates for this size growth argue that massive galaxies may grow in size on average up to a factor of 3 for disk-like galaxies, while for spheroid-like objects this evolution reaches even a factor of 5 since $z=3$ (Buitrago et al. 2008). In this paper we have shown that the effective radius of massive galaxies is altered by the star formation present, growing on average by $16 \\pm 5\\%$ from $z=3$ to $z=0$. This value is only $\\sim 3-5\\%$ of the total increase in the size of massive galaxies from observational studies (e.g. Buitrago et al 2008). This indicates that the star formation has a very minor contribution to the observable overall size evolution at $z<3$. When we apply a simple model of stellar migration to the new stellar mass created via star formation to the present day we find that the size of these massive galaxies is influenced to a greater extent. The effective radius increases by $54 \\pm 19\\%$. This increase would represent $11-18\\%$ of the total size evolution that massive galaxies undergo between $z>1$ and 0. This result shows that the effects of stellar mass added via star formation, and any subsequent stellar migration, plays a minor role in massive galaxy size evolution and only contributes roughly a tenth of the total size growth needed to explain the observed size evolution. This implies that other evolution mechanisms must also be at work to produce the remaining $\\sim80\\%$ of the observed size growth over cosmic time. From also examining the total size growth in the other models of evolution (see \\S 5.4.1) we also find that the maximal size increase we can obtain can only produce $\\sim54\\%$ of the total observed size growth. Recent studies have found that minor and major mergers have a large influence on the size evolution of massive galaxies. These mergers could explain the majority of the remaining $\\sim 80\\%$ of the observed size growth unaccounted for by the SF via increasing the total stellar mass of the galaxies (Bluck et al. 2011). Our results are consistent with this view that something other than SF produces the change in the sizes of massive galaxies. \\subsection{Structural Properties} Recent studies have shown that the massive galaxy population at $z \\ge 1.5$ is dominated by disk like galaxy morphologies with $n<2$ (e.g., Buitrago et al. 2011; Weinzirl et al. 2011). This is in contrast to the local universe where the massive galaxy population is almost entirely dominated by spheroids (e.g. Baldry et al. 2004; Conselice et al. 2006). This transformation is also seen through changes in the S\\'{e}rsic index of these galaxies from a low value of $n$ at $z>1$ to a high value of $n$ at $z<1$. In this study we show that due to the star formation present within the massive galaxies at $z>1.5$ the S\\'{e}rsic index has an insignificant change over cosmic time, $\\Delta n = -0.9 \\pm 0.9$. When we introduce the effects of stellar migration to the mass added via star formation the change in S\\'{e}rsic index is again negligible over cosmic time with, $\\Delta n = -1.1 \\pm 1.3$. In the other methods of SF evolution we find that the change in $n$ is very similar. This implies that with both star formation and stellar migration the change to the S\\'{e}rsic index is minimal. Also, this does not agree with observations of the general increase of $n$ over time. Therefore SF alone cannot account for the observed morphological change which appear to show that $n$ is increasing over time (e.g Buitrago et al. 2011). \\subsection{Spatial Location of Star Formation} In this study we find that the structural properties of our massive galaxies remain largely unchanged after evolution via star formation. This unchanging $n$ shows that the location and magnitude of star formation within massive galaxies largely follows the observed initial stellar mass density profile. This is most pronounced in the case of the inner growth (IG) galaxies. In this class of galaxies the observed stellar mass profile is much smaller than the stellar mass profile added via star formation. Therefore, for this class of galaxy to retain its original S\\'{e}rsic index the stellar mass produced via star formation over evolution to the present day would have to be produced in amounts which largely reflect the already present stellar density i.e. high density regions would have a higher star formation rates than lower density regions. This was also seen in other ways in Trujillo et al.(2007), Buitrago et al. (2008) and Cassata et al. (2010, 2011). The measured $\\sim16\\%$ growth of the effective radii of our massive galaxies due to star formation alone, without any stellar migration, reveals that there is star formation located in the outer regions of our massive galaxies. This is most pronounced in the OG galaxies by definition. In these galaxies the surface stellar mass density of the inner region remains roughly constant over star formation evolution with the outer regions increasing in stellar mass density. Thus in our simulated star formation evolution the observed high redshift galaxy would become surrounded by an envelope of new stellar material over time. With the addition of stellar migration this effect becomes more pronounced with newly created stellar mass migrating outwards. Recent work examining the stellar mass density profiles of high redshift, $z>2$, and low redshift, $z=0$, massive galaxies has shown that the density in the core region of low redshift galaxies is comparable to the density of the compact high redshift galaxies (Hopkins et al. 2009; van Dokkum et al. 2010; Carrasco et al. 2010). The compact high redshift galaxies have become surrounded by an envelope of lower density material from $z>2$ to $0$. This is similar to what we find in the OG class of galaxies. The models that we use in this study do not account for any new gas that can be accreted at later times, at $z<1.5$, and at early times at $z > 3$ where we also do not observe our sample. This new gas and possible new star formation is likely to have a different radial distribution from the current in situ gas, with most of the new gas being at larger radii (Keres et al. 2005; Dekel et al. 2009). Therefore the distribution of star formation that we observe at high redshift is mostly likely the result of previous events of gas accretion (see Conselice et al. 2012). However, not all the gas accreted may convert into stars immediately, and this gas may remain in the outer portions of these galaxies and may form into stars at an epoch later than our observations at $z < 1.5$, which in principle may increase the sizes of these systems at a later time, or alter their S\\'{e}rsic indices. \\subsection{Model Limitations} In this study we have taken a snapshot of our massive galaxy sample over 2 Gyr in time, and derived the resulting evolution based on a derived star formation model. Thus we do not take into account any post\\--observation star formation events in our basic model. However this is likely a fair assumption due to observations of the majority of massive galaxies at $z<1.4$ having old stellar populations and red colours (e.g. McCarthy et al. 2004; Daddi et al. 2005; Saracco et al. 2005; Bundy et al. 2006; Labb\\'{e} et al. 2006; Conselice et al 2007; Mortlock et al. 2011; Gr\\\"{u}tzbauch et al. 2011). This would imply that the SF we observe at $z>1.5$ is the last major burst of SF in massive galaxies. The effect of new star formation events would increase the total amount of stellar mass added to the host galaxy. The galaxy's structural properties and size could also be affected by these events, depending on the location and magnitude of this star formation as discussed in the previous section. Conversely, we also do not take into account any feedback mechanisms that would negatively affect star formation rates. Examples of such processes are AGN and supernovae feedback. Massive galaxies can spend up to 1/3 of their lifetimes in an AGN phase (Hickox et al. 2009, Bluck et al. 2011). This phase introduces energy into the interstellar gas and can expel it from the host galaxy (Schawinski et al. 2006), or heat it such that it cannot cool. Also ongoing star formation results in the creation of many high mass stars which can lose mass during evolution and subsequently die in supernovae, thereby lowering the total stellar mass of the galaxy. When many supernovae are present in a short time the created shock waves introduce vast amounts of energy into interstellar gas. The gas can then can be heated or ejected from the host galaxy (e.g. Bertone et al. 2007). The result of these feedback mechanisms would be a reduction of the star formation rate, and the total stellar mass within the galaxy would be lower. This decreased amount of stellar mass added via SF would also result in the stellar mass added via star formation having a decreased effect on the total size growth and morphological change. We also use a very simple models to describe the stellar migration that is limited to the extent of the $z_{850}$ band profiles. This means that we can not accurately measure how large values of stellar migration would affect the sizes and structural properties of our massive galaxies. However even though we cannot accurately measure the S\\'{e}rsic index or the effective radius of the simulated galaxies with larger values of the stellar migration, we find that the stellar mass begins to be distributed evenly over all radii, with increasing amounts of stellar mass lost outside the confines of the simulation. The amount of stellar mass added via star formation moved by migration is constant for each galaxy but is distributed over wider areas for larger values of stellar migration. This results in the stellar mass density added via star formation to individual regions of the massive galaxies dropping to increasingly smaller values. This implies that with larger values of stellar migration, the stellar mass density added via star formation would have an increasingly smaller effect on the total stellar mass density profile. Therefore even if larger values of stellar migration could be simulated in this study the change in S\\'{e}rsic index and effective radius after star formation evolution and migration would be negligible. Stellar migration has also been found, in simulations, to be most affected by spiral arms in galaxies (Ro{\\v s}kar et al. 2011). $73\\%$ of the sample of massive galaxies have a low S\\'{e}rsic index, $n < 2.5$, implying a disk\\--like morphology. Within these galaxies we may assume therefore that stellar migration via disk features may take place, but this is far from certain. A few of the galaxies in our sample have a high S\\'{e}rsic index, $n >2.5$, implying an early\\--type morphology, and within these galaxies stellar migration is less understood. This does not imply that stellar migration does not take place in these galaxies but it must occur by other processes than those involving disks. Also, as stated in \\S4.3 we cannot reliably distinguish disk-like galaxies in our sample using a S\\'{e}rsic index cut because we cannot rule out that some of the galaxies with $n>2.5$ do not have spiral like features (e.g. Buitrago et al. 2011: Mortlock 2012, private communication). \\subsubsection{Evolutionary models} In this paper we extrapolate the star formation evolution using a exponentially declining star formation model based on SED derived $\\tau$ values. This value can be uncertain so we explore different models of evolution that the star formation could follow down to $z=0$. Firstly we do not investigate an exponentially increasing SFR evolution model because previous studies (e.g. Papovich et al. 2011) show that galaxies at $z<3$ are not well described by this SF history. Therefore we investigate the SF evolution models of: constant SFR to $z=0$, constant SFR to $z=1.5$, maximum valid tau and minimum valid tau. \\begin{itemize} \\item{Constant $\\rm{SFR_0}$ to $z=0$: This model of evolution assumes that the massive galaxies we observe at $z>1.5$ have a very large reservoir of gas and can continue the observed SFR over the next ~10Gyr. This evolutionary method produces galaxies in the local universe with very high star formation rates compared to the galaxies we observe (e.g. Conselice et al 2007). This combined with the fact that over the course of their evolution these galaxies will have accumulated significant amounts of stellar mass with the average massive galaxy in this sample increasing its total stellar mass by $\\sim1500\\%$. This large amount of stellar mass added to the galaxies increases the value of $R_{e}$ by $80\\pm20\\%$. This is an increase of a factor of 5 over the derived tau model in effective radius growth, but still only $16-27\\%$ of the observed size evolution. This model of evolution is highly unlikely due to the many features of this model that we do not observe in the local universe, such as very large stellar mass growth leading to very massive galaxies with stellar masses over $10^{13}M_{\\odot}$ (e.g. Brammer et al. 2011; Conselice et al 2011; Mortlock et al. 2011, all find that the stellar mass growth at the massive end of the luminosity function is on the order of $200\\%$ from $z>1.5$ to 0) and very high star formation rates of 100's of solar masses per year.} \\item{Constant $\\rm{SFR_0}$ to $z=1.5$: This model of SF evolution is based on the observation that the majority of massive galaxies at $z<1.4$ have old stellar populations and red colours (e.g. Conselice et al 2007: Mortlock et al. 2011, Gr\\\"{u}tzbauch et al. 2011). This would imply that these galaxies have turned off their SF before $z=1.5$. To model this we employed a constant observed SFR until $z=1.5$ at which point the SFR is reduced to 0. In this evolution scenario the total stellar mass of the massive galaxies is increased by $126\\pm20\\%$. The effective radii in this model are increased on average by $37\\pm19\\%$. This is a factor of $\\sim2.5$ larger then the increase from the derived tau model. This is still insignificant to the total observed size increase. This model has a very similar effect on the change in $n$, $\\Delta n = -1.1\\pm 1.1$, as the derived tau model.} \\item{Maximum valid tau to $z=0$: In this model of evolution we use the largest value of tau derived for our galaxy sample, $\\tau = 2.71\\times10^{9} \\rm{yr}$. We apply this exponentially declining rate to all the galaxies in the sample. In this scenario we obtain a large average increase in total stellar mass of the sample of $377\\pm172\\%$. The change in the effective radii of this model is on average $R_{e} = 57\\pm33\\%$ a factor of $\\sim3.8$ larger than the derived tau model of evolution. This increase in effective radius is still only $\\sim11-19\\%$ of the observed size evolution. The change in $n$ for this model, $\\Delta n = -1.5\\pm1.7$ is similar to change for the derived tau model.} \\item{Minimum valid tau to $z=0$: This model is similar to the previous model. Except that the minimum valid tau, $\\tau = 1.2\\times10^{8} \\rm{yr}$, is used to extrapolate the SF. This would give the shortest time scale that the SF would occur. In this model the average galaxy in the sample increases its stellar mass by only $\\sim17\\%$. This very small increase in mass is accompanied by an equally small change in $R_{e}$, average $\\Delta R_{e} = 3\\pm1\\%$, and $n$, $\\Delta n = -0.4 \\pm 0.6$. } \\end{itemize} From this investigation of different models of SF evolution to $z=0$ we find that the value in the increase of the effective radii of the massive galaxies can at no point fully explain the total observed size increase. The valid models of SF evolution that we applied can only produce a factor of $\\sim 3.8$ times larger than the size increased we obtained from using the derived tau model at maximum. The change in S\\'{e}rsic index in all the models are within the error consistent with the answer obtained from the derived tau model used in this paper. \\subsubsection{Dust Gradients} In this paper we assume that the dust obscuration is constant across the radius of individual galaxies. From studies of local and distant studies this may not be the case. Colour gradients in the local universe have been shown to correspond to age and dust gradients (e.g. Boquien et al. 2011; Smith et al. 2012). We apply a dust gradient to our sample of massive galaxies that allows the attenuation due to dust to vary within the given error across each galaxy. This is done in two ways. A positive dust gradient with higher attenuation towards the outer regions of the galaxy, and a negative dust gradient with higher dust attenuation towards the central regions of the galaxy. In the positive gradient case we find that the average increase in the effective radius was $68\\pm36\\%$ larger than the original measured effective radius. This is a factor of $\\sim4.5$ larger change than the growth in $R_{e}$ we obtain from using a radially constant dust correction. From this gradient the change in $n$ is largely the same as before but with a much larger scatter, $\\Delta n = -0.9\\pm2.0$. The positive gradient could contribute a maximum of $\\sim23\\%$ to the $300-500\\%$ size growth. In the negative gradient case we find that the average increase in $R_{e}$ is minimal, $\\Delta R_{e}=7\\pm3\\%$. This small increase in the effective radius is accompanied by a change in $n$ that is very similar to most other cases, $\\Delta n = -1.0\\pm1.0$. This negative gradient case would seem to produce a very small increase in the effective radii of our sample and only contribute a maximum of $\\sim2\\%$ to the total observed size growth. Neither of the gradient cases that we applied to the sample are able to fully explain the observed size growth or observed change in S\\'{e}rsic index." }, "1207/1207.3798_arXiv.txt": { "abstract": "\\let\\thefootnote\\relax\\footnotetext{michael.koehn@aei.mpg.de,~~~jlehners@aei.mpg.de,~~~ovrut@elcapitan.hep.upenn.edu} We construct $\\mac{N}=1$ supergravity extensions of scalar field theories with higher-derivative kinetic terms. Special attention is paid to the auxiliary fields, whose elimination leads not only to corrections to the kinetic terms, but to new expressions for the potential energy as well. For example, a potential energy can be generated even in the absence of a superpotential. Our formalism allows one to write a supergravity extension of any higher-derivative scalar field theory and, therefore, has applications to both particle physics and cosmological model building. As an illustration, we couple the higher-derivative DBI action describing a 3-brane in 6-dimensions to $\\mac{N}=1$ supergravity. This displays a number of new features-- including the fact that, in the regime where the higher-derivative kinetic terms become important, the potential tends to be everywhere negative. \\vspace{.3in} \\noindent ", "introduction": "Since its discovery \\cite{Golfand:1971iw,Volkov:1973ix,Wess:1974tw}, supersymmetry has been investigated with enthusiasm by theoretical physicists. Representations of the supersymmetry algebra contain bosonic and fermionic degrees of freedom in equal numbers. Moreover, particles belonging to the same representation have equal mass. Since superpartners with the same mass as conventional particles have not been observed, four-dimensional supersymmetry cannot be an unbroken low energy symmetry. Nevertheless, there are good reasons to take seriously the idea that supersymmetry-- particularly four-dimensional ${\\cal{N}}=1$ supersymmetry --might be relevant at higher energies. For example, when ${\\cal{N}}=1$ supersymmetry is taken into account, the gauge couplings of the electroweak and strong forces unite to good precision at high energies \\cite{Langacker:1983cj}, suggesting the existence of supersymmetric grand unification. Moreover, supersymmetric theories enjoy special finiteness properties that help to explain the hierarchy between the electroweak and the unification/gravitational scales \\cite{Dimopoulos:1981zb,Dimopoulos:1981yj}. Last, but not least, ${\\cal{N}}=1$ supersymmetry is a central feature of phenomenologically realistic string theories--see, for example \\cite{Braun:2005nv,Lukas:1998yy}. All of this motivates studying early universe cosmology within the context of ${\\cal{N}}=1$ supersymmetry. Since cosmology quintessentially involves gravitation, such theories must be constructed using ``local'' supersymmetry-- that is, ${\\cal{N}}=1$ supergravity --and not the ``global'' supersymmetry of low energy particle physics models. This has been done within the context of two-derivative kinetic theories, both in local quantum field theory and superstrings. More recently, however, it has become clear that higher-derivative theories of cosmology are potentially important. These include so-called DBI inflation \\cite{Silverstein:2003hf}, ekpyrotic theories with brane collisions \\cite{Khoury:2001wf,Khoury:2001bz} and ghost-condensation \\cite{Buchbinder:2007ad,Buchbinder:2007tw,Buchbinder:2007at}, as well as other cosmologies constructed on the worldvolume of three-branes \\cite{Khoury:2012dn,Ovrut:2012wn}. Motivated by this, in this paper we will develop a framework for constructing higher-derivative kinetic theories of chiral superfields coupled to ${\\cal{N}}=1$ supergravity. As a first application of this formalism, we present an example of supergravitational DBI inflation. This paper builds on previous work \\cite{Khoury:2010gb,Khoury:2011da} on globally supersymmetric higher-derivative scalar field theories-- extending it to local ${\\cal{N}}=1$ supergravity. We first construct a supergravity version of $(\\pt\\phi)^4$, the square of the usual kinetic energy of a real scalar field. In the present work, we neglect fermions because a) they are typically unimportant in models of early universe cosmology and b) since their inclusion greatly complicates all equations. Instead, we focus on the physics of the scalar bosons and the associated auxiliary fields. We will present the fermionic terms, and discuss their role, in forthcoming publications \\cite{forthcoming}. When the fermions are set to zero, our supergravity extension of $(\\pt\\phi)^4$ has a special-- perhaps unique --property; namely, it can be multiplied by an arbitrary function of the scalar fields and their spacetime derivatives, while not altering the pure supergravity sector of the Lagrangian. Because this multiplicative factor is arbitrary, our formalism allows one to write a supergravity extension of {\\it any} higher-derivative Lagrangian built out of scalar fields and their spacetime derivatives. As always in supergravity, a special role is played by the auxiliary fields. In this paper, we devote considerable attention to their properties. In ordinary two-derivative chiral supergravity, elimination of the auxiliary fields leads to a well-known formula for the potential $V$. In terms of a K\\\"{a}hler potential $K$ and superpotential $W$ \\cite{Cremmer:1978hn}, this is given by \\be V = e^K \\left( K^{,A^{i}A^{j*}} |D_{A^{i}} W|^2 - 3 |W|^2 \\right) \\ , \\ee where $A^{i}$ denotes the complex scalar component of a chiral supermultiplet. In higher-derivative supergravity theories, we find two generic differences. First, the elimination of the auxiliary fields leads to corrections to the above formula. When the higher-derivative terms are important, these corrections can be significant, drastically modifying the dynamics. The second property is that the equation of motion for the auxiliary field $F^{i}$ of a chiral multiplet is now a cubic equation-- whereas previously it was linear. Thus, in general it admits three distinct solutions, which, after substituting back into the Lagrangian, lead to three inequivalent theories. In this paper, we present the basic properties of each of these three branches. The bulk of the paper presents our general formalism. It is useful, therefore, to give an explicit example-- which we do by constructing the supergravity extension of a particular DBI action. This allows us to display the specific corrections to both the kinetic and potential terms induced by the elimination of the auxiliary fields when higher-derivative terms are present. We also analyze one of the new branches of the supergravity DBI theory, commenting on the implications of our results for models of DBI inflation. In particular, we find that in the relativistic regime of the DBI theory, the potential automatically becomes negative-- rendering inflation impossible. These findings illustrate the significance that the auxiliary fields can have on the dynamics of a given model. There are many potential applications of our results, particularly in early universe cosmology. For example, cosmological models that are constructed in-- or inspired by --string theory should admit an effective ${\\cal{N}}=1$ supergravity description in four-dimensions. These theories typically have scalar fields arising as the moduli associated with branes \\cite{Donagi:1999jp}, flux \\cite{Buchbinder:2002ji,Buchbinder:2002pr} or the compactification manifold. For most-- if not all --of these models, whether they are of DBI inflation \\cite{Silverstein:2003hf}, $k$-inflation \\cite{ArmendarizPicon:1999rj}, $k$-essence \\cite{ArmendarizPicon:1999rj}, ekpyrotic/cyclic cosmology \\cite{Khoury:2001wf,Steinhardt:2001st,Buchbinder:2007ad,Lehners:2008vx}, effective theories of Galileons \\cite{Trodden:2011xh} or higher-derivative induced cosmic bounces \\cite{ArkaniHamed:2003uy,Creminelli:2006xe,Lehners:2011kr}, the proper setting is supergravity-- and all contain phases where the dynamic description includes scalar higher-derivative terms. We hope to apply our formalism to these models in the future. The plan of the paper is the following. We begin in Section \\ref{sectionFlat} by reviewing the construction of higher-derivative kinetic terms for chiral multiplets in global supersymmetry; that is, when gravity is neglected. Then, in Section \\ref{sectionX2}, it is shown how this construction can be generalized to supergravity. We proceed by eliminating the auxiliary fields one by one, beginning with $b_{m}$ and $M$ of pure supergravity. . The auxiliary fields $F^i$ of the chiral multiplets require special attention, and Section \\ref{sectionauxiliary} is devoted to them. In Section \\ref{sectionDBI} we apply our formalism to an example of the DBI action. For the benefit of the reader, we include short summaries of our results at the end of each subsection in \\ref{sectionauxiliary} and \\ref{sectionDBI} . After concluding in Section \\ref{sectionconclusion}, we add Appendices describing the difference of our formalism with the framework of Baumann and Green \\cite{Baumann:2011nk,Baumann:2011nm}, as well as comments on K\\\"{a}hler invariance in the present context. The notation and conventions of the book by J. Wess and J. Bagger \\cite{Wess:1992cp} are used throughout the paper. ", "conclusions": "\\label{sectionconclusion} In this paper, we presented a formalism that allows one to obtain an $\\mac{N}=1$ supergravity extension of any scalar field theory with higher-derivative kinetic terms. This was accomplished by constructing a superfield-- quartic in chiral scalars --which contains the term $(\\pt\\phi)^4$ and, when the fermions are set to zero, consists entirely of its top component. Thus, when multiplied by any other superfield, the resulting Lagrangian contains only the lowest component of the multiplicative factor. This property enables one to directly construct a supergravity extension any higher-derivative scalar field term of interest. Moreover, as discussed in the Appendix, our supergravity extension of $(\\pt\\phi)^4$ is likely to be the unique one that does not modify the gravitational sector of the theory-- thus rendering our construction particularly pertinent. For this reason, studying the properties of the auxiliary fields in this context, which are crucial to the structure of supergravity, is important. This was carried out, in detail, in this paper. In our formalism, despite the inclusion of an arbitrarily high number of spacetime derivatives, the auxiliary fields do not have kinetic terms and, therefore, continue to satisfy algebraic equations of motion. We point out that this is a highly non-trivial property, which renders the treatment of the auxiliary fields straightforward. Be this as it may, there is one new, and important, property of our formalism. That is, although the auxiliary fields $F$ satisfy an algebraic equation of motion, that equation is now cubic-- as opposed to the linear equation in the usual second order kinetic theory. Hence, this equation admits up to three distinct solutions. We have shown that these solutions lead to different theories that cannot dynamically transition from one to another. One solution is directly related to the one ordinarily obtained in the absence of higher-derivative terms. This leads to corrections to both the kinetic and potential terms when substituted into the action. We have examined these corrections in different limits. When the higher-derivative terms are small, the corrections are correspondingly small, but need to be taken into account when making precise predictions in phenomenology and cosmology. In the limit that the higher-derivative terms become large, the effect of eliminating the auxiliary field is to suppress certain contributions to the potential. The result is that the negative term $-3e^K|W|^2$ becomes the dominant contribution to the potential energy. Thus, in the large higher-derivative limit, supergravity manifests once more its predilection for negative potentials. This feature implies that the supergravity implementation of inflationary and $k$-essence models-- such as DBI inflation --that rely on higher-derivative kinetic terms in an essential way become more challenging. In addition to this ``usual'' solution for $F$, there exist up to two new solutions. These lead to theories with very unusual properties, which we have only started exploring in the present paper. For example, these new branches seem to prefer solutions with substantial spatial gradients in the scalar fields, and can lead to positive potentials. Moreover they can do this even in the absence of a superpotential. These curious theories, whose physical relevance is not clear yet, form an interesting topic for further research. This work has many foreseeable applications. Most importantly, we hope that our results can be used to bridge the gap between standard model building in cosmology and full-blown string compactifications, leading to well-motivated effective theories of early universe dynamics. In this context, it will be interesting to investigate in more detail models of DBI inflation and $k$-inflation, as well as other models of brane dynamics such as the Galileons and their extensions. Furthermore, it will be enlightening to find out whether null energy violating models, such as the ghost condensate, can be realized in a supergravity context. We hope to explore these topics in the near future." }, "1207/1207.3738.txt": { "abstract": "We look ahead from the frontiers of research on ice dynamics in its broadest sense; on the structures of ice, the patterns or morphologies it may assume, and the physical and chemical processes in which it is involved. We highlight open questions in the various fields of ice research in nature; ranging from terrestrial and oceanic ice on Earth, to ice in the atmosphere, to ice on other solar system bodies and in interstellar space. ", "introduction": "\\begin{quote} The ice was here, the ice was there, \\\\ The ice was all around Samuel Taylor Coleridge, \\\\ \\emph{The Rime of the Ancient Mariner} \\end{quote} Ice is indeed all around us. As the cryosphere, ice or snow covers a small but significant part of the Earth's surface, both land and sea, and it plays a similarly important role in our atmosphere. Moreover ice is present on many other celestial bodies in our solar system and --- surely --- beyond, and it coats grains of dust in interstellar space. Ice is not a static medium but a dynamical one; it shows strong variations of its characteristics with time and place, as we may readily experience at a human scale on any ski slope. A better understanding of ice structures, patterns, and processes is thus a topic of current research in physics. We shall show in the following how progress in understanding these questions is elemental in understanding current questions in astrophysics, atmospheric, cryospheric, and environmental science. Ice research questions are not only tackled separately within distinct fields --- in terrestrial, oceanic, atmospheric, planetary, and interstellar ice research --- but also by researchers with disparate backgrounds: by modelers, field and laboratory experimentalists and theoreticians from both physics and chemistry. We work in these different fields and come from a variety of backgrounds. We came together, some of us initially for a Spanish national project supported by the Spanish CSIC, and then the majority of us for a workshop, Euroice 2008, sponsored by the European Science Foundation, which was organized to connect people working on structures and those working in applied ice fields; to find common ground in the physical and chemical processes at icy surfaces and the physics and chemistry of ice structures from the molecular scale to the macroscale, and to explore whether some of the questions we were asking and some of the answers we were seeking are the same. During the workshop we found that, despite the diversity of ice research, a number of key themes are indeed common between the different fields. The common ground in any field of ice research is the urge to understand better its structure and dynamics. For example: What are the ordering mechanisms of ice as it changes from one of its phases into another? What is the structure at its surface, and how does this differ from the bulk? What is the structure and microenvironment at the contact area of ice crystals? How does ice structure form initially? Are there meta-stable phases present in the environment? This work focuses on this common ground. It thus does not aim to be a comprehensive review; such a review of ice physics and chemistry would be a book; indeed there are excellent books available \\cite{hobbs1974,petrenko1999}. However, many of the issues raised in this article are issues of the 21st century that are not addressed in the textbooks. This article provides a view of the way ahead from some frontiers of research on ice. We set out the main physical and chemical open questions on ice structures, patterns and processes from the fields of ice research in nature: from ice on Earth, in the oceans and the atmosphere, to planetary and interstellar ice. We begin in Section~\\ref{structures} by introducing open questions in the molecular structures of ices; we then examine open issues on dynamical patterns and processes in ice. Following this we look first, in Section~\\ref{astrophysical}, at astrophysical ice. We then focus on ice on Earth, beginning with Section~\\ref{atmospheric} on atmospheric ice, whose precipitation leads to the subsequent formation of terrestrial ice, Section~\\ref{terrestrial}, and, Section~\\ref{oceanic}, sea ice. We conclude in Section~\\ref{perspectives}, in which we attempt to list the important open questions on ice from the perspectives of these different fields of ice research. ", "conclusions": "" }, "1207/1207.5800_arXiv.txt": { "abstract": "{The cosmological peculiar velocity field (deviations from the pure Hubble flow) of matter carries significant information on dark energy, dark matter and the underlying theory of gravity on large scales. Peculiar motions of galaxies introduce systematic deviations between the observed galaxy redshifts $z$ and the corresponding cosmological redshifts $z_{_{\\rm cos}}$. A novel method for estimating the angular power spectrum of the peculiar velocity field based on observations of galaxy redshifts and apparent magnitudes $m$ (or equivalently fluxes) is presented. This method exploits the fact that a mean relation between $z_{_{\\rm cos}}$ and $m$ of galaxies can be derived from all galaxies in a redshift-magnitude survey. Given a galaxy magnitude, it is shown that the $z_{_{\\rm cos}}(m)$ relation yields its cosmological redshift with a $1\\sigma $ error of $\\sigma_z\\sim 0.3$ for a survey like Euclid ($\\sim 10^9$ galaxies at $z\\lesssim 2$), and can be used to constrain the angular power spectrum of $z-z_{_{\\rm cos}}(m)$ with a high signal-to-noise ratio. At large angular separations corresponding to $l\\lesssim 15$, we obtain significant constraints on the power spectrum of the peculiar velocity field. At $15 \\lesssim l\\lesssim 60$, magnitude shifts in the $z_{_{\\rm cos}}(m)$ relation caused by gravitational lensing magnification dominate, allowing us to probe the line-of-sight integral of the gravitational potential. Effects related to the environmental dependence in the luminosity function can easily be computed and their contamination removed from the estimated power spectra. The amplitude of the combined velocity and lensing power spectra at $z\\sim 1$ can be measured with $\\lesssim 5\\%$ accuracy.} ", "introduction": "\\label{sec:int} Dark matter, dark energy, and the theory of gravitation dictate the evolution of large-scale structure in the Universe. The physical conditions allowing for the formation of galaxies, ultimately lead to a bias between the distribution of galaxies and the underlying mass density. However, large-scale motions of galaxies are most certainly locked to the peculiar velocity field associated with the gravitational tug of the total underlying mass fluctuations. This assumes that gravity is the only relevant large-scale force and neglects the contribution from decaying linear modes. Despite the bias between the distribution of galaxies and the underlying matter field, the clustering properties of galaxies have been the main tool for testing cosmological models. Currently planned galaxy surveys will allow us to quantify the clustering of galaxies on hundreds of comoving Mpcs and to even measure coherent distortions of galaxy images which arise from gravitational lensing by the foreground matter. On the other hand, the peculiar motions of galaxies have traditionally been less successful as a cosmological tool, and there are several reasons for that. Neglecting other potentially important effects (see section \\ref{method} for details), peculiar velocities are approximately equal to the redshifts less the corresponding Hubble expansion recession velocities. The latter require direct measurements of galaxy distances which are available for only a small fraction of galaxies. Presently, the number of galaxies with measured distances is several orders of magnitude below that of galaxies in redshift surveys used for clustering studies. Furthermore, although the underlying peculiar velocity is an honest tracer of the general matter flow, the inference of peculiar velocities from observational data is plagued with observational biases \\cite{lyn88}. Traditional peculiar velocity catalogs are expected to improve within the next few years, but it remains questionable how well observational biases will be controlled, especially at large distances. An alternative probe of the peculiar velocity field may be astrometric measurements of galaxies by the Gaia space mission \\cite{perryman01,nbdgaia}. This probe is essentially free of the classic biases contaminating traditional peculiar velocity measurements, but it is also limited to nearby galaxies within $\\sim 100\\hmpc$. Here we describe a method for deriving strong constraints on the power spectrum of the galaxies' peculiar velocity field independent of conventional direct distance measurements, which are prone to systematic errors, and any biasing relation between galaxies and mass. The method is an extension of the approaches we have recently proposed \\citep{NBDL,ND11a,BDN12}, and it relies on using the observed fluxes of galaxies as a proxy for their cosmological distance \\cite{TYS1979}. Although this most basic distance indicator is very noisy, the large number of galaxies available in future surveys will allow one to beat down this noise to a sufficiently low level. Planned galaxy redshift surveys such as Euclid \\cite{euclidL,EuclidRB} will probe the structure of the Universe over thousands of (comoving) Mpcs, comprising $\\lesssim 10^8$--$10^9$ galaxies at $z\\sim 1$ and beyond. These observations will provide redshifts and fluxes, $z$ and $f$, respectively. The observed redshifts deviate from the cosmological redshifts which would be observed in a purely homogeneous universe. The large number of galaxies and the large sky coverage of these surveys can be used to derive a mean global relation between the mean redshift of a galaxy and its apparent magnitude $m=-2.5 \\log f + {\\rm const}$. We interpret this relation as yielding the cosmological redshift $z_{_{\\rm cos}}(m)$ for a given apparent magnitude $m$. Angular power spectra of the difference $z_i- z_{_{\\rm cos}}(m_i)$ between the observed redshift $z_i$ of a galaxy and its expected cosmological redshift $z_{_{\\rm cos}}$ should contain valuable information, mainly on the peculiar velocity field which is the main cosmological source for $z_i- z_{_{\\rm cos}}(m_i)$. In this work, we will show that the velocity power spectrum on large scales of a few 100 Mpcs could be constrained with significant signal-to-noise ratio ($S/N$) at the effective depth of the survey. Another contribution to $z_i- z_{_{\\rm cos}}(m_i)$ results from the time evolution of the gravitational potential along the photon path, but is significantly smaller than that induced by peculiar velocities as we will show below. There are two additional, indirect effects which modify the relation $z_{_{\\rm cos}}(m)$ along a given line of sight. The first effect is related to the environmental dependence between galaxy luminosities and the large-scale structure in which they reside. Since this dependence is closely connected to the underlying density field, it can be self-consistently removed in our analysis. The second effect is caused by gravitational lensing magnification which changes the apparent magnitudes of galaxies in a given direction. This latter contribution could actually be very rewarding since gravitational lensing provides a direct probe of the underlying mass distribution. Considering the analysis presented below, we will therefore treat it as part of the sought signal. The paper is structured as follows: We begin with a detailed description of the method and its application to galaxy redshift surveys in section \\ref{method}. In section \\ref{sec:tpk}, we consider predictions for the standard $\\Lambda$CDM model and discuss the method's viability as well as its expected performance. Finally, we present our conclusions in section \\ref{sec:cnl}. For clarity, some of the technical material is given separately in an appendix. In the following, we adopt the standard notation. The matter density and the cosmological constant in units of the critical density are denoted by $\\Omega$ and $\\Lambda$, respectively. The scale factor $a$ is normalized to unity at the present time ($t=t_0$), and the Hubble function is defined as $H=\\dot{a}/a$. Further, $r=c \\int_{t}^{t_0} dt'/a(t')$ will be the comoving distance to an object and $z$ its corresponding redshift, assuming a homogeneous and isotropic cosmological background. Throughout the paper, the subscript ``$0$'' will refer to quantities given at $t=t_0$, and a dot symbol denotes partial derivatives with respect to time $t$, i.e. $\\dot{A}\\equiv\\partial A/\\partial t$. ", "conclusions": "\\label{sec:cnl} In this paper, we have presented a novel method for deriving direct constraints on the peculiar velocity and gravitational potential power spectra from currently planned galaxy redshift surveys. The large number of galaxies with photometric redshifts in these surveys allows one to exploit apparent galaxy magnitudes as a proxy for their cosmological redshifts since it beats down the large scatter in the $\\zcos(m)$ relation and the uncertainty in the photometric redshifts. The method aims at directly constraining power spectra of the underlying fluctuation fields independent of the way galaxies trace mass. Other methods for extracting cosmological information from redshift surveys rely on accurate measurements of the galaxy power spectrum in redshift space (since galaxy distances remain unknown). The power spectrum and other statistical measures based on the distribution of galaxies have been successful at probing the nature of dark matter and placing important constraints on neutrino masses \\cite{tegmark04,deputter12,sanchez12}. Having said that, however, they depend on a very accurate knowledge of the relation between galaxies and the full underlying matter distribution. The method we have considered here is less precise, but it is completely independent of the galaxy formation process and offers a much more sensitive assessment of the underlying physical mechanism driving cosmic acceleration and structure formation. This approach is particularly worthwhile if such constraints on the velocity field and the gravitational potential are contrasted with local constraints obtained from data at low redshifts. For example, peculiar motions of galaxies within a distance of $\\sim 100\\hmpc$ can be measured using tight relations between intrinsic observables of galaxies \\cite[e.g.][]{TF77,springsixdf}, and also using astrometric observations of the Gaia space mission which is currently scheduled for launch in $2013$. These peculiar motions of galaxies have been useful for constraining cosmological parameters \\citep{DN10} as well as the amplitude of the velocity field in the nearby Universe \\citep{ND11a,bilicki11}. Although we have presented predictions for the Euclid survey, the science proposed in this paper will not have to wait for this space mission. In fact, several ground-based photometric surveys in the optical and near-infrared bands, which will constitute the backbone of Euclid's photometry, will provide photometric redshift catalogs that can be used for our purposes well before the launch of the satellite. On a shorter timescale, the VLT Survey Telescope (VST) will be used to carry out the Kilo Degree Survey (KiDS), one of the ESO public surveys. It will cover 1,500 deg$^2$ to $u=24$, $g=24.6$, $r=24.4$, and $i=23.1$, and will probably contain $\\sim 10^{8}$ galaxies with measured photometric redshifts.\\footnote{\\protect\\url{http://www.eso.org/public/teles-instr/surveytelescopes/vst.html}} Also, the Dark Energy Survey (DES) will start its operations soon, and it will cover 5000 deg$^2$ of the Southern sky within 5 years, reaching magnitudes up to $\\sim 24$ in SDSS $griz$ filters, comparable to the limiting magnitudes of Euclid and with a redshift distribution $dN/dz$ similar to that of Euclid galaxies. DES will measure photometric redshifts of $\\sim 3\\times 10^8 $ galaxies with $\\sigma_{\\rm photo}\\sim 0.12$ at $z\\sim 1$.\\footnote{\\protect\\url{http://www.darkenergysurvey.org/reports/proposal-standalone.ps}} Furthermore, the first of the four planned Pan-STARRS telescopes has been operational since May 2010.\\footnote{\\protect\\url{http://pan-starrs.ifa.hawaii.edu/public/home.html}} The planned 3$\\pi$ area of the sky will be considerably shallower, detecting galaxies below a limiting magnitude of $\\sim 24$ in the $griz$ bands. A deeper survey involving the PS1 and PS2 telescopes is currently being planned. The survey should cover $\\sim 7,500 \\rm deg^2$ with limiting fluxes $g = 24.7$, $r = 24.3$, $i = 24.1$, and $z = 23.6$. Photometric redshifts will then be measured for $\\sim 4.5 \\times 10^8 $ galaxies with similar errors. Finally, on the long run, the Large Synoptic Survey Telescope (LSST) is expected to start operations in 2020. Its main deep-wide-fast survey is expected to observe $\\sim 20,000$ deg$^2$ in the $ugrizy$ bands. After about 10 years of operation, it will reach much deeper depths (down to a co-added magnitude $r=27$), detecting about $3 \\times 10^9$ galaxies \\cite{ivezic08}." }, "1207/1207.3171_arXiv.txt": { "abstract": "{ Close-in planetary systems detected by the Kepler mission present an excess of periods ratio that are just slightly larger than some low order resonant values. This feature occurs naturally when resonant couples undergo dissipation that damps the eccentricities. However, the resonant angles appear to librate at the end of the migration process, which is often believed to be an evidence that the systems remain in resonance. Here we provide an analytical model for the dissipation in resonant planetary systems valid for low eccentricities. We confirm that dissipation accounts for an excess of pairs that lie just aside from the nominal periods ratios, as observed by the Kepler mission. In addition, by a global analysis of the phase space of the problem, we demonstrate that these final pairs are non-resonant. Indeed, the separatrices that exist in the resonant systems disappear with the dissipation, and remains only a circulation of the orbits around a single elliptical fixed point. Furthermore, the apparent libration of the resonant angles can be explained using the classical secular averaging method. We show that this artifact is only due to the severe damping of the amplitudes of the eigenmodes in the secular motion. } ", "introduction": "Dissipation due to tidal interactions is a possible mechanism that explains the abundance of planetary systems that lie near but not at exact mean-motion commensurability. \\citet{papaloizou_dynamics_2010} (in the case of three planets Laplace resonances) and \\citet{papaloizou_tidal_2011} (for two planets resonances) showed that planets that have been temporarily locked in resonance due to differential migration could have their periods ratios to depart from strict commensurability due to the circularization of their orbits by tidal interactions with the star. More recently, \\citet{batygin_dissipative_2012} and \\citet{lithwick_resonant_2012} use similar effects to explain the excess of systems of two planets that lie near resonances but with planets slightly farther from each other than the nominal mean-motion commensurability ratio. One of the most intriguing features that is common in these different studies is the observation that resonant angles continue to librate far from exact commensurability and the authors leave unanswered the question of determining if these systems are in resonance or not. It seems important to clarify this point and to understand why resonant angles can librate so far from exact commensurability. Our analysis is based on the study of the phase space of two planets in mean-motion resonance (MMR) in the conservative case (without any dissipation). This is the object of section \\ref{sec:conservative}. We pay a particular attention to apsidal corotation resonances (ACR) which play a major role in the dynamics of these systems and in the understanding of the topology of the phase space. ACR have been extensively studied both in the asteroidal restricted problem \\citep[e.g.][]{ferraz-mello_symmetrical_1993} and the planetary problem \\citep[e.g.][]{hadjidemetriou_2002,michtchenko_stationary_2006}. Most of the studies on the subject have been made using numerical (or semi-analytical) models that remain valid for arbitrary values of the eccentricities but which do not always provide a global picture of the dynamics. In the present study we are only concerned in the dynamics at low eccentricities since our aim is to understand the motion at the end of the circularization process. A completely analytical model is thus well suited in this case. Analytical studies of planetary MMR have already be done up to degree two in eccentricities in the cases of the 2:1 \\citep{callegari_dynamics_2004} and the 3:2 \\citep{callegari_dynamics_2006} MMR. The dissipative case (studied in section \\ref{sec:dissipation}) is modeled using the conservative case as a basis and very simple and general prescriptions for the dissipation. Our study is mainly aimed at understanding the impact of tides on the dynamics of resonant planets pairs and, in this case, we follow the prescriptions introduced by \\citet{papaloizou_tidal_2011}. However we show that the differential migration process that allows a resonant locking of both planets can also be accounted for in our model. This process has already been widely studied \\citep[e.g.][]{lee_dynamics_2002, ferraz-mello_evolution_2003,lee_diversity_2004,beauge_planetary_2006}, and is also considered in \\citet{papaloizou_dynamics_2010,papaloizou_tidal_2011}. Eventually, we treat this perturbation of the conservative case by following the lines of \\citet{laskar_tidal_2012}. In section \\ref{sec:secular} we show that the final state of the resonant systems that undergo a circularization process is very well characterized by the secular normal form. We explain, with the secular problem, why resonant angles appear to librate even far from resonances. Finally, we present in section \\ref{sec:simulation} the results of two numerical simulations that confirm and illustrate the different mechanisms that we highlight with our analytical model. ", "conclusions": "In this work, we presented a study of planar resonant planetary systems in the conservative case and in presence of a dissipative force. Our main interest was to understand the dynamics at the end of a circularization process in resonant systems. We used a completely analytical model developed in power series of eccentricities which is well suited for the study of these low eccentricity systems. Before introducing the dissipative force in the model we characterized the dynamics in the conservative case. In particular, we highlighted the fact that apsidal corotation resonances (ACR) are a powerful tool to understand the global dynamics of a system. Then, we showed that the introduction of a dissipative force\\footnote{ We consider here tidal effects, but it is clear that other mechanisms could result as well in similar dissipative effects.} in a resonant system has two main effects. On the short-term, the system is attracted toward the libration center if it initially relies in its vicinity. On the long-term, the system tends to follow this ACR solution outside the resonance and both planets tend to move away from each other. These two mechanisms are well illustrated and confirmed by the results of two simulations (one for the 2:1 MMR and the other for the 3:2 MMR) of the innermost planets of the \\object{GJ581} (see section~\\ref{sec:simulation}). Since the ACR solution do not correspond to zero eccentricities even far from the resonance, it is possible to have resonant angles to oscillate outside of resonances. However, we showed that in this case the motion is completely characterized by the secular problem and that the fact that resonant angles appear to librate only means that the secular eigenmodes are (almost) totally damped. The important fact is that the nature of the motion is the same as when the eigenmodes are not damped and \\modif{the separatrix of the resonance does not exist anymore in this region of the phase space.} Thus, it is inappropriate to consider such systems as resonant ones." }, "1207/1207.3989_arXiv.txt": { "abstract": "\\igr is a hard X-ray binary transient discovered recently by \\emph{INTEGRAL}. Here we report on detailed timing and spectral analysis on IGR J18179-1621 in X-rays based on available \\emph{INTEGRAL} and \\emph{Swift} data. From the \\emph{INTEGRAL} analysis, \\igr is detected with a significance of 21.6~$\\sigma$ in the 18--40 keV band by \\emph{ISGRI} and 15.3~$\\sigma$ in the 3--25 keV band by \\emph{JEM-X}, between 2012-02-29 and 2012-03-01. We analyze two quasi-simultaneous \\emph{Swift} ToO observations. A clear 11.82 seconds pulsation is detected above the white noise at a confidence level larger than $99.99\\%$. The pulse fraction is estimated as $22\\pm8\\%$ in 0.2-10 keV. No sign of pulsation is detected by \\emph{INTEGRAL/ISGRI} in the 18--40 keV band. With \\emph{Swift} and \\emph{INTEGRAL} spectra combined in soft and hard X-rays, \\igr could be fitted by an absorbed power law with a high energy cutoff plus a Gaussian absorption line centered at 21.5 keV. An additional absorption intrinsic to the source is found, while the absorption line is evidence for most probably originated from cyclotron resonant scattering and suggests a magnetic field in the emitting region of $\\sim$ $2.4\\times10^{12}$ Gauss. ", "introduction": "One of the most effective techniques to estimate the strength of the magnetic field of a neutron star is the detection of cyclotron resonant scattering features (CRSF) in its X-ray spectrum. The fundamental electron cyclotron resonance energy is $E=11.6\\; B_{12}(1+z_g)^{-1}$ keV, where $B_{12}$ is the magnetic field strength of the neutron star in units of $10^{12}$ G, and $z_g$ is the gravitational redshift. The surface magnetic field of neutron stars in accreting X-ray binaries can then be determined through the observation of the CRSF, which shows absorption lines at the fundamental electron cyclotron resonance and its high-energy harmonics. So far, CRSF are identified in 17 X-ray binaries, and show hints in another 10 (Pottschmidt et al. 2011, Makishima et al 1999, Coburn et al 2001, Yamamoto et al 2011). All the X-ray binaries with identified or possible CRSF are high mass X-ray binaries, except for 4U 1626-67, and all host X-ray pulsars. 7 out of the 10 candidates host X-ray pulsars too, with the exceptions of XTE J1739-302, 4U 1700-377, and IGR 16318-4848, for which pulsations are not yet detected. The magnetic field strength of CRSF-identified X-ray binaries cluster in a relatively narrow range of $(1.1-6.2)\\times10^{12}$ Gauss (assuming the typical neutron star parameters, $z_g$ $\\sim$0.3). \\igr is a newly discovered hard X-ray binary transient found by \\emph{INTEGRAL} during inner Galactic disk observations performed on 2012-02-29 -- MJD 55986 (see the ATel by Tuerler et al. 2012). The significant detection of \\igr by \\emph{INTEGRAL} reveals an absorption line $\\sim$20.8 keV, which may result from cyclotron resonant scattering (Tuerler et al. 2012). In subsequent \\emph{Swift/XRT} and \\emph{Fermi/GBM} observations (see the ATels by Halpern et al. 2012, Li et al.2012, and Finger et al. 2012) a 11.82 s pulsation was discovered. The absorption line and the 11.82~s pulsation suggested that \\igr is a high mass X-ray binary hosting a pulsar. Due to the overlap between the \\emph{Swift} position of the X-ray source and 2MASS J18175218-1621316, Li et al. (2012) proposed the latter as the infrared counterpart of IGR J18179-1621. Such a correlation is compatible with the \\emph{Chandra} determination of the position, later obtained by Paizis et al. (2012). Here, we report on spectra and timing analysis of the \\emph{INTEGRAL} observations as well as on two quasi-simultaneous \\emph{Swift} ToO observations. ", "conclusions": "We present timing and spectral analysis of \\igr from quasi-simultaneous observations made by \\emph{Swift/XRT} and \\emph{INTEGRAL}. A 11.82 s pulsation is discovered in \\emph{Swift/XRT} lightcurve in the 0.2--10 keV band, which is consistent with previous reports by Halpern et al. 2012, Li et al. 2012, and Finger et al. 2012. Because of more data being included in our analysis, a higher significance detection is obtained when compared with that of Tuerler et al. 2012. The 11.82~s pulsation is searched within \\emph{INTEGRAL/ISGRI} data in the 18--40 keV band, but results in a non-detection. Using quasi-simultaneous data from \\emph{Swift} and \\emph{INTEGRAL}, for the first time we carry out a combined spectral analysis in soft and hard X-rays. Above a continuum described by a $\\sim $0.6 power-law with a cut-off at $\\sim$10 keV, a CRSF is significantly detected at 21.5 keV, which is consistent with, but more precise and constrained than the results by Tuerler et al. 2012, obtained using \\emph{INTEGRAL} data only. Among those X-ray binaries which show CRSF and which are fitted by the same model (Coburn et al. 2002), \\igr have a moderate slope. \\ecut and \\efold are low comparing with other sources. 4U 0115+63 is a transient X-ray pulsar too with 3.6~s pulsation and 24-d orbit. It have a similar cutoff energy (10 keV) and folding energy (9.3 keV) than IGR J18179-1621. However, since little is known about IGR J18179-1621, we cannot draw further conclusions on the similarity between these two sources. Though IGR J18179-1621's 2.61 keV width of CRSF is not uncommon, its optical depth of 6.3 is the largest --about ten times that of other sources. A correlation between CRSF relative width (\\ensuremath{\\sigma_{\\rm{c}}}/\\ensuremath{E_{\\rm{c}}}) and optical depth (\\ensuremath{\\tau_{\\rm{c}}}) is observed (Coburn et al. 2002). However, because of the unusually large CRSF optical depth, \\igr is far from the correlation. On the contrary, \\igr follows another two correlations with other sources, CRSF width (\\ensuremath{\\sigma_{\\rm{c}}}) versus centroid energy (\\ensuremath{E_{\\rm{c}}}), and cutoff energy (\\ecut) versus centroid energy (\\ensuremath{E_{\\rm{c}}}) (Coburn et al. 2002). An absorption ($12.3\\times10^{22}$ cm$^{-2}$) which is much larger than the Galactic column density ($1.2\\times10^{22}$ cm$^{-2}$) at IGR J18179-1621's position is obtained. This possibly indicates an additional absorption intrinsic to the source. Li et al. 2012 proposed 2MASS J18175218-1621316 as the infrared counterpart for IGR J18179--1621. This source is obscured and it is well measured only in the Ks band: Ks magnitude=11.14. Both may indicate a complicated surrounding of IGR J18179--1621. Given the value of the spin period, the detection of an absorption feature compatiable with a CRSF, and the proposed optical counterpart, \\igr is most plausibly an accreting pulsar. In general, the transient behavior of X-ray binaries is powered by accretion of matter from the companion to the magnetic poles of the neutron star. The accretion flow onto the magnetic pole will be decelerated in a radiative shock above the neutron star surface when the luminosity reaches $10^{37}$ erg s$^{-1} $ (Basko $\\&$ Sunyaev 1976). Radiation will be modulated as a fan beam coming out from the bottom region of the shock and peaked perpendicular to the magnetic axis. In a lower accretion rate, while luminosity is less than $10^{37}$ erg s$^{-1} $, radiation will be formed into a pencil--beam, which the maximum direction is along the magnetic axis. During a transient outburst similar with the one lead to discovery of IGR J18179--1621, both the fan--beam and pencil beam will influence the pulse profiles in the lightcurve. In a high (low) accretion phase, the fan--beam (pencil) component is dominating the emission region, leading to a double (single) peak pulse profile. The transition point between these two different phases is $\\sim 10^{37}$ erg s$^{-1} $. A demonstration of the pulse profile transformation from two peaks to a single peak accompanying the luminosity evolution is seen in V0332+53 (see, e.g., Zhang et al. 2005). \\igr is characterized with a single peak pulse profile and its unabsorbed flux in 1.5--50 keV is $\\sim 1.3\\times10^{-9}$ erg cm$^{-2}$ s$^{-1}$. If we apply L $<$ $10^{37}$ erg s$^{-1} $ as the threshold of pulse profile shift, we obtain a upper limit on the \\igr distance at d $<$ 8 kpc. From the combined \\emph{Swift} and \\emph{INTEGRAL} spectra CRSF of \\igr is identified at 21.5 keV and no high energy harmonics are discovered. If the fundamental electron cyclotron resonance energy is 21.5 keV, this will indicate a magnetic field of $\\sim$ $2.4\\times10^{12}$ Gauss in the emitting region. Under a magnetic field of $\\sim$ $10^{12}$ Gauss, the ratio of cyclotron absorption coefficient between the fundamental absorption line and first harmonic is $\\sim 10^{-1}-10^{-2}$ (You et al. 1997), which means that the chance of producing the fundamental absorption line is 10--100 times larger than producing the first harmonic. In case of a relatively low accretion rate, there may not be sufficient electrons near the surface of the neutron star to produce the first harmonics in the spectrum. If the accretion rate is high and more electrons are available, the first harmonic may appear in the spectrum, and even the second or third harmonics might. \\igr only shows fundamental absorption line and there is no sight of any harmonics, which hints to a relatively low state of accretion. In addition to its single peak pulse profile, this is another indication that \\igr is not in a very high accretion state during this outburst." }, "1207/1207.3347_arXiv.txt": { "abstract": "We use N-body-spectro-photometric simulations to investigate the impact of incompleteness and incorrect redshifts in spectroscopic surveys to photometric redshift training and calibration and the resulting effects on cosmological parameter estimation from weak lensing shear-shear correlations. The photometry of the simulations is modeled after the upcoming Dark Energy Survey and the spectroscopy is based on a low/intermediate resolution spectrograph with wavelength coverage of $5500 {\\rm \\AA} <\\lambda< 9500 {\\rm \\AA}$. The principal systematic errors that such a spectroscopic follow-up encounters are incompleteness (inability to obtain spectroscopic redshifts for certain galaxies) and wrong redshifts. Encouragingly, we find that a neural network-based approach can effectively describe the spectroscopic incompleteness in terms of the galaxies' colors, so that the spectroscopic selection can be applied to the photometric sample. Hence, we find that spectroscopic incompleteness yields no appreciable biases to cosmology, although the statistical constraints degrade somewhat because the photometric survey has to be culled to match the spectroscopic selection. Unfortunately, wrong redshifts have a more severe impact: the cosmological biases are intolerable if more than a percent of the spectroscopic redshifts are incorrect. Moreover, we find that incorrect redshifts can also substantially degrade the accuracy of training set based photo-z estimators. The main problem is the difficulty of obtaining redshifts, either spectroscopically or photometrically, for objects at $z>1.3$. We discuss several approaches for reducing the cosmological biases, in particular finding that photo-z error estimators can reduce biases appreciably. ", "introduction": "\\label{sec:intro} Large-scale structure surveys benefit enormously from the information about galaxy redshifts. The redshift information reveals the third spatial dimension of a galaxy survey, enabling a much more accurate mapping of the expansion and growth history of the Universe relative to the case when only angular information is available. Unfortunately, obtaining spectroscopic redshifts for all galaxies is typically impossible in wide-field imaging surveys due to the large number ($\\sim 10^8$-$10^9$) of galaxies and the high cost of spectroscopy, especially for the high-redshift galaxies. To circumvent this problem, the current approach in the community is to estimate redshifts using photometric measurements, i.e. fluxes from a few broad band filters. These redshift estimates are known as photometric redshifts, or photo-zs, and are necessarily coarser than spectroscopic redshifts. Because of the intrinsically large errors, photo-zs typically cannot be used directly for cosmological analysis, unless the photo-z error distributions can be quantified precisely. The standard approach to quantify, or calibrate, the photo-z error distributions is to use a small subsample of galaxies with known redshifts. As discussed in detail in \\cite{cun12}, spectroscopic samples used to train photo-zs (cf. Sec. \\ref{sec:train}) need to be locally (in the space of observables) representative subsamples of the photometric samples. For calibration of the photo-z {\\it error distributions}, however, the spectroscopic sample must be globally representative. More specifically, the ideal spectroscopic survey should satisfy the following properties: \\begin{itemize} \\item {\\it Large area:} A spectroscopic survey needs to span a large area to beat down sample variance, and has to have tens of thousands of galaxies to beat down shot-noise in the photo-z error calibration \\citep{cun12}. In addition, the spectroscopic sample needs to be imaged under conditions that faithfully reproduce the variations in the full photometric sample \\citep[see e.g.][]{nak12}. Note that requirements might be alleviated with a correction to the individual galaxy redshift likelihoods \\citep{bor10,bor12}. In the context of dark energy parameter constraints, however, a full analysis that goes beyond the overall redshift distribution and involves the full error matrix $P(z_s|z_p)$ is required \\citep{BH10, Hearin}. \\item {\\it High completeness:} The spectroscopic survey needs to span the same range of redshifts, galaxy types, and other observational selection parameters as the photometric survey. When this is not possible, we say that the survey is {\\it incomplete}. In that case, the photometric survey has to be culled to ensure both surveys have matching selections. Alternatively, the galaxies in the spectroscopic survey can be weighted so as to reproduce the statistical properties of the photometric sample. Achieving high completeness in faint spectroscopic surveys is a major challenge.% \\item {\\it Few wrong redshifts:} We show in this paper that spectroscopic surveys need to have extremely accurate redshifts. As shown by many authors \\citep[e.g. ][]{ma06,hut06,Amara_Refregier_optimal,Abdalla08,Bernstein_Ma,Kitching_sys,Hearin} the photo-z calibration requires exquisite knowledge of the photo-z error distribution. Errors in the spectroscopic redshifts impair the characterization of the photo-z errors and severely degrade our ability to extract cosmological constraints from photometric surveys. \\end{itemize} For fixed observing resources, there is a conflict between accurate redshifts and completeness goals: as we stretch the observational limits (i.e.\\ by observing very faint galaxies) to sample redshifts that would mimic the distribution of the photometric sample, we increase the fraction of incorrect spectroscopic redshifts. As we will show, redshift accuracy is more important for the upcoming surveys. The purpose of this paper is to assess the impact of spectroscopic selection, i.e. completeness and accuracy, on the training and calibration of photometric redshifts and the resulting impact on cosmological constraints derived from weak lensing shear-shear correlations. To achieve this goal, we combine N-body, photometric and spectroscopic simulations patterned after the proposed characteristics of the Dark Energy Survey (DES) and expected spectroscopic follow-up. We then propagate the errors due to imperfect photo-z calibration on the cosmological parameter constraints inferred from the weak gravitational lensing power spectrum observations forecasted for the DES. The paper is organized as follows. In Sec. \\ref{sec:specintro} we provide a pedagogical introduction to the main issues driving completeness and accuracy of a spectroscopic sample. In Sec. \\ref{sec:data} we briefly describe the simulated catalogs we use, leaving the details of the catalog generation to Appendix \\ref{app:sims}. In Sec. \\ref{sec:meth} we give a step-by-step guide describing how we go from the simulated data to the cosmological constraints, detailing the methods used at each step. Results are presented in Sec. \\ref{sec:results}. We discuss the implications of our findings for spectroscopic survey design in Sec. \\ref{sec:design} and present conclusions in Sec. \\ref{sec:concl}. \\begin{figure*} \\centering \\includegraphics[scale=0.4]{plots/SpeczChart2.eps} \\caption{Flowchart describing our step-by-step procedure to go from the simulated observations to cosmological biases.} \\label{fig:chart} \\end{figure*} ", "conclusions": "\\label{sec:concl} We investigated the impact of spectroscopic failures on the training and calibration of photometric redshifts, and the consequent impact on the forecasted dark energy parameter constraints from weak gravitational lensing. Our tests were based on N-body/spectrophotometric simulations patterned after the DES and expected spectroscopic follow-up observations loosely patterned after the VVDS survey. Spectroscopic failures consist of two types of issues: the inability to obtain spectroscopic redshifts for certain galaxies, and incorrect redshifts. The inability to obtain redshifts introduces incompleteness in the spectroscopic sample --- i.e. missing redshifts in some region of parameter space (e.g.\\ at faint magnitudes) represented in the full photometric population of galaxies. This incompleteness must be accounted for before one can use the spectroscopic sample to calibrate photo-zs -- i.e characterize the photo-z error matrices, e.g. the $\\pzs$, of the sample. We studied two approaches to account for the incompleteness in the spectroscopic sample. In the first approach, we used an artificial neural network to estimate the spectroscopic selection function for the photometric sample. This selection function was then used to cull the photometric sample so that its statistical properties matched the spectroscopic sample. We found this approach works extremely well, yielding only insignificant bias in the WL constraints using the culled sample (refer to $\\ztrue$ column in Table \\ref{tab:wstats}). However, the statistical constraints did degrade substantially as, typically, a large fraction of the sample was culled. In the second approach, we accounted for the incompleteness in the spectroscopic sample by applying weights to the galaxies with spectroscopic redshifts, following the approach of \\cite{lim08}, so that the statistical properties of the spectroscopic and photometric samples match. This approach was also successful (cf. $\\ztrue$ column in Table \\ref{tab:wstatswei}) --- as expected, because most of the photometric sample could be used --- yielding tolerable cosmological biases while obtaining the maximum statistical constraints. Overall, we found that the effects of spectroscopic incompleteness are well under control. Unfortunately, on the other hand, we found that wrong redshifts can significantly degrade cosmological constraints and $>99\\%$ of correct spectroscopic redshifts seems to be needed (cf.\\ $\\ssrt$ and $\\zspec$ columns in Tables \\ref{tab:wstats} and \\ref{tab:wstatswei}). We found the results to be independent of the photo-z estimators used, but somewhat dependent on the settings of the spectroscopic pipeline. In particular, we found that attempts to increase the completeness of the spectroscopic sample during the spectral analysis can result in more catastrophic spectroscopic redshift failures, which will increase cosmological biases. We tested a couple of approaches to identify wrong spectroscopic redshifts, finding that the NNE error estimator \\citep{oya08b} is able to reduce the bias in the measured dark energy equation of state by half while removing only $10\\%$ of the photometric sample. Slightly less improvement in the $w$ bias was obtained using the template-fitting error estimator. In summary, we find that wrong redshifts are by far the main issue affecting calibration of photo-z error distributions with spectroscopic samples. Future follow-up spectroscopic observations of the planned and ongoing wide-area photometric surveys must focus primarily on the accuracy of the spectroscopic redshifts even if that implies sacrificing the spectroscopic completeness." }, "1207/1207.2388_arXiv.txt": { "abstract": "{We performed a uniform and detailed abundance analysis of 12 refractory elements (Na, Mg, Al, Si, Ca, Ti, Cr, Ni, Co, Sc, Mn, and V) for a sample of 1111 FGK dwarf stars from the HARPS GTO planet search program. Of these stars, 109 are known to harbor giant planetary companions and 26 stars are exclusively hosting Neptunians and super-Earths.} {The two main goals of this paper are to investigate whether there are any differences between the elemental abundance trends for stars of different stellar populations and to characterize the planet host and non-host samples in terms of their [X/H]. The extensive study of this sample, focused on the abundance differences between stars with and without planets will be presented in a parallel paper.} {The equivalent widths of spectral lines were automatically measured from HARPS spectra with the ARES code. The abundances of the chemical elements were determined using an LTE abundance analysis relative to the Sun, with the 2010 revised version of the spectral synthesis code MOOG and a grid of Kurucz ATLAS9 atmospheres. To separate the Galactic stellar populations we applied both a purely kinematical approach and a chemical method.} {We found that the chemically separated (based on the Mg, Si, and Ti abundances) thin- and thick disks are also chemically disjunct for Al, Sc, Co, and Ca. Some bifurcation might also exist for Na, V, Ni, and Mn, but there is no clear boundary of their [X/Fe] ratios. We confirm that an overabundance in giant-planet host stars is clear for all studied elements.We also confirm that stars hosting only Neptunian-like planets may be easier to detect around stars with similar metallicities than around non-planet hosts, although for some elements (particulary $\\alpha$-elements) the lower limit of [X/H] is very abrupt.} {} ", "introduction": "High-precision radial velocity measurements resulted in the detection of the first extra-solar planetary system surrounding a main-sequence star similar to our own in 1995 (Mayor \\& Queloz \\cite{Mayor-95}). Observational progress in extra-solar planet detection and characterization is now moving rapidly on several fronts. More than 750 planetary companions have already been found orbiting late-type stars\\footnote{http://exoplanet.eu/% }. The total number of planet-harboring systems that are found using Doppler technique is approaching 500. A strong input for this number was made by several dedicated planet-search programs that systematically monitor the sky. Among these programs, the HARPS planet search program made a special contribution. The high spectral resolution and most importantly the long-term stability of the HARPS spectrograph (Mayor et al. \\cite{Mayor-03}) allowed discovering a fairly large number of new planets, including the large majority of the known planets with masses near the mass of Neptune or below (e.g. Santos et al. \\cite{Santos-04a}; Lovis et al. \\cite{Lovis-06}; Mayor et al. \\cite{Mayor-09}, \\cite{Mayor-11}). \\begin{figure*} \\begin{center}$ \\begin{array}{cc} \\includegraphics[width=0.5\\linewidth]{ni_err_trand_typical.pdf}& \\includegraphics[width=0.5\\linewidth]{ni_err_trand_1sigma.pdf} \\end{array}$ \\end{center} \\caption{Ni abundance sensitivity to the stellar parameter variations as a function of model atmosphere parameters. \\emph{Left} - The variation of the atmospheric parameters are the same for all stars and are equal to the typical errors. \\emph{Right} - The variation of the atmospheric parameters are equal to their one-sigma errors taken for each star individually.} \\label{fig-1} \\end{figure*} Shortly after the discovery of the first extra-solar planet, Gonzalez (\\cite{Gonzalez-98}), based on a small sample of eight planet-host stars (PHS), suggested that PHSs tend to be metal-rich compared with the nearby field FGK stars that are known to host no-planet. The metal-rich nature of the PHSs have been confirmed in subsequent papers (e.g. Gonzalez et al. \\cite{Gonzalez-01}; Santos et al. \\cite{Santos-01,Santos-03,Santos-04,Santos-05}; Laws et al. \\cite{Laws-03}; Fischer \\& Valenti \\cite{Fischer-05}; Gilli et al. \\cite{Gilli-06}; Udry et al. \\cite{Udry-06}; Ecuvillon et al. \\cite{Ecuvillon-07}; Sousa et al. \\cite{Sousa-08}; Neves et al. \\cite{Neves-09}; Johnson et al. \\cite{Johnson-10}; Kang et al. \\cite{Kang-11}). This tendency for giant planets that orbit metal-rich stars strongly supports the core-accretion model of planet formation (e.g. Pollack et al. \\cite{Pollack}). This implies that core accretion (Ida \\& Lin \\cite{Ida-04}; Mordasini et al. \\cite{Mordasini-09}) and not disk-instability (Boss \\cite{Boss-97}) is the main working mechanism for the formation of giant planets. Interestingly, recent studies show that Neptune and super-Earth-class planets may easier form in a low-metal-content environment (e.g. Udry et al. \\cite{Udry-06}; Sousa et al. \\cite{Sousa-08,Sousa-11a}; Ghezzi et al. \\cite{Ghezzi-10}; Mayor et al. \\cite{Mayor-11}; Buchhave et al. \\cite{Buchhave-12}). Most spectroscopic studies are in general limited to small samples of a few hundred comparison stars and less than one hundred PHSs at most, and only a few studies have been based on samples as large as 1000 stars (e.g. Gazzano et al. \\cite{Gazzano-10}; Gazzano \\cite{Gazzano-11}; Petigura \\& Marcy \\cite{Petigura-11}). In order to minimize the errors, one needs to have large and homogeneous samples with reliable measurements of their chemical features. In this paper, we present a uniform spectroscopic analysis of 1111 FGK dwarfs observed within the context of the HARPS GTO planet search program. The paper is organized as follows: In Sect. 2, we introduce the sample used in this work. The method of the chemical abundance determination and analysis will be explained in Section 3. This section also includes discussion of the uncertainties and errors in our methodology as well as a comparison of our results with the literature. The calculation of the galactic space velocity data and the selection of different populations of stars, based on their kinematic and chemical properties, are presented in Sect. 4. A discussion of the [X/H] abundances of the exoplanet hosts can be found in Sect. 5. The main conclusions of the paper are finally addressed in Sect. 6. The extensive and full investigation of this sample, focused on the abundance difference between stars with and without planets will be presented in a parallel paper (Adibekyan et al., \\cite{Adibekyan-12}). ", "conclusions": "We have carried out a uniform abundance analysis for 12 refractory elements (Na, Mg, Al, Si, Ca, Ti, Cr, Ni, Co, Sc, Mn, and V) for a sample of 1111 FGK dwarf stars from the HARPS GTO planet search program. Of these stars, 135 are known to harbor planetary companions (26 of them are exclusively hosting Neptunians and super-Earth planets) and the remaining 976 stars do not have any known orbiting planet. The precise spectroscopic parameters for the entire sample were derived by Sousa et al. (\\cite{Sousa-08}, \\cite{Sousa-11a}, \\cite{Sousa-11b}) in the same manner and from the same spectra as were used in the present study. We discussed the possible sources of uncertainties and errors in our methodology in detail, and also we compared our results with those presented in other works to ensure consistency and reliability in our analysis. The large size of our sample allowed us to characterize and remove systematic abundance trends for some elements with \\emph{$T{}_{\\mathrm{eff}}$}. To separate Galactic stellar populations, we applied both purely kinematical approach and chemical method. We showed that both kinematically selected thin- and thick discs are ''contaminated``. The main reason of this ''contamination`` could be the fact that the stars in the local neighborhood have different birth radii and reached the Solar Neighborhood due to their eccentric orbits or via radial migration (e.g. Sch\u00f6nrich \\& Binney \\cite{Schonrich-09}). Inspection of [X/Fe] against [Fe/H] plots suggests us that chemically separated thin- and thick discs, in addition to the Mg, Si, and Ti, are also different for Al, Sc, Co, and Ca. Some bifurcation might also exist for Na, V, Ni, and Mn, but there is no clear boundary of their [X/Fe] ratios. We observed no abundance difference between the thin- and thick discs for chromium. We found that the metal-poor $\\alpha$-enhanced stars and their metal-rich counterparts show different [X/Fe] trends with metallicity for different elements. We confirmed that an overabundance in giant-planet host stars is clear for all studied elements, which lends strong support to the core-accretion model of planet formation (e.g. Pollack et al. \\cite{Pollack}). We also confirmed that stars hosting only Neptunian-like planets may be easier to detect around stars with similar metallicity than non-planet hosts, although for some elements (particularly $\\alpha$-elements) we observed an abrupt lower limit of [X/H], which may indicate that these elements are important in their formation. The maximum abundance difference between Neptunian-like planet hosts and non-host stars is observed for Mg ([Mg/H] $\\approx$ 0.09 dex)." }, "1207/1207.2494_arXiv.txt": { "abstract": "{} {We have investigated the nature of the variability of CHS\\,7797, an unusual periodic variable in the Orion Nebula Cluster.} {An extensive I-band photometric data set of CHS\\,7797 was compiled between 2004-2010 using various telescopes. Further optical data have been collected in R and z$\\arcmin$ bands. In addition, simultaneous observations of the ONC region including CHS\\,7797 were performed in the I, J, K$_{s}$ \\& IRAC $3.6$ and $4.5\\,\\mu m$ bands over a time interval of $\\approx$\\,40\\,d. } {CHS\\,7797 shows an unusual large-amplitude variation of $\\approx$\\,1.7 mag in the R, I, and z$\\arcmin$ bands with a period $17.786\\,\\pm\\,0.03$\\,d ($FAP\\,=\\,1x10^{-15}\\%$). The amplitude of the brightness modulation decreases only slightly at longer wavelengths. The star is faint during $\\approx$\\,2/3 of the period and the shape of the phased light-curves for the seven different observing seasons shows minor changes and small-amplitude variations. Interestingly, there are no significant colour-flux correlations for $\\lambda$\\,$\\lesssim$\\,2\\,$\\mu$m, while the object becomes redder when fainter at longer wavelengths. CHS\\,7797 has a spectral type of M6 and an estimated mass between 0.04-0.1\\,M$_{\\odot}$. } {The analysis of the data suggests that the periodic variability of CHS\\,7797 is most probably caused by an orbital motion. Variability as a result of rotational brightness modulation by a hot spot is excluded by the lack of any colour-brightness correlation in the optical. The latter indicates that CHS\\,7797 is most probably occulted by circumstellar matter in which grains have grown from typical 0.1\\,$\\mu$m to $\\approx$\\,1-2\\,$\\mu$m sizes. We discuss two possible scenarios in which CHS\\,7797 is periodically eclipsed by structures in a disc, namely that CHS\\,7797 is a single object with a circumstellar disc, or that CHS\\,7797 is a binary system, similar to KH\\,15D, in which an inclined circumbinary disc is responsible of the variability. Possible reasons for the typical 0.3\\,mag variations in I-band at a given phase are discussed. } ", "introduction": "\\begin{figure*} \\includegraphics[scale=0.65]{FC_WFI_2004_1.ps} \\caption{I-band images of CHS\\,7797 during maximum and minimum brightness obtained with WFI at the 2.2\\,m telescope on La Silla, Chile. North is up and east is left. The region shown is 2\\,$\\arcmin$x2\\,$\\arcmin$ in size. } \\label{Fig_FC} \\end{figure*} The photometric variability of T Tauri stars (TTS) is attributed in many cases to magnetically induced cool starspots and/or magnetically channelled variable accretion flows, which produce hot spots at the base of the magnetic channels. During the last 10-15 years, various extensive CCD monitoring programs of young clusters like NGC\\,2264 and the Orion Nebular Cluster (ONC) have been carried out. These monitoring programs have significantly increased our knowledge of TTS variability and rotation, and have revealed thousands of mostly low-amplitude periodic pre-main sequence (PMS) variables and a similar number of irregular variable PMS stars \\citep[e.g.][]{Herbst2002, Lamm2004, Rodriguez-Ledesma2009}. As reviewed by \\cite{Herbst2007}, and based on many of these studies, it is possible to distinguish five types of PMS variability. The different types have known properties such as typical amplitude range of the variability at different wavelengths. For example, amplitudes at optical wavelengths of periodic brightness modulation, due to hot spots and/or magnetically channelled accretion, can be a factor 2 to 5 larger than the periodic variations believed to be caused by cool spots, which cannot be larger than 0.5 mag in I \\citep{Herbst1994}. Another relevant difference between these two cases is that cool spots have been observed to be stable over thousands of cycles, while accretion-related variability is much more unstable and irregular, with the observed periodicity often only surviving several tens of cycles. It is important to note that rotational properties and accretion-related variability have been observed also in substellar objects \\cite[e.g.][]{Scholz2004,Mohanty2005,Rodriguez-Ledesma2009}. Because most of these variability studies involve a large number of objects, they are not only valuable for characterising variability, but also for detecting interesting and rare objects. One striking example is KH\\,15D, a unique 48\\,d periodic variable in NGC\\,2264, with extremely deep ($\\approx$3.5 mag) minima. It was first recognized as interesting by \\cite{Kearns1998}. \\cite{Hamilton2001} reported that KH\\,15D is a K7 PMS member of NGC\\,2264 ($\\approx$ 2-4\\,Myr), with a mass of 0.5-1\\,$M_{\\odot}$, and together with \\cite{Herbst2002} proposed that the star was being eclipsed by circumstellar material. Further intense investigations and, in particular, radial velocity studies led to a binary model for KH\\,15D in which the system is surrounded by a nearly edge-on and slowly precessing circumbinary disc, to which the binary system is slightly inclined \\citep[and references therein]{Winn2006, Herbst2008}. Over the past 15 years the circumbinary disc has completely occulted the orbit of component B and has allowed us to see only a diminishing fraction of the orbit of star A. It is the appearing and disappearing of star A behind the disc that causes the strong light modulation. Other TTS binaries are expected to have similar characteristics, but KH15\\,D is not only special because of its particular geometric orientation and the precessing circumbinary disc but also because of the indications of significant grain growth within the disc \\citep{Herbst2008}. This binary is also interesting because it is a jet-driving source, with the jet being most probably a ``common product`` of the whole binary system or the circumbinary disc \\citep{Mundt2010}. Although thousands of PMS stars have been photometrically monitored in the past years, only few objects show a variable behaviour that resembles KH\\,15D, i.e. YLW\\,16A and WL\\,4 in $\\rho$\\,Oph \\citep{Plavchan2010,Plavchan2008}, and V718\\,Per in IC 348 \\citep{Grinin2008}. In the case of WL\\,4 and YLW\\,16A, the authors proposed the existence of a new class of periodic disc-eclipsing binaries to explain the KH\\,15D type variability, in which a third body in these systems is responsible for the warping of the inner circumbinary disc needed to produce the eclipses, while \\cite{Grinin2008} argued that V718\\,Per is most probably a single star surrounded by an edge-on circumstellar disc with an irregular mass distribution at the inner edge of the disc, which causes the observed periodic variability. AATau-like objects show also high amplitude variability, although this variability is found to be only quasi-periodic \\citep{Alencar2010}. As described in \\cite{Alencar2010}, in the fairly common AATau-like objects the quasi-periodic variability is attributed to a magnetically controlled inner disc warp. \\cite{Rodriguez-Ledesma2009} carried out a photometric monitoring in the ONC during 2004 with the primary goal of deriving rotational periods for a large sample of very low-mass PMS stars and brown dwarfs. Several of the objects in this sample show unusual light curves: CHS\\,7797 (i.e. star No.7797 in \\cite{Carpenter2001}) stuck out particularly because of its unusual large-amplitude modulation of $\\approx$1.7\\,mag in I and its periodic variability. Hoping that this newly discovered periodic variable might have a similar potential to KH\\,15D, which has provided considerable insight into the physics of PMS binary systems and their circumbinary discs, we have carried out a photometric follow-up campaign, including multicolour data collected with various telescopes/instruments. This paper is organised as follows: in Section\\,2, we present the optical and near-infrared (NIR) data sets, while in Section\\,3 we describe the photometric data evaluation. Section\\,4 describes the time series analysis procedure. We analyse colour changes in Section\\,\\ref{Colours} and estimate luminosity and mass ranges for CHS\\,7797 in Section\\,\\ref{Luminosity}. A detailed discussion of our results is given in Section\\,\\ref{Discussion}, while final conclusions are presented in Section 7. ", "conclusions": "CHS\\,7797 is undoubtedly a very interesting and unusual young and possibly substellar object that is occulted by circumstellar/circumbinary matter. Whether it is a single object or a binary system, CHS\\,7797 is surrounded by a circumstellar/circumbinary disc that is viewed almost edge-on and in which substantial grain growth has taken place. A comparison of the light-curves and phased light-curves of CHS\\,7797 with the three other known ``unusual`` periodic variables (i.e. KH\\,15D, WL\\,4 and V718\\,Per) suggests some similarities but also discrepancies. In particular, the other three objects show much smoother light curves than CHS\\,7797, which implies that, for as yet unknown reasons, the disc around CHS\\,7797 is much more inhomogeneous, and/or that there is a contribution of accretion-related irregular variations. Clearly, CHS\\,7797 deserves further attention. An upcoming spectroscopic analysis (Rodriguez-Ledesma et al., in prep) will try to provide additional constraints on the nature of CHS\\,7797. Nevertheless, since CHS\\,7797 is very faint and since there is a strong background contribution from the Orion nebula, it will be hard to search for radial velocity variations to confirm the possible binary character of the system even with the instrumentation currently available at 8-10\\,m class telescopes." }, "1207/1207.0943_arXiv.txt": { "abstract": "Since the discovery by Hale in the early 1900s that sunspots harbor strong magnetic field, magnetism has become increasingly important in our understanding of processes on the Sun and in the Heliosphere. Many current and planned instruments are capable of diagnosing magnetic field in the solar atmosphere. Photospheric magnetometry is now well-established. However, many challenges remain. For instance, the diagnosis of magnetic field in the chromosphere and corona is difficult, and interpretation of measurements is harder still. As a result only very few measurements have been made so far, yet it is clear that if we are to understand the outer solar atmosphere we must study the magnetic field. I will review the history of solar magnetic field measurements, describe and discuss the three types of magnetometry, and close with an outlook on the future. ", "introduction": "The importance of magnetic field in astrophysical processes has long been recognised. Magnetic field is found in all parts of the universe and on all scales: planets, stars, galaxies, accretion disks, etc. Even in cases where magnetic field was initially deemed too weak to significantly affect plasma dynamics, it has often turned out to be an important factor in the evolution of the system \\citep[e.g., in accretion disks,][]{1991ApJ...376..214B}. Nowadays a thriving field of research in astrophysics is centered around the life-cycle of magnetic field. How is it created? How does it interact with the plasma it is embedded in? How is it destroyed? The diagnosis of magnetic field is now spreading throughout all parts of astrophysics, e.g., Zeeman-Doppler imaging of stars \\citep{1989A&A...225..456S} and exoplanets interacting with their host star \\citep{2003ApJ...597.1092S}. Indeed, magnetic field has now been inferred to be present or even measured in many astrophysical objects. \\articlefigure[width=0.74\\textwidth]{MaunderButterfly}{dewijn:fig:butterfly}{% The Maunder butterfly diagram showing the latitude variation of active regions over the solar cycle.} Magnetic fields feature perhaps most prominently in solar physics. The definite discovery that sunspots have strong magnetic fields was made by \\cite{1908ApJ....28..315H}, who used a primitive spectro-polarimeter to observe circular polarization in spectral lines resulting from the well-known Zeeman effect \\citep{Zeeman1896}. A more detailed publication \\citep{Zeeman1897} was reproduced in its entirety that same year in The Astrophysical Journal \\citep{1897ApJ.....5..332Z}. These discoveries spurred a century of advancement spectroscopy, polarimetry, instrument development, and research of solar magnetism. It is worth noting that the first detection of widening of spectral lines in sunspots, an effect now unambiguously attributed to magnetism, was first reported by \\cite{1866RSPS...15..256L}, more than 30~years prior to the laboratory discovery of the Zeeman effect, and shown by \\cite{Young1883book} in a graphic illustration of an observation dated back to 1870. This subject is covered in detail in the excellent review by \\cite{1996VA.....40..241D}. On the cutting edge are instruments like the Helioseismic and Magnetic Imager \\citep[HMI,][]{2012SoPh..275..229S} on the Solar Dynamics Observatory \\citep[SDO,][]{2012SoPh..275....3P}, which provides us with the full-disk vector magnetic field in the solar photosphere several times per hour, and the Spectro-Polarimeter (SP) of the Solar Optical Telescope \\citep[SOT,][]{2008SoPh..249..167T} of Hinode \\citep{2007SoPh..243....3K}, which combines unparalleled spatial resolution with high sensitivity to make the most refined measurements of solar photospheric magnetic field to date. ", "conclusions": "" }, "1207/1207.7087_arXiv.txt": { "abstract": "{Limb darkening is an important tool for understanding stellar atmospheres, but most observations measuring limb darkening assume various parameterizations that yield no significant information about the structure of stellar atmospheres.} {We use a specific limb-darkening relation to study how the best-fit coefficients relate to fundamental stellar parameters from spherically symmetric model stellar atmospheres.} {Using a grid of spherically symmetric \\textsc{Atlas} model atmospheres, we compute limb-darkening coefficients, and develop a novel method to predict fundamental stellar parameters. } {We find our proposed method predicts the mass of stellar atmosphere models given only the radius and limb-darkening coefficients, suggesting that microlensing, interferometric, transit and eclipse observations can constrain stellar masses.} {This novel method demonstrates that limb-darkening parameterizations contain important information about the structure of stellar atmospheres, with the potential to be a valuable tool for measuring stellar masses.} ", "introduction": "\\label{sec:intro} Limb darkening, the change of surface intensity from the center to the edge of a stellar disk, is a powerful measure of the physical structure of stellar atmospheres. However, it is difficult to observe the intensity profile across a stellar disk except for the Sun. Most observations of stellar limb darkening are indirect, coming from light curves of eclipsing binaries \\citep{Claret2008} or planetary transits \\citep{Knutson2007}, interferometric visibilities \\citep{Haubois2009} and microlensing light curves \\citep{An2002}, and these indirect methods have limited precision \\citep[e.g.][]{Popper1984, Zub2011}. Because of this, limb darkening is commonly treated as a parameterized function of $\\mu \\equiv \\cos \\theta$, where $\\theta$ is the angle between the direction to the distant observer and the normal direction at each location on the stellar surface. An example of a linear parameterization is \\citep{Schwarzschild1906} \\begin{equation} \\frac{I_\\lambda(\\mu)}{I_\\lambda(\\mu=1)} = 1 - a_\\lambda(1-\\mu). \\end{equation} As numerical model atmospheres became robust, the computed intensity profiles were represented by more elaborate limb-darkening parameterizations that include higher order terms of $\\mu$ or are expressed in terms of powers of $r = \\sin \\theta$ \\citep{Heyrovsky2007}, as well as being normalized with respect to the stellar flux instead of the central intensity \\citep{Wade1985}. These more detailed limb-darkening laws were used to interpret the improving observations noted above. However, even with these advances, current limb-darkening observations are still unable to constrain the model stellar atmospheres. Recent interferometric observations of nearby red giants are not yet precise enough to differentiate between the center-to-limb intensity profiles from \\textsc{Phoenix} and \\textsc{Atlas} model atmospheres or even between plane-parallel or spherically symmetric models \\citep{Wittkowski2004, Wittkowski2006b, Wittkowski2006a, Neilson2008}. However, the combination of the observations and models \\emph{do} provide constraints for fundamental stellar parameters such as the effective temperature and gravity. In this work, we continue our previous analysis \\citep[][hereafter Paper 1]{Neilson2011} to explore the connection between limb darkening and stellar parameters. We adopt the flux-conserving law \\begin{equation} \\label{eq:ldlaw} \\frac{I_\\lambda(\\mu)}{2\\mathcal{H}_\\lambda} = 1 - A_\\lambda \\left (1 - \\frac{3}{2}\\mu \\right ) - B_\\lambda \\left (1 - \\frac{5}{4} \\sqrt{\\mu} \\right ), \\end{equation} where $\\mathcal{H}_\\lambda$ is the Eddington flux defined as \\begin{equation} \\mathcal{H}_\\lambda \\equiv \\frac{1}{2} \\int_{-1}^{1} I_\\lambda(\\mu) \\mu \\mathrm{d} \\mu, \\end{equation} because this law was used by \\citet{Fields2003} to analyze the microlensing event EROS-BLG-2000-5 \\citep{An2002}. We fit this limb-darkening law to the intensities of spherically symmetric model stellar atmospheres computed with the \\textsc{SAtlas} code \\citep{Lester2008} to explore the relation between fundamental parameters and limb-darkening coefficients hinted at in Paper 1. ", "conclusions": "The purpose of this study has been to investigate how limb darkening probes fundamental stellar properties. The limb darkening was parameterized by a flux-conserving linear-plus-square-root law that has two free parameters that are functions of two angular moments of the intensity. The ratio of these two moments provides a measure of the amount of stellar atmospheric extension, represented by $R_\\star/M_\\star$. We tested our method for determining the extension from limb-darkening parameters using a spherically symmetric model atmosphere and found agreement of the mean value of the five spectral bands analyzed to within 8\\%. We also find significant variation of $R_\\star/M_\\star$ due to the definition of the stellar radius. From the context of the model of the stellar atmosphere, we find it natural to define the stellar radius where $\\tau_\\mathrm{Ross} = 2/3$, but other options are possible, such as $\\tau_\\mathrm{Ross} = 1$ or using a specific wavelength, such as 500~nm, to define the radius where $\\tau_{500} = 1$. Observations, of course, are done in specific wavebands, such as $V$, and the stellar radius is assumed to refer to some particular optical depth, such as $\\tau_\\mathrm{V} = 2/3$. Switching to another waveband will lead to another value of the radius. As an example, in the near-infrared the predicted radius is larger than at optical wavelengths. Therefore, the definition of the stellar radius is ambiguous in general for studies of angular diameters \\citep[e.g.][]{Wittkowski2004, Wittkowski2006a, Wittkowski2006b}, which makes the definition of the stellar radius challenging for stars with significant atmospheric extension. A recent example is \\citet{Ohnaka2011}, who measured the $K$-band angular diameter of Betelgeuse to be about 2 mas ($\\approx$ 5\\%) smaller than measured by \\citet{Haubois2009} in the $H$-band. However, simultaneously fitting $R_\\star/M_\\star$ to multiband limb-darkening observations provides a bolometric value for $R_\\star/M_\\star$ analogous to how spectrophotometry can be used to measure a star's effective temperature. The connection between atmospheric extension and stellar limb-darkening is a consequence of the assumption of spherical symmetry for the stellar atmospheres, not a unique feature of our \\textsc{SAtlas} code. The results presented here should also be found using intensity profiles from spherically-symmetric \\textsc{Marcs} models and the \\textsc{Phoenix} model atmospheres used by \\citet{Claret2003} and \\citet{Fields2003}. Unfortunately, in both of these works the authors truncated the intensity profiles to remove the low intensity limb and then redistributed the spacing of $\\mu$-points. Truncating the intensity profile eliminates critical information about the limb darkening, and makes the star's intensity distribution resemble a plane-parallel atmosphere, as we demonstrate in Fig.~\\ref{fig:clipped}. \\begin{figure*}[t] \\centering \\resizebox{\\hsize}{!}{\\includegraphics{clipped_ldr.eps} \\includegraphics{clipped_ldmu.eps} \\includegraphics{fitted_ldmu.eps}} \\caption{(a) Surface-brightness distribution as a function of fractional radius. The solid line is the full set of intensities from a spherical model atmosphere with $T_\\mathrm{eff} = 5000$ K, $\\log g = 2.0$ and $M_\\star = 5~M_{\\sun}$. The dashed line has removed intensity values for $r/R_\\star \\ge 0.995$ and then renormalized the fractional radius to the interval $0-1$. The dotted line has removed intensity values for $r/R_\\star \\ge 0.98$ and then renormalized the fractional radius to the interval $0-1$. (b) Surface-brightness distributions from panel~(a) plotted as a function of $\\mu$, with the lines having the same meaning as in the (a)~panel. (c) Fits to the surface-brightness distributions in panel~(b) using Eq.~(\\ref{eq:ldlaw}). The lines have the same meaning as the other two panels.} \\label{fig:clipped} \\end{figure*} Our method for determining fundamental stellar parameters is based on the general atmospheric properties of flux-conserving limb-darkening laws, not the particular law used here. For instance, \\citet{Heyrovsky2003} and \\citet{Zub2011} found a fixed point in flux-conserving linear limb-darkening laws, which indicates that the linear limb-darkening coefficient is a measure of the mean intensity and flux of the atmosphere. Another example is a quadratic limb-darkening law where the limb-darkening coefficients would provide a measure of the mean intensity and second moment of the intensity, $K$. Because these laws measure the correlation between various moments of the intensity in an atmosphere, they would also measure the atmospheric extension of that star. We conclude that the method outlined in this work is a viable way to use a limb-darkening law, such as the linear-plus-square-root law employed here, to determine the fundamental physical parameters of stars. The method requires knowledge of one of the degenerate parameters effective temperature or luminosity. However, the method is able to constrain stellar parameters using limb-darkening observations even in the case when those limb-darkening observations from microlensing or eclipsing binaries cannot test the model stellar atmosphere directly. It seems as though we still have things to learn from fairly simple representations of limb darkening, and that, as observations continue to improve, this will become an even more powerful tool in the study of stellar astrophysics." }, "1207/1207.0458_arXiv.txt": { "abstract": "{The evidence of a line-like spectral features above 100~GeV, in particular at 130~GeV, recently reported from some parts of the galactic plane poses serious challenges for any interpretation of this surprise discovery. It is generally believed that the unusually narrow profile of the spectral line cannot be explained by conventional processes in astrophysical objects, and, if real, is likely to be associated with Dark Matter.} {In this paper we argue that cold ultrarelativistic pulsar winds can be alternative sources of very narrow gamma-ray lines.} {We demonstrate that Comptonization of a cold ultrarelativistic electron-positron pulsar wind in the deep Klein-Nishina regime can readily provide very narrow ($\\Delta E /E \\leq 0.2$) distinct gamma-ray line features. To verify this prediction, we produced photon count maps based on the Fermi LAT data in the energy interval 100 to 140~GeV.} {We confirm earlier reports of the presence of marginal gamma-ray line-like signals from three regions of the galactic plane. Although the maps show some structure inside these regions, unfortunately the limited photon statistics do not allow any firm conclusion in this regard.} {The confirmation of 130 GeV line emission by low-energy threshold atmospheric Cherenkov telescope systems, in particular by the new 27~m diameter dish of the H.E.S.S. array, would be crucial for resolving the spatial structure of the reported hotspots, and thus for distinguishing between the Dark Matter and Pulsar origins of the `Fermi Lines'. } ", "introduction": "Recent reports on the possible presence of a narrow line-like feature in the spectrum of gamma-ray emission at 130 GeV from the galactic center region, and, presumably, from some other parts of the galactic plane \\citep{brigmann12,weniger12,tempel12,boyarsky12,su12}, have received a prompt and enthusiastic reaction from the astrophysics and particle physics communities. Despite the marginal statistical significance of the reported signals and some outstanding questions and inconsistencies, the implications of these results are hotly debated, basically in the context of Dark Matter (DM). This is motivated not only by the recognition of the potential of gamma-ray observations for indirect searches of DM \\citep[for a recent review see][]{Bergstrom12} and the overall excitement caused by the possible association of the gamma-ray line at 130~GeV with DM \\citep[see e.g.][]{Buckley12}, but also by unusual characteristics of the signal. It is generally believed that the width of the 130~GeV line (of a few tens of GeV) is too narrow to be explained by any physical process except for annihilation of DM\\footnote{Based on the analysis of \\fermi LAT data by \\citet{weniger12}, it has been argued \\citep{Profumo12} that the contamination of a bright featureless power-law background component (related e.g. to the radiation of the interstellar medium) by hard photons arriving from Fermi bubbles with a sharp spectral break between 100 and 150~GeV, can mimic a spurious line. However, a closer look at the morphology shows that the 130 GeV feature is most likely not associated with Fermi Bubbles \\citep[see e.g.][]{tempel12}.}. On the other hand, some other features, in particular the significant shift of the center of gravity of the signal from the position of the galactic center \\citep{su12}, as well as the possible presence of gamma-ray lines at other energies \\citep{boyarsky12} and from other parts of the galactic plane \\citep{tempel12,boyarsky12}, challenge the DM origin of the reported signal. Unfortunately, the marginal gamma-ray photon statistics does not allow definite conclusions in this regard. The situation will be somewhat improved after doubling the photon statistics above 100~GeV by observations of the galactic plane with \\fermi Large Area Telescope (LAT) over the next several years. The development of new dedicated approaches to the data reduction focused on the highest energy domain of \\fermi LAT also may help to clarify some of the current uncertainties and inconsistencies. Finally, there is hope that soon the low-energy threshold imaging Cherenkov telescopes, in particular the new 27m diameter dish of the H.E.S.S. telescope array located in the Southern Hemisphere (see http://www.mpi-hd.mpg.de/hfm/HESS/), will greatly contribute to the clarification of the question concerning the origin of $\\geq 100$~GeV gamma-ray line(s) - are they {\\it instrumental}, {\\it cosmological (DM)}, or {\\it astrophysical}? The last option implies production of gamma-rays by accelerated particles. So far it has been discarded, basically because of the common belief that conventional high energy processes with involvement of ultrarelativistic particles could not reproduce gamma-ray line features. In Sec.\\ref{astrolines} we briefly discuss different gamma-ray production mechanisms in the VHE domain in the context of their ability to produce sharp gamma-ray lines, and argue that the inverse Compton scattering in the Klein-Nishina regime by cold ultrarelativistic electron-positron pulsar winds can result in sharp gamma-ray line emission. The conditions for reproduction of narrow Klein-Nishina line profiles are discussed in Sec.\\ref{KNlines}. Finally, in Sec.\\ref{Data} we describe our study of the spatial distribution of $E \\geq 100$~GeV photons based on the {\\it Fermi} LAT data. We confirm the results reported earlier on the presence of marginal gamma-ray line-like signals from three regions of the galactic plane, but, because of limited photon statistics, we cannot make a strong statement on the existence of localized hot spots inside the `Fermi Line' regions . The results and conclusions are summarized in Sec.\\ref{summary}. ", "conclusions": "\\label{summary} In this paper we argue that there is a viable alternative to the DM origin of the 130~GeV gamma-ray line as recently reported to be present in the galactic gamma-ray emission. We demonstrate that very sharp gamma-ray spectral lines can be produced by pulsar winds through their Comptonizatioin, predominantly by energetic (UV, X-ray) radiation with a relatively narrow spectral distribution, thus the IC scattering proceeds in the deep Klein-Nishina limit. In principle, these conditions can be fulfilled both in isolated pulsars and binary systems. The current paradigm which connects pulsars with their synchrotron nebulae through cold ultrarelativistic electron-positron winds, postulates that the electron-positron wind with a Lorentz factor between $10^4$ to $10^6$, carries away almost the entire rotational energy lost by the pulsar. Thus, under the condition of effective Comptanization of the wind, a substantial fraction of the spin-down luminosity of a pulsar $L_{\\rm rot}$ can be released in a single gamma-ray line. Depending on the intensity of illumination of the pulsar wind by surrounding radiation field(s), the efficiency of formation of such lines can be very high, formally close to 100 \\%. Interestingly, the narrow profile of the 130~GeV line argues against such an extreme efficiency which can be realized in the case of optically thick source; this would imply not only strong attenuation of gamma-rays due to photon-photon pair production, but also significant broadening of the line because of the cooling of electrons. However, both effects become negligible in the case of still very high, $\\leq 10-20 \\% $ efficiency of the wind Comptonization. Conservative estimates show that while such an efficiency can be readily achieved in binary systems, in the case of isolated pulsars the efficiency is expected to be very low unless the acceleration of the wind takes place close to the light cylinder. The efficiency of transformation of kinetic energy of the pulsar wind into Klein-Nishina gamma-ray lines is a key issue the discussion of which is beyond the scope of this paper. Clearly, under certain conditions, the efficiency can be quite high, and, for pulsars with spin-down luminosities exceeding $10^{36}$~erg/s and wind Lorentz factor $\\geq 2 \\times 10^5$, one may expect $\\geq 100$~GeV gamma-ray lines with luminosities exceeding the bolometric luminosities of magnetospheric emission of pulsars. Moreover, one cannot exclude other configurations of compact objects, e.g. magnetars, with (hypothetical) relativistic electron-positron winds, as effective multi-GeV gamma-ray emitters. In this regard a natural question arises: if the pulsars can work as effective generators of VHE gamma-ray lines, why have they not been detected yet? There could be several answers to this question. In particular, in many objects the target radiation field could not be sufficient for effective conversion of the kinetic energy of the wind to IC gamma-radiation. Alternatively, if the wind Lorentz factor is small and/or the target photons have a broad distribution, the IC scattering of electrons in the Thompson regime would lead to a gamma-ray continuum the detection and identification of which could appear not an easy task. The formation of the line is effective only for pulsars with wind Lorentz factor exceeding $10^5$; the corresponding gamma-ray line is formed around or above 100~GeV. At these energies the potential of \\fermi LAT is limited because of the small detection area. On the other hand, the current imaging Cherenkov telescopes operate effectively above 100~GeV, so could simply have missed the lines around 100~GeV. Over the next several years \\fermi LAT can double the photon statistics which will be sufficient, hopefully, for a firm detection of the 130~GeV line, but still not adequate for morphological studies of the gamma-ray line emitting regions. In this regard more promising seem to be forthcoming studies with new, low-energy threshold imaging atmospheric Cherenkov telescope systems, in particular by the new 27~m diameter dish of the H.E.S.S. telescope array which is located in a perfect site in the Southern Hemisphere for observations of the galactic center region. This new instrument with energy threshold as low as 50 GeV, huge collection area exceeding $10^4 \\ \\rm m^2$, and energy resolution close to 20 \\% should allow (in the near future!) deep spectroscopic and morphological studies of the inner galaxy in multi-GeV gamma-ray lines. If confirmed, the existence of such lines may lead to an exciting new research area -- {\\it Klein-Nishina gamma-ray line astronomy} -- that will open the way for future ground-based gamma-ray detectors, in particular the Cherenkov Telescope Array (see http://www.cta-observatory.org/) to probe the physics and astrophysics of relativistic outflows, in particular pulsar winds." }, "1207/1207.5537_arXiv.txt": { "abstract": "We test how well available stellar population models can reproduce observed $u$,$g$,$r$,$i$,$z$-band photometry of the local galaxy population ($0.02\\leq z\\leq 0.03$) as probed by the SDSS. Our study is conducted from the perspective of a user of the models, who has observational data in hand and seeks to convert them into physical quantities. Stellar population models for galaxies are created by synthesizing star formations histories and chemical enrichments using single stellar populations from several groups (Starburst99, GALAXEV, Maraston2005, GALEV). The role of dust is addressed through a simplistic, but observationally motivated, dust model that couples the amplitude of the extinction to the star formation history, metallicity and the viewing angle. Moreover, the influence of emission lines is considered (for the subset of models for which this component is included). The performance of the models is investigated by: 1) comparing their prediction with the observed galaxy population in the SDSS using the ($u$-$g$)-($r$-$i$) and ($g$-$r$)-($i$-$z$) color planes, 2) comparing predicted stellar mass and luminosity weighted ages and metallicities, specific star formation rates, mass to light ratios and total extinctions with literature values from studies based on spectroscopy. Strong differences between the various models are seen with several models occupying regions in the color-color diagrams where no galaxies are observed. We would therefore like to emphasize the importance of the choice of model. Using our preferred model we find that the star formation history, metallicity and also dust content can be constrained over a large part of the parameter space through the use of $u$,$g$,$r$,$i$,$z$-band photometry. However, strong local degeneracies are present due to overlap of models with high and low extinction in certain parts of color space. ", "introduction": "Optical broad band colors have proven to be a powerful tool in studying galaxies. Their dependence on luminosity and environment have greatly increased our knowledge of these systems \\citep{1977ApJ...216..214V,2007ApJ...658..898P,2008AJ....135..380L}. Colors are also used to derive quantities such as star formation histories and stellar masses \\citep{1968ApJ...151..547T,1973ApJ...179..427S,1991ApJ...367..126C,2001ApJ...550..212B,2003MNRAS.344.1000B,2007AJ....133..734B}, both of which are key properties for understanding galaxy formation and evolution. For example, strong correlations between stellar mass and galaxy structure \\citep{2003MNRAS.341...54K}, star formation history \\citep{2007MNRAS.378.1550P}, chemical enrichment \\citep{2008MNRAS.391.1117P} and gas content \\citep{2009MNRAS.397.1243Z} have been presented suggesting that mass is the main driver behind galaxy evolution. These relations reflect the importance of gravity on galactic scales and, moreover, provide further evidence concerning the expected connection between stellar mass and dark matter \\citep[cf.][]{2010ApJ...710..903M}. The derivation of star formation histories and stellar masses are made through a modelling of the light emission from the galaxy and, if necessary, modelling of obscuration by dust. The method enables a comparison between observational quantities and galaxy formation models \\citep[e.g.][]{2006MNRAS.366..499D,2006MNRAS.370..645B,2011MNRAS.413..101G}. The quality of the derivation of star formation histories and stellar masses from colors naturally depends on the quality of the models. A straightforward test is to check how well the models can reproduce the ensemble of observables. In this paper we therefore take a closer look at how successful various models are in reproducing $u$,$g$,$r$,$i$,$z$-band photometry of the $local$\\footnote{By focusing on local galaxies ($0.02\\sim17.0$ as compared to $\\sim15.5$ for the entire sample), which has led us to suspect that a part of this population could simply be an artifact caused by photometric errors. The fraction of galaxies with small $z$-band photometric errors $\\sigma_{m_{r}}<0.05$ in the region ($0.3<(g-r)<0.4$,$-0.05<(i-z)<0.05$) is a factor of two smaller than for the overall sample. However, this is not sufficient to explain the entire population. Another explanation is that these galaxies are even more metal poor than the most metal poor models (see Fig. \\ref{fig:z}) or, alternatively, that the SSP models fail at low metallicity. The star formation histories, along with average ages, metallicities and $r$-band optical depth shown in Fig. \\ref{fig:sfh} look reasonable at first glance. A quiescently looking galaxy, which in fact does not have any emission lines in its SDSS spectra, should according to the models have some amount of recent star formation. However, non star forming models lie within the errors. We thus conclude that the model output is reasonably reliable -- within the limits of broad band photometry -- in providing information about a galaxy's star formation history and dust content." }, "1207/1207.0254_arXiv.txt": { "abstract": "{ Ultra High Cosmic Rays) made by He-like lightest nuclei might solve the AUGER extragalactic clustering along Cen A. Moreover He like UHECR nuclei cannot arrive from Virgo because the light nuclei fragility and opacity above a few Mpc, explaining the Virgo UHECR absence. UHECR signals are spreading along Cen-A as observed because horizontal galactic arms magnetic fields, bending them on vertical angles. Cen A events by He-like nuclei are deflected as much as the observed clustered ones; proton will be more collimated while heavy (iron) nuclei are too much dispersed. Such a light nuclei UHECR component coexist with the other Auger heavy nuclei and with the Hires nucleon composition. We foresaw (2009) that UHECR He from Cen-A AGN being fragile should partially fragment into secondaries at tens EeV multiplet (D,$He^{3}$,p) nearly as the recent twin multiplet discovered ones (AUGER-ICRC-2011), at $20$ EeV along Cen A UHECR clustering. Their narrow crowding occur by a posteriori very low probability, below $3\\cdot 10^{-5}$. Remaining UHECR spread group may hint for correlations with other gamma (MeV-$Al^{26}$ radioactive) maps, mainly due to galactic SNR sources as Vela pulsar, the brightest, nearest GeV source. Other nearest galactic gamma sources show links with UHECR via TeV correlated maps. We suggest that UHECR are also heavy radioactive galactic nuclei as $Ni^{56}$, $Ni^{57}$ and $Co^{57}$,$Co^{60}$ widely bent (tens degree up to $\\geq 100^{o}$) by galactic fields. UHECR radioactivity (in $\\beta$ and $\\gamma$ channels) and decay in flight at hundreds keV is boosted (by huge Lorentz factor $\\Gamma_{Ni}\\simeq 10^{9}- 10^{8}$) leading to PeVs electrons and consequent synchrotron TeVs gamma offering UHECR-TeV correlated sky anisotropy. Moreover also rarest and non-atmospheric $\\tau$, and e neutrinos secondaries at PeVs, as the first two rarest shower just discovered in ICECUBE, maybe the first signature of such expected radioactive secondary tail.} % ", "introduction": " ", "conclusions": "$ decay? } The very rich UHECR map of AUGER and HIRES in celestial coordinate overlapped on TeV anisotropy (North sky by ARGO- South sky by ICECUBE atmospheric muons) is displayed in next figure, see Fig.\\ref{figure3}. The figure shows a clear area around Crab, marked by arrows, that is somehow extending in a wide anisotropy area where few UHECR took place. We also added arrows to remind the Cen A unique clustering as well as the Vela and the Cygnus galactic TeV sources, somehow in connection with UHECR. Surprisingly the UHECR puzzle maybe at a corner stone: the UHECR-Multiplet along Cen A, the absence of Virgo, the hint of correlation with Vela and with galactic TeV anisotropy \\cite{Desiati} \\cite{ARGO}, might be in part solved by an extragalactic lightest nuclei, mainly He, from Cen A ,A partial confirm is the predicted \\cite{Fargion2011} and observed \\cite{Auger11} multiplet clustering by fragments (D,p) at half UHECR edge energy aligned with Cen A: He like UHECR may be bent by a characteristic angle as large as $\\delta_{rm-He} \\simeq 16^\\circ$ \\cite{Fargion2011} while their fragments multiplet follow along a tail spread by a wider angle $\\delta_{rm-p} \\simeq 32^\\circ$ \\cite{Fargion2011},\\cite{Auger11}. also neutron beta decay may feed a TeV correlated anisotropy. Other UHECR spread events, might be due to a dominant heavy radioactive nuclei component $Ni^{56}$, $Ni^{57}$ and $Co^{56}$, $Co^{60}$,originated by galactic sources (old SNR-GRB relics) as suggested also by relic $Al$ nuclei at rest in gamma map. UHECR Ni,Co maybe deflected by $18,7^{o}$ for Vela, $128^{o}$ (or less) for Crab tuning within TeV inhomogeneities, made by boosted hundred keV gamma and beta positrons decay, shining at TeVs . We predict here analogous UHECR traces around Cygnus and Cas A in future TA UHECR maps. Inner galactic core UHECR are widely spread and hidden by magnetic fields in dense magnetic galactic core arms. However more clustering around ($\\geq 20^{o}$) the galactic plane far from the core, is expected in future data. Magellanic cloud and stream may rise in UHECR maps. UHECR should rise around Cas A and Cygnus, seen by T.A. in North sky. Recent doublet toward Aq X1 may be a new galactic source. The UHECR spectra cut off maybe not indebt to the expected extragalactic GZK feature but to the more modest imprint of a galactic confinement and of nuclei spectrography. The UHECR radioactive beta decay in flight may trace in new $\\nu_{\\tau}$ neutrino astronomy or anisotropy, noise free, related to astronomical (parasite oscillated) tau neutrino; boosted tau (\\emph{mini-double bang} \\cite{Learned}, within a 5 meter size) in Deep Core or ANTARES \\cite{antares} may reveal hundred TeV tau decay (seeing similar PeV ones in ICECUBE \\cite{Learned}). Also Tau airshowers may rise in Cherenkov beamed air-showers. \\cite{Fargion1999}, \\cite{FarTau}, as being searched in ASHRA experiment, \\cite{Aita11} or in fluorescence telescopes for higher tau energies \\cite{FarTau},\\cite{Feng02},\\cite{Bertou2002},\\cite{Auger07}, \\cite{Auger08}. The discover of such expected Neutrino astronomy may shed additional light on the UHECR nature, origination and mass composition, while opening our eyes to mysterious UHECR sources. Future gamma data and UHECR correlation, additional multiplet may lead to a more conclusive fit of this unsolved, century old, cosmic ray puzzle. Tau astronomy (observable also by mini-twin double bangs in Deep Core or Tau Airshowers) as well as ICECUBE showering at PeV with no muon may also offer an additional windows for the first extraterrestrial neutrino traces \\cite{ICECUBE-2012}.\\\\" }, "1207/1207.0548_arXiv.txt": { "abstract": "We investigated gravitational microlensing of active galactic nuclei dusty tori in the case of lensed quasars in the infrared domain. The dusty torus is modeled as a clumpy two-phase medium. To obtain spectral energy distributions and images of tori at different wavelengths, we used the 3D Monte Carlo radiative transfer code \\textsc{skirt}. A ray-shooting technique has been used to calculate microlensing magnification maps. We simulated microlensing by the stars in the lens galaxy for different configurations of the lensed system and different values of the torus parameters, in order to estimate (a)\\ amplitudes and timescales of high magnification events, and (b)\\ the influence of geometrical and physical properties of dusty tori on light curves in the infrared domain. We found that, despite their large size, dusty tori could be significantly affected by microlensing in some cases, especially in the near-infrared domain (rest-frame). The very long time-scales of such events, in the range from several decades to hundreds of years, are limiting the practical use of this method to study the properties of dusty tori . However, our results indicate that, when studying flux ratios between the images in different wavebands of lensed quasars, one should not disregard the possibility that the near and mid-infrared flux ratios could be under the influence of microlensing. ", "introduction": "Gravitationally lensed systems with multiple images represent a powerful tool to study the structure of both the galaxy which acts as the lens and the background source. In a number of lensed systems in which a quasar is the source, the flux ratios between the lensed images deviate from those predicted by the simple lens models \\citep[see e.g.][]{koch91,keet03,gold10}. The fluxes, in different wavebands, can be contaminated by different effects such as microlensing by the stars \\citep[e.g.][]{schwam90} or millilensing by a massive structure in the lens galaxy \\citep[e.g.][]{maosch98}, dust extinction \\citep[e.g.][]{elis06} and also by the time delay itself \\citep[e.g.][]{pc05}. Consequently, the flux ratio anomaly observed in some lensed quasars can be caused by extinction and/or gravitational microlensing/millilensing \\citep[see e.g.][]{pc05,yonehara08}. The size of the source has a large effect on the fluctuations due to microlensing. As a large extended source covers a larger area of a microlensing magnification pattern in the source plane at any given time than a small source, its brightness varies less as it moves relative to the lens and observer \\citep{mortonson05}. As a general rule, the variability of a lensed source is significantly affected by microlensing only if the source is smaller than the relevant length scale - the projection of the Einstein radius of a microlens on to the source plane \\citep{courbin02}. Since the sizes of different emitting regions of quasar are wavelength-dependent, microlensing by the stars in the lens galaxy will lead to a wavelength-dependent magnification. The X-ray radiation is coming from the very compact region in the innermost part of the accretion disk, and therefore, it will be magnified more than the radiation in the UV and optical bands, coming from outer, larger parts of the disk. Thus, although the phenomenon of gravitational lensing is achromatic, due to the complex structure of emission regions, ``chromatic'' effect may arise in a lensed quasar system \\citep[see e.g.][]{pc05,jovanovic08,mo09,mo11}. The ``chromaticity'' in lensing effect can be used to investigate both, an unresolved structure of the innermost region of quasars \\citep[see e.g.][]{wyithe02,aba02,pop03,pop06,bate08,mo09,dai10,black11, garsden11} and the structure of the lens galaxy \\citep[see e.g.][]{ic05,chiba05,xu10}. Moreover, comparing flux ratios at different wavelengths makes it possible to constrain the amount of micro- and milli-lensing present in the system, and the sizes of the perturbers \\citep[see e.g.][]{gold10}. Since the X-ray and UV/optical radiations are coming from relatively compact regions (from several light-days to a light-month), the flux ratios in these wavebands are sensitive to both microlensing by the stars and millilensing by a cold dark matter (CDM) substructure \\citep[see e.g.][]{mm01,pc05,dk06,jovanovic08,gold10,xu10}. On the other hand, the infrared (IR) emitting region of a quasar is assumed to be a toroidal structure of dust, with dimensions significantly larger than the the projection of the Einstein radius of a microlens on to the source plane. Therefore one would expect that the IR radiation of lensed quasars would only be affected by the relatively massive subhalos (millilensing) \\citep[see][]{ic05,chiba05,sluse06,yonehara08,min09,xu10,fk11}. However, certain geometrical and physical properties of the dusty torus can conspire to allow non-negligible microlensing effects in the infrared domain. Infrared spectra of most quasars are dominated by thermal emission from hot dust in their tori, or alternatively, by nonthermal synchrotron emission from the regions near their central black holes \\citep{agol00}. Variability in the infrared band due to gravitational microlensing could be used to constrain the size of the infrared emission region, and hence to distinguish between the thermal and synchrotron mechanisms. If the infrared radiation varies on timescales shorter than decades, then its emission region is smaller, located closer to the central black hole, and its emission is nonthermal, while longer timescales indicate a larger, thermal region \\citep{neugebauer99}. Additionally, chromatic effects in the infrared band have been observed in some of the lensed quasars, where the color differences between their multiple images were detected \\citep{yonehara08}. The most realistic scenario that can explain the observed color differences is gravitational microlensing, in contrast to the dust extinction and the intrinsic variability of quasars \\citep{yonehara08}. Some previous theoretical and observational studies suggested that the infrared emission of quasars is not significantly affected by microlensing, implying that it is most likely produced in their dusty tori. For instance, \\citet{agol00} studied the mid-IR emission of Q2237+0305 observed by Keck and found that it was not affected by microlensing, which ruled out the synchrotron mechanism and supported the model with hot dust extended on a length scale of more than 0.03 pc. \\citet{wyithe02} used mid-IR and $V$-band flux ratios for images $A$ and $B$ of Q2237+0305 to infer the size of the mid-IR emission region and found that it was comparable to or larger than the Einstein Ring Radius (ERR) of the microlens, and hence at least two orders of magnitude larger than the optical emission region. They used simple Gaussian and annular intensity profiles of the dusty torus and found that the results were dependent on the assumed source profile \\citep{wyithe02}. Recent Spitzer observations of the same gravitationally lensed quasar \\citep{agol09} showed that a dusty torus model with a small opening angle could satisfactorily explain the shape of the observed infrared SED, excluding an offset in wavelength of the silicate feature. However, the same authors found that the near-IR fluxes are increasingly affected by microlensing toward shorter wavelengths and that this wavelength dependence is consistent with a model in which a dusty torus and an accretion disk both contribute to the infrared radiation near 1 $\\mathrm{\\mu m}$ \\citep{agol09}. In this paper we present simulations of gravitational microlensing of active galactic nucleus (AGN) dusty tori in the infrared domain. We consider microlensing by stars in the lens galaxies, in the case of lensed quasar systems. We modeled the dusty torus as a clumpy two-phase medium. To obtain spectral energy distributions and images of the torus at different wavelengths, we used the 3D radiative transfer code \\textsc{skirt}. For generating microlensing magnification maps, a ray-shooting method was used. To take into account the size of the dusty tori (as they are larger than the Einstein ring radius of the microlens projected on the source plane), the simulated images of the tori are convolved with the magnification maps. We simulated microlensing magnification events for different configurations of the lensed system and different values of the torus parameters. The aims of this paper are to estimate (a)\\ amplitudes and timescales of high magnification events, and (b)\\ the influence of geometrical and physical properties of dusty tori on microlensing light curves in the infrared domain. The paper is organized as follows. In Section 2 we give a description of our dusty torus model, the method we used to calculate microlensing magnification map, and the parameters we adopted in this study. In Section 3 we present and discuss the results of simulated microlensing light curves of the dusty torus. In Section 4 we outline our conclusions. ", "conclusions": "\\label{sec:conc} We investigated gravitational microlensing of AGN dusty tori in the case of lensed quasars. The dusty torus was modeled as a clumpy two-phase medium. The radiative transfer code \\textsc{skirt} was used to obtain SEDs and images of the tori at different wavelengths. The ray-shooting technique has been used to calculate microlensing magnification maps. Due to the large dimensions of dusty tori (compared to the Einstein ring radius of the microlens in the source plane), they must be treated as extended sources. Thus, images of the tori were convolved with the magnification maps. We simulated microlensing by the stars in the lens galaxy, in the case of lensed quasars, for different configurations of the lensed system and different values of the torus parameters, in order to estimate (a)\\ amplitudes and timescales of high magnification events, and (b)\\ the influence of geometrical and physical properties of dusty tori on light curves in the infrared domain. From our investigation, we conclude the following. \\begin{enumerate}[(i)] \\item Despite their large size, we found that AGN dusty torus could be significantly magnified by microlensing in some cases. The amplitude of magnification depends on wavelenght, torus parameters, and configuration of the lensed system. \\item The size of torus is wavelength dependent. As a consequence, the magnification amplitude of microlensing events is also wavelength dependent. The magnification is the highest in the near-infrared, decreases rapidly towards the mid-infrared range, and stays almost constant in the far-infrared part of SED. \\item As microlensing is sensitive to the size of the source, parameters determining the geometry and the apparent size of the torus, have a very important role. Tori with $R_{\\text{out}}\\lesssim10$ pc could be appreciably microlensed. \\hspace{10 pt} More compact dust configurations (e.g. steeper radial density profiles) result in smaller tori and thus in higher magnification amplitudes. However, for primary source (accretion disk) luminosites typical for quasars ($10^{12}\\ L_{\\odot}$), the influence of the dust distribution parameters is diminished, because the radiation is able to penetrate the dust further. \\hspace{10 pt} Tori seen at type 1 (dust-free) inclinations, which provide a direct view of the innermost, hottest region, are more magnified than those at type 2 (dust-intercepting) inclinations. \\item Lensed quasar systems with the lens galaxy closer to the observer, will have higher magnification amplitudes, owing to their larger Einstein ring radius projection on the source plane. \\item Estimated rise times, between the beginning and the peak of HMEs, are in the range from several decades to several hundreds of years. \\end{enumerate} Given such long time-scales, microlensing would hardly prove to be a practical tool to study and constrain the properties of dusty tori, as it is in the case of AGN accretion disks. However, the results presented above should be kept in mind when investigating flux ratio anomally of lensed quasar images in different wavelength bands. In such studies, it is important to determine the true magnification ratios between the images, in the absence of microlensing. In principle, this could be done by looking at the emission-line, infrared, and radio-emitting regions of quasars, as they all should be large enough to safely disregard microlensing effects. However, we have shown that the infrared emission of dusty tori could be significantly microlensed in some cases, and thus, it is a less reliable tool for determining the ``intrinsic'' flux ratios." }, "1207/1207.4477_arXiv.txt": { "abstract": "We present a mathematical method to statistically decouple the effects of unknown inclination angles on the mass distribution of exoplanets that have been discovered using radial-velocity techniques. The method is based on the distribution of the product of two random variables. Thus, if one assumes a true mass distribution, the method makes it possible to recover the observed distribution. We compare our prediction with available radial-velocity data. Assuming the true mass function is described by a power-law, the minimum mass function that we recover proves a good fit to the observed distribution at both mass ends. In particular, it provides an alternative explanation for the observed low-mass decline, usually explained as sample incompleteness. In addition, the peak observed near the the low-mass end arises naturally in the predicted distribution as a consequence of imposing a low-mass cutoff in the true-distribution. If the low-mass bins below 0.02M$_{\\rm{J}}$ are complete, then the mass distribution in this regime is heavily affected by the small fraction of lowly inclined interlopers that are actually more massive companions. Finally, we also present evidence that the exoplanet mass distribution changes form towards low-mass, implying that a single power law may not adequately describe the sample population. ", "introduction": "Samples of exoplanetary systems are increasing rapidly thanks to new ground and space-based dedicated surveys, thus enabling investigation of their statistical properties. One of these properties is the planetary mass distribution, a key aspect needed to understand the origin of exoplanets and its relation to the initial mass function. Currently, radial-velocity (RV) detections (e.g. \\citealp{mayor83}; \\citealp{butler96}; \\citealp{jones10}) have provided the largest sample of unconstrained systems. However, the RV technique does not provide masses directly because the line-of-sight inclination angles, $i$, cannot be measured unless complementary observations are carried out, for instance transit photometry (\\citealp{henry00}) or astrometry (\\citealp{benedict06}). Thus, all masses measured with this technique are indeed 'minimum' planet masses, $\\mo=\\mt\\sin{i}$, where $\\mt$, the 'true' planet mass, is not known a priori. Understanding the \\emph{true} mass distribution rather than the minimum mass distribution will allow modelers to compare their mass distributions against a function that is free from one of the largest sources of uncertainty (see \\citealp{mordasini09}; \\citealp{ida05}). Also, the $\\sin i$ degeneracy that plagues RV signals means we can never be fully sure that any individual signature is planetary in nature from the Doppler data alone. This has consequences for a number of aspects of planetary, brown dwarf, and low-mass star studies that deal with inferences drawn from a RV dominated mass distribution. A prime example of this would be the proper location in mass of the planet-brown dwarf boundary (see \\citealp{sahlmann11}), which allows one to clarify the status of an object as either a planet or a brown dwarf (e.g. \\citealp{jenkins09}), and will help us to better understand the formation mechanisms of both classes of objects. It is thought that the $\\sin{i}$ correction is of order unity and would preserve the power-law shape of the observed mass distribution (e.g., \\citealp{jorissen01}; \\citealp{tabachnik02}; \\citealp{hubbard07}; Morton \\& Johnson 2011). However, no proof has been provided {to substantiate} this. Methods to recover the $\\mt$ distribution from the observed minimum-mass data have been proposed based on: (1) numerically solving an Abel-type integral equation that relates observed and true mass distributions (\\citealp{jorissen01}); (2) analytically finding the distribution that maximizes the likelihood of having a given set of minimum masses (\\citealp{zucker01}); (3) comparing cumulative distributions of projected and de-projected data with non-parametric statistical tools (\\citealp{brown11}); and (4) using the physics of multi-planet systems to resolve the $\\sin{i}$ correction (\\citealp{batygin11}). With the exception of (4), which is a theoretical prediction, these methods are non-parametric, i.e., they do not assume an a priori model of the data. However, they also suffer from some drawbacks like the need to smooth the data (1 and 3), or the complexity of introducing observational limits (2). In this paper we present an alternative method to statistically decouple the $\\sin i$ dependence in observed exoplanet mass functions. The method is based on the expected distribution of the product of two continuous and independent random variables. Its parametric nature requires an assumption on the shape of the underlying ($\\mt$) distribution; but, on the other hand, it offers the possibility of introducing observational constraints in a straightforward fashion. The mathematical problem, applied to planet mass distributions, is stated in \\S~\\ref{problem}, and its solution presented in \\S~\\ref{solution}. In \\S~\\ref{examples} the method is implemented on two example distributions and in \\S~\\ref{predictions} we make a comparison with observational data. In \\S~\\ref{discussion} we look at the significance of the shape of the true mass distribution and how this may affect current models of planet formation and evolution. Finally, in \\S~\\ref{summary} we summarize the results. ", "conclusions": "\\label{discussion} We have shown that if the true mass distribution is described by a power-law with boundaries, there must be a peak in the observed mass distribution of exoplanets, and therefore this peak has implications for planet formation and evolution models. A peak in the planetary mass distribution tells us that most of the mass of the proto-planetary disk that goes into forming planets, gets locked up in the formation of the larger gas and ice giants. This agrees with the mass distribution of the planets in the solar system. The widely accepted planetary formation theory is by core accretion and subsequent planet migration (\\citealp{pollack84}; \\citealp{lin86}) and this model can broadly explain the currently observed population of exoplanets. The peak in the planetary mass distribution needs to be taken into account when comparing the outcomes from core accretion formation models against the observed mass distribution of exoplanets, unless the true mass distribution changes form towards the lowest masses. One interesting question that arises is, what does the position of the peak in the mass distribution tell us and what does it mean? As we have seen in Fig.~\\ref{fig_back}, our true mass distribution model can provide a good fit to the observed mass distribution of the current population of exoplanets, particularly the peak and subsequent decline. The low-mass peak we find that best describes the observed data is located around 0.02M$_{\\rm{J}}$, or $\\sim$6.5M$_{\\oplus}$. When we look at systems that have high inclinations, like an observed mass distribution drawn from transiting planets only, we find that for a complete sample, the observed distribution follows the true distribution with no low-mass peak. Therefore, if we assume that the bins are complete below 0.02M$_{\\rm{J}}$, then this is the regime where we begin to observe the effects of systems with low inclination angles, and hence large mass corrections. Here we are assuming we have sampled all angles above a certain limit, $\\imin$, but our model considers the lower likelihood of observing systems with low inclinations; therefore this tells us something important that even a small fraction of systems with low inclinations can produce large changes in the observed mass distribution in the low-mass regime. This result is in line with the implications of the analysis by Ho \\& Turner (2011). These authors demonstrate, using Bayesian analysis, that the {\\it posterior} distribution of angles is determined by the particular true mass distribution and so the latter cannot be simply obtained from the observed one. Unlike Ho \\& Turner, we deal here with the {\\it prior} distribution of angles and study its effects on the {\\it assumed} shape of the true mass distribution; however, both approaches lead to the conclusion that the low inclination systems modify the low-mass end of the observed distribution. Also, since most of the observed distribution above the 0.02M$_{\\rm{J}}$ boundary broadly follows the true distribution then small to medium values of $\\sin i$ do not affect the overall mass distribution in a large manner. We make it clear that most of the low-mass systems in these bins are genuine rocky planets, but that the small numbers of low-inclination/high-mass interlopers cause a dramatic change to the observed mass distribution, again, assuming that the low-mass bins are complete. Finally, we do caution that the true mass distribution we are discussing here applies to planets with small semimajor axes ($\\le$4~AU's or so). The distribution of mass at ever increasing distances from the central star may change the shape of the true mass distribution, but further analysis on this issue is likely to require many more detections like the directly imaged planets around HR8799 (\\citealp{marois08}) or planets discovered by microlensing techniques (e.g. \\citealp{muraki11}). \\rm" }, "1207/1207.1778_arXiv.txt": { "abstract": "In this paper, we investigate the possibility of significant production of thermal bremsstrahlung radiation at radio continuum frequencies that could be linked to some Galactic supernova remnants (SNRs). The main targets for this investigation are SNRs expanding in high density environments. There are several indicators of radio thermal bremsstrahlung radiation from SNRs, such as a flattening at higher frequencies and thermal absorption at lower frequencies intrinsic to an SNR. In this work we discuss the radio continuum properties of 3 SNRs that are the best candidates for testing our hypothesis of significant thermal emission. In the case of SNRs IC443 and 3C391, thermal absorption has been previously detected. For IC443, the contribution of thermal emission at 1 GHz, from our model fit is 3-57\\%. It is similar to the estimate obtained from the thermal absorption properties (10-40\\% at 1 GHz). In the case of the 3C391 the conclusions are not so clear. The results from our model fit (thermal emission contribution of 10-25\\% at 1 GHz) and results obtained from the low frequency absorption (thermal contribution of 0.15-7\\% at 1 GHz) do not overlap. For the SNR 3C396 we suggest that if previously detected thermal absorption could be intrinsic to the SNR then the thermal emission ($<$47\\% at 1 GHz from our model fit) could be significant enough to shape the radio continuum spectrum at high frequencies. Polarization observations for these SNRs can constrain the strength of a thermal component. Reliable observations at low frequencies ($<100$ MHz) are needed as well as more data at high radio frequencies ($>1$ GHz), in order to make stronger conclusions about the existence of \"radio thermally active\" SNRs. ", "introduction": "The radio continuum emission from supernova remnants (SNRs) is believed to be mainly produced by the non-thermal synchrotron mechanism. The radio continuum spectrum is well fitted by a simple power law. On the other hand, the X-ray radiation from SNRs is produced by thermal bremsstrahlung and line radiation as well as by non-thermal synchrotron radiation (Reynolds 2008, and references therein). In this paper we investigate the possibility of significant production of thermal bremsstrahlung radiation at radio frequencies from SNRs that fullfill certain conditions, as discussed below. {The approximate equation for the volume emissivity of thermal bremsstra\\-hlung radiation $\\varepsilon_{\\nu}^{T}$ for an optically thin ionized gas cloud at radio frequencies has the form:} \\begin{equation} \\varepsilon_{\\nu}^{T}=6.8\\times10^{-38}\\ g_{ff}(\\nu, T)\\ n^{2}\\ T^{-0.5}\\ [\\mathrm{erg\\ cm^{-3}\\ s^{-1}\\ Hz^{-1}}], \\end{equation} where the number densities of electrons and ions are the same and given by $n$ in $\\mathrm{cm}^{-3}$, the temperature of the emitting region $T$ is in $\\mathrm{K}$ and the thermally averaged Gaunt factor $g_{ff}(\\nu, T)$ at radio frequencies is given by Gayet (1970) as: \\begin{equation} g_{ff}(\\nu, T)\\approx\\left\\{ \\begin{array}{ll} 0.55\\ln(4.96\\times10^{-2}\\nu^{-1})+0.82\\ln{T}, & 10^{2}\\ \\mathrm{K}$$10^{-12}$ erg~s$^{-1}$~cm$^{-2}$ \\citep{voasbo1999}, our sample extends the fluxes by two orders of magnitude below theirs. Although we also find that stars generally have lower X-ray-to-IR flux ratios than other types of sources, their result that stars have $J-K_{\\rm s}\\le1.1$ is not well supported by our source sample. Our identified and candidate stars have $J-K_{\\rm s}$ spanning from about 0.0 to 8.0, and about 30\\% and 20\\% of them have $J-K_{\\rm s}\\ge1.1$, respectively. The high value of the $J-K_{\\rm s}$ color for these stars is probably caused by large extinction. The lack of these highly absorbed sources in the RASS BSC is easy to explain, as the energy pass of {\\it ROSAT} is in the very soft X-rays (0.1--2.4 keV). Because of these highly absorbed sources, we choose to use the $K_{\\rm s}$-band flux to calculate the X-ray-to-IR flux ratio instead of the $J_{\\rm s}$-band flux used in \\citet{haru2009}. This is also why stars in our sample are not well separated from other types of sources in terms of the X-ray-to-optical flux ratio, though this is observed to be the case in many studies \\citep[e.g.,][]{voasbo1999}." }, "1207/1207.6118.txt": { "abstract": "The upcoming Wide-Field Upgrade (WFU) has ushered in a new era of instrumentation for the Hobby-Eberly Telescope (HET). Here, we present the design, construction progress, and lab tests completed to date of the blue-optimized second generation Low Resolution Spectrograph (LRS2-B). LRS2-B is a dual-channel, fiber fed instrument that is based on the design of the Visible Integral Field Replicable Unit Spectrograph (VIRUS), which is the new flagship instrument for carrying out the HET Dark Energy eXperiment (HETDEX). LRS2-B utilizes a microlens-coupled integral field unit (IFU) that covers a 7\\arcsec$\\times$12\\arcsec\\ area on the sky having unity fill-factor with $\\sim$300 spatial elements that subsample the median HET image quality. The fiber feed assembly includes an optimized dichroic beam splitter that allows LRS2-B to simultaneously observe $370 < \\lambda \\mathrm{(nm)} < 470$ and $460 < \\lambda \\mathrm{(nm)} < 700$ at fixed resolving powers of $R \\approx \\lambda/\\Delta\\lambda \\approx 1900$ and 1200, respectively. We discuss the departures from the nominal VIRUS design, which includes the IFU, fiber feed, camera correcting optics, and volume phase holographic grisms. Additionally, the motivation for the selection of the wavelength coverage and spectral resolution of the two channels is briefly discussed. One such motivation is the follow-up study of spectrally and (or) spatially resolved \\lya\\ emission from $z \\approx 2.5$ star-forming galaxies in the HETDEX survey. LRS2-B is planned to be a commissioning instrument for the HET WFU and should be on-sky during quarter 4 of 2013. Finally, we mention the current state of LRS2-R, the red optimized sister instrument of LRS2-B. ", "introduction": "\\label{sec:intro} % \\label{} allows reference (\\ref{}) to this section The 10 m Hobby-Eberly Telescope (HET) is currently in the process of undergoing a major wide-field upgrade (WFU; Ref. \\citenum{Hill12b}) that will increase its field of view to 22\\arcmin\\ in diameter and feature improved performance in preparation for the upcoming HET Dark Energy eXperiment (HETDEX; Ref. \\citenum{Hill08a}). HETDEX will amass a sample of $\\sim$0.8 million \\lya\\ emitting galaxies (LAE) to be used as tracers of large-scale structure for constraining dark energy and measuring its possible evolution from $1.9 < z < 3.5$. To carry out the 120 night blind spectroscopic survey covering a 420 square degree field (9 Gpc$^{3}$), the HET will be outfitted with a revolutionary new multiplexed instrument called the Visible Integral Field Replicable Unit Spectrograph (VIRUS; Ref. \\citenum{Hill12a}). VIRUS consists of at least 150 copies (with a goal of 192) of a simple fiber-fed integral field spectrograph and for the first time has introduced industrial-scale replication to optical astronomical instrumentation. VIRUS is one of several instruments being prepared for the new era of HET instrumentation that has been ushered in by the WFU. For the past decade, the current Marcario Low Resolution Spectrograph (LRS; Ref. \\citenum{Hill98}) has been a workhorse instrument for the HET, but its design is incompatible with the WFU. This provides an opportunity to redesign and improve upon the capabilities of LRS in a second generation instrument. The design of the second generation LRS (LRS2) is based on the versatile VIRUS unit spectrograph, which was designed to be adaptable to a range of spectral resolutions and wavelength coverage configurations. The original LRS2 design concept (see Ref. \\citenum{Lee10}) is the first demonstration of the wide range of applications for the basic VIRUS design. The original LRS2 concept is fed by a 7\\arcsec$\\times$12\\arcsec\\ unity fill-factor integral field unit (IFU) and simultaneously covers $350 < \\lambda (\\mathrm{nm}) < 1100$ at a fixed resolving power of $R = \\lambda / \\Delta\\lambda \\approx 1800$. The wide spectral coverage is made possible by utilizing two VIRUS unit pairs (i.e., four total spectrograph channels). Utilizing the VIRUS unit spectrograph as the building block for LRS2 has allowed us to take advantage of the large engineering investment made in VIRUS for optimizing it for mass production. LRS2 will be built on the VIRUS production line, which will greatly reduce the final cost and delivery time to relatively low levels for such a capable instrument. While much of the design concept discussed in Ref. \\citenum{Lee10} remains unchanged in the final LRS2 design, there have been some major adjustments. The largest is that the quad-channel simultaneous coverage has been broken up into two separate double-spectrographs that will independently observe the wavelength ranges of $370 < \\lambda (\\mathrm{nm}) < 700$ and $650 < \\lambda (\\mathrm{nm}) < 1050$, respectively. The blue optimized spectrograph pair (LRS2-B) requires only modest adaptation of the VIRUS design while the red optimized pair (LRS2-R) essentially requires the same modifications as LRS2-B in addition to significant work to adapt the VIRUS camera to the differently packaged red sensitive, thick deep-depletion CCDs. When complete, the latter will significantly improve the performance as compared to the current LRS. A significant factor in dividing LRS2 into two separate double-spectrographs was to ensure the continuity of capable low resolution spectroscopy on the HET through the WFU. As such, LRS2-B will be ready for commissioning during quarter 4 of 2013 so that it is available immediately for the upgraded HET. As stable instruments with rapid setup times and high efficiency, LRS2-B and LRS2-R are designed to be highly complementary to VIRUS and to exploit the HET queue-scheduling for survey follow-up, synoptic observations, and targets of opportunity. Since LRS2-B is planned to be a commissioning instrument for the HET WFU, we focus in this work on presenting the details of its final design. In $\\S$\\ref{sec:science}, we describe the science motivation for LRS2 with an emphasis on the drivers for setting the spectral resolution and wavelength coverage of each of the two LRS2-B channels. In $\\S$\\ref{sec:design}, we outline the design of the instrument and focus on the departures from the nominal VIRUS design. In $\\S$\\ref{sec:status}, we discuss the current status of the instrument, including the progress in procuring parts, its construction, and lab testing. In $\\S$\\ref{sec:lrs2-r}, we change our focus to LRS2-R, and give a brief review of its science drivers as well as the additional details of its design that depart from that of LRS2-B. Finally, we briefly discuss the plan for completing the instrument in $\\S$\\ref{sec:outlook}. ", "conclusions": "" }, "1207/1207.2647_arXiv.txt": { "abstract": "Q1: Why deploy $N$ wavefront sensors on a three mirror anastigmat (TMA) and not $N + 1$?\\\\ Q2: Why measure $M$ Zernike coefficients and not $M + 1$?\\\\ Q3: Why control $L$ rigid body degrees of freedom (total) on the secondary and tertiary and not $L + 1$?\\\\ The usual answer: ``We did a lot of ray tracing and $N, M,$ and $L$ seemed OK.''\\\\ We show how straightforward results from aberration theory may be used to address these questions. We consider, in particular, the case of a three mirror anastigmat. ", "introduction": "\\label{sec:intro} % \\subsection{New Wide Field Telescopes and Wavefront Sensing} The next decade will see the construction of several billion-dollar-class wide field telescopes, ground-based and in space, for which the measurement of cosmological weak gravitational lensing is a major programmatic goal.\\footnote{LSST, Euclid, WFIRST} The image quality delivered by these telescopes must be understood to a level never before achieved by astronomical telescopes. At least four different effects will contribute to the observed point spread function (henceforth the PSF): telescope tracking errors, optical misalignments, mirror figure errors, and diffusive and transfer effects within the detector. For ground-based telescopes, atmospheric seeing makes an additional contribution. Wavefront sensors may be deployed to disentangle these effects, and in particular, to measure mirror figure errors and the misalignments of the telescope optics. A recent study of the design for the Large Synoptic Survey Telescope (henceforth the LSST) describes four wavefront sensors measuring Zernike polynomials up to 36th order\\cite{Manuel2010}. But why four, and why 36th order? The answer would appear to be that ray tracing was carried out for many combinations, and that these choices sufficed. While ray tracing simulations may identify wavefront sensing configurations that are effective and efficient, one might still want to understand {\\it why} they are effective and efficient. And insofar as ray tracing simulations are only approximate in their modeling of telescopes and detectors, an understanding of what lies behind the answers can be quite useful. The purpose of this paper is to present a coherent and concise development of such an understanding. \\subsection{Outline of Paper} In \\S2 we describe the generic 3rd order Seidel misalignment aberration patterns that arise in circularly symmetric telescopes. In \\S3 we consider the particular case of a three mirror anastigmat, and show that the 3rd order misalignment aberration patterns do not suffice to produce perfect telescope alignment. In \\S4 we describe the generic 5th order Seidel misalignment aberration patterns that arise in circularly symmetric telescopes. In \\S5 we introduce the {\\it subspace of benign misalignment} defined by the subset of misalignments that produce none of the third order asymmetric aberration patterns described in \\S 3. In \\S6 we answer, in part, the questions of the number of wavefront sensors needed and the order to which the wavefront must be measured. ", "conclusions": "" }, "1207/1207.0642_arXiv.txt": { "abstract": "{Stellar shells observed in many giant elliptical and lenticular as well as a few spiral and dwarf galaxies presumably result from galaxy mergers. Line-of-sight velocity distributions of the shells could, in principle, if measured with a sufficiently high signal-to-noise ratio, constitute a method to constrain the gravitational potential of the host galaxy.} {Merrifield \\& Kuijken (1998) predicted a double-peaked line profile for stationary shells resulting from a nearly radial minor merger. In this paper, we aim at extending their analysis to a more realistic case of expanding shells, inherent to the merging process, whereas we assume the same type of merger and the same orbital geometry.} {We used an analytical approach as well as test particle simulations to predict the line-of-sight velocity profile across the shell structure. Simulated line profiles were convolved with spectral PSFs to estimate peak detectability.} {The resulting line-of-sight velocity distributions are more complex than previously predicted due to nonzero phase velocity of the shells. In principle, each of the Merrifield \\& Kuijken (1998) peaks splits into two, giving a quadruple-peaked line profile, which allows more precise determination of the potential of the host galaxy and contains additional information. We find simple analytical expressions that connect the positions of the four peaks of the line profile and the mass distribution of the galaxy, namely, the circular velocity at the given shell radius and the propagation velocity of the shell. The analytical expressions were applied to a test-particle simulation of a radial minor merger, and the potential of the simulated host galaxy was successfully recovered. Shell kinematics can thus become an independent tool to determine the content and distribution of the dark matter in shell galaxies up to $\\sim$100\\,kpc from the center of the host galaxy.} {} ", "introduction": "Several methods have been used to measure the gravitational potentials and their gradients of elliptical galaxies, including strong and weak gravitational lensing \\citep[e.g.,][]{gavazzi07,mandelbaum08,auger10}, X-ray observations of hot gas in the massive gas-rich galaxies \\citep[e.g.,][]{fukazawa06,nagino09,churazov08,das10}, rotation curves of detected disks and rings of neutral hydrogen \\citep[e.g.,][]{weijmans08}, stellar-dynamical modeling from integrated light spectra \\citep[e.g.,][]{weijmans09,delorenzi09,churazov10,thomas11}, and the use of tracers such as planetary nebulae, globular clusters, and satellite galaxies \\citep[e.g.,][]{coccato09,nierenberg11,deason12,norris12}. All these methods have their limits, such as the redshift of the observed object, the luminosity profile, and gas content. In particular, the use of stellar dynamical modeling is plausible in the wide range of galactic masses, as long as spectroscopic data are available. However, it becomes more challenging beyond a few optical half-light radii. Other complementary gravitational tracers or techniques are required to derive mass profiles in outer parts of the galaxies. When comparing independent techniques for the same objects at similar galactocentric radii, discrepancies in the estimated circular velocity curves were revealed together with several interpretations \\citep[e.g.,][]{churazov10,das10}. The compared techniques usually employ modeling of the X-ray emission of the hot gas (assuming hydrostatic equilibrium) and dynamical modeling of the optical data in the massive early-type galaxies. Therefore, even for the most massive galaxies with X-ray observations available, there is a need for other methods to independently constrain the gravitational potential at various radii. Shell galaxies are galaxies that contain arc-like fine features, which were first noticed by \\citet{arp66}. These structures are made of stars and form open, almost concentric arcs that do not cross each other. Shells are relatively common in elliptical or lenticular galaxies. At least 10\\,\\% of all these galaxies in the local universe possess shells. Nevertheless, shells occur markedly most often in regions of low galaxy density, and perhaps up to half of E and S0 galaxies in these environments are shell galaxies \\citep{malin83,schweizer83,schweizer85,colbert01}. Shells can also be associated with dust \\citep{sikkema07,stickel04} and neutral hydrogen emission \\citep{schiminovich94,schiminovich95,balcells96,petric97,horellou01}. In addition, \\citet{charmandaris00} detected the presence of dense molecular gas in the shells of NGC 5128. Shells are thought to be by-products of minor mergers of galaxies \\citep{quinn84}, although they can also be formed during major mergers \\citep{hernquist92}. The most regular shell systems are believed to result from nearly radial mergers \\citep{dupraz86,hernquist88}. When a small galaxy enters the sphere of influence of a big elliptical galaxy on a radial or close-to-radial trajectory, it disintegrates and its stars begin to oscillate in the potential of the big galaxy. At their turning points, the stars have the lowest speed and thus tend to spend most of the time there, where they pile up and produce arc-like structures in the luminosity profile of the host galaxy when viewed perpendicular to the axis of the collision. Measurement of the number and distribution of shells can, in principle, yield to an approximate estimate of the mass distribution of the host galaxy and the time since the merger \\citep{quinn84,dupraz86,hernquist87a,hernquist87b,canalizo07}. But both of these observables are sensitive to details such as the dynamical friction and the gradual decay of the cannibalized galaxy during the merger \\citep{dupraz87,james87,heisler90,ebrova10}. Moreover, if the core of the cannibalized galaxy survives the merger, new generations of shells are added during each successive passage. This was predicted by \\citet{dupraz87} and successfully reproduced by \\citet{katka11} in self-consistent simulations. All these effects complicate the simulations to such an extent that the interest in shell galaxies largely faded by the end of the 1980s. Recently, this topic has raised interest again, thanks to the discovery of shells in a quasar host galaxy \\citep{canalizo07} and shell structures in M31 \\citep{fardal07,fardal08} and in the Fornax dwarf \\citep{coleman04}. \\citet{helmi03} suggested that ring-like stellar structures, including the one observed in the outer disk of the Milky Way (the so-called Monoceros ring), could be analogous to shells. A significant number of shells is also contained in the early-type galaxy sample of the ongoing ATLAS$^{\\mathrm{3D}}$ project, including images of galaxies with a surface brightness down to 29\\,mag$/$arcsec$^2$ \\citep[see, e.g.,][]{krajnovic11,duc11}. \\citet{kim12} identified shells in about 6\\% of a sample of 65 early-type galaxies from the Spitzer Survey of Stellar Structure in Galaxies (S$^4$G). Shells also appear to be suitable for indirect detection of dark matter via gamma-ray emission from dark matter self-annihilations \\citep{sanderson12}. About 70\\,\\% of a complete sample of nearby (15--50\\,Mpc) luminous ($M_{\\mathrm{B}}<-20$\\,mag) elliptical galaxies were found to show tidal features by Tal et al. (2009). Faint structures, including shells and other signatures of recent gravitational interaction (tidal tails and streams), were found in the Sloan Digital Sky Survey (SDSS). \\citet{kaviraj10} identified 18\\,\\% of early-type galaxies (ETGs) in the SDSS Stripe82 sample as having disturbed morphologies; similarly, \\citet{miskolczi11} found tidal features in 19\\,\\% of their sample of galaxies from SDSS DR7. Observations of warm gas by \\citet{rampazzo03} in five shell galaxies showed irregular gaseous velocity fields (e.g., a double nucleus or elongated gas distribution with asymmetric structure relative to the stellar body), and in most cases, gas and stellar kinematics appear decoupled. \\citet{rampazzo07}, \\citet{marino09}, and \\citet{trinchieri08} investigated star formation histories and hot gas content using the NUV and FUV Galaxy Evolution Explorer (GALEX) observations (and in the latter case also X-ray ones) in a few shell galaxies. The results support accretion events in the history of shell galaxies. \\citet{mk98}, hereafter \\citetalias{mk98}, studied theoretically the kinematics of a stationary shell, a monoenergetic spherically symmetric system of stars oscillating on radial orbits in a spherically symmetric potential. They predicted that spectral line profiles of such a system exhibit two clear maxima, which provide a direct measure of the gradient of the gravitational potential at the shell radius. The first attempt to analyze the kinematical imprint of a shell observationally was made by \\citet{romanowsky12}, who used globular clusters as shell tracers in the early-type galaxy M87. \\citet{fardal12} obtained radial velocities of giant stars in the so-called western shelf in M31 Andromeda galaxy. They successfully analyzed the shell pattern in the space of velocity versus radius. Nevertheless, real-world shells are not stationary features. The stars of the satellite galaxy have a continuous energy distribution, and so at different times, the shell edge is made of stars of different energies, as they continue to arrive at their respective turning points. Thus, the shell front appears to be traveling outwards from the center of the host galaxy and shell spectral-line profiles are more complex (\\citealt{jilkova10}, \\citealt{ebrova11}, see also \\citealt{fardal12}). In this paper, we derive spectral-line profiles of nonstationary shells. We assume that shells originate from radial minor mergers of galaxies, as proposed by \\citet{quinn84}. We find that both of the original \\citetalias[][]{mk98} peaks in the spectral line are split into two, resulting in a quadruple shape, which can still be used to constrain the host galaxy potential and even bring additional information. We outline the simplified theoretical model and derive the shell velocities in Sect.~\\ref{sub:rad_osc}, and describe the origin of the quadruple line profile in Sect.~\\ref{sec:4peak}. In Sect.~\\ref{sec:app}, we derive equations connecting the observable features of the quadruple-peaked line-of-sight velocity distribution (LOSVD) with parameters of the host galaxy potential in the vicinity of the shell edge. We compare these analytical predictions with the theoretical model (Sect.~\\ref{sec:Compars}) and with results of test-particle simulations of the radial minor merger (Sect.~\\ref{sec:N-Simulations}). Section~\\ref{sec:N-Simulations} also demonstrates the derivation of the galactic potential from the simulated spectral data. ", "conclusions": "Kinematics of regular shells produced during nearly radial minor mergers of galaxies can be used to constrain their gravitational force field and thus the dark matter distribution. \\citet{mk98} showed that the LOSVD measured near the edges of a shell has a double-peaked shape, and found a relation between the values of the two LOS velocity peaks and the circular velocity. Their approximation is limited to stationary shells. We have extended their theoretical analysis to traveling shells. We find that in two-component giant galaxies with realistically massive dark matter halos, shell propagation velocity is significantly higher, typically 30--150 km$/$s, compared to values quoted in the theoretical studies in the literature. We show that such large speeds have considerable impact on the LOS kinematics of shells. We demonstrate that each peak of the double-peaked profile is split into two, producing a quadruple-peaked LOSVD. We derive a new approximation, relating the circular velocity of the host galaxy potential at the shell edge radius, as well as the current phase velocity of the shell, to the positions of the four peaks. In galaxies with multiple shells, we can use circular velocities measured by these methods to determine the potential of the host galaxy over a large span in radii, whereas the measured shell phase velocity carries information on the age of the shell system, and the arrival direction of the cannibalized galaxy. The potential observation of multigeneration shell systems contains additional limits on the shape of the potential of the host galaxy." }, "1207/1207.1191_arXiv.txt": { "abstract": "{Radio observations using the Very Long Baseline Interferometry (VLBI) technique typically have fields of view of only a few arcseconds, due to the computational problems inherent in imaging larger fields. Furthermore, sensitivity limitations restrict observations to very compact and bright objects, which are few and far between on the sky. Thus, while most branches of observational astronomy can carry out sensitive, wide-field surveys, VLBI observations are limited to targeted observations of carefully selected objects. However, recent advances in technology have made it possible to carry out the computations required to target hundreds of sources simultaneously. Furthermore, sensitivity upgrades have dramatically increased the number of objects accessible to VLBI observations. The combination of these two developments have enhanced the survey capabilities of VLBI observations such that it is now possible to observe (almost) any point in the sky with milli-arcsecond resolution. In this talk I review the development of wide-field VLBI, which has made significant progress over the last three years.} ", "introduction": "Since the invention of the VLBI technique in the 1960s, observations using it have almost exclusively been limited to the study of carefully selected, small samples of objects. There are two fundamental reasons which have led to this situation. First, the observing bandwidth is restricted by the need to record the raw data on tape or disk and then replay and process the data afterwards on a reasonable timescale. Since the sensitivity of an interferometer is linked to the bandwidth of the observations (i.e., the number of measurements being made), this limitation directly affected the sensitivity. Connected-element interferometers such as the VLA\\footnote{Very Large Array} and ATCA\\footnote{Australia Telescope Compact Array} are not affected by this, since there is no need to store the raw antenna signals before the correlation is carried out. Second, the field of view of a radio interferometer is limited by the spectral and temporal resolution with which the data are correlated (which can also be expressed as the channel width, $\\Delta\\nu$, and integration time, $\\Delta\\tau$, over which the data are averaged). In an often-used analogy the limitation arising from bandwidth averaging can be compared to chromatic aberration in a lens, and the limitation from time averaging to motion blur in a photograph when the exposure time has been too long. Whilst both effects also affect connected-element interferometers, the long baselines and consequently high spatial resolution of VLBI observations limit the field of view to around one arcsecond. Therefore traditional VLBI observations are helplessly unsuitable for surveys of large fractions of the sky. ", "conclusions": "Wide-field VLBI observations are now practical and relatively easy to carry out. The new multi-phase centre mode of the DiFX software correlator allows one to observe a large number of sources simultaneously in a single observation, and the multi-source self-calibration strategy we present here enables noise-limited calibration of these data. We point out that wide-field VLBI observations have a number of applications outside the study of galaxy populations in sensitive surveys: a sample of (relatively) faint background sources can serve as a reference in astrometric observations to measure proper motions and parallaxes; and observations of faint targets are no longer limited by the accuracy of phase corrections derived from interleaved calibrator observations. The full potential of wide-field VLBI observations, however, is expected to be reached in even more sensitive observations of large fields such as the COSMOS survey, or in less sensitive observations of very large areas (tens or hundreds of square degrees)." }, "1207/1207.6064_arXiv.txt": { "abstract": "We consider the astrophysical reaction rates for radiative neutron capture reactions ($n,\\gamma$) in the crust of a neutron star. The presence of degenerate neutrons at high densities (mainly in the inner crust) can drastically affect the reaction rates. Standard rates assuming a Maxwell-Boltzmann distribution for neutrons can underestimate the rates by several orders of magnitude. We derive simple analytical expressions for reaction rates at a variety of conditions with account for neutron degeneracy. We also discuss the plasma effects on the outgoing radiative transition channel in neutron radiative capture reactions and show that these effects can also increase the reaction rates by a few orders of magnitude. In addition, using detailed balance, we analyze the effects of neutron degeneracy and plasma physics on reverse ($\\gamma,n$) photodisintegration. We discuss the dependence of the reaction rates on temperature and neutron chemical potential and outline the efficiency of these reactions in the neutron star crust. ", "introduction": "\\label{S:intro} Nuclear reactions in the atmosphere and the crust of accreting neutron stars affect important observational manifestations such as X-ray bursts and superbursts (e.g., Refs.\\ \\cite{sb06,schatz03,cummingetal05,gu07}) as well as deep crustal heating of neutron stars in X-ray transients (e.g., Refs.\\ \\cite{hz90a,hz03,bbr98,gu07}). In the vicinity of the neutron drip density ($\\rho\\sim 4\\times 10^{11}$~g~cm$^{-3}$ for the cold-catalyzed crust and $\\rho\\sim 6\\times 10^{11}$~g~cm$^{-3}$ for the accreted crust \\cite{hz90a}) and beyond in the inner crust the dense matter contains an increasing amount of free degenerate neutrons (see, e.g., Ref.\\ \\cite{hpyBOOK}). Neutron capture and reverse reactions are important components of nuclear burning under these conditions \\cite{gkm08}. Standard thermonuclear neutron capture rates, which are used in reaction network simulations of nucleosynthesis in stars or supernova explosions, are obtained (e.g., Ref.\\ \\cite{rt00}), assuming the classical Maxwell-Boltzmann distribution of neutrons. However, free neutrons in the neutron star crust can be degenerate, in particular when the density exceeds the neutron drip point~\\cite{hpyBOOK}. For instance, ground-state (cold-catalyzed) matter at $\\rho=6.2 \\times 10^{12}$~g~cm$^{-3}$ has a neutron Fermi energy of $\\approx$ 2.6~MeV \\cite{nv73}. Consequently neutron degeneracy needs to be taken into account for neutron capture rates under such conditions. In addition, the dense stellar plasma of the neutron star crust strongly affects emission, absorption, and propagation of photons \\cite{sy09} and therefore modifies radiative capture and photodisintegration reactions, like ($n,\\gamma$) and ($\\gamma,n$). Because of the high density, the electron plasma frequency $\\omega_\\mathrm{p}$ can be of the order of or higher than characteristic frequencies of radiative transitions in nuclei. Under these conditions, well-defined elementary electromagnetic excitations (photons or plasmons) become either suppressed or forbidden (e.g., Ref.~\\cite{abr84eng}) although radiative transitions are not suppressed because they can be realized by emission (or absorption) of excess energy to (from) the plasma as a collective system \\cite{sy09}. These plasma physics effects can be important since they may enhance the radiative transition strength. In Sec.\\ \\ref{S:deg_n} we discuss the effects of neutron degeneracy on ($n,\\gamma$) radiative neutron capture reactions in dense matter. In Sec.\\ \\ref{S:plasma} we analyze plasma effects on the outgoing radiative transition channel of ($n,\\gamma$) reactions. In Sec.\\ \\ref{S:reverse} we consider the same neutron degeneracy and plasma physics effects on inverse ($\\gamma,n$) photodisintegration reactions. We discuss our results in Sec.\\ \\ref{S:discuss} and summarize them in Sec.\\ \\ref{S:concl}. For brevity, we use the units in which the Boltzmann constant $k_B=1$. ", "conclusions": "\\label{S:concl} We have considered neutron captures ($n,\\gamma$) in dense stellar matter, taking into account the effects of neutron degeneracy and plasma physics. The effects of neutron degeneracy increase the amount of high-energy neutrons and mainly enhance the reaction rates; plasma physics effect enhance the radiative transition in the outgoing channel and enhance the reaction rates as well. The effects of neutron degeneracy on neutron capture reaction rates can be quantified by introducing the ratio $R_n$, Eq.~(\\ref{eq:R}), of rates calculated for given conditions to those for non-degenerate neutrons. We have described this ratio by a simple analytic expression (\\ref{eq:R_powerlaw}) assuming the power-law energy dependence of the reaction cross section (\\ref{eq:sigma_pl}) at energies that are not too high. The derived expression (\\ref{eq:R_powerlaw}) seems sufficient for many applications. Furthermore, approximating $\\sigma(E)$ by a power-law function (\\ref{eq:sigma_pl}), one also obtains the power-law index $\\nu$ and the maximum energy $E_\\mathrm{max}$ to which the power-law approximation of $\\sigma(E)$ is valid. $E_0$ and $\\nu$ are needed in Eq.~(\\ref{eq:R_powerlaw}), while $E_\\mathrm{max}$ controls the validity of Eq.~(\\ref{eq:R_powerlaw}). Our conclusions are as follows: \\begin{enumerate} \\item Neutron degeneracy can significantly affect ($n,\\gamma$) reactions in deep neutron star crust (Sec.\\ \\ref{S:deg_n}). In many cases the effects of neutron degeneracy are well described by the factor $R_n$ given by Eq.~(\\ref{eq:R_powerlaw}). For threshold reactions, strong neutron degeneracy enhances the reaction rate by many orders of magnitude. \\item Plasma physics effects can additionally enhance ($n,\\gamma$) rates (Sec.\\ \\ref{S:plasma}), that is described by the factor $R_\\mathrm{pl}$. These effects are less dramatic but can reach a few orders of magnitude. \\end{enumerate} Furthermore, in Sec.\\ \\ref{S:reverse} we have used the detailed balance principle and calculated the rates of inverse ($\\gamma,n$) reactions taking into account neutron degeneracy and plasma effects. Finally, in Sec.\\ \\ref{S:discuss} we discussed the efficiency of ($n,\\gamma$) reactions in a neutron star crust, with the conclusion that neutron degeneracy can be most important. Finally it should be noted that free degenerate neutrons in a neutron star crust can be in a superfluid state. Critical temperature $T_\\mathrm{cn}$ for the appearance of neutron superfluidity is very model dependent. Numerous calculations using different techniques (e.g., Ref.~\\cite{ls01}) give density-dependent $T_\\mathrm{cn}(\\rho)$ with maximum values ranging from $\\sim 0.2-0.3$ MeV to $\\sim 2$~MeV, indicating that superfluidity is most likely. Superfluidity produces a gap in the energy spectrum of neutrons near the Fermi level and modifies matrix elements of neutron capture reactions. Both effects on neutron captures are not explored but may strongly modify the reaction rates. Our consideration of neutron degeneracy and plasma effects on neutron capture rates is simplified. A more rigorous (and complicated) analysis of these effects (including also neutron superfluidity) would be desirable. It would be instructive to perform self-consistent calculations of the structure of atomic nuclei immersed in a Fermi sea of free neutrons, taking into account a compression of the nuclei by free neutrons (e.g., Refs.\\ \\cite{st83,hpyBOOK}). For simplicity, we have used a model of free neutrons which occupy the space between atomic nuclei. At densities $\\rho$ not much higher than the neutron drip density, it is sufficiently accurate (as follows, for instance, from results of Ref.\\ \\cite{nv73}). In a self-consistent approach, this model should be replaced by a more elaborated unified treatment of neutrons bound in nuclei and free outside. In any case one should bear in mind that neutron capture reactions in a deep neutron star crust can be affected by neutron degeneracy, plasma physics, and neutron superfluidity. These effects may have important impacts on nuclear burning and nucleosynthesis in the deep neutron star crust. The effects should be taken into account to correctly simulate and interpret various observational phenomena in accreting neutron stars such as X-ray bursts and superbursts as well as quiescent thermal emission of neutron stars in X-ray transients (e.g., Refs.\\ \\cite{sb06,schatz03,gu07} and references therein)." }, "1207/1207.6408_arXiv.txt": { "abstract": "It is believed that first-order phase transitions at or around the GUT scale will produce high-frequency gravitational radiation. This radiation is a consequence of the collisions and coalescence of multiple bubbles during the transition. We employ high-resolution lattice simulations to numerically evolve a system of bubbles using only scalar fields, track the anisotropic stress during the process and evolve the metric perturbations associated with gravitational radiation. Although the radiation produced during the bubble collisions has previously been estimated, we find that the coalescence phase enhances this radiation even in the absence of a coupled fluid or turbulence. We comment on how these simulations scale and propose that the same enhancement should be found at the Electroweak scale; this modification should make direct detection of a first-order electroweak phase transition easier. ", "introduction": "As the bubbles expand and begin to collide, individual collisions are seen to contribute to the power spectrum. These are no longer apparent by $\\tau \\approx \\beta^{-1}/2$, after which time $h^2 \\Omega_{gw}$ rises more steadily, as shown in Fig. \\ref{earlyheightbinv}. We begin by considering the time interval $0 < \\tau < \\beta^{-1}$. \\subsection{Collisions: $0 < \\tau < \\beta^{-1}$ } \\begin{figure}[htbp] \\centering \\includegraphics[width=3.25in]{heightvsbetainv.eps} % \\caption{The maximum intensity of the gravitational wave spectrum for a $\\mu = 10^{-4}\\,m_{\\rm pl}$ simulation initialized with 40 (red, solid), 32 (blue, dotted), 24 (green, dashed), or 16 (black, dot-dashed) bubbles per Hubble volume, $\\tau<\\beta^{-1}$.} \\label{earlyheightbinv} \\end{figure} During this stage, the peak frequency of the gravitational radiation should correspond to the scale $a_e\\beta^{-1}=L_*$ \\cite{Kamionkowski:1993fg}, the mean distance between bubbles on the initial slice. This corresponds to a physical wavenumber $k_{\\rm{phys}} = 2 \\pi L_*^{-1} \\approx 11.34 n^{1/3} H_* $. The associated frequency observed at the present day is given by Eq.~\\ref{freqtransfer}, where $H_e$ is the Hubble constant at the time when we calculate the spectrum; $H_e = \\frac{H_0}{a^2_e} \\sim H_0$, where $a_e$ is the scale factor at the time when we take the spectrum. Putting this together, we expect the peak frequency to occur at \\begin{equation} f_{\\rm peak} \\approx 6.8\\times 10^{11}n^{1/3}\\sqrt{\\frac{H_0}{m_{\\rm pl}}} \\,{\\rm Hz} \\approx 1.16\\times10^{12}\\mu\\,n^{1/3}\\,{\\rm Hz} \\end{equation} This corresponds to a peak around $f_{\\rm peak}\\sim 10^8\\,{\\rm Hz}$ when $\\mu=10^{-4}\\, m_{\\rm{pl}}$ and $n\\sim10$. These numbers are consistent with the location of a low-frequency peak visible in the spectrum at early times; see Fig.~\\ref{separations1}. \\begin{figure}[htbp] \\centering \\includegraphics[width=3.25in]{separations1.eps} % \\caption{The present-day gravitational wave spectrum produced by time $\\tau=\\beta^{-1}$. This is for a simulation with $\\mu=10^{-4}\\,m_{\\rm pl}$ and $n=16$ bubbles per Hubble volume (red, solid), $n=24$ (blue, dotted), $n=32$ (green, dashed), and $n=40$ (black, dot-dashed). The bump at high frequencies is a numerical artifact.} \\label{separations1} \\end{figure} The frequency at which the maximum intensity of gravity waves is expected to be found is also given in \\cite{Kamionkowski:1993fg} as \\begin{equation} f_{\\rm{peak}} \\approx 5.2\\times10^{-8}\\,{\\rm Hz}\\left(\\frac{\\beta}{H_*}\\right)\\left(\\frac{T_*}{1\\,{\\rm GeV}}\\right)\\left(\\frac{g_*}{100}\\right)^{1/6} \\end{equation} where $g_*$ is the number of ultra-relativistic degrees of freedom at the time of the transition, $\\beta$ is, again, the nucleation rate, and $T_*=\\mu$ is the energy density when the phase transition occurs. Taking $g_*=400$, \\begin{equation} f_{\\rm{peak}}=1.18 n^{1/3} \\mu \\times 10^{12}, \\end{equation} which predicts the spectrum to peak at frequencies three or four times $\\mu \\times 10^{12}$, depending on the number of bubbles. Fig.~\\ref{separations1} confirms that for $\\mu=10^{-4}$, peak frequencies occur around this value. A more recent calculation of the peak frequency \\cite{Huber:2008hg} predicts \\begin{equation} f=16.5 \\times 10^{-6} \\, \\rm{Hz} \\left( \\frac{f_*}{\\beta} \\right)\\left(\\frac{\\beta}{H_*}\\right)\\left(\\frac{T_*}{100\\,{\\rm GeV}}\\right)\\left(\\frac{g_*}{100}\\right)^{1/6} \\end{equation} where we use the same values as before and approximate $f_*/\\beta$ by \\begin{equation} \\frac{f_*}{\\beta}=\\frac{0.62}{1.8-0.1v+v^2}=0.22963 \\end{equation} when the bubble wall velocity $v \\approx 1$ (the scalar bubbles in our simulation accelerate quickly to $v\\approx 1$). This prediction of the frequency simplifies to \\begin{equation} f_{\\rm{peak}}=8.60 n^{1/3} \\mu \\times 10^{11}, \\end{equation} which also predicts a few times $10^{8}$, also consistent with Fig.~\\ref{separations1}. In addition to calculating the peak frequency, \\cite{Kamionkowski:1993fg} also gives the fraction of critical density found in gravity waves as \\begin{equation} \\begin{split} \\Omega_{gw}h^2 =& 1.1 \\times 10^{-6} \\kappa^2 \\left(\\frac{H_*}{\\beta}\\right)^2\\left(\\frac{\\alpha}{1+\\alpha}\\right)^2\\\\ &\\left(\\frac{v^3}{0.24+v^3}\\right)\\left(\\frac{100}{g_*}\\right)^{1/3} \\end{split} \\end{equation} where the efficiency factor $\\kappa$ measures energy lost to the motion of a coupled fluid. We set $\\kappa = 1$, as our simulations do not include a fluid. The fraction $H_*/\\beta$ is the mean initial bubble separation as a fraction of the Hubble distance, $v$ is the velocity of the bubble walls, and $g_*$ is, again, the number of ultra-relativistic degrees of freedom. In our simulations we set $\\alpha = 1/3$. Assuming $v=1$ and $g_* = 400$, the expected intensity of the spectrum reduces to \\begin{equation} \\Omega_{gw}h^2=1.42n^{-2/3} \\times 10^{-9}. \\end{equation} This estimate varies between 90 and 1250 times greater than simulation results at $\\tau = \\beta^{-1}$, becoming more consistent in the large $n$ limit. The more recent prediction \\cite{Huber:2008hg} of intensity as \\begin{equation} \\begin{split} \\Omega_{gw}h^2 = &1.67 \\times 10^{-5} \\left(\\frac{0.11v^3}{0.42+v^2}\\right) \\kappa^2 \\left(\\frac{H_*}{\\beta}\\right)^2 \\\\ &\\left(\\frac{\\alpha}{1+\\alpha}\\right)^2\\left(\\frac{100}{g_*}\\right)^{1/3} \\\\ = & 2.07n^{-2/3} \\times10^{-9} \\end{split} \\end{equation} falls slightly farther from our results. There are two main factors that describe the discrepancies: (1) we have implemented an expanding background in which Hubble friction depletes the energy of the source by a small fraction and (2) our models realistically thicken the bubble walls. ``Thick\" walls prolong collisions and dilute the gradient terms that source strong gravitational waves. At the same time the inclusion of these effects strengthens the validity of the current model. \\subsection{Coalescence: $\\beta^{-1} < \\tau < 2 \\beta^{-1}$:} Although most of the volume of the simulations is in the true minimum by $\\tau=\\beta^{-1}$, there is still a lot of kinetic, gradient and even potential energy in the fields. The phase transition is, more or less, complete, but the production of gravitational radiation has not ceased. Fig.~\\ref{lateheightbinv} shows how the peak amplitude rises from $\\tau=\\beta^{-1}$ until $\\tau= 2.5\\beta^{-1}$. \\begin{figure}[htbp] \\centering \\includegraphics[width=3.25in]{heightvsbetainv2.eps} % \\caption{The maximum intensity of the gravitational wave spectrum for a $\\mu = 10^{-4}\\,m_{\\rm pl}$ simulation initialized with 40 (red, solid), 32 (blue, dotted), 24 (green, dashed), or 16 (black, dot-dashed) bubbles per Hubble volume, $\\beta^{-1}<\\tau<2.5\\beta^{-1}$.} \\label{lateheightbinv} \\end{figure} We expect that a period of turbulence can increase the magnitude of the spectrum by several orders of magnitude, e.g. \\cite{Kamionkowski:1993fg,Caprini:2006jb,Gogoberidze:2007an,Caprini:2007xq,Kahniashvili:2008pf,Caprini:2009fx}. One of these estimates \\cite{Kamionkowski:1993fg} say that the intensity of gravitational radiation after turbulence is predicted to be \\begin{equation} \\Omega_{gw}h^2 = 10^{-5}\\left(\\frac{H_0}{\\beta}\\right)^2v v_0^6\\left(\\frac{100}{g_*}\\right)^{1/3}, \\end{equation} which reduces to \\begin{equation} \\Omega_{gw}h^2=2.55n^{-2/3} \\times 10^{-7}. \\end{equation} Thus, the intensity of the gravitational wave spectrum is expected to be around order $10^{-8}$ if we take $v_0 \\approx 1$. We can do slightly better if we try to assign a sound speed, $v_0$, by estimating the speed of perturbations in the true vacuum. Near $\\phi = -\\phi_0$, the effective mass of the field, $m^2_{\\rm eff} = \\lambda \\phi_0^2$, and \\begin{equation} v_0 = \\frac{\\partial \\omega}{\\partial k} = \\sqrt{\\frac{1}{1+\\lambda \\phi_0^2/k^2}}. \\end{equation} This ranges between $10^{-2}$ for low frequency modes, $k_{\\rm low} \\sim H_* \\approx \\sqrt{\\lambda}\\phi_0/500$, and almost 1 for higher frequency modes, $k\\sim\\sqrt{\\lambda}\\phi_0$. The final amplitudes that we present here, see Fig.~\\ref{lateheightbinv}, are some three orders of magnitude lower, but this is understandable as there is no turbulence {\\sl per se} in our simulation. There is, however, significant post-collisionary amplification of the gravitational wave spectrum. This period of {\\sl coalescence} after $\\tau = \\beta^{-1}$ of the simulation, {\\sl amplifies the spectrum by more than an order of magnitude}, depositing energy in higher frequency modes. The frequency of the turbulence peak is \\cite{Kamionkowski:1993fg} \\begin{equation} \\begin{split} f_{\\rm{peak}} &\\simeq 2.6\\times10^{-8}\\,{\\rm Hz}\\,v_0v^{-1}\\left(\\frac{\\beta}{H_*}\\right)\\left(\\frac{T_*}{1\\,{\\rm GeV}}\\right)\\left(\\frac{g_*}{100}\\right)^{1/6}\\\\ &=5.91 \\times 10^{11}n^{1/3} \\mu \\times 10^{11} \\end{split} \\end{equation} This predicts the peak to shift downward during turbulence, while we see a higher-frequency peak at the end of the simulation, with frequency $f \\sim \\mathcal{O}(\\mu) \\times 10^{13}$. This peak comes from the amplification of higher frequency modes during the coalescence period. Fig.~\\ref{separations2} shows the integrated gravitational wave spectrum for $\\mu = 10^{-4}$ after this period. In this plot the peak at higher frequencies is quite apparent at $\\tau=2.5\\beta^{-1}$. \\begin{figure}[htbp] \\centering \\includegraphics[width=3.25in]{separations2.eps} % \\caption{The present-day gravitational wave spectrum at $t=2\\beta^{-1}$, nearing the end of the coalescence phase. This is for a simulation with $\\rho_*=(10^{-4}\\,m_{\\rm pl})^4$ and 16 bubbles per Hubble volume $H_*^{-3}$ (red, solid), 24 bubbles per Hubble volume (blue, dotted), 32 bubbles per Hubble volume (green, dashed), 40 bubbles per Hubble volume (black, dot-dashed). The bump at high frequencies is a numerical artifact.} \\label{separations2} \\end{figure} It should be noted that Fig.~\\ref{separations2} does not explicitly show the $k^3$ low-frequency tail of the gravitational wave spectrum. This is due to the lack of resolution at the relevant scales; the longest wavelength that we can resolve is well within the horizon at the time of the phase transition. The leftmost few points in all of our spectra are averaged only over a small number of modes and should not be used to extrapolate to very low frequencies. Lastly, it's important to check how the spectrum varies with energy scale. In Fig.~\\ref{scales} we vary the energy scale, $\\mu$, between $10^{-6}\\,m_{\\rm pl}$ and $10^{-4}\\,m_{\\rm pl}$. The amplitude of the gravity wave spectrum should be independent of scale of the simulation, $\\mu$ (also $T_*$ in \\cite{Kamionkowski:1993fg,Huber:2008hg} among others), and should only depend on $\\beta^{-1}$ and dynamical factors. Indeed, we recover the scale-independent behavior for the three orders of magnitude that we test. \\begin{figure}[htbp] \\centering \\includegraphics[width=3.25in]{scales.eps} % \\caption{The present-day gravitational wave spectra at $\\tau=2.5\\beta^{-1}$ for cases with 32 bubbles per Hubble volume at three energy scales: $\\mu = 10^{-6}\\,m_{\\rm pl}$ (green, dashed), $10^{-5}\\,m_{\\rm pl}$ (blue, dashed) and $10^{-4}\\,m_{\\rm pl}$ (red, solid). The bump at high frequencies is a numerical artifact. } \\label{scales} \\end{figure} We anticipate that our simulations would continue to produce similar spectra even at much lower energy scales (modified only slightly when the number of ultra-relativistic degrees of freedom, $g_e$, decreases). ", "conclusions": "Gravitational radiation should be the most obvious relics of first-order phase transitions that may have existed in the history of the Universe. To the best of our knowledge, these results represent the first 3-dimensional simulations of first-order cosmological phase transitions and the highest-resolution lattice gravitational wave predictions to date. Using these simulations, without a coupled fluid and with only scalar degrees of freedom, we have confirmed previous analytic and numerical simulations of gravity waves from first-order processes, and have reproduced the predicted scalings of the results both in frequency and amplitude. We had identified the two relevant stages of the process: (1) the stage during which the bubbles collide and (2) a coalescence phase during with the field settles into the true vacuum. During the first stage, we precisely reproduce the location of the peak of gravitational radiation from previous estimates. The amplitude at this time is lower than expected from these estimates, due to the inclusion of friction and by realistically ``thickening\" the walls. Primarily, though, we have discovered that a coalescence phase following the phase transition amplifies the gravitational wave signal for a first-order phase transition by about an order of magnitude and increases the peak frequency by about a decade. This is most likely a consequence of the persistence of energy in domain walls, even after regions have collided, e.g. \\cite{Hawking:1982ga}, residual anisotropic stress-energy produced as the Universe relaxes to a thermal state. This compensates for the lack of power in gravity waves at $t\\approx \\beta^{-1}$. The electroweak phase transition will occur at a much lower frequency than those presented here. If we estimate this energy scale as $\\sim 200\\,{\\rm GeV}$, then we expect the peak frequency of gravitational radiation from the phase transition to occur at a few times $10^{-5}\\,{\\rm Hz}$, assuming $\\beta/H_*\\approx 5$ as we have here. Coalescence will both amplify the signal and raise the peak frequency about an order of magnitude. We believe that this is an important effect to be considered in the next generation of detector experiments. We intend to follow up on this model by adding dynamical fluids to our 3-dimensional simulations. The evolution of fluids alongside our scalar fields will allow us to confirm the parametric dependence of the gravitational wave signal on the terminal velocity of the bubble walls, check for the amplification of the signal due to the existence of turbulence and confirm that a coalescence phase exists in the presence of a viscous fluid." }, "1207/1207.4916_arXiv.txt": { "abstract": "Extended UltraViolet (\\xuv) disks have been found in a substantial fraction of late-type --S0, spiral and irregular-- galaxies. Similarly, most late-type spirals have an extended gas disk, observable in the 21cm radio line (\\hi). The morphology of galaxies can be quantified well using a series of scale-invariant parameters; Concentration-Asymmetry-Smoothness (CAS), Gini, \\m20, and $G_M$ parameters. In this series of papers, we apply these to \\hi\\ column density maps to identify mergers and interactions, lopsidedness and now \\xuv\\ disks. In this paper, we compare the quantified morphology and effective radius ($R_{50}$) of the Westerbork observations of neutral Hydrogen in Irregular and Spiral galaxies Project (\\whisp) \\hi\\ maps to those of far-and near-ultraviolet images obtained with \\galex, to explore how close the morphology and scales of \\hi\\ and \\uv\\ in these disks correlate. We find that \\xuv\\ disks do not stand out by their effective radii in \\uv\\ or \\hi. However, the concentration index in \\fuv\\ appears to select some \\xuv\\ disks. And known \\xuv\\ disks can be identified via a criterion using Asymmetry and \\m20; 80\\% of \\xuv\\ disks are included but with 55\\% contamination. This translates into 61 candidate \\xuv\\ disk out of our 266 galaxies, --23\\%-- consistent with previous findings. Otherwise, the \\uv\\ and \\hi\\ morphology parameters do not appear closely related. Our motivation is to identify \\xuv\\ disks and their origin. We consider three scenarios; tidal features from major mergers, the typical extended \\hi\\ disk is a photo-dissociation product of the \\xuv\\ regions and both \\hi\\ and \\uv\\ features originate in cold flows fueling the main galaxy. We define extended \\hi\\ and \\uv\\ disks based on their concentration ($C_{HI} > 5$ and $C_{FUV} > 4$ respectively), but that these two subsamples never overlap in the \\whisp\\ sample. This appears to discount a simple photo-dissociation origin of the outer \\hi\\ disk. Previously, we identified the morphology space occupied by ongoing major mergers. Known \\xuv\\ disks rarely reside in the merger dominated part of \\hi\\ morphology space but those that do are Type 1. Exceptions, \\xuv\\ disks in ongoing mergers, are the previously identified UGC 4862 and UGC 7081, 7651, and 7853. This suggests cold flows as the origin for the \\xuv\\ complexes and their surrounding \\hi\\ structures. ", "introduction": "Introduction} Interest in the outskirts of spiral galaxy disks has increased over recent years as these regions are the site of the most recent acquisition of gas for these systems \\citep[e.g.][]{Sancisi08}, as well as low-level star-formation \\cite[e.g.,][for recent results]{Dong08,Bigiel10b, Alberts11}, making these faint outskirts the interface between the island universes --the galaxies themselves-- and the cosmic web of primordial gas. The low-level star-formation was first discovered in H$\\alpha$ emission by \\cite{Ferguson98a} and \\cite{Lelievre00}. After the launch of the Galaxy Evolution EXplorer \\citep[{\\sc galex},][]{galex}, initial anecdotal evidence pointed to ultraviolet disks of spiral galaxies extending much beyond their optical radius \\citep{Thilker05a,Thilker05b, Gil-de-Paz05, Gil-de-Paz07, Zaritsky07}. Subsequent structural searches for these Extended Ultraviolet (\\xuv) Disks by \\cite{Thilker07b} and \\cite{Lemonias11} find that some 20--30\\% of spirals posses an \\xuv\\ disk and 40\\% of S0s \\citep{Moffett11}, making this type of disk common but not typical for spiral and S0 galaxies. These \\xuv\\ disk complexes are generally $\\sim100$ Myr old, explaining why most lack H$\\alpha$ \\citep{Alberts11}, as opposed to a top-light IMF \\citep[as proposed by][]{Meurer09}, and sub-solar but not excessively low metallicities \\citep[][$0.1-1 ~ Z_\\odot$, based on emission lines]{Gil-de-Paz07,Bresolin09a, Werk10}. Additionally, it has been known for some time now that atomic hydrogen (\\hi) as observed by the 21cm fine structure line also extends well beyond the optical disk of spiral galaxies \\citep[e.g.,][]{Begeman89, Meurer96, Meurer98, Swaters02, Noordermeer05, Walter08, Boomsma08, Elson11, Heald11a, Heald11b, Zschaechner11b}. In those few cases where both high-quality \\hi\\ and deep {\\sc galex} data are available, a close relation in their respective morphology was remarked upon \\citep{Bigiel10b}. While we have to wait for the all-sky surveys in \\hi\\ to catch up to the coverage of the {\\sc galex} surveys (e.g., the wide survey with WSRT/APERTIF or the WALLABY survey with ASKAP), we can compare the morphology for those galaxies for which uniform \\hi\\ information is available. The canonical view of the origin of the \\xuv\\ disks, is that the Kennicutt-Schmidt law \\citep{Kennicutt98} needs to be extended to low global surface densities of gas and the formation of individual O-stars in the very outskirts of disks \\citep{Cuillandre01,Bigiel10c}. The recent accretion of cold gas flows \\citep[][]{Keres05} into the \\hi\\ disk is the origin for the young stars \\citep[the fueling rate implied by XUV disks is explored in][]{Lemonias11} and \\hi\\ warps\\citep[][]{Roskar10a}. The fraction of spirals that have an \\xuv\\ disk ($\\sim20-30$\\%) supports this scenario as the remaining spirals may simply have no current cool gas inflow. In this case, one would expect \\uv\\ and \\hi\\ morphology to follow each other reasonably closely for the \\xuv\\ disks but not for many of the others, as the star-formation in the inner disk is more closely related to the molecular phase \\citep{Bigiel08}. However, the existence of \\xuv\\ disks pose an intriguing alternate possibility for the origin of the atomic hydrogen disk. Instead of primordial gas accreting onto the disk, the \\hi\\ disk could also be the byproduct of photodissociation of molecular hydrogen on the `skins' of molecular clouds by the ultraviolet flux of the young stars in the \\xuv\\ disk \\citep[see][]{Allen97, Allen02, Allen04}. This explanation has been explored in the stellar disks of several nearby galaxies \\citep{Heiner08a, Heiner08b, Heiner09,Heiner10}. \\cite{Gil-de-Paz07} calculate the time-scales (molecular gas dissociation and re-formation) involved but these are inconclusive regarding the origin of the \\xuv\\ disk. In this scenario, one would expect the \\uv\\ and \\hi\\ morphologies to follow each other closely in all cases; in the outer disk, the \\uv\\ flux from a few young stars would reach out to large areas of the low-column density gas to dissociate enough hydrogen to form the outer \\hi\\ disk. Thus, the low-flux and low \\hi\\ column density morphology --those defining the limits and extent of the \\xuv\\ and \\hi\\ disks-- should show a close relation in parameters such as concentration, Gini, \\m20\\ and the effective radius. A third explanation is in terms of recent tidal interaction. A major merger often pulls gas out of the planes of galaxies and triggers star-forming events. Some anecdotal evidence \\citep[e.g., UGC 04862 is a late-stage major merger][]{Torres-Flores12} does point to this possible origin. In this scenario, one would expect that most \\hi\\ disks hosting an \\xuv\\ disk would be seriously tidally disrupted. In this series of papers, we have explored the quantified morphology of available \\hi\\ maps with the common parameters for visible morphology; concentration-asymmetry-smoothness, Gini and \\m20 and $G_M$. In \\cite{Holwerda11a}, we compare the \\hi\\ morphology to other wavelengths, noting that the \\hi\\ and ultraviolet morphologies are closely related. In subsequent papers of the series, we use the \\hi\\ morphology to identify mergers \\citep{Holwerda11b}, their visibility time \\citep{Holwerda11c} and subsequently infer a merger rate from \\whisp\\ \\citep{Holwerda11d} as well as identify phenomena unique to cluster members \\citep{Holwerda11e}. In this paper, we explore the morphological link between the \\hi\\ and \\xuv\\ disks in the Westerbork \\hi\\ Survey Project \\citep[\\whisp,][]{whisp,whisp2}, a survey of several hundred \\hi\\ observations of nearby galaxies. We complement this data with {\\sc galex} images to explore the morphology and typical scales of these maps. A direct and quantified comparison between the gas and ongoing star-formation morphology could help answer open questions regarding the origin and nature of \\xuv\\ disks: how do their respective sizes relate? Are their morphologies closely related in every case? Do their respective morphologies point to a dominant formation mechanism; gas accretion, photodissociation or tidal? Are \\xuv\\ disks in morphologically distinct or typical \\hi\\ disks? Are \\xuv\\ disks embedded in \\hi\\ disks that appear to be in an active interaction? What is the relation between \\uv\\ flux and \\hi\\ column density in the \\xuv\\ disks, especially the outer disk? The paper is organized as follows; \\S \\ref{s:morph} gives the definitions of the quantified morphology parameters we employ, and \\S \\ref{s:data} describes the origin the data. We describe the application of the morphological parameters and the results in \\S \\ref{s:app} and \\ref{s:analysis}, and we discuss them in \\S \\ref{s:disc}. We list our conclusions and discuss possible future work in \\S \\ref{s:concl}. ", "conclusions": "\\label{s:concl} Based on the quantified morphology of the \\hi\\ and \\uv\\ maps of the \\whisp\\ sample of galaxies, we conclude the following: \\begin{itemize} \\item[1.] There are distinct galaxy populations that stand out by their high concentration values in \\fuv\\ or \\hi\\ (\\xuvc\\ and \\xhi\\ thoughout the paper). These population do not overlap. Some known \\xuv\\ disks are in the \\xuvc\\ sample (Figures \\ref{f:morphfuv}-a \\ref{f:morphfuv}-e, \\ref{f:thingsfuv}-a and \\ref{f:thingsfuv}-c).To compute this concentration index, the outer \\hi\\ contour is needed to delineate the extent of the disk. \\item [2.] The fact that relative dilute \\uv\\ disks ($C_{FUV} > 4$) and \\hi\\ disks ($C_{HI} > 5$) population do not overlap at all and the lack of close morphological relations in any other parameters together suggest that a common small-scale origin for the \\uv\\ and \\hi\\ disks, such as a photodissociation of molecular hydrogen scenario, is less likely for \\xuv\\ disks (but may still very much hold for the inner disks). \\item [3.] Asymmetry and \\m20\\ can be used in combination to select \\xuv\\ disks with reasonable reliability (80\\% included) but substantial contamination (55\\%), when properly calibrated for the survey's spatial resolution (equation \\ref{eq:things:M20A} and \\ref{eq:whisp:M20A}) and the morphologies computed over a large enough aperture. With this selection, one can find the number of \\xuv\\ disks in a survey or candidates for visual classification (Figure \\ref{f:fuvmorph}). \\item [4.] Based on the morphology of the \\hi\\ disk in which they occur, \\xuv\\ disks appear not to occur often in tidally disturbed gas disks (Figures \\ref{f:himorph} and \\ref{f:hithings}). \\item [5.] In a few cases, the \\xuv\\ disk {\\em is} the product of a major merger; for example UGC 04862, identified by \\cite{Thilker07b} and UGC 7081, 7651, and 7853, identified by their large \\fuv\\ concentration. This appears to be an avenue to form {\\em some} of the Type 1 \\xuv\\ disks. \\item [6.] The \\hi\\ morphology and anecdotal evidence for small satellite cannibalism all point to a third mechanism for the origin of most \\xuv\\ disks; cold flow accretion. \\end{itemize} With the emergence of new and refurbished radio observatories in preparation for the future Square Kilometre Array \\citep[SKA;][]{ska}, a new window on the 21 cm emission line of atomic hydrogen gas (\\hi) is opening. The two SKA precursors, the South African Karoo Array Telescope \\citep[MeerKAT;][]{MeerKAT,meerkat1,meerkat2}, and the Australian SKA Pathfinder \\citep[ASKAP;][]{askap2, askap1, ASKAP, askap3,askap4} stand poised to observe a large number of Southern Hemisphere galaxies in \\hi\\ in the nearby Universe (z$<$0.2). In addition, the Extended Very Large Array \\citep[EVLA;][]{evla} and the APERture Tile In Focus instrument \\citep[APERTIF;][]{apertif,apertif2} on the Westerbork Synthesis Radio Telescope (WSRT) will do the same for the Northern Hemisphere. The surveys conducted with these new and refurbished facilities will add new, high-resolution, \\hi\\ observations on many thousands of galaxies. In the case of WALLABY (Koribalski et al. {\\em in preparation}), these will be of similar quality spatial resolution as the \\whisp\\ survey. Combining these with existing \\uv\\ observations from \\galex, we can explore the relation between the morphology of the atomic hydrogen and ultraviolet light for much greater samples and in much greater detail. An application of the Asymmetry-\\m20\\ identification of \\xuv\\ disks in the \\galex\\ Nearby Galaxy Atlas \\citep{nga} or a sample similar to \\cite{Lemonias11} could reveal additional examples but the combination with deep \\hi\\ observations can prove the link with cold flows for the majority of \\xuv\\ disks." }, "1207/1207.4769_arXiv.txt": { "abstract": "We present for the first time Washington $CT_1$ photometry for 11 unstudied or poorly studied candidate star clusters. The selected objects are of small angular size, contain a handful of stars, and are projected towards the innermost regions of the Small Magellanic Cloud (SMC). The respective Colour-Magnitude Diagrams (CMDs) were cleaned from the unavoidable star field contamination by taking advantage of a procedure which makes use of variable size Colour-Magnitude Diagram cells. This method has shown to be able to eliminate stochastic effects in the cluster CMDs caused by the presence of isolated bright stars, as well as, to make a finer cleaning in the most populous CMD regions. Our results suggest that nearly 1/3 of the studied candidate star clusters would appear to be genuine physical systems. In this sense, the ages previously derived for some of them mostly reflect those of the composite stellar populations of the SMC field. Finally, we used the spatial distribution in the SMC of possible non-clusters to statistically decontaminate that of the SMC cluster system. We found that there is no clear difference between both expected and observed cluster spatial distributions, although it would become more important at a 2$\\sigma$ level between $a$ $\\approx$ 0.3$\\degr$ and 1.2$\\degr$ (the semi-major axis of a ellipse parallel to the SMC bar and with $b/a$ = 1/2), if the asterisms were increased up to 20\\%. ", "introduction": "The different catalogues of Small Magellanic Cloud (SMC) star clusters have been compiled on the basis of star counts, either by visually inspecting photographic plates \\cite[for example]{b75,h86,bs95} or by automatic algorithmic searches \\cite[for example]{petal99}. As far as we are aware, the most recent catalogue which puts all the previous ones together is that of Bica et al. \\shortcite[hereafter B08]{bietal08}. Although it is expected that most of the catalogued objects are indeed genuine physical systems, it was beyond the scope of Bica et al. (2008) to verify the physical nature of such faint objects. The task of cleaning cluster catalogues from non physical systems or asterisms is far from being an exciting job. Most of the astronomers desire to deal with prominent clusters. For this reason studies concluding about the asterism or overdensity nature of faint objects in the Clouds are rare or absent. However, those works would be very important and are also required if any statistical analysis about the cluster formation and disruption rates, the cluster spatial, age and metallicity distributions, etc. is attempted. As it is commonly accepted, an apparent concentration of stars in the sky does not necessarily lead to the conclusion that such concentration constitutes a physical system. The presence of such star concentration implies that we are dealing with a physical cluster in the case of typical globular clusters or very populous open clusters. For most of the apparent star concentrations in the sky, however, it may be necessary to have supplementary information available about proper motions, radial velocities, spectral types and photometry to confirm their physical reality. The photometric data are often the only information at our disposal from which the existence of a star cluster may be inferred. Even though photometric data are indeed valuable, the steps to conclude on the physical nature of a star aggregate from its Colour-Magnitude Diagram (CMD) might not be a straighforward task. This usually happens when dealing with small objects or sparse clusters projected or inmersed in crowded star fields. An example of such situations are those clusters located in the inner regions of the SMC. In such cases, simple circular CMD extractions around the cluster centre could lead to a wrong conclusion, since the CMDs are obviously composed of stars of different stellar populations \\cite{p12}. Consequently, it is hardly possible to assess whether the bright and young Main Sequence (MS) or the subgiant and red giant branches trace the fiducial cluster features. Glatt et al. \\shortcite[hereafter G10]{getal10} have studied 324 SMC clusters using data from the Magellanic Cloud Photometric Surveys \\cite{zetal02}. They show from isochrone fittings on to the CMDs that the studied objects are clusters younger than 1 Gyr mostly distributed in the main body of the galaxy, which is highly crowded. Although they mention that field contamination is a severe effect in the extracted cluster CMDs and therefore influences the age estimates, no decontamination from field CMDs were carried out. It would not be unexpected that some of the studied objects are not real star clusters, particularly those with very uncertain age estimates ($\\sigma$(age) $\\ge$ 0.5 for log(age) $\\la$ 9.0). This possibility alerts us that the sole circular extraction of the observed CMDs of clusters located in highly populated star fields is not enough neither for an accurate isochrone fitting to the cluster MSs nor for confirming their physical natures. Different statistical procedures have been proposed with an acceptable success, in order to avoid as much as possible the field contamination in the cluster CMDs. Chiosi et al. \\shortcite[hereafter C06]{cetal06} studied 311 clusters in the central part of the SMC from OGLE data \\cite{uetal98} and other own data. They used equivalent cluster areas of fields close to the clusters, but outside the cluster radii, to build field CMDs. Subsequently, they divided the CMDs of both clusters and fields in boxes of size $\\Delta$($V$)= 0.5 mag and $\\Delta$($V-I$) = 0.2 mag, and subtracted for every field star in any box the closest cluster star in the respective box. The cluster ages derived from isochrone fittings on to the cleaned CMDs for 136 clusters also included in the study of G10, resulted to be $\\sim$ 0.2-0.3 in log(age) younger than those by G10 ($\\sigma$($\\Delta$log(age)) = 0.13, age $\\la$ 9.0). The main reason for this systematic shift is probably the different metallicty of the isochrones involved. While C06 used the isochrone set of Girardi et al. \\shortcite[Z = 0.008]{getal02}, G10 fitted the cluster CMDs with isochrones computed by Giradi et al. \\shortcite[Z = 0.004]{getal95}. Note that for those clusters with the most significant age deviation large age uncertainties are obtained in both C06 and G10 works. Furthermore, bearing in mind the large age uncertainties quoted by C06 and G10 for some clusters, and that nobody has confirmed that they are genuine physical systems, the doubt about their cluster reality might arise unavoidably. In this paper we present an analysis of 11 candidate star clusters from new CCD Washington $CT_1$ photometry, in combination with a computational tool for cleaning the star field signature in the cluster CMDs. Our main aim is to be able to confirm the physical reality of the studied objects, once their photometric data have been properly cleaned from field contamination. Indeed, the proposed computational tool for estimating the probability of a star of being an intrinsic feature of the cluster field shows to be able to produce reliable CMDs revealing the genuine nature of the considered object. Note that the studied objects were catalogued as clusters on the basis of star counts on less deep images than those using in this study. We present the data set in Section 2, while we describe the data handling in Section 3. Section 4 deals, on the one hand, with available star field decontamination procedures and, on the other hand, with the presently used method. In Section 5, our analysis shows that most of these stellar groups are likely genuine star clusters. Finally, we summarize the main results in Section 6. ", "conclusions": "To date, the catalogue by Bica et al. \\shortcite{bietal08} has been the most complete compilation of star clusters in the SMC. Most of these objects have not been studied yet. Here, we present for the first time Washington $CT_1$ photometry for 11 unstudied or poorly studied candidate star clusters. As compared with the data sets from previous photometric surveys, the present $CT_1$ photometry turns out to be deeper and more accurate. In general the selected objects appear to be of small angular size and contain a handful of stars. They are projected towards the most crowded star field regions in the SMC, at distances shorter than $\\sim$ 1$\\degr$ from its centre. We have designed a procedure for cleaning the cluster CMDs from the unavoidable star field contamination which makes use of variable cells in the CMDs. The cells are adjusted in such a way that they result bigger in CMD regions with a scarce number of field stars, and viceversa. This way, we reproduce the field CMD as closely as possible on to the cluster CMD. The method does not need to know whether a star is placed close to the cluster centre nor the cluster radial density profile to infer a membership probability. However, it takes into account the star field density, since the more populous a star field the larger the number of stars subtracted from the cluster CMD. As a result, the intrinsic spatial star distribution is uncovered within the object region. Once the field CMD is adopted, the method defines a free path for each field star as the distance to the closest star in the field CMD. The method has shown to be able to eliminate stochastic effects in the cluster CMDs caused by the presence of isolated bright stars, as well as, to make a finer cleaning in the most populous CMD regions. When applying the cleaning procedure to the CMDs of the 11 selected candidate star clusters, we found that nearly 1/3 of them would appear to be genuine physical systems. We estimated their ages from the matching of the isochrone which best reproduces the CMD cluster features. In this sense, the ages previously derived for some of them mostly reflect those of the composite stellar populations of the SMC field. The present analysis tools applied to faint poorly populated clusters or candidates in the Magellanic Clouds points to the need of better scale deep observations with e.g. the 8m class telescopes. Finally, we used the spatial distribution in the SMC of possible non-clusters to statiscically decontaminate that of the SMC cluster system. By assuming that the area covered by 11 studied fields (36$\\arcmin$$\\times$36$\\arcmin$ each) represents an unbiased subsample of the SMC as a whole and by using an elliptical framework centred on the SMC centre ($b/a$ = 1/2), we found that there is no significant difference between the expected and the observed cluster spatial distributions. However, a difference at a 2$\\sigma$ level would become visible between $a$ $\\approx$ 0.3$\\degr$ and 1.2$\\degr$, if we doubled the amount of possible non-clusters." }, "1207/1207.7076_arXiv.txt": { "abstract": "We present detailed elemental abundances of 12 subgiants in the open cluster IC 4756 including Na, Al, Mg, Si, Ca, Ti, Cr, Ni, Fe, Zn and Ba. We measure the cluster to have [Fe/H] = $-0.01\\pm0.10$. Most of the measured star-to-star [X/H] abundance variation is below $\\sigma < 0.03$, as expected from a coeval stellar population preserving natal abundance patterns, supporting the use of elemental abundances as a probe to reconstruct dispersed clusters. We find discrepancies between {Cr\\,{\\sc i}} and {Cr\\,{\\sc ii}} abundances as well as between {Ti\\,{\\sc i}} and {Ti\\,{\\sc ii}} abundances, where the ionized abundances are larger by about 0.2~dex. This follows other such studies which demonstrate the effects of overionization in cool stars. IC 4756 are supersolar in Mg, Si, Na and Al, but are solar in the other elements. The fact that IC 4756 is supersolar in some $\\alpha$-elements (Mg, Si) but solar in the others (Ca, Ti) suggests that the production of $\\alpha$-elements is not simply one dimensional and could be exploited for chemical tagging. ", "introduction": " ", "conclusions": "The derived abundances and total uncertainty per star per element are listed in Tables~\\ref{table:EW_analysis_1}--\\ref{table:EW_analysis_3}. We plot our results in Fig.~\\ref{fig:xfe_feh}, and overplot field stars from \\citet{red03,red06}, \\citet{ben03} and recent open cluster compilation from \\citet{car11} for comparison. For the open cluster compilation, we have 78 clusters with $6.4 \\leq R_G \\leq 20.8$ kpc. For each cluster, we take the mean abundance of the cluster for each element. \\begin{table*} \\caption{Comparison with previous studies on IC 4756 giants.\\label{table:previous_studies}} \\begin{tabular}{lccccccccccccccc} \\hline Source & Type & Resolution & [Fe\\,{\\sc i}/H]& [Na/Fe] & [Al/Fe] & [Mg/Fe] & [Si/Fe] \\\\ \\hline This study & 12 giants & 30\\,000 & $-0.01\\pm0.10$ & $0.21\\pm0.15$ & $ 0.12\\pm0.12$ & $ 0.12\\pm0.13$ & $ 0.13\\pm0.13$ \\\\ \\citet{pac10} & 3 giants & 100\\,000 & $ 0.08\\pm0.11$ & $0.11\\pm0.12$ & $-0.12\\pm0.12$ & --- & $ 0.02\\pm0.11$ \\\\ \\citet{smi09} & 5 giants & 48\\,000 & $ 0.05\\pm0.11$ & $0.02$ & --- & $-0.05$ & $ 0.05\\pm0.08$ \\\\ \\citet{jac07} & 6 giants & 15\\,000 & $-0.15\\pm0.18$ & $0.57\\pm0.19^a$& $ 0.29\\pm0.21$ & --- & $ 0.34\\pm0.25$ \\\\ \\citet{luc94} & 3 giants & 18\\,000 & $-0.05\\pm0.21$ & $0.14\\pm0.22$ & $ 0.05\\pm0.28^a$ & $ 0.22\\pm0.40^a$ & $ 0.23\\pm0.29$ \\\\ \\vspace{-0.15cm} \\\\ Source & [Ca/Fe] & [Ti\\,{\\sc i}/Fe] & [Cr\\,{\\sc i}/Fe] & [Cr\\,{\\sc ii}/Fe] & [Ni/Fe] & [Zn/Fe] & [Ba/Fe] \\\\ \\vspace{-0.15cm} \\\\ This study & $ 0.05\\pm0.16$ & $0.00\\pm0.19$ & $0.03\\pm0.15$ & $0.25\\pm0.14$ & $ 0.03\\pm0.13$ & $0.06\\pm0.13$ & $ 0.00\\pm0.14$ \\\\ \\citet{pac10} & $-0.02\\pm0.12$ & $0.03\\pm0.13$ & $0.01\\pm0.13$ & --- & $-0.04\\pm0.12$ & --- & --- \\\\ \\citet{smi09} & $ 0.02\\pm0.09$ & $-0.05\\pm0.09$ & $0.04\\pm0.11$ & $0.19\\pm0.11$ & $-0.01\\pm0.06$ & --- & --- \\\\ \\citet{jac07} & $ 0.07\\pm0.24$ & --- & --- & --- & $ 0.08\\pm0.21$ & --- & --- \\\\ \\citet{luc94} & $-0.07\\pm0.35$ & $-0.08\\pm0.31$ & $0.08\\pm0.26$ & --- & $ 0.08\\pm0.30$ & --- & $ 0.02\\pm0.26$ \\\\ \\hline \\end{tabular} \\flushleft $^a$[X/H] does not have uncertainty estimation; we have no choice but to assume that all uncertainties come from [Fe/H]. The stated uncertainty is therefore smaller than it should be. \\end{table*} \\subsection[]{Light odd-Z elements: Na and Al}\\label{sec:results_and_discussion} For [Al/Fe], all stars from IC 4756 show homogeneous abundances. On the other hand, [Na/Fe] shows more intracluster star-to-star dispersion. As shown in Fig.~\\ref{fig:teff}, abundances of Na are strongly dependent on the effective temperature with the two coolest stars (Her 6 with [Na/H] = 0.13 dex and Her 90 with [Na/H] = 0.04 dex) being the reason for the trend. As discussed in Section~\\ref{subsection:ew}, the inhomogeneity may be explained by NLTE effects. This could, however, also be a sign of internal mixing in giants. The larger abundance scatter in Na is also consistent with the open cluster compilation values. As shown in Fig.~\\ref{fig:xfe_feh}, among light odd-Z elements, [Na/Fe] is subjected to more scatter, presumably due to the choice of dwarfs or giants in each survey. This enhanced Na is also seen in field star surveys of clump and red giant stars \\citep{siv09}, but this is not the case for [Al/Fe]. This observation is also consistent with field stars results from \\citet{and08} and \\citet{bon09} where they argued that Al does not suffer from significant internal mixing. \\subsection[]{$\\alpha$-elements: Mg, Si, Ca and Ti}\\label{subsection:alpha_elements} Our analysis shows that [Mg/Fe] and [Si/Fe] are supersolar for IC 4756, but [Ca/Fe] and [Ti/Fe] seem to be solar. This is not new. Several high-resolution analyses of various open clusters show overabundances in lighter $\\alpha$-elements \\citep*[e.g.][]{yon05}. \\citet{ting12} also showed that, compared to field stars, open cluster compilation shows more variation in $\\alpha$-elements and more independent dimensions. This is consistent with the more interclusters scatter in lighter $\\alpha$-elements, especially Mg, as shown in Fig.~\\ref{fig:xfe_feh}. However, since Mg is usually derived with very few lines, the results should be viewed with caution. We chose to use neutral Ti for our analysis. It is important to note that there is a discrepancy between Ti\\,{\\sc i} and Ti\\,{\\sc ii} abundances. Similar effects have previously been reported for stars cooler than \\teff$<5200~K$ \\citep{sch03}. The mechanism responsible for this, known as overionization effect, is the pumping of the electrons into the ionized state due to strong UV flux of the hot chromospheres \\citep[refer to][for a more detailed discussion on the topic]{sch03,dor09}. Therefore we adopt the Ti\\,{\\sc i} abundance as a more reliable value. \\subsection[]{Fe-peak elements: Cr, Ni and Zn}\\label{subsection:Fe_peak_elements} All our Fe-peak elements are homogeneous, solar-like and show little intracluster star-to-star dispersion. As observed with Ti\\,{\\sc ii} abundances, we see the effects of overionization in Cr\\,{\\sc ii} as well. Therefore, we adopt the Cr\\,{\\sc i} abundance, which is in the solar proportions, in agreement with the other Fe-peak elements. \\subsection[]{Neutron-capture elements: Ba}\\label{subsection:Neutron-capture} Our study shows that Ba abundance is homogeneous and solar-like. This is strikingly different from what has been observed for the similar aged Hyades open cluster, where Ba is enhanced over 0.2 dex \\citep{siv06,siv11}. It is often thought that different analysis methods and line lists might contribute to the scatter in [Ba/Fe]-[Fe/H], as shown in Fig.~\\ref{fig:xfe_feh}. However our analysis adopts a very similar method compared to the studies of \\citet{siv06,siv11}. It is therefore possible that the cluster-to-cluster dispersion (and star-to-star dispersion for field stars) in Ba abundance is real. This supports the idea that the metallicity-dependence nature of $s$-process in AGB phase \\citep{bus01,cri09,bis10} creates a unique neutron-capture elements imprint on the gas from which the stars form \\citep{ting12}. \\subsection[]{Comparison with previous studies}\\label{subsec:previous_studies} Her 85, Her 87 and Her 176 were previously studied by \\citet{luc94}. Similar to our analysis, Her 85 also showed some elemental abundances that do not fit into the cluster mean. As discussed in Section~\\ref{subsec:observation}, Her 85 is likely a non-member of IC 4756; therefore, we exclude Her 85 to calculate the cluster mean. For Her 87, they derived \\teff~$=5100$, $\\log g= 2.8$, $v_t=2.40$, and for Her 176, they deduced \\teff$=5200$, $\\log g=3.0$, $v_t=1.90$. Furthermore, \\citet{gil89} studied Her 87, Her 144, Her 176, Her 228, Her 249, Her 296 and Her 314. Other than some discrepancies in microturbulence, the comparison with our atmospheric parameters is very satisfactory. The comparison of elemental abundances with previous studies is listed in Table~\\ref{table:previous_studies}. Note that, for a better comparison, we do not directly adopt the cluster mean and its uncertainty quoted in each study. {\\it Instead, we calculate the cluster mean with the same method as described in Section~\\ref{subsec:mean_abundances}}. All previous studies agree on [Fe/H], [Ca/Fe], [Ti/Fe], [Cr/Fe], [Ni/Fe] and [Ba/Fe], including [Fe/H] $= 0.04\\pm0.20$ from seven giants of \\citet{gil89} and [Fe/H] $=0.03\\pm0.07$ from three giants of \\citet{san09} that are not listed in Table~\\ref{table:previous_studies}. All these abundances appear to be solar. Furthermore, [Cr\\,{\\sc ii}/Fe] is consistently about $0.2$ dex enhanced compared to [Cr\\,{\\sc i}/Fe], which mimics our results as well. Our results are consistent with other studies that [Na/Fe] is above solar. Our analysis shows Na abundance in between the results of \\citet{jac07} and \\citet{pac10}. \\citet{jac07} claimed that their Na abundance can be brought down if different line lists were considered and therefore might be more consistent with ours and \\citet{pac10} results. However, \\citet{smi09} claimed that there is no enhancement in Na if NLTE corrections is taken. On the other hand, there are discrepancies in term of [Al/Fe], [Mg/Fe] and [Si/Fe]. Our study agrees with \\citet{jac07} that both Al and lighter $\\alpha$-elements, such as Mg and Si, are enhanced. \\citet{luc94} also reported the same for Mg and Si. \\citet{pac10} reported a [Na/Fe] enhancement of 0.11~dex, which is similar to the level of enhancement in this study for [Al/Fe], [Mg/Fe] and [Si/Fe]." }, "1207/1207.2186_arXiv.txt": { "abstract": "In Dirac-Born-Infeld inflation, changes in the sound speed that transiently break the slow roll approximation lead to features in the power spectrum. We develop and test the generalized slow roll approximation for calculating such effects and show that it can be extended to treat order unity features. As in slow-roll, model independent constraints on the potential of canonical inflation can be directly reinterpreted in the DBI context through this approximation. In particular, a sharp horizon scale step in the warped brane tension can explain oscillatory features in the WMAP7 CMB power spectrum as well as features in the potential. Differences appear only as a small suppression of power on horizon scales and larger. ", "introduction": "\\label{sec:intro} In Dirac-Born-Infeld (DBI) inflation \\cite{Silverstein:2003hf,Alishahiha:2004eh}, transient but rapid changes in the sound speed leave their imprint as features on the power spectrum. For string-motivated DBI examples, such features might arise from duality cascades which impart steps in the warped brane tension \\cite{Hailu:2006uj,Bean:2008na}. Annihilation of branes during DBI inflation has also been shown to lead to particle production and to the imprint of features on the warp \\cite{Firouzjahi:2010ga}. More generally, within the context of effective field theory \\cite{Cheung:2007st} it has been shown that a sharp step in the sound speed leads to oscillatory features in the power spectrum of fluctuations \\cite{Park:2012rh}. Power spectrum features from sudden changes in the warped brane tension of DBI inflation are closely related to those from sudden changes in the potential for canonical single field inflation. Measurements of the CMB temperature power spectrum from WMAP place observational constraints on the latter. Recently, the generalized slow roll approach (GSR) \\cite{Stewart:2001cd,Dvorkin:2009ne} has been used to extract model-independent constraints from the WMAP data on features as sharp as 1/4 of an efold \\cite{Dvorkin:2010dn,Dvorkin:2011ui}. Even sharper features lead to highly oscillatory power spectrum features which can evade these constraints due to projection effects. Indeed there is a special case where a sharp step in the potential on scales near the current horizon can fit the WMAP data better that a smooth model in the acoustic regime \\cite{Adshead:2011jq}. The GSR approach remains valid for single field inflation with non-canonical kinetic terms \\cite{ArmendarizPicon:1999rj}, including DBI inflation, with a suitable reinterpretation of the source of deviations from slow-roll \\cite{Hu:2011vr}. In this Paper, we develop the GSR approach for DBI inflation and show how observational constraints on potential features translate to constraints on warp features. In \\S \\ref{sec:DBI}, we briefly review the phenomenology of DBI inflation and the exact computation of its power spectrum. In \\S \\ref{sec:GSR} we develop and test the GSR approach in the DBI context and establish the correspondence between potential features and warp features. In \\S \\ref{sec:step}, we consider the special case of a sharp step in the warp analytically and show that it can explain the WMAP data as well as a sharp step in the potential. We discuss these results in \\S \\ref{sec:discussion}. ", "conclusions": "\\label{sec:discussion} We have shown that the GSR approximation can be applied to DBI inflation to constrain features in the warped brane tension $T(\\phi)$ from observational data. The approximation accurately recovers corresponding features for up to order unity deviations. Previous work on constraining the GSR source function $G'$ and hence second derivatives of the potential $V(\\phi)$ for canonical fields can be directly reinterpreted in the DBI context as limits on the second derivative of $T(\\phi)$ \\cite{Dvorkin:2010dn,Dvorkin:2011ui}. The main difference between the two is that features in $T(\\phi)$ once traversed can strongly affect the slow-roll attractor for modes that cross the sound horizon later. The correspondence between features in $V(\\phi)$ and $T(\\phi)$ is especially close in the limit of extremely sharp features, for example a step feature. In both cases the power spectrum exhibits constant amplitude oscillations for modes that cross the sound horizon after the step. Consequently, the preference for a horizon-sized step in the potential in WMAP7 implies a corresponding preference for a step in $T(\\phi)$. The main difference is a reduction of power for the low $k$ modes that cross before the feature. The large cosmic variance of these modes prevents a significant distinction between the two. On the other hand, features in $V(\\phi)$ for canonical inflation and $T(\\phi)$ for DBI inflation should induce very different bispectra. We leave these considerations to a future work." }, "1207/1207.2465_arXiv.txt": { "abstract": "I suggest a simple signature for new particles which are unstable partners of a dark matter particle. The suggested mass range is from 8 TeV to 3 PeV,\\ the former being the mass of the dark matter particle and the latter being the knee energy mass scale from the cosmic ray energy spectrum. It can be the energy spectrum of a specific particle such as a muon, a neutrino, jets or any other particles produced in cosmic ray showers, as long as the spectrum is measued. As for the detection of a 3 PeV particle by the neutrino energy spectrum, all dark matter targets throughout the galaxy that are bombarded by high energy cosmic rays and high energy dark matter particles contribute to the process. This is new in the study of dark matter physics. ", "introduction": "LHC experiments are underway with the hope that new particles in a dark matter family might be observed. So far no new partitcles have been discovered. From high energy gamma ray searches the prognosis for the existence of a relatively low mass (less than 1 TeV) dark matter particle (DMP) seems unlikely. HESS\\cite{hess1} concluded that there is no gamma ray peak below 2 TeV, while above 2 TeV up to 10 TeV, there is gamma ray excess above the power law extension from the lower energy data. This suggests that the discovery of new particles in a DMP family might require various kinds of cosmic ray detectors, as was the case for the discovery of strange particles in the pre-1950 era, a golden age of cosmic ray physics. All new particles were found exclusively in cosmic ray detectors, before the arrival of the accelerator era. In the next section, I summarize the expected mass range for a DMP family to see what we are dealing with. The method which we develop for detection is of more general application, however. ", "conclusions": "" }, "1207/1207.0656_arXiv.txt": { "abstract": "{The low-mass star formation evolutionary sequence is relatively well-defined both from observations and theoretical considerations. The first hydrostatic core is the first protostellar equilibrium object that is formed during the star formation process.} {Using state-of-the-art radiation-magneto-hydrodynamic 3D adaptive mesh refinement calculations, we aim to provide predictions for the dust continuum emission from first hydrostatic cores.} {We investigated the collapse and the fragmentation of magnetized 1 M$_\\odot$ prestellar dense cores and the formation and evolution of first hydrostatic cores using the {\\ttfamily{RAMSES }}\\rm code. We used three different magnetization levels for the initial conditions, which cover a wide variety of early evolutionary morphology, e.g., the formation of a disk or a pseudo-disk, outflow launching, and fragmentation. We post-processed the dynamical calculations using the 3D radiative transfer code {\\ttfamily{RADMC-3D}}. We computed spectral energy distributions and usual evolutionary stage indicators such as bolometric luminosity and temperature.} {We find that the first hydrostatic core lifetimes depend strongly on the initial magnetization level of the parent dense core. We derive, for the first time, spectral energy distribution evolutionary sequences from high-resolution radiation-magneto-hydrodynamic calculations. We show that under certain conditions, first hydrostatic cores can be identified from dust continuum emission at 24 $\\mu$m and 70 $\\mu$m. We also show that single spectral energy distributions cannot help in distinguishing between the formation scenarios of the first hydrostatic core, i.e., between the magnetized and non-magnetized models.} {Spectral energy distributions are a first useful and direct way to target first hydrostatic core candidates but high-resolution interferometry is definitively needed to determine the evolutionary stage of the observed sources.} \\keywords {Magnetohydrodynamics (MHD), radiative transfer - Methods: numerical- Stars: low mass, formation, protostars} \\titlerunning{} \\authorrunning{B. Commer\\c con et al.} ", "introduction": "It is well established that low-mass stars form from the collapse of dense prestellar cores. The star formation process can be divided both observationally and theoretically into an evolutionary sequence. From observations, protostars are classified according to the parametrized properties of the observed spectral energy distribution (SED) and appearance. For instance, a Class 0 source \\citep{Andre_et_al_1993}, which corresponds to the phase where the mass of the protostar is less than half the mass of the surrounding envelope, is defined according to its bolometric temperature $T_\\mathrm{bol}$ \\citep{Chen_et_al_1995} or to the ratio of bolometric to submillimeter luminosities $L_\\mathrm{bol}/L_\\mathrm{smm}$ \\citep{Andre_et_al_1993}. At later stages, the Class I, Class II, and Class III sources can be classified according to the slope $\\alpha_\\mathrm{IR}=\\mathrm{dlog}(\\lambda F_\\lambda)/\\mathrm{dlog}(\\lambda)$ between 2.2 $\\mu$m and $10-25$ $\\mu$m \\citep{Lada_1987}. These classifications attempt to replicate the different evolutionary stages of a protostar. Observational ambiguities (e.g., projection effects, episodic accretion), however, imply that these classifications may not always reflect the true evolutionary stage of an object \\citep[e.g., ][]{Dunham_et_al_2010}. The protostellar collapse phase is divided into two phases. First the prestellar dense core (gravitationally bound, no protostar, and no outflow) collapses isothermally and is able to freely radiate its gravitational energy into space, until it becomes dense ($\\rho\\sim 10^{-13}$ g cm $^{-3}$) and opaque enough for that radiation to be trapped. At this time, the gas begins to heat up and the first hydrostatic core (FHSC) is formed \\citep{Larson_1969}. The FHSC is optically thick and its temperature increases adiabatically until it reaches $\\sim2000$ K, where H$_2$ dissociation starts. This endothermic reaction allows the gas to undergo a second collapse phase until stellar densities are reached, i.e., the formation of the second hydrostatic core (SHSC), the protostar that corresponds to the protostellar Class 0 evolutionary stage. The FHSC is thus a transition phase from prestellar cores to Class 0 objects and has not been yet identified conclusively because its observational appearance is unclear. The predicted timescales of the FHSC phase range from a few hundred years in spherical models \\citep{Masunaga_Inutsuka_2000} to a few thousand years in very-low-mass non-magnetized rotating dense cores \\citep{Tomida_2010b}. In recent years, significant progress has been achieved in both theory and observation thanks to increasing computing power and new observational capabilities provided by {\\it{Spitzer }}\\rm and {\\it{Herschel}}\\rm. For instance, the large {\\it{Spitzer }}\\rm \"Core to Disks\" (c2d) Legacy Program survey \\citep{Evans_et_al_2003} helped to increase the number of known protostars up to the point where first detections of FHSC can now be expected, given the short duration of the FHSC phase \\citep{Enoch_et_al_2010,Dunham_et_al_2011}. {\\it{Spitzer }}\\rm also helped to identify a new class of low-luminosity objects: the Very Low Luminosity Objects \\citep[VeLLOs, e.g., ][]{DiFrancesco_et_al_2007}. These are embedded in a dense envelope and have low intrinsic luminosity, $L<0.1$ L$_\\odot$. It is still a matter of debate whether VeLLOs are dense cores undergoing first collapse or low-mass protostars with a low-accretion rate. In a recent study, \\cite{Kim_et_al_2011} showed that at least in one case (the Bok globule CB130-1), the latter is the most likely explanation. Prestellar dense core populations have been identified over the past twenty years using space- and ground-based telescopes \\citep[e.g., ][]{Ward-Thompson_et_al_1994,Motte_et_al_1998,Johnstone_et_al_2000,Nutter_2007}. More recently, {\\it{Herschel }}\\rm observations have identified hundreds of prestellar cores and Class 0 sources \\citep[e.g., ][]{Andre_et_al_2010,Bontemps_et_al_2010,Henning_et_al_2010}. Combining different instrument data, \\cite{Launhardt_et_al_2010} presented the results of a comprehensive infrared, submillimeter, and millimeter continuum emission study of isolated low-mass star-forming cores, and identified nine Class 0 sources. Finally, a few FHSC candidate detections have been reported in recent years thanks to the increasing spectral and spatial scale coverage \\citep[][]{Belloche_et_al_2006, Chen_et_al_2010,Enoch_et_al_2010,Dunham_et_al_2011,Pineda_et_al_2011,Pezzuto_et_al}. Their confirmation remains controversial, however, because the expected observable signatures of FHSC are still uncertain. In parallel, enormous work has been put into developing numerical hydrodynamical models that include magnetic fields and radiative transfer \\citep[e.g., ][]{Price_Bate_2009,Commercon_et_al_2010,Tomida_2010} but relatively little has been done in producing synthetic observations that are directly comparable with actual observations. Most of the observables used to confront theory and observation remain statistical in nature, such as the initial mass function (IMF) or the binary distribution, whereas observations report mostly dust continuum and molecular line emission data. Therefore, it is vital that the interpretation of the data is sound, necessitating the use of non-symmetric models incorporating complex physics to derive predictions of dust continuum and molecular line emission observations. For the early stages of star formation, \\cite{Boss_Yorke_1995}, \\cite{Masunaga_et_al_1998}, \\cite{Omukai_2007}, \\cite{Yamada_2009}, \\cite{Tomida_2010b}, \\cite{Tomisaka_2011}, and \\cite{Saigo_Tomisaka_2011} made FHSC observational predictions of either dust emission or molecular line emission. \\cite{Young_Evans_2005} and \\cite{Dunham_et_al_2010} also presented evolutionary signatures of the formation of low-mass stars, but used a crude approximation of the FHSC and the evolution of the envelope during collapse. To date, there has been no systematic study of observational predictions for FHSCs that explore physical conditions, and in particular on the magnetization level. Magnetic fields have indeed been found to control the first collapse and fragmentation phases \\citep[e.g., ][]{Hennebelle_Teyssier_2008,Commercon_et_al_2010} and strong magnetization seems to be a favored scenario for reproducing observational data of Class 0 multiplicity \\citep{Maury_et_al_2010}. The ability of a disk to fragment depends ion the magnetization level, leading to potential changes in the SED. In this paper, we derive SED evolutionary sequences and classical observational indicators such as $L_\\mathrm{bol}$ and $T_\\mathrm{bol}$ from state-of-the-art radiation-magneto-hydrodynamic (RMHD) 3D calculations of collapsing 1 M$_\\odot$ dense cores with different initial magnetic field strengths. We address the questions of whether or not SEDs can help in distinguishing physical conditions (e.g., magnetic field strengths) and how SEDs can help in targeting FHSC candidates. This paper is the first of a series and is associated with a companion paper in which we produce synthetic ALMA dust continuum emission maps (Commer\\c con et al. {\\it{in prep}}, hereafter Paper II). The paper is organized as follows: in Sect. 2, we present the numerical codes we used and our post-processing methodology. The RMHD calculations are qualitatively presented in Sect. 3. In Sect. 4, the RMHD calculations are post-processed to produce synthetic SEDs. In Sect 5., we discuss our results and potential observational tests. Sect. 6 concludes our paper. ", "conclusions": "\\subsection{Comparison to previous work} Spectral energy distributions of FHSCs embedded in their parent cores have already been studied for instance by \\cite{Boss_Yorke_1995}, \\cite{Masunaga_et_al_1998}, \\cite{Omukai_2007}, \\cite{Tomida_2010b}, and \\cite{Saigo_Tomisaka_2011}. While they all found that the radiation emitted by the FHSC is essentially completely reprocessed by the dust in the surrounding core to 50 - 200 $\\mu$m with no observable emission below 30 $\\mu$m, we found that there is indeed observable emission at shorter far-infrared wavelengths down to 20 $\\mu$m. These differences may be explained by the different physical models and post-processing of each study. While we directly post-processed 3D AMR RMHD calculations, \\cite{Boss_Yorke_1995} used a spherical grid and integrated the RHD equations with a low numerical resolution. Their FHSC temperature is also relatively low ($<200$ K). Their results are therefore similar to ours shortly after FHSC formation. In addition, as mentioned previously, \\cite{Masunaga_et_al_1998} and \\cite{Omukai_2007} used spherical models and their SEDs are then similar to ours with $\\theta=90^{\\circ}$. Next, \\cite{Saigo_Tomisaka_2011} post-processed 3D high-resolution collapse calculations with plane-parallel approximation but neglected magnetic fields and used a barotropic EOS. They observed an increase of the total energy in the SED but their SEDs were calculated over a very small area of 40 AU $\\times$ 40 AU and therefore excluded the large-scale emission dominant at wavelengths $> 100$ $\\mu$m. Last, \\cite{Tomida_2010b} computed SEDs from 3D RHD calculations but their cores do not fragment (slow rotation). In addition, they used Bonnor-Ebert initial density profiles in which accretion rates are lower and therefore FHSCs are fainter.% \\subsection{What can SEDs tell about physical processes and evolutionary stages?} Figures \\ref{sed} and \\ref{SED_evol} show that the evolution of an FHSC throughout its lifetime can be reflected in observable changes of the SED, but only under certain conditions. Only the intermediate MU10 model does not show any evolutionary sequence because of the thick envelope structure consisting of a system disk + pseudo-disk + outflow. The MU2 and MU200 models show an evolutionary sequence very different than that of the MU10 model, but very similar to each other. The MU2 model is strongly magnetized and exhibits a pseudo-disk + outflow system, whereas the MU200 model fragments and exhibits a rotationally supported disk. An FHSC with a clear evolutionary sequence can then be observed in simple geometrical configurations and with inclination angle $\\theta<60^{\\circ}$. Nevertheless, it is impossible to distinguish between a magnetized and a non-magnetized scenario from a strict SED point of view. High angular resolution images of, at least, dust emission are required to distinguish between these scenarios. This kind of diagnostic can only be provided by interferometry at longer wavelengths (see Paper II). We have shown than it can be relatively straightforward to distinguish a starless core from an FHSC thanks to flux received at mid- to far-infrared wavelengths, i.e., at 24 $\\mu$m and 70 $\\mu$m, but it remains hard to distinguish between an FHSC and a young stellar object (YSO) with a disk \\citep{Pineda_et_al_2011}, i.e., an accreting second hydrostatic core (SHSC). This difficulty is due to the thick envelope in which the FHSC and SHSC are each embedded throughout the Class 0 phase that reprocesses short-wavelength radiation to longer wavelengths. Here, models clearly fail to predict any detectable differences. Using only dust continuum emission is not the appropriate way to differentiate the two core stages, since almost all radiation emanating from both the FHSC and the SHSC will be reprocessed by the thick envelope. Another approach could consist of focusing on the outflow properties (collimation, velocity, momentum, etc.) and on the chemical properties of the sources, such as gas phase and ice molecules abundances, which reflect the local thermal history \\citep[e.g., ][]{Kim_et_al_2011}. We will study these properties in upcoming papers. \\subsection{Example of methodology to target first core} We showed in Sec. \\ref{new_FHSC} that the fluxes at 24 $\\mu$m and 70 $\\mu$m can be good indicators of FHSCs in comparison with starless cores. We also did not find any significant emission below 10 $\\mu$m. We infer that the combination of a point source detection at 70 $\\mu$m or even at 24 $\\mu$m with no detection at wavelength $<10$ $\\mu$m can be a good indicator of FHSCs from a strict SED point of view. Under the current telescope sensitivity limits, a detection at 70 $\\mu$m alone is most likely, since the emission at 24 $\\mu$m is faint, and reaches the {\\it{Spitzer }} sensitivity limit only in the MU2 model in the second half of the FHSC lifetime. Figure \\ref{mu2_images} shows dust continuum emission maps for the MU2 model at three different evolutionary stages and for four wavelengths (24 $\\mu$m, 70 $\\mu$m, 160 $\\mu$m, and 1200 $\\mu$m). It clearly shows that the emission at 70 $\\mu$m and 24 $\\mu$m should be point-like after the FHSC formation. Note that at all evolutionary stages, a flux at 70 $\\mu$m can be detected but its emission would not look like a point source for a starless core. It will instead be much more extended (i.e., the size of the initial dense core) because of the ISRF reprocessing. Thanks to the recent {\\it{Herschel }}\\rm and {\\it{Spitzer }}\\rm observations, observers should be able to combine data from both telescopes and make comparisons with our predicted color-color plot in Fig. \\ref{L70}. For instance, \\cite{DiFransesco_et_al_2010} already used this procedure and combined maps of {\\it{Herschel }}\\rm data with {\\it{Spitzer }}\\rm ones to distinguish between prestellar and protostellar cores. Once FHSC candidates are identified from these flux examinations, the next step will be to target them with high-resolution interferometry such as ALMA (see Paper II) to constrain the physical properties (e.g., strength of the magnetic field). As mentioned earlier, a complementary approach would be to use chemical analysis to separate the FHSC and the SHSC stages. At temperatures above $\\sim 1000$ K, grain evaporation starts and refractory chemical species such as C and Si are released into the gas phase, considerably changing the chemical network and the opacities \\citep[e.g., ][]{Lenzuni_et_al_1995}. Deuterium-related chemistry could also trace an SHSC because deuterium-bearing species should be under-abundant in high-temperature regions \\citep{Albertsson_et_al_2011}. Another more direct and more promising way to stage differentiation would be to find evidence of a low-velocity outflow without its high-velocity counterparts. Recent theoretical studies have shown that a low-velocity outflow can be launched during the FHSC stage \\citep[e.g., ][]{Hennebelle_Fromang_2008, Commercon_et_al_2010,Tomida_2010} without the high-velocity outflow that is always present at the SHSC stage \\citep{Machida_et_al_2008}. Not only may the morphology and energetics of the outflow help in identifying the cores, but also the chemical composition of the gas in the outflow. Since temperatures are different during the FHSC and SHSC stages, we may expect that the low-velocity outflow would show evidence of low-temperature gas-grain chemistry and that the high-velocity outflow would show evidence of high-temperature chemistry. In addition, the low-velocity outflow is launched from regions with temperatures $<1000$ K where grains are still not fully evaporated. Grains should then be present in the FHSC outflow and since the velocity is relatively low (a few km s$^{-1}$), they are not expected to be destroyed in the shock at the envelope interface. Therefore, we may expect FHSC outflows to be seen in thermal dust emission if they are dense enough. It is worth mentioning that some dust continuum observations may be also contaminated by CO emission from outflow gas, making them easier to see \\citep[e.g., ][]{Drabek_et_al_2012}. Recent work by \\cite{Price_2012} reported the formation of a collimated jet driven from the FHSC, supporting recent observations of an FHSC candidate in Per-Bolo 58 \\citep{Dunham_et_al_2011}. The timescale needed to produce such a collimated jet ($\\sim 2$ kyr), however, is longer than the expected lifetime of FHSCs in a magnetized collapsing dense core. In addition, \\cite{Price_2012} used sink particles (scales smaller than 5 AU are not described) and a barotropic approximation with an adiabatic index of $\\gamma=7/5$, for which FHSC lifetimes are even shorter than those we predict in this study with $\\gamma=5/3$ (the FHSC contracts adiabatically faster with $\\gamma=7/5$). It is therefore very likely that FHSCs reported in \\cite{Price_2012} should have already undergone second collapse and formed an SHSC within $\\sim 2$ kyr after their formation. \\subsection{Limits of the model} In this study, we used the gray radiative transfer approximation for the RMHD calculations which may not be the most appropriate method at later stages of the evolutions. As mentioned earlier, however, \\cite{Vaytet_et_al_2011} showed that multi-frequency collapse calculations were well-reproduced by the gray FLD calculations of \\cite{Commercon_et_al_2011b}, at least for the first collapse and FHSC stages. In addition, we used a constant ratio of specific heats $\\gamma$, set to $5/3$, although it is known that $\\gamma$ should change to a value of $7/5$ as temperature increases, and H$_2$ rotational degrees of freedom become available \\citep[e.g., ][]{Machida_et_al_2008}. With $\\gamma=7/5$, FHSCs compress and cool more efficiently and thus reach the second collapse more quickly \\citep[e.g., ][]{Commercon_et_al_2010}, which reduces the FHSC lifetime. For the {\\ttfamily{RADMC-3D}} post-processing, we used one specific opacity model of \\cite{Semenov_et_al_2003A&A} for all temperatures and densities. This choice may not be appropriate since grain mantle compositions change as temperature increases. Nevertheless, the changes in opacity remain relatively minor at temperatures $<1000$ K and we do not expect them to have a significant impact on our results. Finally, we have neglected scattering although it has been found to be important for wavelengths $<10$ $\\mu$m \\citep{Young_Evans_2005,Dunham_et_al_2010}. Since we did not find any emission from FHSCs at these wavelengths and we do not consider this wavelength range significant for our conclusions, scattering effects are not relevant for our purposes." }, "1207/1207.1273.txt": { "abstract": "Scalar modifications of gravity have an impact on the growth of structure. Baryon and Cold Dark Matter (CDM) perturbations grow anomalously for scales within the Compton wavelength of the scalar field. In the late time Universe when reionisation occurs, the spectrum of the 21cm brightness temperature is thus affected. We study this effect for chameleon-f(R) models, dilatons and symmetrons. Although the $f(R)$ models are more tightly constrained by solar system bounds, and effects on dilaton models are negligible, we find that symmetrons where the phase transition occurs before $z_{\\star} \\sim 12$ will be detectable for a scalar field range as low as $5\\ {\\rm kpc}$. For all these models, the detection prospects of modified gravity effects are higher when considering modes parallel to the line of sight where very small scales can be probed. The study of the 21 cm spectrum thus offers a complementary approach to testing modified gravity with large scale structure surveys. Short scales, which would be highly non-linear in the very late time Universe when structure forms and where modified gravity effects are screened, appear in the linear spectrum of 21 cm physics, hence deviating from General Relativity in a maximal way. ", "introduction": "A major challenge for theoretical cosmology is the explanation of the recent acceleration of the Universe's expansion~\\cite{Riess:1998cb}. In the standard $\\Lambda$-CDM scenario, it is the consequence of the existence of a cosmological constant, although it could be also due to a dark energy fluid whose origin has yet to be determined~\\cite{Copeland:2006wr}. Models of modified gravity~\\cite{Clifton:2011jh} complement dark energy scenarios and provide an explanation of the absence of long range fifth force effects in the solar system and laboratory experiments. Indeed most involve at least one scalar field coupled to matter, and eventually an environmental dependence leading to a screening mechanism of the scalar field in high density regions~\\cite{Khoury:2010xi}. This mechanism is an essential ingredient for the models to pass the stringent constraints on the possible modifications of gravity in the laboratory~\\cite{Adelberger:2002ic}, the solar system~\\cite{1990MNRAS.247..510S}, and the galactic environments~\\cite{Pourhasan:2011sm}. Moreover, the scalar fields are required to sit at the minimum of the density dependent effective potential prior to Big Bang Nucleosynthesis (BBN), so that catastrophic modifications in the formation of light elements are avoided. Numerous models of this type have been proposed. Let us mention the chameleons~\\cite{Khoury:2003rn,Mota:2006fz,Brax:2004px,Brax:2005ew,Brax:2004ym,Brax:2010kv,Gannouji:2010fc,Hees:2011mu}, involving a thin shell shielding the scalar field in dense bodies, the symmetrons~\\cite{Pietroni:2005pv,Olive:2007aj,Hinterbichler:2010es,Hinterbichler:2011ca,Brax:2011pk,Davis:2011pj,Clampitt:2011mx}, involving a symmetry breaking potential so that the scalar field is decoupled from matter at high densities, the dilatons~\\cite{Brax:2010gi,Brax:2011ja}, where the coupling to gravity turns off in dense environments, and the $f(R)$ models~\\cite{Starobinsky:1980te,Carroll:2003wy,Carroll:2004de,Faulkner:2006ub,Navarro:2006mw,Amendola:2006we,Carloni:2007br,Song:2006ej,Li:2007xn,Sawicki:2006jj,Brax:2011ja}, which are a sub class of chameleon models~\\cite{Brax:2011ja}. At the homogeneous level, all these scenarios coincide with the $\\Lambda$-CDM model. However, the evolution of linear perturbations differ. As a consequence, models of modified gravity can induce observable signatures in the matter power spectrum at redshifts $z \\lesssim 2$, and less importantly in the cosmic microwave background~\\cite{Hojjati:2011ix,Brax:2011ja} at $z \\simeq 1100$. Modified gravity models can also be probed with weak lensing (see e.g.~\\cite{Schmidt:2008hc}). During the dark ages and the reionisation period, \\textit{i.e.} in the range $1100 > z \\gtrsim 6 $, no cosmological signal have yet been observed. Such observations would be however of great interest for cosmology, especially for the study of modified gravity through the time evolution of the matter perturbations. In the near future, this gap is expected to be partially filled with the observation of the 21cm signal from reionisation~\\cite{Madau:1996cs,Ciardi:2003hg,Loeb:2003ya,Furlanetto:2006jb,Pritchard:2011xb}, and maybe in the more distant future, from the dark ages. During reionisation, transitions between the fundamental hyperfine levels of neutral hydrogen atoms are possible, via the Wouthuysen-Field effect involving the absorption and re-emission of Lyman-$\\alpha$ photons from the first stars (for a review, see Ref.~\\cite{Furlanetto:2006jb}). These induce the 21cm signal corresponding to a stimulated emission of 21cm photons against the Cosmic Microwave Background (CMB) radiation. The 3D power spectrum of the 21cm radiation maps the baryon distribution and thus is sensitive to modifications of gravity. In this paper, we explore for the first time the effects of modified gravity on the 21cm power spectrum at reionisation, and discuss the detectability of such effects with instruments of future generation giant Fast Fourier Transform radio-telescopes~\\cite{Tegmark:2008au,Tegmark:2009kv}. Several parameterisations of modified gravity have been proposed~\\cite{Bertschinger:2006aw,Hu:2007pj,Jain:2007yk,Amendola:2007rr,Zuntz:2011aq,Bertschinger:2008zb,Song:2008vm,Bean:2010zq,Daniel:2010ky,Pogosian:2010tj,Zhao:2011te,Skordis:2008vt,Ferreira:2010sz} for the evolution of linear perturbations such as, for instance, through the Poisson equation, $- k^2 \\Phi = 4 \\Pi (1 + \\nu ) G a^2 \\delta \\rho_{\\rr m}$ and $ \\Psi = (1+ \\gamma) \\Phi$, where $\\delta \\rho_{\\rr m}$ is the matter density perturbation and where $\\Phi$ and $\\Psi$ are the potentials in the Newtonian gauge. This parametrization involves two functions $\\nu(k,a)$ and $\\gamma(k,a)$ that depend both on time and on the perturbation wavenumber $k$. For $f(R)$ models, those functions depend on $B = (f_{RR} / f_R ) H \\dd R / \\dd H $ together with $f_{R0} $ today~\\cite{Song:2006ej}. In this paper, we adopt the parameterisation proposed in Refs.~\\cite{Brax:2011aw,Brax:2012gr}, for which the action and dynamics can be fully and uniquely reconstructed from two time-dependent functions: the coupling to matter $\\beta(a)$ and the scalar field mass $m(a)$. Of course, in the $f(R)$ case, it coincides with the usual approach. Moreover, for all these models, the $\\nu$ and $\\gamma$ functions can be explicitly obtained as a function of $m(a)$ and $\\beta(a)$. Using the $(m(a),\\beta(a))$ parameterisation, we calculate the signatures of $f(R)$, dilaton, chameleon and symmetron models on the 21cm power spectrum at reionisation, as well as on the present matter power spectrum in the linear approximation. For each model, we discuss the range of parameters that could be probed by 21cm experiments and compare it to the constraints from local experiments. Finally, we evaluate how the coupling to photons could be bounded via 21-cm constraints on the variation of the fine structure constant $\\alpha$. In all cases we find that the 21 cm signal obtained by varying the modes parallel to the line of sight is the most relevant. In the $f(R)$ case, we expect the constraints to be less stringent than the ones from the gravity tests in the solar system. For dilatons, the signal is found to be negligible while symmetron models with a transition at a redshift larger than $z_{\\star} \\sim 12$ could be detected even when the range of the symmetron interaction now is as low as $5$ kpc. This paper is organised as follows: In Sec.~\\ref{sec:21cm}, we briefly summarise the physics of the 21cm signal at the period of reionisation and derive its 3D power spectrum. We also give the evolution of the baryon and dark matter perturbations after the time of last scattering for the $\\Lambda$-CDM model. In Sec.~\\ref{sec:MGmodels}, we show how these equations are modified in the context of scalar models and describe the dynamics of $f(R)$, symmetron, chameleon and dilaton models using the reconstruction of the scalar field dynamics from the parametrization of Ref~\\cite{Brax:2011aw,Brax:2012gr}. In Sec.~\\ref{sec:21cmeffects} we give the specifications of the considered FFTT radio-telescope as well as the forecast errors on the 21cm power spectrum for single redshift measurements. We then evaluate the range of parameter values leading to observable effects and compare to the constraints from local tests of gravity. Our results are summarised in the conclusion. %These data, combined with galaxy, solar system, and laboratory experiments, ", "conclusions": "The 21 cm line can in principle be used to probe the evolution of the matter perturbations over a wide range of redshifts, typically from the dark ages up to the completion of the reionisation. Observing the 21cm cosmological signal should therefore further our understanding of the evolution of the Universe and it is thus important to investigate the predictions of different cosmological models on the 21cm three-dimentional power spectrum. In this paper we have considered modified gravity models with a screening mechanism and study their signatures on the 21cm power spectrum at reionization. Our archetypical experiment is the Fast Fourier Transform radio-Telescope (FFTT), consisting of a one kilometer side square of dipole antennas, that is designed especially for the detection of the 21cm power spectrum at reionisation. While the current and next generation of giant radio-telescope are expected to have a limited interest for cosmology, the ability of a FFTT-type experiment to put strong constraints on the various cosmological parameters has been demonstrated in~\\cite{Mao:2008ug}. We have investigated modified gravity models with a screening mechanism using a unified parametrisation whereby the models are characterised by the scale dependence of the coupling to matter and the mass of the scalar field. This has previously been shown to encapsulate all the effects of modified gravity with a screening mechanism~\\cite{Brax:2011aw} and has been used to investigate the effects of such models on large scale structure \\cite{Brax:2012gr}. For a fixed redshift, the consequences of modified gravity are important on small scales, within the Compton wavelength of the scalar field mediating the deviation from General Relativity. At the time of reionisation, such small scales could be in the linear regime of perturbations, where modifications of gravity enhance the growth of structure in a maximal way, whereas they are in the non-linear regime at the higher redshifts of the large scale structure probes, where screening effects take place and therefore reduce the magnitude of modified gravity effects. Hence 21cm cosmology offers a prime possibility of observing modification of gravity unhampered by screening effects. However, it is important to notice that with our specifications of the FFTT, such small scales can be only probed through wavelength modes parallel to the line of sight. The parametrisation of~\\cite{Brax:2011aw} is used here in the context of generalised chameleon, dilaton and symmetron models, as well as f(R) gravity, which is a special chameleon case. For all these models (except for dilatons), we find that 21cm observations at reionisation should constrain them more tightly than the large scale structures and CMB observations. The 21cm signal appears to be also a good discriminator of modified gravity models. Some models are already tightly constrained by the tests of gravity in the solar system, in the laboratory and in galactic environments. If these constraints are imposed, predictions for 21 cm cosmology are very similar to that of $\\Lambda$CDM and signatures of modified gravity should be hardly observable. However, in the case of generalised symmetron models, strong constraints could be established with 21cm observations. We have also considered the effect of the scalar field coupling to photons in modified gravity models. This coupling arises naturally from the conformal anomaly \\cite{Brax:2010uq} and gives rise to a time variation of the fine-structure constant. Since the 21 cm line is very sensitive to any variation in the fine structure constant, it can be used to probe the coupling to photons. This has enabled us to forecast tight bounds on this coupling on the basis of a 21cm experiment covering redshifts going from the beginning to the completion of the reionization. More accurate predictions will require sophisticated MMCM and Fisher matrix techniques. This will allow us to consider the whole ensemble of wavelength modes accessible to the experiment and to perform a multi-redshift analysis in order to set precise bounds on the various model parameters, taking account the possible degeneracies with other cosmological and reionisation parameters. This is left for future work." }, "1207/1207.1136_arXiv.txt": { "abstract": "{During an [O {\\sc iii}] survey for planetary nebulae, we identified a region in Sagittarius containing several candidate Supernova Remants and obtained deep optical narrow-band images and spectra to explore their nature. The images of the unstudied area have been obtained in the light of \\hnii, \\sii\\ and \\oiii. The resulting mosaic covers an area of $1.4\\degr \\times 1.0\\degr$~where filamentary and diffuse emission was discovered, suggesting the existence of more than one supernova remnants (SNRs) in the area. Deep long slit spectra were also taken of eight different regions. Both the flux calibrated images and the spectra show that the emission from the filamentary structures originates from shock-heated gas, while the photo-ionization mechanism is responsible for the diffuse emission. Part of the optical emission is found to be correlated with the radio at 4850 MHz suggesting their association, while the WISE infrared emission found in the area at 12 and 22 $\\mu$m marginally correlates with the optical. The presence of the \\oiii\\ emission line in one of the candidate SNRs suggests shock velocities into the interstellar \"clouds\" between 120 and 200 km s$^{-1}$, while the absence in the other indicates slower shock velocities. For all candidate remnants the [S {\\sc ii}] $\\lambda\\lambda$ 6716/6731 ratio indicates electron densities below 240 cm$^{-3}$, while the \\ha\\ emission has been measured to be between 0.6 to 41$\\times$\\flux. The existence of eight pulsars within 1.5$\\degr$ away from the center of the candidate SNRs also supports the scenario of many SNRs in the area as well as that the detected optical emission could be part of a number of supernovae explosions.} ", "introduction": "Supernova explosions belong to the most spectacular events in the Universe. Observations of galaxies reveal several events every year \\citep{manu05}, where the supernova is of comparable brightness to the entire galaxy for days up to weeks. Supernova remnants (SNRs) which are the consequent result of such events are some of the strongest radio sources observed. SNRs have a major influence on the properties of the interstellar medium (ISM) and on the evolution of galaxies as a whole. They enrich the ISM with heavy elements, release about 10$^{51}$ ergs and heat the ISM, compress the magnetic field and efficiently accelerate in their shock waves energetic cosmic rays as observed throughout the Galaxy. The majority of known SNRs have been discovered by their non-thermal radio emission \\citep{gre09} while a smaller number of them are observed in other wavelengths (e.g optical; \\citealt{bou08, bou09}, X-rays; \\citealt{rey09}, infrared; \\citealt{rea06}). In this paper, we report the optical detection of many filamentary and diffuse structures (possibly more than one SNR) in the region of Sagittarius constellation. During an \\OIII\\ survey for planetary nebulae \\citep{bou03,bou06}, we identified a very strong \\sii\\ source designated as candidate supernova remnant instead of planetary nebula. Following that detection, a number of images in \\hnii, \\sii\\ and \\oiii\\ were taken in order to explore the SNR candidate area and many filamentary structures were discovered. Only one known SNR was found in the area (G 16.2--2.7; \\citealt{tru99}), hence all the other filamentary structures were examined in detail in order to identify their origin. Information about the observation and data reduction is given in Sect. 2. In Sect. 3 and 4 the results of the imaging and spectra observations are presented, while in Sect. 5 we report on observations in wavelengths other than the optical. Finally, in Sect. 6 we discuss the properties of the new candidate SNRs. ", "conclusions": "The newly discovered candidate SNRs towards the Sagittarius constellation show up as incomplete circular or elliptical structures in the optical and in most of the cases in the radio and without any X--ray and H {\\sc i} emission detected so far. The absence of soft X--ray emission may indicate a low shock temperature and/or a low density of the local interstellar medium. Its optical emission marginally correlates with the infrared. The elliptical shape of the candidate SNRs is unusual compared to most known SNRs suggesting that the surrounding medium is very irregular with not constant interstellar density. However, it should be mentioned that a significant number of known SNRs show not circular structure (bilateral, barriel, elliptical, cilindrical etc.) depending on their position to the line of sight so this might also be a reason of their appearance. Detailed optical observations have been performed in an attempt to understand the nature of the candidate SNRs. The lower ionization images in \\hnii\\ and \\sii\\ reveal several filamentary and diffuse structures while the higher ionization image in \\oiii\\ shows emission only in one region. The \\hnii\\ image best describes the newly detected structures. Sulfur line emission is also detected and generally appears less filamentary and more diffuse than in the \\hnii\\ image with their position and shape in agreement with that of the \\hnii. The \\oiii\\ flux production depends mainly on the shock velocity and the ionization state of the preshocked gas. Therefore, as mentioned in Sect. 3.1, the absence of \\oiii\\ emission in almost all of the areas may be explained by slow shocks propagating into the ISM. The presence of [O {\\sc i}] 6300 \\AA\\ line emission is also consistent with the emission being from shocked material. The \\oiii/\\hbeta\\ ratio is a very useful diagnostic tool for complete or incomplete shock structures \\citep{ray88}. However, the absence of \\oiii\\ does not allow us to suggest for complete or incomplete shock structures apart from pos.1 where shocks with incomplete recombination zones should be present. Both the calibrated images and the long--slit spectra suggest that the detected emission results from shock heated gas since the \\sii/\\ha\\ ratio exceeds the empirical SNR criterion value of 0.4--0.5, while the measured \\nii/\\ha\\ ratio also confirms this result. \\par The SNR origin of the proposed candidate remnants is strongly suggested by the positions of the line ratios in Fig.~\\ref{fig4} compared with those of Herbig--Haro objects, H{\\sc ii} regions and planetary nebulae (PNe). They follow closely the shape of those observed for those of shock ionized evolved SNRs. \\par The \\ha/\\hbeta\\ ratios in Table \\ref{table4} can be used to estimate the variations in logarithmic extinction coefficient c over these sources, assuming an intrinsic ratio of 3 and the interstellar extinction curve as implemented in the nebular package \\citep{sha95} within the IRAF software. An interstellar extinction c (see Table~\\ref{table4}) between 0.2 and 1.1 or an A$_{\\rm V}$~between 0.4 and 2.2 were measured, respectively. We have also determined the electron density measuring the density--sensitive line ratio of \\siirat. The measured densities lie below 240 \\dens. % \\par The candidate remnants under investigation have not been studied in the past hence the current stage of their evolution is unknown. Our aim is to provide a first indication of their stage of evolution by estimating basic SNR parameters, assuming that the temperature is close to 10$^{4}$ K. Estimated values of N$_{\\rm H}$ between 4.8 and 6.8 $\\times 10^{21}$~cm$^{-2}$ and N$_{\\rm H}$ between 3.2 and 7.4 $\\times 10^{21}$~cm$^{-2}$ are given by \\citet{dic90} and \\citet{kal05} respectively, for the column density in the direction of the candidate remnants. Using the relations of \\citet{ryt75} and \\citet{pre95}, we obtain N$_{\\rm H}$ between 0.9 and 4.9 $\\times 10^{21}~{\\rm cm}^{-2}$~for the minimum and maximum c values calculated from our spectra. In Table~\\ref{table3}, we present the estimated values of N$_{\\rm H}$ for each candidate remnant, where it can be seen that the values based on the optical data and the statistical relations are consistent with the estimated galactic N$_{\\rm H}$~ from \\citet{kal05} and less by that estimated from \\citet{dic90}. However, it should be noted that the slightly higher values calculated by the latter method can be explained by the fact that it also covers gas beyond the area of interest. Assuming that they are still in the adiabatic phase of their evolution the preshock cloud density n$_{\\rm c}$ can be measured by using the relationship \\citep{dop79} \\begin{equation} {\\rm n_{[SII]} \\simeq\\ 45\\ n_c V_{\\rm s}^2}~{\\rm cm^{-3}}, \\end{equation} where ${\\rm n_{[SII]}}$ is the electron density derived from the sulfur line ratio and V$_{\\rm s}$ is the shock velocity into the clouds in units of 100 \\vel. Furthermore, the blast wave energy can be expressed in terms of the cloud parameters by using the equation given by \\citet{mck75} \\begin{equation} {\\rm E_{51}} = 2 \\times 10^{-5} \\beta^{-1} {\\rm n_c}\\ V_{\\rm s}^2 \\ {\\rm r_{s}}^3 \\ \\ {\\rm erg}. \\end{equation} The factor $\\beta$ is approximately equal to 1 at the blast wave shock, ${\\rm E_{51}}$ is the explosion energy in units of 10$^{51}$ erg and {\\rm r$_{\\rm s}$} the radius of the remnant in pc. By using the upper limit on the electron density of 240 \\dens, which was derived from our spectra, we obtain from Eq. (1) that ${\\rm n_c} V_{\\rm s}^2 < 5.3$. Then Eq. (2) becomes ${\\rm E_{51}} < \\alpha \\times 10^{-3}~{\\rm D_{1 kpc}^3}$, where $\\alpha$ is a value depends on the diameter of each candidate remnant and ${\\rm D_{1 kpc}}$~the distance to the remnant in units of 1 kpc. For the different diameters of the remnants and assuming the typical value of 1 for the supernova explosion energy (E$_{51}$), we derive that the candidate SNRs lie at distances greater than 8 kpc. Since, there are no other measurements of the interstellar density n$_{0}$, values of 0.1 and 1.0 will be examined. Then, the lower interstellar density of $\\sim$0.1 cm$^{-3}$~suggests that their distance is between 10 and 22 kpc while for n$_{0} \\approx 1~{\\rm cm}^{-3}$~it is between 1.0 and 2.2 kpc, for the lower and higher column densities calculated above. Combining the previous results, values between 0.1 and 0.2 cm$^{-3}$~for the interstellar density seem to be more probable. It should also be noted that the ambient density of gas around the candidate SNRs is not the same so this might also change the estimated vaues. We also searched for pulsars in the region using the ATNF Pulsar Catalogue \\citep{man05}. In total, we found 8 pulsars within a 1.5$\\degr$ diameter circle away from the center of each candidate SNR. In Table 5, we present their names, coordinates and rotation period as well as which candidate SNRs fulfill the 1.5$\\degr$ limit. The closest one is PSR1826-1526 to G 15.8$-$1.9 at a distance of 0.7$\\degr$. It is not clear at the moment if any of these pulsars are related to the candidate SNRs, however, the existence of a significant number of pulsars very close to the area of the candidate SNRs it is another strong indication of the existence of more than one SNRs in the region. It is more plausible that due to their distance from the candidate SNRs they are not associated with them since the closest pulsar is about 4 radii from the nearest SNR or $>$ 75 pc in the plane of the sky at a distance of 8 kpc. However, if they are not in the plane of the sky which is probably the case, then these numbers might change. Also, they might be at different distances which change the numbers, too. Therefore, their correlation cannot be confirmed or ruled out and it should be further examined in the future in detail. It is possible that this irregular group of filaments is part of a wider structure but is being seen through holes in intervening clouds, leading to patchy optical interstellar extinction. The current data are not sufficient to claim strong a correlation. Furthermore, there is a possibility that the detected optical emission could be part of a number of supernova explosions in the area. However, since neither the distance nor the interstellar medium density are accurately known, we cannot confidently determine the current stage of evolution of the candidate remnants and more observations are needed. Further study of the area would benefit from higher resolution multiwavelength observations (optical, radio and X--ray) and will help to clarify the current uncertainties. In particular, higher resolution imaging observations will verify their filamentary structure and confirm their morphological appearance, while kinematic observations will help to determine their 3-D morphology and measure expansion velocities. X--ray observations will help in order to clarify if there is any correlation of the faint and diffuse X--ray emission with the optical filaments and provide more information about their evolutionary stage, while radio observations would also be useful to examine their non--thermal spectral index and confirm their SNR nature." }, "1207/1207.1070_arXiv.txt": { "abstract": "We subject the steady solutions of a spherically symmetric accretion flow to a time-dependent radial perturbation. The equation of the perturbation includes nonlinearity up to any arbitrary order, and bears a form that is very similar to the metric equation of an analogue acoustic black hole. Casting the perturbation as a standing wave on subsonic solutions, and maintaining nonlinearity in it up to the second order, we get the time-dependence of the perturbation in the form of a Li\\'enard system. A dynamical systems analysis of the Li\\'enard system reveals a saddle point in real time, with the implication that instabilities will develop in the accreting system when the perturbation is extended into the nonlinear regime. The instability of initial subsonic states also adversely affects the temporal evolution of the flow towards a final and stable transonic state. ", "introduction": "\\label{sec1} The classical model of steady spherically symmetric accretion~\\cite{bon52} is a mathematical problem of conservative and compressible hydrodynamics. This model has acquired a paradigmatic status in studies on astrophysical accretion, which, as a fluid flow, falls under the general class of nonlinear dynamics. The mathematical description of such fluid systems involves a momentum balance equation (with gravity as an external force in accretion), the continuity equation and a polytropic equation of state~\\cite{fkr02}. Set in full detail, the condition of momentum conservation in a fluid is a balance of dynamic effects, nonlinear effects and the effects of the pressure in a continuum system~\\cite{ll87}. Some early studies on astrophysical accretion considered only the interplay between dynamics and nonlinearity (see~\\cite{bon52} and references therein), but in present studies, it is customary to ignore the dynamics and instead consider the effects of pressure, in what becomes a stationary flow. In either case, however, nonlinearity endures. From the plethora of mathematical solutions of the stationary spherically symmetric compressible fluid flow, the ones of physical relevance are identified to be locally subsonic very far away from the accretor. Within the class of inflows obeying this outer boundary condition, there is an infinitude of globally subsonic solutions, along which a fluid element may reach the accretor with a low subsonic velocity. For the same outer boundary condition, a single critical solution stands out in a class by itself, and allows matter to reach the accretor with a high supersonic velocity, crossing the sonic horizon along the way. This is the unique transonic~\\citet{bon52} accretion solution. The exact fashion in which accreting matter reaches the accretor is related to the inner boundary condition of the inflow problem. If the accretor is a black hole, the infall process must be transonic~\\cite{nt73,shateu83}. This is because a black hole has an event horizon instead of a physical surface, and thus precludes all possibility of a pressure build-up at small radii, which could otherwise have dominated over the free-fall conditions close to the accretor. The situation is not so clearly understood if the accretor has a hard surface like a neutron star or a white dwarf. For such an accretor, it is supposed that the accumulated matter will build up pressure near the surface and cause the supersonic flow to be shocked down to subsonic levels, although for a neutron star, all accreted matter might be ``vacuum cleaned\" away efficiently, making it easier for the flow to remain supersonic~\\cite{pso80}. Evidently then, the question of the inner boundary condition and an inflow trajectory in relation to it, is by no means a trivial one. Nevertheless, working with the stationary problem itself,~\\citet{bon52} invoked the physical criteria of the maximization of the accretion rate and the minimization of the total energy of the flow, to propose that the transonic solution would be the one selected by a fluid element to reach the accretor from a distant outer boundary. \\citet{bon52}, however, left a definitive conclusion regarding the realizability of the transonic solution to its stability. The trouble with the transonic solution in the stationary regime is that its realizability is extremely vulnerable to even an infinitesimal deviation from the precisely needed boundary condition to generate the solution~\\citep{rb02}. This difficulty may be overcome by considering the temporal evolution of global solutions towards the transonic state~\\citep{rb02,rrbon07}, but there is no analytical formulation to solve the nonlinear partial differential equations governing the temporal evolution of the flow. So, much of all time-dependent studies in spherically symmetric accretion is perturbative and linearized in character, although some non-perturbative studies are also known~\\citep{rb02,rrbon07}. The commonly accepted view to have emerged from the linearized approach is that perturbations on the flow do not produce any linear mode with an amplitude that gets amplified in time~\\citep{gai06}, and that the perturbative method does not indicate the primacy of any particular class of solutions~\\citep{gar79}. This is as far as could be said, working in the linear regime, but our general experience of any nonlinear system is that an understanding gained about it under linearized conditions can scarcely be imposed on circumstances dominated by nonlinearity. In this work we attempt to bridge the gap. First we adopt a time-dependent radial perturbation scheme implemented originally by~\\citet{pso80} and retain all orders of nonlinearity in the resulting equation of the perturbation. A most striking feature of this equation is that even on accommodating nonlinearity in full order, it conforms to the structure of the metric equation of a scalar field in Lorentzian geometry (Section~\\ref{sec3}). This fluid analogue (an ``acoustic black hole\"), emulating many features of a general relativistic black hole, is a matter of continuing interest in fluid mechanics from diverse points of view~\\citep{monc80,un81,jacob91,un95,vis98,bil99,su02,das04,sbr05, us05,vol05,vol06,rbcqg07,rbpla07,rrbon07,ncbr07,dbd07,macmal08, blv11,rob12,sbbr13}. Then we apply our nonlinear equation of the perturbation to study the stability of globally subsonic stationary solutions under large-amplitude time-dependent perturbations. Our motivation to do so lies in certain dynamic features of the flow. For the non-perturbative time evolution of the accreting system, the initial condition of the evolution is a globally subsonic state, with gravity subsequently driving the system to a transonic state, sweeping through an infinitude of intermediate subsonic states. So, to ensure an unhindered temporal convergence to a stable transonic trajectory, the stability of the subsonic states is crucial. In a numerical study, an instability in the subsonic states was observed by~\\citet{sb78}, who were consequently of the view that the subsonic flows would quickly change into the transonic flow. Our nonlinear perturbative analysis does agree with the fact that there is an instability in the subsonic states, but short of an exact non-perturbative analysis of all the nonlinear flow equations concerned (for which there exists no analytical prescription yet), it would be hasty to claim that the transonic state is indeed the final stable attractor state for the unstable subsonic states. The feasible analytical alternative, therefore, is to study the behaviour of the system under progressively higher orders (nonlinear orders) of time-dependence in the perturbative approach. We truncate all orders of nonlinearity beyond the second order in our equation of the perturbation. Following this, we integrate out the spatial dependence of the perturbation with the help of well-defined boundary conditions on globally subsonic flows~\\citep{pso80,td92}. After this, we extract the time-dependent part of the perturbation and note with intrigue that it has the mathematical form of a Li\\'enard system~\\citep{stro,js99} (Section~\\ref{sec4}). On applying the standard analytical tools of dynamical systems to study the equilibrium features of this Li\\'enard system, we discover the existence of a saddle point in real time, whose implication is that the stationary background solutions will be unstable, if the perturbation is extended into the nonlinear regime (Section~\\ref{sec5}). We also provide independent numerical support in favour of our analytical findings on the dynamics~(Sections~\\ref{sec5}~\\&~\\ref{sec6}). ", "conclusions": "\\label{sec7} It will be well in order now to make some general remarks about our work, to put it in perspective. First, accretion being a fluid problem, is very much within the realm of nonlinear dynamics. Our work addresses the nonlinear aspects of the dynamics of an accretion process, by having recourse to the usual analytical tools of nonlinear dynamics. One salient outcome of the nonlinear approach is obtaining an acoustic metric, in spite of accommodating nonlinearity completely. This marks a significant departure from the linear treatment (small perturbation) of the problem. Another new result of this work is the discovery of a Li\\'enard system (a nonlinear oscillator) in a very common and basic model of accretion. A noteworthy aspect of all these new results is that they have been extracted from a system that has been known to the astrophysical community for more than sixty years --- a conservative, non-self-gravitating, compressible fluid inflow, driven by Newtonian-like external gravitational fields, with coupled density and velocity fields. It is a very simple system in essence, and yet it continues to offer novel insights. Going by the form of the Li\\'enard system derived here, it is easy to see that the number of equilibrium points will depend on the order of nonlinearity that we may wish to retain in the equation of the perturbation. In practice, however, the analytical task becomes formidable with the inclusion of every higher order of nonlinearity. Going up to the second order, an instability in real time appears undeniable, but then we must realize that this conclusion has been made regarding a purely inviscid and conservative flow. Real fluids have viscosity as another important physical factor to influence their dynamics. In fact, fluid flows are usually affected by both nonlinearity and viscosity, occasionally as competing effects, and apropos of this point, we note that for a linearized radial perturbation in spherically symmetric inflows, viscosity helps in decaying the amplitude of the standing waves on globally subsonic solutions~\\citep{ray03}. So the instability that arises because of nonlinearity can very well be tempered by viscosity in the flow~\\citep{sb78}. Closely related to viscous dissipation, the stability of spherically symmetric accretion is expected to be affected by turbulence as well~\\citep{rb05}. A suitable mechanism that favours stability may be found in accretion onto black holes, where the coupling of the flow with the geometry of space-time acts in the manner of a dissipating effect. General relativistic effects have been known to enhance the stability of the flow~\\citep{ncbr07}. And the stability of fluids may also be studied by constraining a perturbation to behave like a travelling wave~\\citep{pso80,rbpla07,ncbr07,sbbr13}. At times, we encounter the surprising situation of a fluid flow being stable under one type of perturbation, but unstable under the effect of another~\\citep{rbpla07}. Formally involving nonlinearity, all these features merit a close examination." }, "1207/1207.5186_arXiv.txt": { "abstract": "The surface tension of quark matter plays a crucial role for the possibility of quark matter nucleation during the formation of compact stellar objects, because it determines the nucleation rate and the associated critical size. However, this quantity is not well known and the theoretical estimates fall within a wide range, $\\gamma_0\\approx5-300{\\rm\\ MeV/fm^2}$. We show here that once the equation of state is available one may use a geometrical approach to obtain a numerical value for the surface tension that is consistent with the model approximations adopted. We illustrate this method within the two-flavor linear $\\sigma$ model and the Nambu--Jona-Lasinio model with two and three flavors. Treating these models in the mean-field approximation, we find $\\gamma_0 \\approx7-30{\\rm\\ MeV/fm^2}$. Such a relatively small surface tension would favor the formation of quark stars and may thus have significant astrophysical implications. We also investigate how the surface tension decreases towards zero as the temperature is raised from zero to its critical value. ", "introduction": "Lattice-gauge calculations yield a non-vanishing value of the quark condensate $\\langle{\\overline \\psi} \\psi \\rangle$ in the QCD vacuum \\cite{Aoki}, indicating that chiral symmetry is broken. This general feature of the vacuum remains present even for massless quarks because the symmetry is then broken spontaneously. On the other hand, chiral symmetry is expected to become restored at sufficiently high values of the net-baryon density $\\rho$ or/and the temperature $T$. The character of this phase change is not yet well understood but it has significant implications in areas such as cosmology and astrophysics and it is a focal point for current experimental and theoretical research in nuclear physics. Nuclear collision experiments carried out with the Relativistic Heavy Ion Collider (RHIC) at the Brookhaven National Laboratory and with the Large Hadron Collider at CERN explore systems having relatively small net-baryon densities $\\rho$ and the associated chemical potentials $\\mu$ are negligible. Lattice calculations can readily be carried out at vanishing $\\mu$ and they indicate that a cross-over transformation from the chirally broken phase to the restored phase occurs as the temperature is increased from below to above the cross-over temperature $T_\\times\\approx160\\,{\\rm MeV}$ \\cite{Aoki,Aoki:2006br,Bazavov:2011nk}. The other extreme region of the QCD phase diagram, namely low temperatures and high chemical potentials, cannot be addressed by current lattice-QCD methods, due to the fermion sign problem, and studies of this phase region must therefore rely on less fundamental models. Most investigations suggest that there is a first-order phase transition which, for $T\\approx0$, sets in at baryon densities several times that of the nuclear saturation density, $\\rho_0\\approx0.153/\\fm^3$. The properties of strongly interacting matter in this phase region are important for our understanding of compact stars. If indeed such a first-order phase transition exists at $T=0$, then, as the temperature is raised, one would expect it to remain present but gradually weaken and eventually terminate at a critical point $(\\mu_c,T_c)$. The existence and location of such a critical point is a subject of intense theoretical investigation with a variety of models, including in particular effective-field models, such as the linear $\\sigma$ model (LSM), and effective quark models, such as the Nambu--Jona-Lasinio (NJL), at different levels of sophistication considering up to three quark flavors and possibly including the Polyakov loop to account for confinement \\cite {pedropoly,mao}. Experimentally the corresponding region of density and excitation may be produced in current nuclear collisions at the low-energy end of RHIC and in the future with FAIR at GSI and NICA at JINR which are being constructed with such investigations in mind. In the present work, we concentrate on the high-$\\mu$ and low-$T$ part of the phase diagram with the aim of exploring the expected chiral phase transition which has significant implications for the possible existence of quark stars \\cite {ed1,bie}. It should be noted that chiral symmetry may be restored already during the early post-bounce accretion stage of a core-collapse supernova event and the associated neutrino burst might then provide a spectacular signature for the presence of quark matter inside compact stars \\cite {sagert}. However, as pointed out in Refs.\\ \\cite{bruno,bombaci}, the possibilities depend on the dynamics of the phase conversion and especially on the time scales involved. When the phase diagram of bulk matter exhibits a first-order phase transition, the two phases may coexist in mutual thermodynamic equilibrium and, consequently, when brought into physical contact a mechanically stable interface will develop between them. The associated interface tension $\\gamma_T$ (which we shall often refer to simply as the surface tension of quark matter) depends on the temperature $T$; it has its largest magnitude at $T=0$ and approaches zero as $T$ is increased to $T_c$. This quantity plays a key role in the phase conversion process and it is related to various characteristic quantities such as the nucleation rate, the critical bubble radius, and the favored scale of the blobs generated by the spinodal instabilities \\cite{PhysRep,RandrupPRC79}. (As we shall see, the surface tension is essentially proportional to the effective interaction range, which determines the width of the surface region, and the spatial size of the most rapidly amplified density irregularity is also proportional to this quantity.) Unfortunately, despite its central importance, the surface tension of quark matter is rather poorly known. Estimates in the literature fall within a wide range, typically $\\gamma_0\\approx10-50~\\MeV/\\fm^2$ \\cite{heiselberg,sato} and values of $\\gamma_0\\approx30\\,\\MeV/\\fm^2$ have been considered for studying the effect of quark matter nucleation on the evolution of proto-neutron stars \\cite{constanca}. But the authors in Ref.\\ \\cite {russo}, taking into account the effects from charge screening and structured mixed phases, estimate $\\gamma_0\\approx50-150~\\MeV/\\fm^2$, without excluding smaller values, and an ever higher value, $\\gamma_0\\approx300~\\MeV/\\fm^2$, was found by Alford {\\em et al.}\\ \\cite {alford} on the basis of dimensional analysis of the minimal interface between a color-flavor locked phase and nuclear matter. The surface tension for two-flavor quark matter was evaluated within the framework of the LSM by Palhares and Fraga \\cite{leticia}. In that work the authors considered the one-loop effective potential and then fitted its relevant part, which included both the chirally symmetric and broken state, by a quartic polynomial. The surface tension was evaluated using the thin-wall approximation for bubble nucleation and the estimated values cover the $5-15\\,\\MeV/\\fm^2$ range, depending on the inclusion of vacuum and/or thermal corrections. In principle, this range makes nucleation of quark matter possible during the early post-bounce stage of core-collapse supernovae and it is thus a rather important result. It is also worth noting that a small surface tension would facilitate various structures in compact stars, including mixed phases \\cite {kurkela}. The present work is devoted to the evaluation of the surface tension for quark matter using both the LSM (with two flavors) and the NJL model (with two and three flavors) following the procedure employed in Ref.\\ \\cite {RandrupPRC79}. Here, the LSM is mainly included to check the consistency of our procedure by comparing our present results with those obtained by the thin-wall approximation of Ref.\\ \\cite{leticia} (we find the agreement to be very good). The NJL model is considered with two and three flavors because the latter, which contains strangeness, is one of the most popular effective quark models used in studies related to compact stars. As explained below, the method described in Ref.\\ \\cite{RandrupPRC79} makes it possible to express the surface tension for any subcritical temperature in terms of the free energy density for uniform matter in the unstable density range. Because the models employed readily provide the equation of state (EoS) for the full density range, they are well suited for our purpose and we may directly employ the method without any further approximations. In practice, the procedure is rather simple to implement and it provides an estimate for the surface tension that is consistent with the EoS implied by the model employed, with its specific approximations and parametrizations. The paper is organized as follows. In Sect.\\ \\ref{method} we review the method for extracting the surface tension from the equation of state. In Sec.\\ \\ref{models} we then present the two-flavor versions of the two models considered and discuss how to extract the surface tension. Section \\ref{NJL3} is devoted to the treatment for the more realistic $SU(3)$ version of the NJL model and our numerical results are presented in Sec.\\ \\ref{results}, both for cold matter and for temperatures up to the critical value. The conclusions and final remarks are presented in Sec.\\ \\ref{conclude}. ", "conclusions": "\\label{conclude} In this work we have shown that the interface tension related to a first-order phase transition may be evaluated once the uniform-matter equation of state is available for the unstable regions of the phase diagram. It is a convenient feature of the method employed that knowledge of the interface profile functions is is not required, because their determination can be quite complicated, as is the case for NJL model \\cite{russos} (although it is easy for the LSM). In addition to the EoS, the geometrical approach also requires a proper setting of three input parameters, namely the characteristic densities $\\rho_g$ and ${\\cal E}_g$ together with the length scale $a$. While this does encumber the numerical results with some degree of uncertainty, our zero-temperature LSM result, $\\gamma_0=13.18\\,\\MeV/\\fm^2$, agrees within a few percent with the approximate value obtained in Ref.\\ \\cite{leticia}, thus suggesting that those parameters were chosen reasonably. The surface tension determined in the present fashion is entirely consistent with the employed model, including the approximations and parametrizations adopted. For the effective quark models employed here, this amounts to considering {\\it all} the solutions to the gap equation (stable, metastable and unstable) and determine the relevant effective quark masses. In most non-perturbative approximations (large $N_c$, mean field, {\\em etc.}) the various quantities of interest, such as the free energy density, become functions of this effective mass and will therefore also reflect the metastable and unstable character of the configuration considered. As a cross check on our procedure, we have evaluated $\\gamma_0$ for the LSM obtaining a result that differs by only about $2\\%$ from estimates based on the thin-wall approximation \\cite{leticia}. We have investigated the two-flavor NJL model as well as its more realistic three-flavor version. Our main conclusion is that all these effective models generate relatively low values for the the surface tension. This would favor the formation of quark matter and may thus have important astrophysical consequences regarding the existence of pure quark stars. Of particular interest is the three-flavor NJL result, $\\gamma_0=20.34\\,\\MeV/\\fm^2$, because this model is widely used in studies related to neutron stars. Here, for simplicity, we have considered pure quark matter where all flavors share the same chemical potential, but it is just a technical matter to generalize our procedure so as to include leptons ($e,\\mu$) in order to enforce $\\beta$ equilibrium ($\\mu_d=\\mu_s=\\mu_u +\\mu_e \\, , \\, \\mu_e=\\mu_\\mu$), although the additional chemical potential introduces an increased degree of complexity into the features of the phase transition. In principle, more refined treatments, such as the Polyakov-NJL model, can also be considered within the same framework. However, because the effects of the Polyakov loop become more important above $100\\,\\MeV$ \\cite{costa} we believe that our results, especially the three-flavor ones, can be considered as reasonably accurate, although numerical variations may arise due to the parametrizations and approximations adopted." }, "1207/1207.2135_arXiv.txt": { "abstract": "The similarity of the observed baryon and dark matter densities suggests that they are physically related, either via a particle physics mechanism or anthropic selection. A pre-requisite for anthropic selection is the generation of superhorizon-sized domains of different $\\Omega_{B}/\\Omega_{DM}$. Here we consider generation of domains of different baryon density via random variations of the phase or magnitude of a complex field $\\Phi$ during inflation. Baryon isocurvature perturbations are a natural consequence of any such mechanism. We derive baryon isocurvature bounds on the expansion rate during inflation $H_{I}$ and on the mass parameter $\\mu$ which breaks the global $U(1)$ symmetry of the $\\Phi$ potential. We show that when $\\mu \\lae H_{I}$ (as expected in SUSY models) the baryon isocurvature constraints can be satisfied only if $H_{I}$ is unusually small, $H_{I} < 10^{7} \\GeV$, or if non-renormalizable Planck-suppressed corrections to the $\\Phi$ potential are excluded to a high order. Alternatively, an unsuppressed $\\Phi$ potential is possible if $\\mu$ is sufficiently large, $\\mu \\gae 10^{16} \\GeV$. We show that the baryon isocurvature constraints can be naturally satisfied in Affleck-Dine baryogenesis, as a result of the high-order suppression of non-renormalizable terms along MSSM flat directions. ", "introduction": "The cosmological dark matter and baryon mass densities are observed to be within an order of magnitude of each other, $\\Omega_{B}/\\Omega_{DM} \\approx 1/5$ \\cite{wmap}. However, baryogenesis and dark matter production are often physically unrelated in particle physics models. So why are these densities similar? It is possible to produce both dark matter and baryon number simultaneously, thereby directly relating their number densities. Such models are usually based on an overall conserved charge which is shared by the baryons and dark matter particles. This implies that dark matter is asymmetric with a small dark matter particle mass, $m_{DM} \\sim 1-10 \\GeV$. However, in the case where thermal relic WIMPs are the explanation for dark matter, a particle physics mechanism cannot simply relate the baryon and dark matter number densities directly to each other (as in the charge conservation models), but must specifically relate the baryon asymmetry to the thermal relic WIMP density. This implies a connection between the weak annihilation freeze-out process responsible for the thermal relic WIMP density and the mechanism determining the observed baryon asymmetry. Recently there have been some proposals which make this connection, based either on the modification of a pre-existing baryon asymmetry ({\\it baryomorphosis}) \\cite{bm,bm2} or on the generation of the baryon asymmetry via annihilation of dark matter ({\\it WIMPy baryogenesis}) \\cite{wimpy}. Such mechanisms require a number of additional particles and are strongly constrained by B washout. Since the new particles are necessarily at the TeV scale, these models may be testable at the LHC. The alternative is anthropic selection. Anthropic selection models have two components: (i) a mechanism to generate domains\\footnote{By 'domain' we mean any patch of the Universe with different conditions from ours. This could also include domains in different inflationary patches, although we will focus on a single inflated patch.} with varying $\\Omega_{B}/\\Omega_{DM}$ and (ii) the assumption that domains with $\\Omega_{B}/\\Omega_{DM} \\sim 1$ are favoured by the evolution of observers. An example of such a model was proposed in \\cite{lindeax}. In this model dark matter is due to a condensate of axions with a domain-dependent density, while the baryon number density is assumed fixed. Domains with average dark matter densities larger than in our domain result in the formation of galaxies with baryon and dark matter densities which are strongly enhanced relative to the average. The enhancement is due to perturbations becoming non-linear earlier \\cite{lindeax}. The enhanced dark matter and baryon densities in galaxies are then assumed to provide the required anthropic cut-off. However, if the dark matter density is fixed throughout the Universe, as in the case of thermal relic WIMPs, we need an alternative way to vary $\\Omega_{B}/\\Omega_{DM}$. Here we consider varying the baryon density between domains. It is, in principle, easy to vary the baryon density on superhorizon scales. All that is necessary is that the CP-violating phase or strength of B-violation depends on a field which is effectively massless until the onset of baryogenesis. During inflation the field can take random values on scales much larger than the horizon when the observed Universe exits the horizon at $N = 60$ e-foldings before the end of inflation. Therefore superhorizon domains with different baryon number will exist at present. It is therefore likely that there will exist some domains with $\\Omega_{DM} \\sim \\Omega_{B}$. It is also likely that the largest field value (magnitude or phase) will have the largest probability, in which case we will most likely live in a domain with the largest possible baryon asymmetry up to anthropic selection effects. As in the axion model, an average baryon density which is larger than the observed baryon density will be enhanced to a much larger baryon density in galaxies, which may then serve as an anthropic cut-off. We will refer to such models as {\\it anthropic baryogenesis} models in the following. In order for the baryon density in a domain to be random, it should not be determined purely by the parameters of the $\\Phi$ potential. For example, suppose the CP-violating phase $\\theta$ of a complex field $\\Phi = \\phi e^{i \\theta}/\\sqrt{2}$, which is effectively massless during inflation ($m_{\\Phi}^{2} \\ll H_{I}^{2}$), determines the baryon asymmetry. (The CP-conserving direction can be defined to be $\\theta = 0$.) In a 'typical' domain we expect $\\theta \\sim \\pi$. The baryon asymmetry will then be near maximal and will be essentially determined by the parameters of the $\\Phi$ potential. As a result, a coincidence between the maximal baryon asymmetry and the DM density is required; there is no real anthropic selection. In order to have a randomly-varying baryon asymmetry, the domain which has the observed baryon asymmetry must be {\\it atypical}, with $\\theta \\ll 1$. In this case there is no direct connection between the baryon asymmetry and the parameters of the potential and so no element of coincidence. Moreover, the baryon density in neighbouring domains can then be much larger or smaller than in our domain, allowing anthropic selection to function, whereas in the case where $\\theta \\sim \\pi$ only O(1) fractional increases in the baryon asymmetry relative to our domain are possible. However, the dependence on a massless complex scalar has a consequence that will impose a strong constraint on any anthropic selection mechanism of this type; quantum fluctuations of the massless field will produce baryon isocurvature pertubations. We will show that the atypically small value of $\\theta$ enhances the baryon isocurvature perturbations, resulting in strong constraints on anthropic baryogenesis models. The paper is organized as follows. In Section 2 we discuss the effect of a varying average baryon density on the properties of galaxies in neighbouring domains. In Section 3 we consider general constraints on anthropic baryogenesis models. In Section 4 we consider the case of Affleck-Dine baryogenesis. In Section 5 we present our conclusions. ", "conclusions": "The similarity of the observed baryon and dark matter densities suggests that there is a physical process connecting them. This similarity is particularly difficult to understand in the case of thermal relic WIMP dark matter, since this requires an explanation of why the baryon abundance is within an order of magnitude of the thermal relic dark matter density, ruling out simple co-production of baryons and dark matter. Here we have considered an anthropic selection mechanism based on superhorizon-sized domains of varying baryon density. We have discussed a general framework, {\\it anthropic baryogenesis}, in which the domains are generated by variations of a complex scalar field $\\Phi$ and anthropic selection is assumed to disfavour domains in which galaxies have a baryon density which is much larger than in our domain. Baryon isocurvature perturbations impose strong constraints on anthropic baryogenesis. In the case where the $\\Phi$ mass during inflation satisfies $|m_{\\Phi}| \\lae H_{I}$, either an inflation model with an unusually small expansion rate during inflation, $H_{I} < 10^{7} \\GeV$, or a high-order suppression of Planck-suppressed terms in the $\\Phi$ potential is necessary to suppress the baryon isocurvature perturbation. The need to suppress non-renormalizable terms to a high order rules out models with only simple symmetries. This case is relevant to SUSY models, since in that case the mass-squared terms are at most of order $H^2$. Alternatively, an unsuppressed $\\Phi$ potential is possible if the symmetry-breaking mass term $\\mu$ in the $\\Phi$ potential is sufficiently large, $\\mu \\gae 10^{16} \\GeV$. The necessary suppression of potential terms is natural in the case of Affleck-Dine baryogenesis, where the combination of the SM gauge symmetry, SUSY and R-parity provides a sufficiently complex symmetry to suppress the non-renormalizable terms to a high order. We have considered Affleck-Dine baryogenesis for the case of a $d = 6$ $(u^{c}d^{c}d^{c})^2$ flat-direction in the context of F-term hybrid inflation. With inflaton superpotential coupling $\\kappa = 0.005$, the value of the CP-violating phase in our domain must be in the range $0.001 \\lae \\theta \\lae 0.03$, where the lower bound is the isocurvature constraint and the upper bound is the value below which the baryon density can be considered to be anthropically selected. The existence of this range allows $\\theta$ in our domain to be small enough for anthropic selection to function but large enough to evade large baryon isocurvature perturbations. Since $\\theta$ in our domain, which is determined anthropically, can take any value within this range, it is possible that baryon isocurvature perturbations will be large enough to be observed in the future. In our model we have considered all the parameters of the Universe to be fixed to their observed values except the baryon density. In particular, we have considered the dark matter density to be fixed and equal to its value in the observed Universe. This raises an important issue for the class of anthropic selection model considered here. The underlying assumption is that there is a critical baryon density above which life is anthropically disfavoured. The baryon density in a domain will take the largest value possible up to anthropic selection effects, therefore the baryon density will be close to this critical density. It is therefore assumed that this critical baryon density is close to the observed baryon density. But this does not explain why the observed dark matter density, which is assumed to be a fixed parameter, is also close to the critical baryon density. (A similar problem arises in the model of \\cite{lindeax}, where the baryon number is assumed to be fixed and the axion dark matter density varies between domains. In this case it is not explained why the fixed baryon density is close to the critical density.) In order to achieve a complete solution, it may be necessary for both the baryon and dark matter densities to vary independently between domains. In this case the baryon density in a domain will have the highest probability when it is close to the critical density. The dark matter density will then have the highest probability when it is close to baryon density, since there will be a rapid increase in the baryon and dark matter densities in galaxies once $\\Omega_{DM} > \\Omega_{B}$ \\cite{lindeax}. In the case of thermal relic WIMP dark matter, this suggests that a domain-dependent Higgs expectation value and so domain-dependent weak scale is necessary. Alternatively, axion dark matter combined with an anthropic baryogenesis model, such as Affleck-Dine baryogenesis, could provide the basis for such a model. We will return to this possibility in future work." }, "1207/1207.0306_arXiv.txt": { "abstract": "A northern subsample of 89 {\\it Spitzer GLIMPSE} extended green objects (EGOs), the candidate massive young stellar objects, are surveyed for molecular lines in two 1-GHz ranges: 251.5-252.5 and 260.188-261.188\\,GHz. A comprehensive catalog of observed molecular line data and spectral plots are presented. Eight molecular species are undoubtedly detected: H$^{13}$CO$^+$, SiO, SO, CH$_3$OH, CH$_3$OCH$_3$, CH$_3$CH$_2$CN, HCOOCH$_3$, and HN$^{13}$C. H$^{13}$CO$^+$\\,3-2 line is detected in 70 EGOs among which 37 ones also show SiO\\,6-5 line, demonstrating their association to dense gas and supporting the outflow interpretation of the extended $4.5\\,\\mu$m excess emission. Our major dense gas and outflow tracers (H$^{13}$CO$^+$, SiO, SO and CH$_3$OH) are combined with our previous survey of $^{13}$CO, $^{12}$CO and C$^{18}$O\\,1-0 toward the same sample of EGOs for a multi-line multi-cloud analysis of line width and luminosity correlations. Good log-linear correlations are found among all considered line luminosities, which requires a universal similarity of density and thermal structures and probably of shock properties among all EGO clouds to explain. It also requires that the shocks should be produced within the natal clouds of the EGOs. Diverse degrees of correlation are found among the line widths. However, both the line width and luminosity correlations tend to progressively worsen across larger cloud subcomponent size-scales, depicting the increase of randomness across cloud subcomponent sizes. Moreover, the line width correlations among the three isotopic CO\\,1-0 lines show data scatter as linear functions of the line width itself, indicating that the velocity randomness also increases with whole-cloud sizes and has some regularity behind. ", "introduction": "Formation and evolution of massive young stars is still a matter of debate. As reviewed by \\citet{zinn07}, the high luminosity, high protostar temperature, and much shorter formation time scales make the formation of massive stars distinct from that of low mass stars. Particularly, many observed massive protostars reside in proto-clusters, which further complicates their formation processes by introducing possible interplay between neighboring forming massive stars and their lower mass siblings. The paucity of known massive star forming regions is one of the major obstacles in observational studies. Thus, the recently uncovered new sample of over 300 massive young stellar object (MYSO) candidates -- Extended Green Objects (EGOs) -- that are identified from the {\\it Spitzer} Galactic Legacy Infrared Mid-Plane Survey Extraordinaire ({\\it GLIMPSE}) at 3.6, 4.5, 5.8, 8.0, and {\\it MIPSGAL} at 24\\,$\\mu$m, have significantly extended the working sample for massive star formation studies \\citep{cyga08}. The EGOs are so named because they show excess emission in extended structures in the Infrared Array Camera ({\\it IRAC}) $4.5\\,\\mu$m band images that are conventionally coded as green in the {\\it IRAC} false color images. EGOs are candidate birth places of MYSOs because, as shown by \\citet{cyga08}, many of them are associated with Class II CH$_3$OH masers and/or infrared dark clouds (IRDCs). Both the Class II CH$_3$OH masers \\citep{crag05,elli06} and the IRDCs \\citep{simo06a,simo06b,rath06,rath07} are signposts of massive star formation sites. A more detailed interferometric thermal millimeter line mapping of two EGOs by \\citet{cyga11} has lent further support to EGOs being MYSOs or massive protoclusters. EGOs very possibly possess prominent outflows, because the characteristic extended green structures in them are possibly shock originated \\citep{nori04,smit05,debu10}. Studies have confirmed that it is the $4.5\\,\\mu$m excess instead of the $4.5\\,\\mu$m total flux that traces the outflow shocks \\citep[e.g.,][]{simp12,lee12}. The outflow nature of the EGO samples is further confirmed by the high detection rate of SiO\\,5-4 line in a selected subsample of 10 EGOs \\citep{cyga09}. Furthermore, the interferometric mapping of two EGOs by \\citet{cyga11} has also revealed their outflow nature. These EGOs are possibly still in early cloud collapse stage, because their {\\it Spitzer} [3.6]-[5.8] and [8.0]-[24] colors mimic that of the youngest massive star formation models that are still in early cloud collapse phase \\citep{cyga08}. Furthermore, infall tracers such as HCO$^{+}\\,1-0$ have been surveyed by \\citet{chen10} toward a subsample of 88 EGOs with a single dish telescope in the northern hemisphere. They found more blue skewed HCO$^{+}\\,1-0$ line profiles than red skewed ones, which statistically supports the existence of infall motions in these EGOs. The early stage of star formation in EGOs is also confirmed by the non-detection of continuum emission at 3.6 or 1.3\\,cm (before the onset of UC\\,H\\,II region) by \\citet{cyga11b}. The existing methanol (CH$_3$OH) maser data have provided rich information to diagnose the physical conditions in EGOs. Many EGOs associate with both Class\\,I and II methanol masers. On the basis of methanol maser surveys in literature, \\citet{chen09} found that about two third of the observed EGOs are spatially coincident with known Class\\,I methanol masers at 44 and/or 95\\,GHz. The high detection rate of Class\\,I methanol masers has been confirmed by the 55 per cent detection rates of the 95\\,GHz methanol maser in a recent Mopra survey of almost all EGOs in the \\citet{cyga08} catalog by the same group \\citep{chen11}. Thus, more than half of the EGOs should be undoubtedly associated with active outflow shocks. The VLA mapping of both classes of methanol masers by \\citet{cyga09} illustrated that the two classes of masers actually occur in different spatial components of the objects: Class\\,II 6.7\\,GHz maser spots concentrate on the peak of the $24\\,\\mu$m emission sources, tracing the warm star forming cores, whilst the Class\\,I 44\\,GHz masers usually scatter around the extended $4.5\\,\\mu$m emission structures, tracing the interface between the proposed outflows and interstellar medium. The coexistence of both classes of methanol masers in some EGOs may complicate the interpretation of methanol thermal lines observed by single dishes, because part of the detected emission could be mainly excited by strong IR radiation in the hot star-forming cores while the rest part could be contributed by collisional excitation in the hot and dense shock regions around the outflows. Up to now, these new objects are still lack of large scale survey of thermal molecular lines beside our previous survey of the three isotopic CO\\,1-0 and HCO$^+$\\,1-0 lines \\citep{chen10}. To further explore the nature of the clouds around EGOs (EGO clouds hereafter), we have performed a survey of outflow tracers, SiO\\,6-5 and CH$_3$OH\\,J$_3$-J$_2$\\,A line series, toward a sample of 89 EGOs mainly in the northern sky, nearly the same sample as in \\citet{chen10}. Our observations also simultaneously cover some rotational transitions of H$^{13}$CO$^+$, SO, and other interesting complex species such as CH$_3$OCH$_3$, CH$_3$CH$_2$CN and HCOOCH$_3$. In this paper, we mainly present the spectral plots, line parameters and discuss the observed line width and luminosity correlations among the considered molecules, while more detailed analysis of the results will be presented in future papers. Sects.~\\ref{sample} and \\ref{obs} of this paper present target selection and observations. Some notes on line identification are given in Sect.~\\ref{lineiden}. We present the results in Sect.~\\ref{results} and also discuss the dense gas, outflows, distance and line width and luminosity correlations in Sect.~\\ref{discuss}. In the end, the main points are summarized in Sect.~\\ref{summary}. \\section[]{Sample selection} \\label{sample} We have selected 89 EGOs with DEC $>-38\\,\\deg$ from the over 300 EGOs in \\citet{cyga08} as our working sample. The infrared positions from the same paper are used. It is nearly identical to the sample of the isotopic CO\\,1-0 observations by \\citet{chen10}. During the object selection, we found five pairs of EGOs that are so close to each other that our telescope beam (29$\\arcsec$) can hardly get them resolved. Thus, only one source in each pair was selected and the observed lines should be deemed as contributed by both EGOs. The five dropped sources were \\object[EGO G24.11-0.18]{G24.11-0.18}, \\object[EGO G43.04-0.45(b)]{G43.04-0.45(b)}, \\object[EGO G54.45+1.02]{G54.45+1.02}, \\object[EGO G54.11-0.05]{G54.11-0.05} \\footnote{Later we recognized that this EGO is actually far enough away from G54.11-0.04 so that it should be included in our sample. We realized this only after the observations were done.} and \\object[EGO G58.78+0.65]{G58.78+0.65}. In addition, two EGOs, \\object[EGO G19.01-0.03 O-N]{G19.01-0.03\\,O-N} and \\object[EGO G19.01-0.03 O-S]{G19.01-0.03\\,O-S}, were also excluded from our sample, because they were covered by the beam toward their associated point source, G19.01-0.03. \\section[]{Observations and data reduction} \\label{obs} The observations were done with the Arizona Radio Observatory 10-m Submillimeter Telescope (AROSMT) in several nights during two epochs: 2009 April 24 to May 11 and 2010 March 29 to April 27. The ALMA band-6 sideband-separating receiver was used with two Filter Banks to record the upper and lower sidebands (USB and LSB) separately. The observed frequency ranges were 251.5-252.5\\,GHz (USB) and 260.188-261.188\\,GHz (LSB). The LRS velocities of most of our EGOs from the C$^{18}$O\\,1-0 observations of \\citet{chen10} had been used for tuning. The spectral resolution was 1\\,MHz in the 1\\,GHz bandwidth in both sidebands. We tested several off positions that were selected on the basis of the CO\\,1-0 survey result of \\citet{dame01} at locations several arc minutes to degrees away from some targets but found no line emission. However, our test runs in position-switch mode were found to suffer from strong ripples in baseline. Thus, we adopted beam switch mode with a 2-arcmin throw at 1\\,Hz for all our targets to improve the spectral baseline. This was viable perhaps because our observed lines are mostly high density tracers. This strategy was justified by the fact that no obvious absorption like features were seen in the final data. Although the existence of weak emission at the off-positions can not be excluded, their effects to our data should be hopefully small. The telescope pointing and focus was checked roughly every two hours with a planet. The representative beam width is $29\\arcsec$ (at 260\\,GHz), which is corresponding to a linear resolution of roughly 0.56\\,pc at a representative distance of 4\\,kpc. Main beam efficiencies of 0.65 and 0.55 were applied to the LSB and USB, respectively, to obtain the main beam temperature $T_{\\rm mb}$ from the antenna temperature $T_{\\rm A}^{\\rm *}$. Line flux can be computed from $T_{\\rm mb}$ with the nominal conversion factor of 46.5\\,Jy/K (at 260\\,GHz) for the ARO SMT. The typical T$_{\\rm sys}$ during most of the observations were $\\sim\\,280$\\,K in the LSB and $\\sim\\,311$\\,K in the USB. Sideband leakage can produce false line features to make trouble to line identification. Fortunately, it can be checked in our data by comparing the USB and LSB data which are obtained simultaneously and separately. We did not find leakage in most of our data, except for the cases of \\object[EGO G14.33-0.64]{G14.33-0.64} and \\object[EGO G24.33+0.14]{G24.33+0.14} of which the strong SO line in the LSB leaked into the USB as a small artificial line feature. The rareness of the leakage is due to the image rejection ratio of about 10 to 20\\,dB during most of our observations. Gildas/CLASS software\\footnote{\\url{http://www.iram.fr/IRAMFR/GILDAS}} was used to reduce the data and analyze the line profiles. A linear baseline was removed from almost all spectra, except the LSB spectrum of \\object[EGO G11.92-0.61]{G11.92-0.61} for which a cubic baseline was subtracted to correct for the gently curvy baseline. With an on+off exposure time of 10 minutes, the average baseline RMS (in mainbeam temperature) are $\\sim\\,25$\\,mK in the LSB and $\\sim\\,33$\\,mK in the USB at a resolution of 1\\,km\\,s$^{-1}$. The rms of individual objects are listed in Table~\\ref{tab2} (together with their LRS velocity and distance that will be discussed in later sections). ", "conclusions": "" }, "1207/1207.4486.txt": { "abstract": "We present an X-ray stacking analysis of a sample of 38 submillimeter galaxies with $\\left< z \\right>=2.6$ discovered at $\\geq 4\\sigma$ significance in the Lockman Hole North with the MAMBO array. We find a $5\\sigma$ detection in the stacked soft band (0.5--$2.0\\,\\rm keV$) image, and no significant detection in the hard band (2.0--$8\\,\\rm keV$). We also perform rest-frame spectral stacking based on spectroscopic and photometric redshifts and find a $\\sim 4\\sigma$ detection of $\\rm Fe\\,K\\alpha$ emission with an equivalent width of $\\rm EW\\gtrsim1\\,\\rm keV$. The centroid of the $\\rm Fe\\,K\\alpha$ emission lies near $6.7\\,\\rm keV$, indicating a possible contribution from highly ionized Fe\\,XXV or Fe\\,XXVI; there is also a slight indication that the line emission is more spatially extended than the X-ray continuum. This is the first X-ray analysis of a complete, flux-limited sample of SMGs with statistically robust radio counterparts. ", "introduction": "Submillimeter galaxies (SMGs) are distant star-forming systems with tremendous infrared luminosities ($L_{\\rm IR}$ [8--1000$\\,\\mu{\\rm m}] \\gtrsim 10^{12}\\, L_{\\odot}$). In the (sub)millimeter waveband they are observable out to high redshifts due to the strong negative $K$-correction in the Rayleigh-Jeans regime of their thermal spectrum \\citep[see, e.g., ][]{blai02}. The prevalence of SMGs at $z> 1$ \\citep{chap05} in combination with their high rates of dust-obscured star-formation imply that they may be responsible for the production of a significant fraction of all the stellar mass in present-day galaxies. X-ray \\citep{alex03,alex05b} and mid-infrared \\citep{vali07, mene07, mene09, pope08} spectroscopy shows that SMGs frequently contain active galactic nuclei (AGN) as well as powerful starbursts. This connection between star formation and accretion at high redshift may help explain the black hole mass-bulge mass relation in present-day galaxies \\citep[e.g., ][]{alex05a}. However, it remains hard to determine the relative importance of accretion and star formation for the SMG population as a whole because of the challenge of assembling large, statistically unbiased SMG samples. Studying the X-ray properties of SMGs is difficult for two main reasons. First, the X-ray counterparts to SMGs are extremely faint. The count rate is so low that even the deepest $Chandra$ and $XMM-Newton$ spectra of SMGs cannot resolve features that serve as sensitive diagnostics of the physical conditions inside galaxies, like the Fe\\,K$\\alpha$ emission line. Fe\\,K$\\alpha$ emission is a ubiquitous feature in spectra of optically-selected AGN up to $z \\simeq 3$ \\citep[e.g., ][]{brus05, chau10, iwas11}, but the Fe\\,K$\\alpha$ emission properties of SMGs remain considerably more uncertain \\citep{alex05b}. Second, the requirement that SMGs need radio, (sub)millimeter, or mid-IR counterparts capable of nailing down their positions in high-resolution X-ray maps can lead to concessions of inhomogeneously-selected samples \\citep[e.g., including radio-selected galaxies; ][]{alex05b}, yielding results that conflict with X-ray studies of purely submillimeter-selected SMG samples \\citep{lair10, geor11}. To disentangle the relationship between SMGs and X-ray selected AGN, we need to overcome the uncertainty introduced by inhomogeneously selected samples, requiring X-ray spectral analyses of large, flux-limited samples of (sub)millimeter-selected SMGs with robust counterparts. In this work, we report on an X-ray stacking analysis of a sample of 38 SMGs detected in a $1.2\\,\\rm mm$ map of the Lockman Hole North (LHN), one of the fields in the $Spitzer$ Wide-Area Infrared Extragalactic (SWIRE) Survey \\citep{lons03}, using data from the $Chandra$-SWIRE survey \\citep{poll06,wilk09}. The high radio counterpart identification rate of the LHN SMG sample \\citep[93\\%; ][]{lind11} is afforded by the extremely deep $20\\,\\rm cm$ map of the same field \\citep{owen09}, and allows for reliable X-ray photometry. The sample benefits from spectroscopic \\citep{poll06, owen09, fiol10} and optically-derived photometric \\citep{stra10} redshifts. Additionally, analyses of $Herschel$ observations of the LHN \\citep{magd10, rose12} have delivered reliable photometric redshifts and infrared luminosities for a large fraction of the sample by fitting far-IR photometry with thermal-dust spectral energy distribution (SED) models. In \\S 2, we describe the observations used in our analysis. \\S 3 outlines our X-ray stacking technique, and our method for deriving rest-frame luminosities. In \\S 4, we compare our results to previous X-ray studies of SMGs, and discuss the possible origins of the Fe\\,K$\\alpha$ emission seen in our stacked spectrum. In \\S 5, we present our conclusions. In our calculations, we assume a $WMAP$ cosmology with $H_0=70\\,\\rm km\\,s^{-1}\\,Mpc^{-1}$, $\\Omega_M=0.27$, and $\\Omega_\\lambda=0.73$ \\citep{koma11}. ", "conclusions": "We analyze the X-ray properties of a complete sample of SMGs with radio counterparts from the LHN. This sample's X-ray detection rate of $2^{+6}_{-2}\\%$ is consistent with those for other uniformly-mapped, submillimeter-detected samples, considering the depth of our X-ray data. The X-ray undetected SMGs show a strong stacked detection in the $S_C$ band, and no significant detection in the $H_C$ band, similar to results from SMG stacking in the CDF-N \\citep{lair10} and CDF-S \\citep{geor11}. We also use the available redshift information of our SMGs to compute the rest-frame, stacked count-rate spectrum of our sample. The rest-frame spectrum shows strong ($ \\rm EW> 1\\,\\rm keV$) emission from Fe\\,K$\\alpha$, possibly with contributions from Fe\\,XXV and Fe\\,XXVI. A comparison with other high-ionization Fe\\,K$\\alpha$-emitting systems from the literature indicates that accretion onto obscured AGNs is the likely explanation for the strong Fe\\,K$\\alpha$ emission line. In our sample, the Fe\\,K$\\alpha$ emission is responsible for $\\sim 20\\%$ of the observed soft-band X-ray flux. Therefore, if strong Fe line emission is a common feature in other SMG samples, it would significantly decrease the measured values of HR and lead to overestimates of the continuum spectral index $\\Gamma$. We find a tentative indication (71\\% confidence) that our sample's stacked distribution of Fe\\,K$\\alpha$ photons is more spatially extended than that of the X-ray continuum. If confirmed by future studies, this result can help determine the physical origin of the prominent Fe\\,K$\\alpha$ emission in SMGs." }, "1207/1207.4438_arXiv.txt": { "abstract": "We introduce the Marenostrum-MultiDark SImulations of galaxy Clusters (MUSIC) dataset. It constitutes one of the largest sample of hydrodynamically simulated galaxy clusters with more than 500 clusters and 2000 groups. The objects have been selected from two large N-body simulations and have been resimulated at high resolution using Smoothed Particle Hydrodynamics (SPH) together with relevant physical processes that include cooling, UV photoionization, star formation and different feedback processes associated to Supernovae explosions. In this first paper we focus on the analysis of the baryon content (gas and star) of clusters in the MUSIC dataset both as a function of aperture radius and redshift. The results from our simulations are compared with a compilation of the most recent observational estimates of the gas fraction in galaxy clusters at different overdensity radii. We confirm, as in previous simulations, that the gas fraction is overestimated if radiative physics is not properly taken into account. On the other hand, when the effects of cooling and stellar feedbacks are included, the MUSIC clusters show a good agreement with the most recent observed gas fractions quoted in the literature. A clear dependence of the gas fractions with the total cluster mass is also evident. However, we do not find a significant evolution with redshift of the gas fractions at aperture radius corresponding to overdensities smaller than 1500 with respect to critical density. At smaller radii, the gas fraction do exhibit a decrease with redshift that is related the gas depletion due to star formation in the central region of the clusters. The impact of the aperture radius choice, when comparing integrated quantities at different redshifts, is tested. The standard, widely used definition of radius at a fixed overdensity with respect to critical density is compared with a definition of aperture radius based on the redshift dependent overdensity with respect to background matter density: we show that the latter definition is more successful in probing the same fraction of the virial radius at different redshifts, providing a more reliable derivation of the time evolution of integrated quantities. We also present in this paper a detailed analysis of the scaling relations of the thermal SZ (Sunyaev Zel'dovich) Effect derived from MUSIC clusters. The integrated SZ brightness, $Y$, is related to the cluster total mass, $M$, as well as, the $M-Y$ counterpart which is more suitable for observational applications. Both laws are consistent with predictions from the self-similar model, showing a very low scatter which is $\\sigma_{\\log Y}$ $\\simeq$ 0.04 and even a smaller one ($\\sigma_{\\log M}$ $\\simeq$ 0.03) for the inverse $M-Y$ relation. The effects of the gas fraction on the $Y-M$ scaling relation is also studied. At high overdensities, the dispersion of the gas fractions introduces non negligible deviation from self-similarity, which is directly related to the $f_{gas}-M$ relation. The presence of a possible redshift dependence on the $Y-M$ scaling relation is also explored. No significant evolution of the SZ relations is found at lower overdensities, regardless of the definition of overdensity used. ", "introduction": "\\label{sec:introduction} Galaxy clusters are the biggest gravitationally bound objects of the Universe and constitute one of the best cosmological probes to measure the total matter content of the Universe. However, the total mass of these objects cannot be directly measured. It must be inferred from other observational quantities (X-ray or SZ Surface Brightness, Lensing distortions or number of galaxies). In all cases, one has to relate these quantities with the total mass of the system. Due to the complex physics involved in the processes of cluster formation, hydrodynamical numerical simulations have been a fundamental tool to calibrate mass proxies, define new ones, and to study the systematics affecting observational measurements. They are also indispensable to deeply study the formation and evolution of clusters of galaxies and all their gas-dynamical effects \\citep{BK09}. The big progresses achieved in the last years by numerical simulations are well represented by their use to describe and to study X-ray temperatures and their relation with cluster gas mass (\\citealt{ETTORI04}; \\citealt{MUA06}; \\citealt{NKV07}) as well as to compare numerical predictions with observed temperatures (\\citealt{LOKEN02}; \\citealt{BORG04}; \\citealt{LECCA08}) , gas profiles (\\citealt{RONCA06}; \\citealt{CROS08}), or pressure profiles \\citep{ARNAUD10}. \\textbf{A detailed review on cluster simulations can be found in \\cite{BK09} and \\cite{KB12}.} In an ideal scenario one would need to have a large sample of simulated galaxy clusters with enough numerical resolution (both in mass and in the gravity and pressure forces) to accurately resolve the internal substructures and with a detailed modelling of the most relevant physical processes. The best way to achieve this goal would be by simulating large cosmological boxes (\\citealt{BORG04}; \\citealt{MN07}; \\citealt{BURNS08}; \\citealt{HAR08}; \\citealt{BOY08}). Unfortunately, due to the large computational demand of these simulations, one needs to find a compromise between the three main ingredients: volume size, mass resolution and physical processes included . A possible solution to the computational problems related with scalability of the present-day hydro codes is to proceed, mimicking the observations, by creating a catalogue of resimulated galaxy clusters that are extracted from low resolution N-body simulations. The regions containing clusters of galaxies are then resimulated with very high resolution, adding only gas physics in the resimulated areas (\\citealt{PUCH08};\\citealt{DOLAG09}; \\citealt{LAU09}; \\citealt{FABJAN10}). This so-called 'zooming' technique permits to simulate thousands of clusters basically independently from each other with less computational cost than a full box hydrodynamical simulation of the same resolution. By selecting all the objects formed in a given volume above a given mass threshold, mock volume limited sample catalogues can be generated and used in the study of the properties and interrelations of the different scaling laws of galaxy clusters. Following this procedure, we have generated the MUSIC (Marenostrum-MultiDark SImulation of galaxy Clusters) dataset, a large sample of simulated clusters of galaxies, composed of objects extracted from two large box cosmological simulations: the MareNostrum Universe and the MultiDark simulation. We selected all the clusters using criteria based on mass (selecting all the clusters having a total mass larger than 10$^{15}$ \\hMsun) or on morphology (selecting groups of clusters corresponding to different morphology classes, bullet-like clusters and relaxed clusters). All these objects have been resimulated with SPH particles, radiative physical processes and star formation prescriptions, improving by an order of magnitude the resolution with respect to the original simulations, as described in the next section. From this database, we will obtain mock observations for X-rays, SZ, lensing as well as optical galaxy counts. This will allow us to study the interrelations between the scaling laws associated to the different observables. In this paper we base our attention on the properties of the baryon content and the SZ effects and will leave a more detailed analyzes of the relations with the other observed properties for a further work. In the last few decades the SZ effect \\citep{SZ70} has become one of the most powerful cosmological tool to study clusters of galaxies, as well as the nature of the dark matter and dark energy components of the Universe. The physical process of the SZ effect is the diffusion of CMB (Cosmic Microwave Background) photons with a hot plasma due to inverse Compton scattering. The thermal component of the SZ effect is largely enhanced by the presence of clusters of galaxies, the most massive bound objects in the Universe, where plasma is in hydrostatic equilibrium inside the gravitational potentials of dark matter. The Intra Cluster Medium (ICM) composed by high energy electrons constitutes an ideal laboratory to investigate the SZ effect. The brightness of the SZ effect turns out to be independent of the diffuser position, thus making it the best tool to find galaxy clusters at high redshift. Moreover, the SZ flux collected from the cluster region is proportional to the total thermal energy content, with a weak dependence on the complex physical processes acting at the inner regions (e.g. cooling flows, galaxy feedbacks etc) which mostly affect the X-ray luminosity. Therefore, these two properties: redshift independence and low scatter mass proxy makes the integrated Compton $Y$-parameter an efficient high-$z$ mass-estimator. Under the hypothesis that the evolution of galaxy clusters is driven mainly by gravitational processes \\citep{KAI1986} and assuming hydrostatic equilibrium and an isothermal distribution of dark matter and ICM \\citep{BN1998} it is possible to derive simple scaling power-laws linking cluster properties: the so-called self-similar scaling relations. In the case of SZ science, the relation linking the SZ brightness with the cluster total mass, the $Y-M$ scaling law, is continuously under analysis to test for its robustness, allowing the application of the SZ effect as a mass-finder. Observational studies started to collect data of a few clusters, mostly those with high X-ray luminosities (and therefore high masses) (\\citealt{BENS04}; \\citealt{MORANDI07}; \\citealt{BONA2008}; \\citealt{VI09}; \\citealt{ARNAUD10}). Recent large surveys have shown the possibility of detecting undiscovered clusters only through SZ effect observations, as in the claims by South Pole Telescope (SPT, \\citealt{SPT2009}; \\citealt{SPT2010}; \\citep{SPT2011A}; \\citealt{SPT2011B}), Atacama Cosmology Telescope (ACT, \\citealt{ACT2011A}; \\citealt{ACT2011B}) and Planck (\\citealt{PLANCKa}; \\citealt{PLANCKb}; \\citealt{PLANCKc}; \\citealt{PLANCKd}). The possibilities to explore more distant objects or to deeply map single cluster morphology are planned with the ongoing higher angular resolution projects like AMI \\citep{AMI08}, AMiBA \\citep{AMIBA}, MUSTANG \\citep{MUSTANG}, OCRA \\citep{OCRA}, CARMA \\citep{CARMA} or the in-coming projects like the ground based C-CAT \\citep{CCAT}, the upgraded with new spectroscopic capabilities MITO \\citep{MITO} and the balloon-borne OLIMPO \\citep{OLIMPO} or the proposed satellite mission {\\it Millimetron} ({\\ttfamily http:// www.sron.rug.nl/millimetron}). More recent blind-surveys carried out by SPT, ACT and Planck enlarged the existing dataset confirming the self-similarity in the sample at least for the massive clusters (\\citealt{ACT2011A}; \\citealt{AND11}; \\citealt{PLANCKc}). Numerical simulations have shown that $Y$ (the integral of the Compton $y$-parameter over the solid angle of the cluster) is a good mass proxy (\\citealt{DASILVA04}; \\citealt{MOTL05}; \\citealt{AGH09}) and that the slope and the evolution of SZ scaling relations are apparently not affected by redshift evolution and cluster physics. In order to find an X-ray equivalent of the SZ integrated flux, the numerical simulations also led to the introduction of a new mass proxy: the $Y_X$ parameter, defined as the product of the cluster gas mass and its tempearture \\citep{YX06}. Many works based on large $N$-body cosmological simulations have already studied the impact of gas physics on SZ scaling relations (\\citealt{DASILVA00}; \\citealt{HS02}; \\citealt{DASILVA04}; \\citealt{MOTL05}; \\citealt{NAGAISZ}; \\citealt{BONALDI07}; \\citealt{HALL07}; \\citealt{AGH09}; \\citealt{BATTAGLIA11}; \\citealt{KAY12}), such as the effects introduced by clusters with disturbed morphologies (\\citealt{POOL07}; \\citealt{WIK08}; \\citealt{YANG10}; \\citealt{KRAUSE12}), showing that the self-similar model is valid at least up to cluster scales (even if with some differences introduced by the different models used in the simulations). In what follows, we will present an analysis of the relation between the SZ integrated flux and the total mass, $Y$-$M$, in MUSIC clusters and will compare it with the predictions of the self-similar model. We also study the possible biases introduced by the common assumption of considering quantities ($M$ or $Y$) integrated inside a radius defined by a fixed overdensity with respect to the critical density instead of a more suitable definition with respect to the background density whose value depends on redshift. We will focus our analysis of the $Y$-$M$ scaling relation on the most massive objects of the MUSIC dataset which constitutes an almost complete volume limited sample. Therefore, only clusters with virial masses $M_v>$5$\\times$10$^{14}$\\hMsun, are considered in this paper. We will extend our analysis to a wider range of masses in an upcoming work. The paper is organized as follows. In Section 2 the MUSIC database is described. The baryon content in the clusters at different overdensities is presented in Section 3, together with a study of its evolution with redshift and a comparison of our numerical results with the most recent observational estimates. In Section 4 the $Y$-$M$ scaling relation is computed from the MUSIC clusters and compared with the predictions from the self-similar model. We also discuss the validity of the integrated $Y$ as a proxy for the cluster mass. In Section 5 we focus on the impact of the gas fraction on the $Y$-$M$ relation together with the dependence of the gas fraction on the cluster mass. Finally in Section 6 the redshift evolution of the $Y$-$M$ relation is discussed. In Section 7 we summarize and discuss our main results. ", "conclusions": "\\label{sec:conclusions} Recent new large scale surveys on the far infrared and millimeter bands confirmed the validity of the Sunyaev-Zel'dovich Effect as a useful cosmological tool to detect and to study galaxy clusters. Today, one of its most relevant applications is to infer the total cluster mass using the $Y-M$ scaling relation. Recent observations (Planck, SPT and ACT) provide the SZ-measurements of a large number of clusters from which the $Y-M$ relation can be reliably inferred. In this paper we have presented the MUSIC dataset, a large sample of gasdynamical resimulated massive clusters, that have been selected from the MareNostrum Universe simulation and from the MultiDark simulation. Our dataset comprises two different subsamples. The MUSIC-1 which contains 164 objects that have been selected according to their dynamical state (i.e. bullet-like clusters and relaxed clusters) and the MUSIC-2 which contains a total of more than 2000 galaxy groups and clusters. The MUSIC-2 constitutes a complete volume limited sample for clusters with virial masses above $8\\times 10^{14} h^{-1} M_\\odot$ at $z=0$. To our knowledge, there are no other works about resimulated clusters that have shown such a large number of massive objects. These objects have been simulated using two physical processes: non-radiative (NR, gravitational heating only) and a radiative model (CSF) including cooling, UV photoionisation, multiphase ISM and star formation and supernovae feedbacks. No AGN feedback has been included in the radiative simulations. We plan to generate mock catalogues from the MUSIC clusters for the most common observables such as: X-ray and SZE surface brightness maps, lensing maps and galaxy luminosity functions. All these deliverables, together with the initial conditions for all the MUSIC simulations, will be made publicly available from the website http://music.ft.uam.es. Here we have presented the first results from MUSIC based on the analysis of the baryon content and the scaling relations of the thermal SZ effect. The analysis of the integrated cluster quantities, such as the fractions of the different baryon components, or $Y$ parameter, have been done using two definitions of the aperture radius. The most commonly used approach assumes that the integration domain is defined by a fixed value of the mean overdensity inside the cluster region compared to the critical overdensity of the Universe. It has been shown (see Appendix A) that this definition introduces a bias when any integrated quantity is compared at different redshifts. On the contrary, by using a redshift-dependent overdensity compared to the background density of the Universe we can rescale the aperture radius such that the integrated region is always the same volume fraction of the total cluster volume defined by the virial radius. Our numerical results on the baryon content and the SZ scaling relations have been explored to check possible dependences on overdensity, redshift and cluster physics modelling. We have also made a comparison between MUSIC clusters and available observational results. The scaling relation between SZ brightness, quantified with the integrated Compton $y$-parameter, $Y$, and the cluster total mass, $M$, confirmed the robustness of the self-similarity assumption for this class of objects present in MUSIC sample. We evaluated the effect on the $Y-M$ of having the cluster gas fraction, $f_{gas}$, as a free parameter that can vary with the cluster mass. Finally, we looked for a possible redshift evolution of the $Y-M$ relation using two different approaches. Our main conclusions can be summarized as follows: \\begin{itemize} \\item The mean value of the gas fraction of MUSIC-2 CSF clusters at $\\Delta_c$=500 is $f_{gas}$=0.118$\\pm$0.005. This value is compatible with observational results shown in \\cite{MGAN04}, \\cite{LAROQUE06}, \\cite{MGAN06}, \\cite{VI09}, \\cite{ZHANG10}, \\cite{ET10}, \\cite{JU10}, and \\cite{DJF12}. The values of the baryon fraction at virial radius, for both simulation flavours (CSF and NR), are consistent, within the errors, with the cosmological ratio $\\Omega_b/\\Omega_m$ (according to WMAP7 cosmology). At higher overdensities (i.e. $\\Delta_c$=2500), we also find agreement for the $f_{gas}$ with observational results from \\cite{MGAN06}, \\cite{VI09} and \\cite{ZHANG10} and we are marginally consistent with \\cite{AL08} and \\cite{DJF12}. This means that the possible effects of numerical overcooling do not appear to affect dramatically our simulations, despite the fact that we have not included strong AGN feedback in our model. \\textbf{The star fraction measured in MUSIC-2 CSF clusters does not agree with the one estimated by observations, which for massive clusters predict a star fraction considerably smaller than one extracted from MUSIC simulations.} Concerning the evolution with redshift of the different baryon components, we find significant differences in the evolution depending on the definition of aperture radius used. This is more evident at large overdensities, where the considered fraction of virial radius increases with redshift when a fixed critical overdensity is used to the define it (see Fig. \\ref{rad}). \\item The $Y-M$ scaling relation derived from our clusters, assuming a fixed $f_{gas}$, agrees very well with the predictions of the self-similar model and shows a very low dispersion ($\\sigma_{logY}$ $\\simeq$ 0.04). The resulting best fit $Y_{500}-M_{500}$ relation can then be expressed as: \\begin{equation} Y_{500} = 10^{-28.3\\pm 0.2}\\bigg(\\frac{M_{500}}{\\hMsun}\\bigg)^{1.66\\pm 0.02}E(z)^{-2/3}[h^{-2}Mpc^{2}] \\end{equation} The $M-Y$ relation, which is more suitable to infer cluster masses from $Y$ measurements, is also compatible with the self-similar model (which predicts a slope of 3/5) with an even lower dispersion ($\\sigma_{logM}\\simeq$0.03). The corresponding best-fit $M_{500}-Y_{500}$ relation can be expressed as: \\begin{equation} M_{500} = 10^{17.0\\pm 0.1}\\bigg(\\frac{Y_{500}}{h^{-2}Mpc^2}\\bigg)^{0.59\\pm 0.01}E(z)^{-2/5}[\\hMsun] \\end{equation} We have compared the above fits with the recent results on the $Y-M$ from the Planck Collaboration. The agreement is very good, provided that the masses of Planck clusters are overestimated by 22 per cent due to the biases of the X-ray mass estimations with respect to the lensing measurements. We have tested whether this bias could be due to the lack of hydrostatic equilibrium hypothesis. By comparing the hydrostatic masses with the true one in our MUSIC samples, we find a negative HMB (an opposite trend than in Planck clusters). In our case, we find 25 per cent underestimation of the true mass by the hydrostatic mass, \\textbf{in agreement with other simulations studying the same topic which already found that the HSE hypothesis introduce a negative bias on the true mass estimation.} \\item The dependence of the gas fraction on the cluster total mass has been also studied along the cluster aperture radius. There is a linear relation between $f_{gas}$ and total mass which is more evident, although with larger scatter, when we approach the cluster core. This effect is rather insensitive on the adopted physics in the numerical simulations: both CSF and NR clusters present a similar behavior. \\item Leaving $f_{gas}$ as a free parameter in the $Y-M$ scaling relation leads to a deviation from self-similarity. The modified $Yf_{gas}^{-1} -M $ relation and the $f_{gas}-M$ relation are directly related with the $Y-M$ if both of them are expressed as power laws. We confirm this hypothesis and provide the fits to the $f_{gas} -M $ for two overdensities $\\Delta_c=2500$ and $\\Delta_c=500$ \\begin{equation} f_{gas}=10^{-3.5\\pm0.2}\\bigg(\\frac{M_{2500}}{\\hMsun}\\bigg)^{0.17\\pm0.01} \\end{equation} \\begin{equation} f_{gas}=10^{-1.6\\pm0.1}\\bigg(\\frac{M_{500}}{\\hMsun}\\bigg)^{0.04\\pm0.01} \\end{equation} \\item We have studied the redshift evolution of the slope, and of the normalization, in the $Y-M$ scaling law. We did not find any evolution for NR clusters. On the contrary, CSF clusters show a marginal deviation of the slope from the self-similar prediction at high overdensity for redshifts larger than 0.5. This result is confirmed using two different methods to study the evolution with redshift. \\item No significant differences have been found when comparing results at different redshifts using the two alternative definitions of aperture areas based on different overdensity criteria: the standard, redshift independent, critical overdensity and the redshift-dependent background overdensity show the same results, at least for the gas fraction analysis, with the latter appearing to be more sensitive to the redshift evolution of the sample properties ($Y-M$ slope, gas fraction, etc.). We can therefore conclude that even if the redshift-dependent background overdensity leads to results which better identify a possible evolution, no significant errors are introduced in the analysis of the SZ effects when using the standard critical overdensity criteria. The same conclusions in the case of X-ray analyzes have been raised by \\cite{MGAN06} . \\item The use of radiative physics on galaxy cluster simulations introduces fundamental improvements respect to non-radiative simulations: CSF clusters show better agreement with observations in the estimate of the gas fraction as well as in the study of the $Y-M$ scaling relation. On the other hand, NR clusters overestimate the gas fraction and do not seem to be compatible with $Y-M$ scaling relations from observational results. \\end{itemize} In summary, we have shown that the MUSIC dataset is well suited for the study of massive cluster properties and provides a reasonable description of observed objects with similar mass. Therefore, the MUSIC clusters can be a good cosmological tool. In upcoming papers we are going to extend the analysis to other complementary scaling relations in X-ray, lensing and optical in order to have a full set of observables for a complete volume limited sample of $\\Lambda$CDM simulated galaxy clusters. \\appendix" }, "1207/1207.5799.txt": { "abstract": "We use the SAURON and GMOS integral field spectrographs to observe the active galactic nucleus (AGN) powered outflow in NGC\\,1266. This unusual galaxy is relatively nearby (D=30 Mpc), allowing us to investigate the process of AGN feedback in action. We present maps of the kinematics and line strengths of the ionised gas emission lines H$\\alpha$, H$\\beta$, [OIII], [OI], [NII] and [SII], and report on the detection of Sodium D absorption. We use these tracers to explore the structure of the source, derive the ionised and atomic gas kinematics and investigate the gas excitation and physical conditions. NGC\\,1266 contains two ionised gas components along most lines of sight, tracing the ongoing outflow and a component closer to the galaxy systemic, the origin of which is unclear. This gas appears to be disturbed by a nascent AGN jet. We confirm that the outflow in NGC\\,1266 is truly multiphase, containing radio plasma, atomic, molecular and ionised gas and X-ray emitting plasma. The outflow has velocities up to $\\pm$900 \\kms\\ away from the systemic velocity, and is very likely to be removing significant amounts of cold gas from the galaxy. The LINER-like line-emission in NGC\\,1266 is extended, and likely arises from fast shocks caused by the interaction of the radio jet with the ISM. These shocks have velocities of up to 800 \\kms, which match well with the observed velocity of the outflow. Sodium D equivalent width profiles are used to set constraints on the size and orientation of the outflow. The ionised gas morphology correlates with the nascent radio jets observed in 1.4 Ghz and 5 Ghz continuum emission, supporting the suggestion that an AGN jet is providing the energy required to drive the outflow. ", "introduction": "In recent years the idea feedback from an active galactic nucleus (AGN; e.g. \\citealt{2005ApJ...620L..79S,2006MNRAS.365...11C}) could be responsible for the quenching of star formation has grown in popularity. Such quenching seems to be required to create the red-sequence galaxies we observe today \\citep[e.g.][]{2004ApJ...600..681B}. There is circumstantial evidence to support AGN-driven quenching, such as the study by \\cite{2007MNRAS.382.1415S} suggesting that AGN are predominantly found in green valley galaxies, but direct evidence for removal/heating of cold star-forming gas is rare. The physical mechanism by which an AGN could drive molecular gas out of a galaxy is still debated. Radiation pressure is thought to be important in star formation-driven outflows (e.g. \\citealt{2005ApJ...618..569M}), and is potentially implicated in AGN-powered `quasar mode' outflows (e.g. \\citealt{1999ApJ...516...27A,2009MNRAS.397.1791K}). Kinetic feedback from an AGN jet can provide sufficient power to directly push through the ISM of a galaxy and entrain or destroy it (e.g. \\citealt{2010ApJ...716..131R,McNul2012}), but it is unclear if the geometry of a bipolar jet, which often emerges perpendicular to the nuclear disk, will allow the jet to remove the ISM from an entire galactic disk. Broad-line winds can deposit significant momentum into gas surrounding an AGN, which could also lead to large outflows (e.g. \\citealt{2010ApJ...722..642O}). Alternatively, heating by X-rays and cosmic rays could destroy/alter the molecular clouds close to an AGN, removing the need to expel them from the galaxy (e.g. \\citealt{1991ApJ...382..416B,Ferland}). These processes should be distinguishable if we can identify and study local galaxies where AGN feedback is ongoing. Our recent Combined Array for Research in Millimetre-wave Astronomy (CARMA) and Sub-Millimetre Array (SMA) observations of the nearby lenticular galaxy NGC\\,1266 suggest that it harbors a massive AGN-driven molecular outflow, providing an excellent local laboratory for studying AGN-driven quenching \\citep[][hearafter A2011]{2011ApJ...735...88A}. {NGC\\,1266} is a nearby (D= 29.9 Mpc; derived from recession velocity in \\citealt{2011MNRAS.413..813C}; hereafter \\atlas\\ Paper I), early-type galaxy (ETG) in the southern sky ($\\delta=-$2$^{\\circ}$), which was studied as part of the \\atlas\\ project. A three colour image of this galaxy (from \\citealt{SINGS}) is presented in Figure \\ref{ngc1266color}. While typical CO spectra from early-type galaxies reveal the double-horned profile characteristic of gas in a relaxed disk with a flat rotation curve, the spectrum of NGC\\,1266 shows a narrow central peak (FWHM $\\approx$120 \\kms) with non-Gaussian wings out to at least $\\pm$400 \\kms\\ with respect to the systemic velocity \\citep{2011MNRAS.414..940Y}. Imaging of the high-velocity components using the SMA revealed that the wings resolve into redshifted and blueshifted lobes (A2011), coincident with H$\\alpha$ emission \\citep{2003PASP..115..928K}, 1.4 GHz continuum \\citep{2006A&A...449..559B}, and thermal bremsstrahlung emission (detected with Chandra; A2011; Fig. 3). Molecular gas observations suggest that 3$\\times$10$^8$ M$_{\\odot}$ of molecular gas is contained within the central 100 pc of NGC\\,1266, and that at least 5$\\times$10$^7$ M$_{\\odot}$ of this gas is involved in a molecular outflow (A2011). This is thus the first observed large-scale expulsion of molecular gas from a non-starbursting ETG in the local universe, and this presents a unique opportunity to study this powerful process in action. In this paper we present SAURON (Spectrographic Areal Unit for Research on Optical Nebulae) and Gemini Multi-Object Spectrograph (GMOS) integral-field unit (IFU) observations of the ionised gas in NGC\\,1266. {By investigating the ionised gas kinematics and line ratios we hope to constrain the outflow parameters and ionisation mechanisms and thus shed light on the mechanism driving gas from the galaxy}. In Section \\ref{dataredux} we present the data, and describe our reduction processes. We then present the derived maps of the gas kinematics and line fluxes. In Section \\ref{results} we discuss the kinematic structure of the system, gas excitation mechanisms and the driving force behind the outflow. Finally we conclude and discuss prospects for the future in Section \\ref{conclude}. \\begin{figure} \\begin{center} \\includegraphics[width=0.48\\textwidth,clip,trim=10cm 3.0cm 11cm 1.5cm]{NGC1266_sings2.eps} \\end{center} \\caption{\\small SINGs \\citep{SINGS} $B$,$V$ and $R$ band composite three colour image of S0 galaxy NGC\\,1266. The white bar shows a linear scale of 1 Kpc (6\\farc94 at an adopted distance of 29.9 Mpc; \\atlas\\ Paper I). Overlaid are the total field of view of our \\sauron\\ IFU (red) and GMOS IFU (blue) observations. } \\label{ngc1266color} \\end{figure} ", "conclusions": "\\label{conclude} In this paper we have been able to shed some light on the outflow activity in NGC\\,1266. This unusual galaxy is relatively nearby, allowing us to investigate the process of AGN feedback in action. Using the SAURON and GMOS IFUs we detected strong ionised gas emission lines (H$\\alpha$, H$\\beta$, [OIII], [OI], [NII], [SII] and HeI), as well as NaD absorption. We use these tracers to explore the structure of the source, derive the ionised and atomic gas kinematics and investigate the gas excitation and physical conditions. NGC\\,1266 contains two ionised gas components along each line of sight, tracing the ongoing outflow and a component closer to the galaxy systemic, the origin of which is unclear. The gas appears to be being disturbed by a nascent AGN jet. We have presented further evidence the outflow in this object is truly multiphase, containing radio emitting plasma, ionised and atomic gas, as well as the molecular gas and X-ray emitting plasma (as detected in A2011). With outflow velocities up to 900 \\kms, the outflow is very likely to be removing significant amounts of cold gas from the galaxy. The ionised gas morphology correlates well with the radio jets observed in 1.4 Ghz and 5 Ghz continuum emission, supporting the suggestion of A2011 that an AGN jet is the most likely driving mechanism for the ionised gas outflow. The line emission in NGC\\,1266 causes it to be classified as a LINER in optical diagrams. We show here that although NGC\\,1266 undoubtably hosts an AGN (see A2011) the line emission in this object is extended, and is most consistent with excitation from fast shocks caused by the interaction of the radio jet with the ISM. These shocks have velocities of up to 800 \\kms, which match well with the observed velocity of the outflow. Using the observed NaD absorption we are able to set further constraints on the size and orientation of the outflow. We show that we are able to detect atomic gas entrained in the outflow out to a (deprojected) distance of $\\approx$400$\\pm$50 pc, and that the outflow has an inclination (between the galaxy plane and the outflow) of 53$\\pm$8$^{\\circ}$. This suggests the outflow is misaligned from the stellar body. Furthermore we were able to provide an independent estimate of the column density of neutral material in the outflow N$_{\\mathrm{\\hi}}$=(1.2$\\pm$0.6)$\\times$10$^{21}$ cm$^{-2}$. This estimate is consistent with that derived from \\hi\\ absorption in A2011. The neutral and molecular outflows are well correlated, but appear to be outflowing along a slightly different axis to the ionised gas. The cause of this affect is unclear. NGC\\,1266 is a highly complex object, and it is clear that further observations will be required to fully understand it. Observations of single emission lines (either with an IFU or with a Fabry-P\\'erot instrument) will be important to overcome problems with line blending inherent in this high velocity dispersion source. Higher spatial resolution would also be advantageous to obtain clear diagnostics of AGN activity, and better understand the shock structure. Sensitive interferometric radio observations at high angular resolution would also enable us to understand the morphology and orientation of the nascent radio jet. The Atacama Large Millimeter/submillimeter Array (ALMA) will allow us to study the molecular component of the outflow in greater detail, and further determine the driving mechanism. For instance if molecular shock tracers are found predominantly in the outflow then a kinetic driving mechanism would be favoured, while if photon dominated region tracing species were detected this would argue for a radiation driven component to the outflow. This galaxy is one of the few currently known in which we can witness ongoing active feedback, where a central AGN is disrupting its star-forming reservoir. It is clear that understanding the processes removing the ISM will have widespread applications to both theoretical and observation attempts to understand AGN feedback, its affect on the ISM, and role in building up the red-sequence. \\vspace{10pt} \\noindent \\textbf" }, "1207/1207.3959_arXiv.txt": { "abstract": "In this paper, we examine the cosmological viability of a light mass galileon field consistent with local gravity constraints. The minimal, $L_3=\\Box\\phi(\\partial_\\mu \\phi)^2$, massless galileon field requires an additional term in order to give rise to a viable ghost free late time acceleration of Universe. The desired cosmological dynamics can either be achieved by incorporating an additional terms in the action such as $(L_4,L_5)$ $-$ the higher order galileon Lagrangians or by considering a light mass field {\\it \\`a la} galileon field potential. We analyse the second possibility and find that: (1) The model produces a viable cosmology in the regime where the non-linear galileon field is subdominant, (2) The Vainshtein mechanism operates at small scales where the non-linear effects become important and contribution of the field potential ceases to be significant. Also the small mass of the field under consideration is protected against strong quantum corrections thereby providing quantum stability to the system. ", "introduction": "Modified theories of gravity pose a serious alternative to dark energy, an exotic cosmic fluid, needed to account for late time cosmic acceleration in the framework of standard lore\\cite{review1, vpaddy,review2,review3,review3C,review3d,review4,review5}. The task of building an alternative to Einstein's theory is very challenging as the latter fits with the observation with a great accuracy locally thereby a large scale modification is either felt locally or the framework gets reduced to $\\Lambda CDM$. A viable alternative theory of gravity should satisfy important requirements: (1) It should be close to $\\Lambda CDM$ but yet distinguishable from it, (2) Theory should be free from ghost and tachyon instabilities and (3) Theory should not conflict with local physics. The third requirement is often very stringent and its compliance needs the invoking of special mechanisms. The modified theories necessarily include additional scalar degree(s) of freedom in some form or the other. And to be relevant to late time cosmic acceleration, it should be light mass entity which, on the other hand, could cause a havoc as a fifth force never seen in the laboratory or in the solar neighborhood. It therefore necessary to screen out the effect of the fifth force in a delicate manner. Broadly, there are two methods of mass screening, the chameleon mechanism and the Vainshtein effect. The chameleon scenario \\cite{Khoury:2003aq} relies on the direct coupling of matter to scalar field with potential. The mass of the field becomes dependent on the density of environment which justifies the designation \"chameleon\" for such a field. The chameleon potential is chosen such that the effective mass of the field increases with local matter density thereby leading to suppression of fifth force. The chameleon mechanism is an extremely powerful tool for mass screening but as noticed by many authors (see e.g. \\cite{Babichev:2009fi}) the numerical integration of the system for a spherically symmetric background hits singularity because of the form of the chameleon potential. Hence an extreme fine-tuning of the initial conditions is required to avoid the problem. The chameleon theories are also plagued with the problem of large quantum corrections because of the large mass of the chameleon field required to pass the local gravity tests \\cite{Upadhye:2012vh}. The Vainshtein mechanism \\cite{Vainshtein:1972sx} is a superior field theoretic method of mass screening. It was invented by Vainshtein in 1972 to address the discontinuity problem in massive gravity of Pauli-Fierz. It relies on non-linear derivative term of the type $L_3=(\\partial_\\mu \\phi)^2\\Box \\phi$ where $\\phi$ is scalar degree of freedom of helicity zero graviton in this case. The dynamics of non-linear term gives rise to a miraculous phenomenon: Around a massive body, in a large radius dubbed Vainshtein radius, fifth force is suppressed switching off any modification to gravity locally. The field is strongly coupled to itself and hence becomes weakly coupled to matter sector. The DGP model \\cite{Dvali:2000hr} contains such a non-linear term in the so called decoupling limit responsible for the compliance of the model with local gravity constraints \\cite{Luty:2003vm}. The non-standard kinetic term also occurs in Kaluza-Klein reduction of Gauss-Bonnet gravity to four dimensional space time \\cite{Gannouji:2011qz} which makes clear why the theory is free from Ostrogradsky ghosts. The role of the scalar degree of freedom is played by dilaton in this case. The field $\\phi$ due to the presence of an underlying symmetry in flat space time was termed as galileon. There exist higher order galileon Lagrangians $L_4$ and $L_5$ that contain higher order non linear derivative terms in which four and five galileon fields participate respectively. Recently, more general galileon action were constructed \\cite{deRham:2010eu,Goon:2011qf} and their cosmological implications were investigated \\cite{DeFelice:2010pv}. They belong to a more general class of models first introduced in \\cite{Fairlie:1991qe}. The galileon field with the lower order term $L_3$ is sufficient to take care of the local gravity constraints whereas $L_5$ does not contribute to mass screening and though $L_4$ can effect the numerics of Vainstein radius but adds nothing to the underlying physics of Vainshtein mechanism. However, galileon system with $L_3$ term alone can not give rise to late time acceleration of universe. It was first demonstrated in \\cite{Gannouji:2010au} that at least one higher order Lagrangian, say $L_4$ \\cite{Nicolis:2008in} be added to action in order to produce a stable de Sitter solution. This is the analogue of the DGP model, where the self-accelerating branch is unstable because of the presence of ghost \\cite{Koyama:2005tx}. Also physical implications of the cubic Galileon term coupled non-minimally to the metric was first studied in \\cite{Chow:2009fm,Silva:2009km}. Since we are working in a phenomenological setting, we could restrict ourselves to $L_3$ but add a potential to galileon to produce late time cosmic acceleration. We do not assign a field theoretic mechanism to produce the potential for galileon, perhaps one would need some non-perturbative machinery to do the job. Our proposal can be seen as an attempt to replace chameleon mechanism by the Vainshtein effect: In chameleon scenario, the choice of specific form of chameleon potential is crucial to the model which needs to be extremely fine-tuned locally for the successful implementation of the underlying chameleon ideology. ", "conclusions": "In this paper, we have investigated a particular case of a more general class of models which do not produce a gravitational slip. This model includes the $L_3$-term derived in the decoupling limit of DGP and a galileon field potential added to the Lagrangian on phenomenological grounds. The potential breaks the galilean symmetry in a flat spacetime but serves as an important toll at large scales for obtaining a viable cosmology. The scale in the model and the parameters are carefully set such that the galileon term is dominant at small scales which suppress the effect of the chameleon effect leaving the Vainshtein mechanism to operate. This was one of the motives of our proposal to replace the chameleon mechanism as the latter is plagued by the problems of fine tuning and large quantum corrections due to the large mass of the chameleon field. In contrast, the domination of the galileon term at small scales successfully generates a screening effect without requiring a large mass of the field and the quantum corrections remain small. We have, therefore shown that models including light mass scalar fields coupled to matter become viable as soon as we introduce a lower order galileon, term, $(\\nabla \\phi)^2\\Box\\phi$ in the action. We should also emphasize that this term does not significantly modify the background cosmology of the model. Our model is built in away as not to produce large deviation from standard model for cosmological perturbations due to the vanishing of the gravitational slip. This approach can be generalized to a large class of light mass scalar field models which are otherwise ruled out by local gravitational constraints.\\\\ We have demonstrated that the model passes the local tests and can give rise to viable cosmology at late times. It produces a small deviation of the growth factor compared to $\\Lambda$CDM model. It interesting to note that in case of pure coupled quintessence, $G_{eff}=G(1+2\\beta^2)$ thereby requiring coupling to be small in order to respect the BBN constraint. On the other hand, in the model under consideration, $G_\\text{eff}=G$ as in General Relativity which certainly alleviates the constraints on coupling at large redshifts. In principle, our scenario involves cosmological constraints at low redshift which depends on the form of the potential as usual." }, "1207/1207.1447_arXiv.txt": { "abstract": "In several gamma-ray bursts (GRBs) excess emission, in addition to the standard synchrotron afterglow spectrum, has been discovered in the early time X-ray observations. It has been proposed that this excess comes from black body emission, which may be related to the shock break-out of a supernova in the GRBs progenitor star. This hypothesis is supported by the discovery of excess emission in several GRBs with an associated supernova. Using mock spectra we show that it is only likely to detect such a component, similar to the one proposed in GRB~101219B, at low redshift and in low absorption environments. We also perform a systematic search for black body components in all the GRBs observed with the \\emph{Swift} satellite and find six bursts (GRB~061021, 061110A, 081109, 090814A, 100621A and 110715A) with possible black body components. Under the assumption that their excess emission is due to a black body component we present radii, temperatures and luminosities of the emitting components. We also show that detection of black body components only is possible in a fraction of the \\emph{Swift} bursts. ", "introduction": "Gamma-ray bursts (GRBs) emit extreme amounts of $\\gamma$-rays on a short time scale; typically $10^{50}-10^{54}$ erg are released in $0.1-100$ seconds. Only violent processes, such as a compact object merger or the collapse of a massive star, can explain these large energy releases, which have made GRBs observable out to high redshifts of $z\\approx 8-9$ \\citep{2009Natur.461.1254T,2009Natur.461.1258S,2011ApJ...736....7C}. There exists strong evidence, that the collapse of massive stars can produce long GRBs \\citep[$T_{90}>2$ s;][]{1993ApJ...413L.101K}, since spectroscopic features from supernovae (SNe) have been detected in optical follow-up observations of GRBs (e.g. \\citet{1998Natur.395..670G,2003Natur.423..847H,2011MNRAS.411.2792S}. Also see review by \\citet{2011arXiv1104.2274H}). All these SNe are of type Ic with broad lines and no signs of Hydrogen or Helium. Besides these spectroscopic detections, evidence for SN Ic features is also found in light curves of some GRBs \\citep[][]{2001ApJ...555..900P,2004ApJ...606..381L,2010ApJ...718L.150C,2011MNRAS.413..669C}. One burst of particular interest was GRB~060218 \\citep{2006Natur.442.1014S,2006Natur.442.1018M}. It had an associated supernova \\citep{2006Natur.442.1011P}, and its X-ray afterglow could best be described by a combination of synchrotron emission, which is usual for afterglows, and black body emission \\citep{2006Natur.442.1008C}. \\citet{2007ApJ...667..351W} showed that this black body emission could origin in a shock generated by the breakout of a supernova through the surface of the GRBs progenitor star. Subsequently, thermal X-ray emission which may be described by a black body has been suggested in GRB~090618, GRB~100316D and GRB~101219B, which all have associated SNe (\\citet{2011MNRAS.416.2078P}, \\citet{2011MNRAS.411.2792S}, \\citet{Starling2012}, respectively. See also \\citet{IAU}). This supports the connection of the black body component with emission from a supernova. Deviations from a single power law in the early X-ray spectra in GRBs, was also found by \\citet{2007ApJ...656.1001B}, who identified a soft emission component in $5-10$ \\% of the bursts in the studied sample. In this series of papers, we search for more bursts with X-ray black body components, and derive the conditions under which such components may be reliably recovered. In Paper I \\citep{Starling2012} black body components were identified in bursts with spectroscopic or photometric signatures in the optical. The aim of this paper is to perform a systematic search for more bursts with X-ray black body components in the \\emph{Swift} sample, and to derive the conditions under which such components may be reliably recovered. In Section~2 the sample and model fitting are described, and in Section~3 we create simulated spectra to set constraints on the detectability of black body components. In Section~4 bursts are selected as candidates for having black body emission, and Section~5 presents and discusses the final list of candidates. Section~6 extracts physical parameters, assuming that the excess emission is black body emission, and the fraction of GRBs with probable excess emission is examined. For the cosmological calculations we assume a $\\Lambda$CDM-universe with $h_0=0.71$, $\\Omega_\\mathrm{m} = 0.27$, and $\\Omega_\\Lambda = 0.73$. All stated errors and error bars are 90\\% confident. In the plots $n_H$ is in units of $10^{22}$ cm$^{-2}$ unless stated otherwise. We will use the words \\emph{thermal components} and \\emph{black body components} interchangeably. ", "conclusions": "We have first examined under which conditions a black body component, like the one proposed to be in GRB~101219B, can be recovered. At high redshift and in environments with intermediate or high column densities a detection of such a component will not be possible. We also find that it will be hard to recover a black body component, when a bright afterglow emission is present. We show that detection of a black body component only will be possible in a small fraction of all GRBs in our sample. We have searched for black body components in all the \\emph{Swift} bursts with known redshift and created a list of bursts with possible black body components (GRB~061021, 061110A, 081109, 090814A, 100621A and 110715A). They have temperatures, radii and luminosities similar to those found in previous studies of black body components in GRBs." }, "1207/1207.4781_arXiv.txt": { "abstract": "We study the effects of current systematic errors in Type Ia supernova (SN Ia) measurements on dark energy (DE) constraints using current data from the Supernova Legacy Survey (SNLS). We consider how SN systematic errors affect constraints from combined SN Ia, baryon acoustic oscillations (BAO), and cosmic microwave background (CMB) data, given that SNe Ia still provide the strongest constraints on DE but are arguably subject to more significant systematics than the latter two probes. We focus our attention on the temporal evolution of DE described in terms of principal components (PCs) of the equation of state, though we examine a few of the more common, simpler parametrizations as well. We find that the SN Ia systematics degrade the total generalized figure of merit (FoM), which characterizes constraints in multi-dimensional DE parameter space, by a factor of two to three. Nevertheless, overall constraints obtained on more than five PCs are very good even with current data and systematics. We further show that current constraints are robust to allowing for the finite detection significance of the BAO feature in galaxy surveys. ", "introduction": "\\label{sec:intro} Since the discovery of the accelerating universe in the late 1990s \\cite{Riess_1998,Perlmutter_1999}, a tremendous amount of effort has been devoted to improving measurements of dark energy (DE) parameters. As constraints on these parameters improved, controlling the systematic errors in measurements became critical for continued progress. The systematics come in many flavors, including a multitude of instrumental effects and astrophysical effects. Type Ia supernovae (SNe Ia) were used to discover DE and still provide the best constraints on DE. The advantage of SNe Ia relative to other cosmological probes is that {\\it every} SN provides a distance measurement and therefore some information about DE. More recently, SN Ia observations have been joined by measurements of baryon acoustic oscillations (BAO), which provide exceedingly accurate measurements of the angular diameter distance in redshift bins. Cosmic microwave background (CMB) anisotropies come mostly from high redshift and are thus not particularly effective in probing DE, but they do provide one measurement of the angular diameter distance to redshift $z \\simeq 1100$ very accurately. Galaxy clusters also constrain DE usefully, while weak gravitational lensing is expected to become one of the most effective probes of DE in the near future. For recent comprehensive reviews of DE probes, see \\cite{FriTurHut,Weinberg:2012es}. In this work, we are interested in studying the effect of SN Ia systematics on DE constraints by including the {\\it covariance} of measurements between different SNe. The covariance includes primarily systematic errors, and for the first time it has been quantified in depth by \\citet{Conley2011}. Including the effects of the systematic errors, represented by nonzero covariance, weakens the overall constraints on model parameters. Here we wish to explore the effect of systematic errors for general models of DE described by a number of principal components (PCs) of the equation of state, though we first consider these effects for simpler, more commonly used descriptions of the DE sector. We choose to combine the SN Ia data with BAO and CMB measurements and estimate the effects of {\\it current} systematic errors in SN Ia observations. We then proceed to study another systematic concern that is particularly relevant for BAO: whether the finite significance of the detection of the BAO feature in various surveys, when taken into account, weakens the constraints imposed on DE parameters. While we closely follow the accounting for the SN Ia systematics from \\citet{Conley2011}, we note that several other analyses have considered the effect of SN systematics. However, most of these analyses only studied the effects of the systematic errors on the constant equation of state (e.g.\\ \\cite{WoodVasey_2007,Constitution,SDSS_SN,Conley2011}) or included the additional parameter $w_a$ to describe the variation of the equation of state with time (e.g.\\ \\cite{Sullivan2011}). Notable exceptions are studies by \\citet{Davetal07} and \\citet{Rubin_SCP}, which considered a number of specific DE models with non-standard behavior, and \\citet{Amanullah_Union2} and \\citet{Suzuki_SCP}, which parametrized the DE density in several redshift bins. Here our goal is to go beyond any specific models and study the effects of systematic errors in current data on DE constraints in the greatest generality possible. While a truly model-independent description of the DE sector is of course impossible, a description of the expansion history in terms of 10 or so parameters -- which we adopt in this paper -- comes close\\footnote{We do not, however, consider allowing departures from general relativity; doing so would further generalize the treatment.}. In this sense, our paper complements the recent investigations by Mortonson et al.\\ \\cite{Mort_current,MHH_FoM} (see also \\cite{Huterer_Cooray,Wang_Tegmark_2005,Zunckel_Trotta,Zhao_Huterer_Zhang,Hojjati:2009ab,Ishida:2010nk,Shafieloo:2012ht,Seikel:2012uu,Zhao:2012aw}), which studied constraints on very general descriptions of DE using (a slightly different set of) current data but without specific study of the effects of systematic errors. The paper is organized as follows. In Sec.~\\ref{sec:data}, we describe the SN Ia, BAO, and CMB data (and for BAO and CMB, the distilled observable quantities) that we use in our analysis. In Sec.~\\ref{sec:results}, we discuss useful parametrizations of DE and compare constraints on the DE parameters with and without systematic errors included in the analysis. In Sec.~\\ref{sec:BAO_detect}, we investigate the effects of the finite detection significance of the BAO feature in galaxy surveys on the cosmological parameter constraints. In Sec.~\\ref{sec:conclude}, we summarize our conclusions. ", "conclusions": "\\label{sec:conclude} In this paper, we have investigated the effects of systematic errors in current SN Ia observations on DE parameter constraints. We accounted for the systematic errors in SN Ia observations, including the effects of photometric calibration, dust, color, gravitational lensing, and other systematics by adopting a fully off-diagonal covariance matrix between $\\sim 500$ SNe from the SNLS compilation (see Fig.~\\ref{fig:covmatrixplot}). We extended the similar analysis from \\citet{Conley2011} by constraining the temporal evolution of the equation of state of DE described by the pair of parameters $(w_0, w_a)$ as well as a much richer description in terms of 10 PCs of the equation of state (shown in Fig.~\\ref{fig:PCs}). We combined the SN Ia constraints with data from BAO from four different surveys (see Fig.~\\ref{fig:Ameas}) as well as the principal information on DE given by the acoustic peak measurements of the CMB anisotropies measured by the WMAP experiment. The constraints on the simple parametrizations of DE are affected by the systematics, but the overall constraints are still strong even after their inclusion (see Figs.~\\ref{fig:Omw} and \\ref{fig:w0wa}). More importantly, we found that systematic errors affect the contraints somewhat, reducing the DETF FoM by a factor of about three (see Table~\\ref{tab:w0wafom}), while the generalized PC-based FoM is degraded by a factor of two (see Fig.~\\ref{fig:fomvary}). However, as the PC analysis shows, this degradation is mainly restricted to the first two numbers (PC amplitudes) describing DE. In fact, what is particularly impressive about current data is that more than five PCs are well-constrained even in the presence of systematic errors (see Figs.~\\ref{fig:gridPCs} and \\ref{fig:PCconstraints}). In the spirit of testing for systematic effects in current data constraining DE, we also wondered if the relatively low detection significances of BAO features, ranging from about 2.4-$\\sigma$ to 5.0-$\\sigma$ in various surveys, change the overall cosmological constraints. While not a systematic error {\\it per se}, a small but non-negligible probability that the BAO feature has not been detected in some of these surveys implies that the posterior probability of cosmological parameter values asymptotes to a small but nonzero value far from the likelihood peak \\cite{Bassett_Afshordi}. We find that, while the BAO-only constraints are somewhat affected, the combined constraints are not (see Fig.~\\ref{fig:BAO_nondetect}). From all this, we conclude that current systematic errors do degrade DE constraints and FoMs, but not in a major way. Given that future constraints are forecasted to be much better, continued control of current systematic errors remains key for progress in characterizing DE." }, "1207/1207.2672_arXiv.txt": { "abstract": "In this paper, gravothermal oscillations are investigated in multi-component star clusters which have power law initial mass functions (IMF). For the power law IMFs, the minimum masses ($m_{min}$) were fixed and three different maximum stellar masses ($m_{max}$) were used along with different power-law exponents ($\\alpha$) ranging from $0$ to $-2.35$ (Salpeter). The critical number of stars at which gravothermal oscillations first appear with increasing $N$ was found using the multi-component gas code SPEDI. The total mass ($M_{tot}$) is seen to give an approximate stability condition for power law IMFs with fixed values of $m_{max}$ and $m_{min}$ independent of $\\alpha$. The value $M_{tot}/m_{max} \\simeq 12000$ is shown to give an approximate stability condition which is also independent of $m_{max}$, though the critical value is somewhat higher for the steepest IMF that was studied. For appropriately chosen cases, direct N-body runs were carried out in order to check the results obtained from SPEDI. Finally, evidence of the gravothermal nature of the oscillations found in the N-body runs is presented. ", "introduction": "The condition for the onset of gravothermal oscillations is best understood for the case of one-component star clusters, clusters consisting of stars of equal mass. \\cite{Goodman1987} found that gravothermal oscillations first appear when the number of stars $N$ is greater than 7000. This condition has also been confirmed with Fokker-Planck calculations \\citep{Cohn_et_al1989} and by direct N-body simulations \\citep{Makino1996}. However, the multi-component case is more complicated. This is due to the fact that the presence of several components introduces new dynamical processes into the system, and several additional parameters in addition to $N$. Even for the two-component case, which is the simplest kind of mass spectrum, the condition for the onset of gravothermal oscillations is not so simple. Two-component models can be subdivided into Spitzer stable and Spitzer unstable cases depending on whether or not the two components can achieve equipartition of kinetic energy during core collapse \\citep{Spitzer}. \\cite{KimLeeGood1998} studied a range of Spitzer stable two-component models. Their research supported the applicability of the Goodman stability parameter $\\epsilon$ \\citep[see][]{Goodman1993} as a stability criterion. \\cite{breenheggie1}, whose research focused on the more general Spitzer unstable two-component case, indicated that the occurrence of gravothermal instability depends approximately on the number of stars in the heavier component. \\cite{breenheggie1} also found that the critical value of $\\epsilon$ depended on the parameters of the mass function (e.g. stellar mass ratio). However, by using a slightly modified version of $\\epsilon$, one with a modified definition of the half mass relaxation time, they found a nearly constant critical value. \\citet{Murphyetal1990} found that the post-collapse evolution of multi-component models was stable to much higher values of $N$ than in one-component models and that the value of $N$ at which gravothermal oscillations appeared varied with different mass functions. They studied seven-component systems constructed to approximate evolved power law IMFs with masses ranging from $0.1$ to $1.2M_{\\sun}$. The power law exponent that they considered ranged from $-2$ to $-4.5$. They found that gravothermal oscillations appeared when the total mass of the system ($M_{tot}$) was of order $8 \\times 10^4 M_{\\sun}$ \\citep[see][ Figure 6]{Murphyetal1990} and that the critical value of $M_{tot}$ increased with decreasing power law exponent. They suggested that the appearance of oscillations depends on the number of heavier stars. However, this leads to the issue that in a multi-component system it is not clear what the definition of a heavy star should be (this point is discussed in Section 2). The main aim of the present paper is to provide a theoretical understanding of the onset of gravothermal oscillations in multi-component systems. As this present paper follows on from the research of \\cite{breenheggie1} it is worthwhile attempting to extend the concepts developed in that paper to the multi-component case. Although two-component systems may be realistic approximations of multi-component systems \\citep{KimLee1997}, it is best to have a better understanding of gravothermal oscillations in multi-component systems as real globular clusters contain a continuous mass spectrum. What is of particular importance is the effect of varying the maximum stellar mass ($m_{max}$) on the onset of instability as this was not studied by \\cite{Murphyetal1990}. The rest of this paper is structured as follows. In Section 2, we state the results concerning gravothermal oscillations in gaseous models. This section also contains subsections on the Goodman stability parameter and a variant which used a modified relaxation timescale. This is followed by Section 3 in which the results of N-body simulations are given. Finally, Section 4 consists of the conclusion and discussion. ", "conclusions": "\\label{sec:conanddis} The focus of this paper has been on the conditions for the onset of gravothermal oscillations in multi-component systems. We have investigated power law IMFs with different exponents and three different stellar mass ranges $(3.0,0.1)$, $(2.0,0.1)$ and $(1.0,0.1)$. A multi-component gas code has been used to obtain the values of $N_{crit}$. In order to verify the validity of the results direct $N$-body runs were carried out on appropriately chosen cases. The values of $N_{crit}$ found ranged from $2\\times 10^4 $ to $10^5$, and varied with $\\alpha$ and the stellar mass range. Motivated by \\cite{Murphyetal1990}, who found that the total mass of the systems they studied could be used as an approximate stability condition, the value of $M_{crit}$ (the total mass of the system at $N_{crit}$) for each system was calculated (see Table \\ref{table:mcrit}). While for a fixed mass range $M_{crit}$ does provide an approximate stability condition, the value of $M_{crit}$ varied by roughly a factor $m_{max}$. $M_{crit}$ can also be used as an approximate stability condition for the two-component models of \\cite{breenheggie1} so long as the stellar mass ratio is fixed (see Appendix \\ref{A}). In order to find a more general stability condition we applied an extension of an idea first employed in \\cite{breenheggie1}. They used a quantity called the effective particle number ($N_{ef}$). The value was useful because the two-component system that was being considered was expected to behave in roughly the same manner as a one-component system with $N_{ef}$ stars. In the present paper this idea has been extended to multi-component systems. The values of $N_{ef}$ for the multi-component models in this paper are given in Table \\ref{table:mcritNef}. The variation in Table \\ref{table:mcritNef} is significantly less then that in either Table \\ref{table:tab1} or Table \\ref{table:mcrit}. A stability condition of $N_{ef} \\sim 10^4$ covers most of the values of Table \\ref{table:mcritNef} and indeed the two-component models of \\cite{breenheggie1} (see Table \\ref{table:aptab3}). The Goodman Stability Parameter was also tested for the multi-component case (see Table \\ref{table:goodmanepsilon1}). The critical values in Table \\ref{table:goodmanepsilon1} were found to be lower than the value for a one-component model ($\\log_{10} \\epsilon = -2$) and also varied with $\\alpha$ and, to a much lesser extent, with $m_{max}$. By modifying the Goodman Stability Parameter using a slightly different definition for the half-mass relaxation time (based on the effective particle number) a critical value was found which was more consistent with the critical value for a one-component model (see Table \\ref{table:goodmanepsilon2}). \\cite{Goodman1987} used a gas model to find the value of $N_{crit}$ ($=7000$) for a single component system. Technically what he showed was that steady post-collapse expansion was possible in a gas model for all $N$, but that it was unstable for $N>7000$. While the gas model used by \\cite{Goodman1987} is similar in form to the model used here and in \\cite{breenheggie1}, there are two notable differences. Firstly \\cite{Goodman1987} used a larger energy generation rate than the one used here. Secondly, the parameter of the coulomb logarithm that was used was $\\lambda=0.4$. A value of $\\lambda=0.4$ \\citep{Spitzer} was a reasonable choice at the time, but it has since been shown that $\\lambda=0.11$ is a better choice for a single-component model \\citep{GierszHeggie1994}. (For multi-component models the value of $\\lambda=0.02$ was found to provide a good fit \\citep{GierszHeggie2}). These two differences affect the stability in opposite ways$:$ by arguments similar to those given in the last two paragraphs in Appendix \\ref{A}, a larger energy generation rate will increase stability, whereas a larger value of $\\lambda$ tends to reduce stability. For example for $N=7000$ with $\\lambda=0.4$ the increase in the relaxation rate is $20\\%$ compared with $\\lambda=0.11$. % In the present paper, we have made the assumption that multi-component systems will be depleted in stars with stellar mass greater than $3 M_{\\sun}$. This neglects the possibility of systems containing a population of stellar mass Black Holes, which would require a value of $m_{max}$ about an order of magnitude greater than what is considered here. These systems are outside the parameter space studied by \\cite{breenheggie1} and \\cite{KimLeeGood1998}, as the total mass ratio is lower than the range considered by \\cite{breenheggie1} and the stellar mass ratio is higher than the values considered by \\cite{KimLeeGood1998}. The onset of gravothermal oscillations and the more general evolution of systems containing a population of stellar mass black holes are the topics of the next paper in this series. To conclude, a stability condition of $N_{ef} \\sim 10^4$ does apply to the multi-component systems in this paper and the two-component systems of \\cite{breenheggie1}. This condition is expected to apply to any multi-component system provided that there is a sufficient number of stars with stellar mass $\\sim m_{max}$." }, "1207/1207.0993_arXiv.txt": { "abstract": "\\baselineskip=0.6 cm The properties of the ergosphere and energy extraction by Penrose process in a rotating non-Kerr black hole are investigated. It is shown that the ergosphere is sensitive to the deformation parameter $\\epsilon$ and the shape of the ergosphere becomes thick with increase of the parameter $\\epsilon$. It is of interest to note that, comparing with the Kerr black hole, the deformation parameter $\\epsilon$ can enhance the maximum efficiency of the energy extraction process greatly. Especially, for the case of $a>M$, the non-Kerr metric describes a superspinning compact object and the maximum efficiency can exceed $60\\%$, while it is only $20.7\\%$ for the extremal Kerr black hole. ", "introduction": "In 4-dimensional general relativity, no-hair theorem \\cite{noh} guarantees that a neutral rotating astrophysical black hole is uniquely described by the Kerr metric which only possesses two parameters, the mass $M$ and the rotational parameter $a$. For the Kerr black hole, the fundamental limit is the bound $a\\leq M$, and the central singularity is always behind the event horizon due to the weak cosmic censorship conjecture \\cite{WCC}. However, the hypothesis that the astrophysical black-hole candidates are described by the Kerr metrics still lacks the direct evidence, and the general relativity has been tested only for weak gravitational fields \\cite{CMW}. In the regime of strong gravity, the general relativity could be broken down and astrophysical black holes might not be the Kerr black holes as predicted by the no-hair theorem \\cite{FCa,TJo,CBa}. Several parametric deviations from the Kerr metric have been suggested to study observational signatures in both the electromagnetic \\cite{Reviews} and gravitational-wave \\cite{EMRIs} spectral that differ from the expected Kerr signals. Recently, Johannsen and Psaltis proposed a deformed Kerr-like metric \\cite{TJo} suitable for the strong field of the no-hair theorem, which describes a rotating black hole (we named it the non-Kerr black hole) in an alternative theory of gravity beyond Einstein's general relativity. The non-Kerr black hole possesses the following novel features: there is no restriction on the value of the rotational parameter $a$ due to the existence of the deformation parameter $\\epsilon$. Interestingly, for a positive parameter $\\epsilon$, the non-Kerr black hole becomes more prolate than the Kerr black hole and there are two disconnected horizons for high spin parameters, but there is no horizon when $a>M$. Therefore, for a negative parameter $\\epsilon$, the non-Kerr black hole is more oblate than the Kerr black hole, and the horizon always exists for an arbitrary $a$ and the topology of the horizon becomes toroidal \\cite{CBa,CBa1}. The non-Kerr metric is an ideal spacetime to carry out strong-field tests of the no-hair theorem. Therefore, a lot of effort has been dedicated to the study of the rotating non-Kerr black hole recently \\cite{FCa,TJo,CBa, CBa1,CBa2, VCa1,TJo1}. In Ref. \\cite{csbchen}, we studied the properties of the thin accretion disk in the rotating non-Kerr spacetime and found that the presence of the deformation parameter $\\epsilon$ changes the inner border of the disk, energy flux, conversion efficiency, radiation temperature, spectral luminosity and spectral cut-off frequency of the thin accretion disk. Moreover, for the rapidly rotating black hole, the effect of the deformation parameter $\\epsilon$ on the physical quantities of the thin disk becomes more distinct for the prograde particles and more tiny for the retrograde ones. These significant features in the mass accretion process may provide a possibility to test gravity in the regime of the strong field in the astronomical observations. The power energy for a active galactic nuclei, X-ray binaries and quasars has always been concerned in the high energy astrophysics. Several mechanisms (i.e. the accretion disk model \\cite{Kozfowski,shakura} and Blandford-Znajek mechanism \\cite{Blandford}) have been proposed to interpret how to extract energy from a black hole and the formation of the power jet. Furthermore, the Penrose process \\cite{WCC,chandrasekhar,efficiency} also provides an important method to extract energy from a black hole. The Penrose process was also extended to the five-dimensional supergravity rotating black hole \\cite{K. Prabhu}, higher dimensional black holes and black rings \\cite{Nozawa}, and Ho\\v{r}ava-Lifshitz Gravity \\cite{Abdujabbarov}. In this paper, we will investigate in detail the ergosphere of the non-Kerr black hole and how the deformation parameter $\\epsilon$ of the non-Kerr black hole affects the negative energy state and the efficiency of the energy extraction. The paper is organized as follows: in Sec. II, we review briefly the metric of the rotating non-Kerr black hole proposed by Johannsen and Psaltis \\cite{TJo} to test gravity in the regime of the strong field and then analyze the ergosphere structure. In Sec. III, we investigate the efficiency of the energy extraction by using the Penrose process. Sec. IV is devoted to a brief summary. ", "conclusions": "In this paper, we present a detail analysis of the properties of the ergosphere in the rotating non-Kerr black hole proposed recently by Johannsen and Psaltis \\cite{TJo} to test gravity in the regime of the strong field in the future astronomical observations. We now summarize our results as follows: (1) We present the restricted conditions on the deformation parameter $\\epsilon$ to guarantee that the non-Kerr black hole has the connected horizons (see Eq. \\ref{region2} and Fig. \\ref{htt}). (2) We show that the ergosphere is sensitive to the deformation parameter $\\epsilon$ (see Fig. \\ref{ergosphere11}) and the shape of the ergosphere becomes thick with increase of the deformation parameter $\\epsilon$. (3) We find that, comparing with the Kerr black hole, the deformation parameter $\\epsilon$ not only enlarges the negative energy $E$ (see Fig. \\ref{negativeff}) but also enhances the maximum efficiency of the energy extraction process (see tables \\ref{table1}-\\ref{table2} and Fig. \\ref{effffffffffff}). Moreover, the maximum efficiency can exceed $60\\%$ for the non-Kerr compact objects with $a>M$. The influence of the deformation parameter $\\epsilon$ on the maximum efficiency presents a good theoretical opportunity to distinguish the non-Kerr black hole from the Kerr one and to test whether or not the current black-hole candidates are the black holes predicted by Einstein's general relativity. However, we think such a test is not possible at present." }, "1207/1207.0488_arXiv.txt": { "abstract": "We present the discovery of planet-mass companions to two giant stars by the ongoing Penn State-Toru\\'n Planet Search (PTPS) conducted with the 9.2 m Hobby-Eberly Telescope. The less massive of these stars, K5-giant BD+20 274, has a 4.2 $M_{J}$ minimum mass planet orbiting the star at a 578-day period and a more distant, likely stellar-mass companion. The best currently available model of the planet orbiting the K0-giant HD 219415 points to a $\\gtrsim$ Jupiter-mass companion in a 5.7-year, eccentric orbit around the star, making it the longest period planet yet detected by our survey. This planet has an amplitude of $\\sim$18 m s$^{-1}$, comparable to the median radial velocity (RV) ``jitter'', typical of giant stars. ", "introduction": "Searches for planets around giant stars offer an effective way to extend studies of planetary system formation and evolution to stellar masses substantially larger than 1 $M_{\\odot}$ \\citep{sat03,sat08a,omi11b,hat06,dol09a,nie07,nie09a,qui11}. Unlike their progenitors on the main-sequence, the cool atmospheres of evolved stars produce many narrow spectral lines, making a radial velocity (RV) measurement precision of $<$ 10 m s$^{-1}$ possible. The most conspicuous result of the RV surveys of giants is the absence of short-period planets around these stars. While hot Jupiters around main sequence dwarfs are relatively common \\citep[eg.][]{wri11}, the closest planets in orbit around giants detected so far have orbital radii of $\\sim$0.6 AU \\citep{dol09b, nie09a}. This effect is not seen in subgiant systems, which have at least one known hot Jupiter \\citep{joh10}, and is most likely related to stellar evolution \\citep{vil09,nor10}. Typical periods for planets around giants are on the order of one year or longer, requiring multiple years of observations to characterize their orbits. The longest such period announced to date, detected by \\citet{dol09a} is just over three years. In this paper, we report the detection of a planet in a nearly 6 year orbit around HD 219415. In this case, we are close to encountering the wide-orbit limit of planet detection, which results from the sensitivity of the Doppler velocity method reaching the noise level determined by the intrinsic RV jitter of giant stars. Another consequence of the systematically lengthening time baseline of our survey is the growing frequency of planet detection in binary systems. In addition to one such case recently reported by \\citet{get12}, here we describe a planet around BD+20 274, which has another, most likely low-mass stellar companion. This paper is organized as follows. An outline of the observing procedure and a description of the basic properties of the stars are given in Section 2, followed by the analysis of RV measurements in Section 3. The accompanying analysis of rotation and stellar activity indicators is given in Section 4. Finally, our results are summarized and further discussed in Section 5. ", "conclusions": "In this paper, we report detections of periodic RV variations in two K-giant stars from the list of targets monitored by the PTPS program. The accompanying analyses of the photometric data, as well as of the time variability of line bisectors indicate that the most likely origin of the observed periodicities is the Keplerian motion of planet-mass companions. The selection criteria of our survey have been designed to reject targets that show a RV ramp of $>1$ km s$^{-1}$ in preliminary measurements over the period of 2-3 months. However, stars with a ramp of $<1$ km s$^{-1}$ over that time are continued to be observed, if they exhibit $>20$ m s$^{-1}$ RV scatter around this trend. This has the net effect of rejecting close binaries, but allowing the detection of planets in wide binary systems. Indeed, several of the planetary systems detected by PTPS, including BD+20 274, show long-term radial velocity trends with linear drift velocities of order 1 m s$^{-1}$ day$^{-1}$, suggesting that these stars have binary companions. While the dynamical constraints from these partial orbits are not highly restrictive, further information about the secondaries can be obtained by examining the stellar spectra. The spectra of BD+20 274 do not show signs of lines from a second star, even in the stellar template which has S/N = 355. If we assume that the companion must have S/N $\\sim$10 to be detectable, this suggests that the luminosity ratio between the two objects is of order $\\sim$1300, or about 8 magnitudes. As the absolute magnitude of a typical giant star is $M_{bol} = 0.08$, the maximum mass of a MS companion is 0.5 M$_{\\odot}$. By examining the cross-correlation function used in the line bisector measurements in the manner of \\citet{que95}, this estimate can be further restricted. Using the many lines in the CCF template increases the signal, giving a detectable line intensity ratio for the two stars of $\\sim 3200$. This corresponds to a secondary that is $\\sim$9 magnitudes fainter, or about 0.3 M$_{\\odot}$ for a main sequence star. Both these limits are consistent with those discussed in Section 3. % The range of orbital radii of planets discoverable around giants is restricted by stellar evolution and the $\\sqrt{a}$ minimum mass scaling of the Doppler velocity method. Indeed, no planet around a K-giant with an orbital radius smaller than 0.6 AU has been detected so far \\citep{dol09b, nie09a}, in general agreement with the theoretical estimates based on the influence of tidal effects and stellar mass-loss on orbital evolution \\citep{vil09,nor10,kun11}. At the other end of the range of orbital sizes, the HD 219415 planet reported in this paper illustrates the detectability limits of wide orbit, long-period planets imposed by the enhanced intrinsic RV jitter in giants. This effect has been studied by \\citet{hek06}, who have demonstrated that the rms RV noise of K-giants has a median value of $\\sim$20 m s$^{-1}$ and it tends to increase toward later spectral types. Both these limits are outlined in Figure \\ref{fig5}. Evidently, planets down to the Saturn-mass should be easily detectable around early giants over the 0.1-0.5 AU range of orbital radii, but their existence is apparently impaired by the dynamical effects of stellar evolution. For wide orbits, a 1 M$_{J}$ planet around a 2 M$_{\\odot}$ star would have a RV signal of $<$20 m s$^{-1}$ for $a \\geq$ 1 AU, and could be buried in the intrinsic RV noise of a late-type giant. This detection threshold will become more important as the time baseline of the ongoing giant surveys continues to expand, and it will eventually place a practical upper limit on long-period planet detection around the evolved stars. \\bigskip We thank the HET resident astronomers and telescope operators for support. SG and AW were supported by the NASA grant NNX09AB36G. AN, MA, PZ and GN were supported in part by the Polish Ministry of Science and Higher Education grant N N203 510938 and N N203 386237. GM acknowledges the financial support from the Polish Ministry of Science and Higher Education through the Juventus Plus grant IP2010 023070. The HET is a joint project of the University of Texas at Austin, the Pennsylvania State University, Stanford University, Ludwig-Maximilians-Universit\\\"at M\\\"unchen, and Georg-August-Universit\\\"at G\\\"ottingen. The HET is named in honor of its principal benefactors, William P. Hobby and Robert E. Eberly. The Center for Exoplanets and Habitable Worlds is supported by the Pennsylvania State University, the Eberly College of Science, and the Pennsylvania Space Grant Consortium. \\clearpage" }, "1207/1207.5863_arXiv.txt": { "abstract": "The spectrum and amplitude of the stochastic background of relic gravitons produced in a bouncing universe is calculated. The matter content of the model consists of dust and radiation fluids, and the bounce occurs due to quantum cosmological effects when the universe approaches the classical singularity in the contracting phase. The resulting amplitude is very small and it cannot be observed by any present and near future gravitational wave detector. Hence, as in the ekpyrotic model, any observation of these relic gravitons will rule out this type of quantum cosmological bouncing model. ", "introduction": "\\label{sec:introduction} Bouncing cosmological models \\cite{bouncebn} are being widely investigated because, besides solving the singularity problem in cosmology by construction, they can also solve the horizon and flatness puzzles, and lead to an almost scale invariant spectrum of scalar perturbations if the contracting phase is dominated by dust at large scales \\cite{pert}. One of the calculations that have been done was the evaluation of the spectral index of long wavelength tensor perturbations, $n_T$, and in most of bouncing models they were found to be also scale invariant. With respect to their amplitudes, specific models must be worked out. For instance, in the cyclic ekpyrotic scenario, they were evaluated and it was shown that the amplitudes are too small to be detected by present gravitational wave detectors \\cite{cyclic}. In this paper we will calculate the spectrum and amplitude of relic gravitons in a different bouncing cosmological model. It consists of a Friedmann-Lema\\^{\\i}tre-Robertson-Walker (FLRW) universe filled by dust and radiation, which is contracting classically. As it approaches the singularity, quantum cosmological effects on the background, here described through the Wheeler-DeWitt equation interpreted along the lines of the Bohm-de Broglie quantum theory, avoids the singularity and ejects the universe to the expanding phase we are now experiencing. Note that our bouncing model is very conservative: there is nothing else than dust and radiation, we are working with a $3+1$ dimensional space-time, and we are performing a canonical quantization of the second order perturbed (with respect to the FLRW background) Einstein-Hilbert action of general relativity, interpreted along the lines of a quantum theory appropriate to quantum cosmology, namely, one which does not need any external agent to the quantum system to give a meaning to the quantum calculations. Evolution of quantum perturbations (scalar, vector and tensor) on these quantum backgrounds can be described by simple equations, as it was demonstrated in Refs.~\\cite{Peter:2006id,nelson}. With these equations, we were able to calculate the spectrum (analytically and numerically) and amplitude (numerically) of this stochastic background of relic gravitons. Although, the spectrum and amplitude differ considerably from the cyclic ekpyrotic scenario, the main conclusion remains the same, namely, that the amplitudes are too small to be detected by any present and near future gravitational wave detectors. The paper is divided as follows: in the next section we review the main aspects of the quantum cosmological bouncing model on which the relic gravitons evolves, and we obtain the dynamical equations that the tensor perturbations which describe these relic gravitons must obey. In section III we derive the expression for the critical fraction of the relic gravitons energy density from the tensor perturbations described in section II. In section IV we calculate the spectrum and amplitude of this energy density and the graviton strain either analytically (for the spectrum) and numerically. We end up with the conclusions. ", "conclusions": "In this paper we have calculated the amplitude and spectrum of energy density and strain of relic gravitons in a quantum bouncing cosmological model. The strain spectrum of this quantum bouncing model is different from the cyclic and inflationary scenarios. While these two models have spectra $\\approx k^{-2}, \\approx k^{-1}$ at dust domination, and $\\approx k^{-1}, \\approx k^{0}$ at radiation domination, respectively, our model have spectra $\\approx k^{-2}$ and $\\approx k^{0}$ at the same eras. One possible different scenario in which the power of relic gravitons could be enhanced could be obtained by adding to the matter content of the model some amount of stiff matter, which has primordial spectral index $n_T =2$. This will be the subject of future investigations. As a final point, as in the cyclic ekpyrotic model, the resulting amplitude is too small to be detected by any gravitational wave detector. In particular, the sensitivity of the future third generation of gravitational wave detectors, as for example the Einstein Telescope, could reach $\\Omega_{\\rm GW} \\sim 10^{-12}$ at the frequency range $10-100\\,{\\rm Hz}$ to an observation time of $\\sim 5$ years and with a signal-to-noise ratio $({\\rm S/N})\\sim 3$. Therefore, any detection of relic gravitons, in this frequency range, will rule out this type of quantum bouncing model as a viable cosmological model of the primordial universe." }, "1207/1207.3923_arXiv.txt": { "abstract": " ", "introduction": "Astronomy is as old as human culture. Early agricultural civilisations required reliable predictions of the positions and motions of the Sun and Moon, in order to predict in turn seasons, tides, and river risings. Even in the absence of an extensive scientific model, these predictions relied on careful observations, preserved in the form of almanacs or ephemerides. Documents such as these associate astronomy with not only the first data archives but, since these artifacts still exist, also the oldest data archives in the world. Long-term digital preservation in astronomy is nothing new. We cannot resist saying more about this, in \\prettyref{s:babylon}. Astronomical archiving does however evolve, and in the last few decades both astronomy and particle physics have had to become leaders in large-scale data management. Although astronomical images (now all born digital) have always been substantial in size, they have generally been reasonably manageable. Newer astronomical techniques\\dash and we are thinking of 21st century radio astronomy and gravitational astronomy\\dash are capable of generating truly challenging quantities of data; and particle physics has been generating, and addressing, intimidating data problems for decades. These problems cover both the management and preservation of large data volumes, as technical problems, and the preservation of the data's information content, on substantially varying timescales. \\subsection{Project Background} The Managing Research Data/Gravitational Waves project (MRD-GW) is concerned with the data management arrangements of the \\gls{LSC}, and of the broader \\gls{GW} community. It is one of the six projects in the RDMP strand of the JISC \\gls{MRD} programme~\\cite{jisc-mrd-programme}. The GW community was selected by the \\gls{STFC}, at JISC's invitation, as a representative example of big-science data management practice\\dash as we elaborate below, it has features of both the traditional astronomy and HEP communities, without being identifiable with either of them. While many of the specifics, below, relate to this community, we believe much of the discussion is relevant to the other disciplines. Here\\label{s:curatingbigscience}, we are focusing on the big-science projects which receive strategic support from \\gls{STFC}, rather than the smaller-scale projects funded by specific research grants, since it is these large-scale projects that are distinctive about STFC-funded research. We assume that the outputs of the smaller projects will be managed through disciplinary repositories, in a manner which more closely resembles that of other research councils. The MRD-GW project exists to inform three sets of stakeholders: \\begin{itemize} \\item Although the \\gls{JISC} and the \\gls{DCC} have extensive experience with digital libraries and digital curation in general, there are problems specific to \\q{\\gls{big science}} data which JISC would like to better understand. \\item The Research Councils have recently started to require bidders to include a \\q{data management plan} within project proposals. However there is no consensus on what such a plan should look like for science funded by the \\gls{STFC}. The US \\gls{NSF} has recently placed binding requirements on projects to produce data management plans~\\cite{nsf11}. \\item The LSC community has considerable internal software and administration experience, and has solved a large number of data management problems focused on large-scale data storage and transport. However there is an awareness that (partly because there have been no immediate imperatives to do so) there was until recently no published plan for a long-term data archive. \\end{itemize} The existence of these three groups is reflected in the overall structure of the document. This project's context also includes the broad \\gls{VO} movement, which aims to develop standards and areas of consensus which help scientists have ready access to astronomical data across sub-disciplines and wavelengths. All the stakeholder groups have interests in the success of the \\gls{VO} movement. The project aims to bring together two sets of practice, namely the long-term digital preservation perspectives represented by the OAIS reference model\\index{OAIS} in the abstract and the \\gls{DCC} in particular, and the very considerable experience of practical large-scale data management, embedded within the LSC community.\\footnote{For OAIS, see~\\cite{std:oais} and \\prettyref{s:oais}; in this report the \\q{DCC} is the JISC Digital Curation Centre, not the LIGO Document Control Center.} \\subsection{How to read this document} \\label{s:howtoread} This document is organised into three main sections, broadly corresponding to the three audiences we are addressing. \\prettyref{s:education} is about data management in \\gls{big science}. It is addressed to the \\gls{JISC} and to the data preservation community in general, and is intended to illuminate the ways in which scientists in these areas have distinctive data management requirements, and a distinctive data culture, which contrasts informatively with other disciplines. \\prettyref{s:responsibilities} is primarily addressed to \\gls{STFC} and other similar funders of this type of science. It is concerned with the responsibilities which are imposed on funders by the wider society, and which are passed on to the funded through requirements on the governance of projects and the availability of data. The recommendations here are concerned with how best to express these responsibilities. Finally, \\prettyref{s:practicalities} is primarily addressed to the \\gls{LSC}, as a proxy for similar big-science projects. The explicit recommendations here are intended to be of as much interest to projects, as actions they may wish to take, as to funders, as behaviour it may be prudent or productive to require. \\subsection{Working with communities\\dash pragmatics} This report is the result of a fruitful collaboration with the \\gls{GW} community. It may be useful to note some of the features of the project, and the community, which contributed to this. \\begin{itemize} \\item The project team, as part of Glasgow University, has current involvement in the community, and the project director (Woan) is a senior figure there. \\item The \\gls{LIGO} community is already aware of the general need for data management, and the specific need for preservation (see~\\cite{anderson11}). \\item The project personnel have relevant scientific background, and are to some extent in the position of being informatics-for-astronomy specialists (ie we're \\q{insiders}). \\item The community is large and (via studies such as~\\cite{collins04}) has some experience of being \\q{studied}. \\item The existing \\gls{LVC} workshop series meant that we could contact relevant people easily in a context where newcomers were expected, and we didn't have to add our own data management workshop. \\end{itemize} ", "conclusions": "In this report, we have described some of the ways in which \\q{big science} manages its data, as part of a broader data culture which is characterised by large collaborations, and which has decades of experience in agreeing how, and when, and when not, to share data. We can say with some confidence that the big science data culture manages its data well (and this seems to be corroborated by the AIDA assessment discussed in \\prettyref{s:aida}), but we are not suggesting that other disciplines could or should simply copy this culture, since there are various reasons (cf, \\prettyref{s:big-science-easy}) why this culture is particularly natural in some areas. There are however some practices which we do believe are straightforwardly portable to other disciplines. As we discuss in \\prettyref{s:reification}, the notions of \\emph{data products} and \\emph{proprietary periods} very naturally concretize otherwise diffuse arguments about data management and sharing, transforming them from \\q{whether} and \\q{why} to \\q{which} and \\q{how long}. As well, we believe that embedding data management in the day-to-day practice of researchers lowers costs in both the short term (researchers can easily re-find their own data, and interpret others') and the long term (since preservation becomes a technical problem of conserving an in-use repository). We discuss the costing of data management at slightly greater length in \\prettyref{s:preservation-costs}. We repeat our explicit recommendations below. \\recommendations \\subsection*{Acknowledgements} This project was funded by JISC, as part of the \\q{Managing Research Data} programme. We are most grateful to the numerous people who have commented on various drafts of this report, or provided us with information or resources. In particular, we thank Stuart Anderson (LIGO), Paul Butterworth (NASA), Harry Collins (Cardiff), Fernando Comer\\oacute n (ESO), Joy Davidson (DCC), Fran\\ccedilla oise Genova (CDS), Magdalena Getler (DCC), Simon Hodson (JISC), Sarah Jones (DCC), Dorothea Salo (Wisconsin), Angus Whyte (DCC), and Roy Williams (LIGO). \\appendix \\begingroup" }, "1207/1207.4397_arXiv.txt": { "abstract": "{Shocks in jets and hot spots of Active Galactic Nuclei (AGN) are one prominent class of possible sources of very high energy cosmic ray particles (above $10^{18}$eV). Extrapolating their spectrum to their plausible injection energy from some shock, implies an enormous hidden energy for a spectrum of index $\\sim -2$. Some analyzes suggest the particles' injection spectrum at source to be as steep as -2.4 to -2.7, making the problem much worse, by a factor of order $10^{6}$. Nevertheless, it seems implausible that more than at the very best 1/3 of the jet energy, goes into the required flux of energetic particles thus, one would need to allow for the possibility that there is an energy problem, which we would like to address in this work.} {Sequences of consecutive oblique shock features, or conical shocks, have been theorized and eventually observed in many AGN jets. Based on that, we use by analogy the 'Comptonisation' effect and we propose a scenario of a single injection of particles which are accelerated consecutively by several oblique shocks along the axis of an AGN jet.} {We use detailed test-particle approximation Monte Carlo simulations in order to calculate particle spectra by acceleration at such a shock pattern while monitoring the efficiency of acceleration, calculating differential spectra.} {We find that the first shock of a sequence of oblique shocks, establishes a low energy power-law spectrum with $\\sim E^{-2.7}$. The following consecutive shocks push the spectrum up in energy, rendering flatter distributions with steep cut-offs, and characteristic depletion at low energies, an effect which could explain the puzzling apparent extra source power.} {Our numerical calculations show a variation of spectral indexes, starved spectra and a general spectral flattening, connecting to multiple shock acceleration conditions, the relativistic nature of the shocks and the steepness of the magnetic field to the shock normal, shedding further light into understanding the jet-magnetic field geometry and the irregular or flat spectra observed in many AGN jets (e.g. Cen A, 3C279, PKS 1510-089). Furthermore, the $E^{-2.4}-E^{-2.7}$ UHECRs injected source spectra claimed by many authors, could be explained by the superposition of several, perhaps many emission sources, all of which end their particle shock acceleration sequence with flatter, starved spectra produced by two or more consecutive oblique shocks along their jets, or could also imply a mixed component of the accelerated particles above $10^{19}$eV.} ", "introduction": "The flux of ultra high energy cosmic rays (UHECRs) at $E\\ge10^{9}$~GeV is believed to arise in plasma shock environments in extragalactic sources; most favourable being the hot spots or strong shocks in the jets and radio lobes of Active Galactic Nuclei (AGN) (e.g. Ginzburg \\& Syrovatskii 1963, % Biermann \\& Strittmatter 1987, Rachen \\& Biermann 1993) or alternatively, shocks in Gamma Ray Burst environments (e.g. qualitatively by Biermann 1994, quantitatively by Waxman 1995, Vietri 1995). Nevertheless, An absence of neutrinos associated with cosmic-ray acceleration Gamma Ray Bursts (IceCube Collaboration, 2012) shows that no neutrinos are associated with GRBs In the mechanism of diffusive shock acceleration, particles repeatedly gain energy in multiple crossings of an astrophysical shock discontinuity, due to collisions with upstream and downstream magnetic scattering centers, resulting in a power-law spectrum extending up to the energy of the observed UHECRs events. Initially, theoretical developments on particle acceleration were made by Fermi (1949, 1954), Darwin (1954) and during 1970s, the work of Axford et al. (1977), Krymsky (1977), Bell (1978) and Blandford \\& Ostriker (1978) established the 1st order Fermi mechanism for diffusive particle shock acceleration. Ultra-high energy measurements by the Auger experiment (Auger Collaboration 2007, 2010a, 2010b), have shown that AGN seem a quite plausible candidate source for UHECRs. Auger also reported a suppression of cosmic ray flux above $4 \\cdot 10^{19}$ eV. The HiRes experiment reported an observation of the cutoff (HiRes Collaboration 2009). Moreover Auger measurements indicated a weak statistical trend which supports the idea that very high energy events are localized (Auger Collaboration 2010b). In addition, the detection of very hard TeV-energy spectra of high-redshift ($z>0.15$) Blazars, like the case of 3C279 (Teshima et al., 2009) detected by the MAGIC telescope at a redshift of $z=0.536$, poses one of the most interesting questions in astroparticle physics and shock acceleration efficiency. It seems implausible that more than at the very best of 1/3 of an AGN jet energy goes into energetic particles (Falcke and Biermann 1995, Falcke et al. 1995); and so equating the power in UHECRs, adding the additional power hidden at lower energies, and equating it with 1/3 of the observed jet-power, it yields a condition on the spectrum: that in turn this would imply a spectrum flatter than $E^{-2.2}$. Thus, one needs to allow for the possibility that there is an energy problem, which we would like, among other aspects, to address in the present study. The best simple power-law spectral fit to the UHECR ($\\leq 10^{9}$ GeV) data suggests a spectral index between -2.4 and -2.7 (de Marco \\& Stanev, 2005, Berezinsky et al. 2006, 2009). Thus a total power of $10^{48.5} \\times D_{10}^{2}$ erg/s is required, where $D_{10}$ denotes the distance in units of 10 Mpc. The prime candidates seem to be Cen A and M87 (Ginzburg \\& Syrovatskii 1963, Biermann \\& Strittmatter 1987). Very recently, Gopal-Krishna et al. (2010) obtained an excellent fit for an UHECR signal contribution by Cen A, but not just a simple power-law, but a spectrum with a kink downwards. Assuming that up to 1/3 of the total jet power can be supplied, then the maximal power possible to be provided in UHECR is $10^{42.5}$ erg/s for Cen A, and $10^{44.5}$ erg/s for M87 (e.g. Whysong \\& Antonucci 2003, Abdo et al. 2009, Gopal-Krishna et al. 2010), which falls far below than what is required to explain the data. Therefore, one solution would be to consider spectra, that are maybe starved at lower energies. We note here that, it seems plausible to assume that at around a few $10^{19}$ eV there are many candidates and only at higher energies there are either very few or perhaps only one source contribution, which would worsen the energy problem (i.e. the distance of either Cen A or M87 and the interaction with the microwave background is not quite so important, see Greisen 1966, Zatsepin and Kuzmin 1966). The observed standing shock features in AGN jets (e.g. Sanders 1983, Gawthorne 2006, Marscher et al. 2008), and the system of repeated shocks observed in jet structure at vastly discrepant spatial resolutions, like in the radio galaxy NGC6251 (Bridle and Perley, 1984) motivated us for this work. Moreover, it is well known from normal quantum statistics, that prescribing the number of photons in a box filled with hot gas of temperature $T$, their total energy, and forbidding both creation and destruction of photons, gives highly distorted Planck spectra. These spectra as a function of photon energy $\\epsilon$ can be described using a \"chemical potential\", $\\mu$ (e.g. Leighton 1959), $N(\\epsilon) \\; \\sim \\; e^{-\\mu -\\epsilon/kT}$. This is called incomplete Comptonisation effect. The process of incomplete Comptonisation in the disks of accreting compact objects is one well known astrophysical application (e.g. Sunyaev 1970, Katz 1976, Rybicki and Lightman 1993), see figure 1. If there is an insufficient number of photons available, the generated spectrum is depressed at low energies. This entails, that a low energy extension of high energy spectra may not be realistic, and therefore, we need to allow for depressed low energy spectra, thus we wish to draw here an analogy to the possibility that cosmic ray spectra produced by acceleration in multiple oblique shock structures in AGN jets, may also be depressed (starved) in lower energies while extending to higher energies. In the present study, we will show that the first shock from a shock sequence, establishes a power-law spectrum with a spectral index of $\\sim 2.7$, while the following shocks of the sequence push the particle spectrum up in energy, with flatter distributions manifesting a characteristic flux depletion at lower energies. \\subsection{Jets and multiple-oblique shocks} First attempts to simulate relativistic hydrodynamical jets were made by Marti et al. (1994) and by Duncan and Hughes (1994) for low Mach numbers and by Marti et al. (1995) for high Mach numbers. Their studies showed the same trends with the global phenomenology of non-relativistic jet evolution, nevertheless it was shown that relativistic jets propagate more efficiently into the ambient external medium comparing to a non-relativistic head jet velocity. A year later, Massanglia et al. (1996a) presented more detailed two-dimensional simulations. Three-dimensional simulations (e.g. Arnold and Arnett 1986, Clarke 1996) also reveal internal oblique shock structures in light or moderately supersonic jets. Among many insightful findings regarding the structure of a relativistic jet, it was found that when a high pressure cocoon develops, it squeezes the jet and drives shock waves into it, which reflect on the axis and form a conical (in other words an oblique shock sequence) shape. The aspect of this interaction depends on the Mach number. Furthermore, one could establish a relationship between an optimum angle ($a$) of the cone between its tangential shock surface and the flow parallel to the jet axis expressed in Lorentz factor $\\Gamma$ as, $sin~ $a$~\\propto 1/\\Gamma$. The critical parameter was found to be the inclination of the conical shocks that determined the thrust behind the jet head, meaning in planar jet geometry: \\textit{Oblique shocks must have a small inclination angle to the axis in order to produce a strong acceleration effect}. The physics of magneto-hydrodynamic (MHD) shocks topology in magnetized flows is still not well understood in the plasma and astrophysical community. Efforts to understanding the jet physics via MHD simulations in 1D, 2D and 3D are currently under investigation. Despite the fact that jet dynamics is inherently three-dimensional (3D), much of the physics of these jets can be obtained from simple 1D simulations of flow along a cylindrical shell. For example Cheung et al. (2007) demonstrated that observed proper motions of forward (superluminal) and reverse (subluminal) knots can be reproduced precisely by a 1D relativistic MHD simulation model. A very interesting work was presented in a series of papers by Jones et al. (2001), Tregillis et al. (2001ab, 2004), O'Neil et al. (2006), etc where the authors have conducted extensive studies of high-resolution 2D and 3D MHD simulations of supersonic jets, exploring the influence of the jet Mach number and the ambient medium on jet propagation and energy deposition over long distances. Obviously, only 3D simulations can hope to reproduce the obliquity of the features in some jets observed with very high angular resolution, e.g. NGC6251, 3C66B, M87, etc During the course of this work, we will assume a sequence of oblique shocks along the jet axis (as depicted in figure 4), following the \\textit{reconfinement} mechanism (Sanders, 1983). Sanders showed by the method of characteristics, that the reconfinement process is a typical property of jets observed transversely. The internal pressure in a free jet decreases rapidly with distance from the AGN nucleus, and therefore the jet would be expected to come into pressure equilibrium with the ambient medium. The reconfinement shocks allow the jet flow to adjust to the outside pressure staying supersonic. These shocks shocks are almost periodically repeated, contacting the jet circumference and they are mostly subluminal as long as the flow stays supersonic. The reconfinement mechanism and the solutions of a hydrodynamic wind (e.g. Fukue \\& Okada 1990) show that a jet flow goes through multiple critical points, passing from subsonic to supersonic and vice versa several times, with formation of shock structures. This mechanism is actually an attempt to recover the strength of a weak jet flow, which would mean that one would have repeatedly formations of superluminal to subluminal shocks and so forth, until the critical point were the outside pressure totally overpowers the pressure of the jet in the deceleration phase. Observations of polarized synchrotron radiation in AGN jets (e.g. Agudo et al., 2001), indicate stationary conical components in the observer frame in the parsec-scale jets, which could be identified as recollimation shocks. We note that systems of repeated oblique (conical) shocks are observed in jet structures at vastly discrepant spatial resolutions, like in the radio galaxy NGC6251 (Bridle \\& Perley 1984), interpreting this as a quasi-self-similar pattern, with the local scales all proportional to radial distance along the jet. Marscher et al. (2008, 2010) showed high-resolution radio images and optical polarization measurements of the BL Lac type object PKS 1510-089, revealed a bright feature in its jet that caused a double flare of radiation from optical frequencies to TeV gamma-ray energies, as well as a delayed outburst at radio wavelengths. It was concluded that the observational event started in a region with a \\textit{helical} magnetic field which was identified with the acceleration and collimation zone predicted by theoretical works, and the variability in the observed brightness was due to emission by the plasma excited by a conical standing shock wave. Furthermore, observations of the jet structure of M87 (Walker et al., 2008) and PKS 1510-089 (Marscher et al., 2008) near the central black hole indicate that, while there can be moving shocks between 10 and 1000 Schwarzschild-radii ($r_s$), the first strong stationary shock occurs at $\\sim$ 3000 $r_s$, as already proposed by Markoff et al. (2005) and confirmed by Britzen et al. (2008) in the case of the BL Lac type object S5 1803+784. The evidence of helical jets transverse rotation measure gradients across AGN jets, was predicted by Blandford (1993), and among others found by Gabuzda, Murray and Cronin (2004) in BL Lacs. Brown et al. (2009) found similar evidence in an the FRI radio galaxy 3C78. Asada et al. (2008) performed multi-frequency polarimetry for the quasar NRAO 140 using the VLBA. The observations revealed the existence of helical magnetic components associated with the jet itself. Evidence for a co-existence of shocks and helical magnetic fields is given by Keppens et al. (2008). It is important to note that using MHD simulations they explored the morphology of AGN jets by studying propagation characteristics of a series of highly relativistic, helically magnetized jets and among other they showed that the magnetic helicity changes at internal cross-shocks, which then act to repeatedly re-accelerate the jet. Moreover, Nakamura et al. (2010), showed that extragalactic jets are the result of MHD shocks produced in helically twisted, magnetized relativistic outflows. Additionally, we note that state-of-the art laboratory experiments of supersonic highly conductive plasma flows (e.g. Lebedev 2005, Ampleford et al. 2008) have also shown development of conical shocks forming at its base and along the axis of the developed jet. The topology we assume here considers the axisymmetric reconfinement of shocks as a sequence of shocks with inclination angles $a, b, c, d$ within a supersonic AGN jet, with a helical magnetic field, as shown in figure 4. Specifically we assume that particle acceleration can take place in such internal shocks as theorized or observed in e.g. PKS 1510-089, NGC6251, etc. The paper is structured as follows: In section 2 we present our numerical method, and we discuss results of individual and multiple shock acceleration studies. In section 3 we conclude discussing our findings. ", "conclusions": "We performed Monte Carlo simulations allowing acceleration at relativistic \\textit{multiple} oblique shocks as an application to muliple shocks observed in various AGN jets. We based our work on the re-confinement mechanism in jets and incomplete Comptonisation effect, injecting test-particles once upstream the first shock, and we let them accelerate through a sequence of subluminal and superluminal shocks. Specifically, we studied an exemplary case for six possible combinations of a set of four consecutive relativistic oblique shocks, with an aim to facilitate our understanding on acceleration efficiency and potential particle spectra. In principle, our investigation was initiated by the fact that particle energies with $\\ge 10^{18.5}$ eV imply a lot more total source power. We showed that when flat and starved cosmic ray spectra are attainable, then the puzzling extra power to be provided for UHECRs does not seem necessary. By injecting particles once towards a sequence of shocks it means that with a smaller number of low energy particles one could have a final particle spectrum extending up to very high energies with depletion at lower energies. We numerically showed that the first shock from the multiple-shock sequence, establishes a low energy power-law spectrum with $\\sim E^{-2.7}$, while the next consecutive shocks push the particle spectrum up in energy with flatter distributions, leaving a flux deficiency at low energies. We have shown that for two or more shocks the spectrum becomes flatten than -2.2. Especially, the case where only identical subluminal shocks are involved, the spectra are the flattest and the highest maximum particle energy of a spectacular $10^{11}$eV is attained, confirming as well the expected high efficiency of subluminal shocks shown previously by Meli \\& Quenby (2003b), Meli et al. (2008). On the contrary when only superluminal shocks are present, they render the spectra steeper and the acceleration less efficient. Within our model, other particles except protons (e.g. electrons, heavy nuclei), would naturally suffer losses in addition to all the effects from oblique shocks that could accentuate the sharpness of the final spectrum, as we discussed above, which is of interest to our astrophysical interpretation. Consequently, taking under consideration the work of de Marco \\& Stanev (2005) and Berezinsky et al. (2006, 2009) which require an injected source spectrum for UHECRs between $E^{-2.4}$ and $E^{-2.7}$ before propagation, our model can explain the latter on the basis of either a single source such as Cen A, accelerating a composite population of protons and heavy nuclei, initially flat at source (justifying source power requirements) with additional propagation energy losses and a significant steepening, or 2), due to the superposition of several sources, all of which end their acceleration shock sequence with two or more relativistic subluminal shocks with \\textit{starved} flat spectra (justifying as well source power requirements). Moreover, with the current model we could also explain irregular or very flat gamma-ray spectra by various flaring high red-shifted extragalactic sources. Specifically, inverted spectra may be required to comprehend the tendency to detect TeV sources at redshifts such as 0.5 (3C279; The MAGIC coll., 2008); sources with normal spectra would be undetectable. This remains to be examined in more detail in a follow-up work. Intensive future observations may yet to test or reveal one or more of the above predictions." }, "1207/1207.5558.txt": { "abstract": "We propose a transform for signals defined on the sphere that reveals their localized directional content in the spatio-spectral domain when used in conjunction with an asymmetric window function. We call this transform the directional spatially localized spherical harmonic transform (directional SLSHT) which extends the SLSHT from the literature whose usefulness is limited to symmetric windows. We present an inversion relation to synthesize the original signal from its directional-SLSHT distribution for an arbitrary window function. As an example of an asymmetric window, the most concentrated band-limited eigenfunction in an elliptical region on the sphere is proposed for directional spatio-spectral analysis and its effectiveness is illustrated on the synthetic and Mars topographic data-sets. Finally, since such typical data-sets on the sphere are of considerable size and the directional SLSHT is intrinsically computationally demanding depending on the band-limits of the signal and window, a fast algorithm for the efficient computation of the transform is developed. The floating point precision numerical accuracy of the fast algorithm is demonstrated and a full numerical complexity analysis is presented. ", "introduction": "\\IEEEPARstart{S}{ignals} that are inherently defined on the sphere appear in various fields of science and engineering, such as medical image analysis~\\cite{Chung:2010}, geodesy~\\cite{Simons:2006}, computer graphics~\\cite{Han:2007}, planetary science~\\cite{Audet:2011}, electromagnetic inverse problems~\\cite{Colton:1998}, cosmology~\\cite{Spergel:2007}, 3D beamforming~\\cite{Ward:1995} and wireless channel modeling~\\cite{Pollock:2003}. In order to analyze and process signals on the sphere, many signal processing techniques have been extended from the Euclidean domain to the spherical domain~\\cite{Antoine:1999,Khalid:2012,Khalid_icassp:2012,Khalid2:2012,Marinucci:2008,McEwen:2006,McEwen:2007,McEwen:2008,Sadeghi:2012,Simons:1997,Simons:2006,Starck:2006,Wiaux:2005,Wiaux:2006,Wiaux:2008,Yeo:2008}. Due to the ability of wavelets to resolve localized signal content in both space and scale, wavelets have been extensively investigated for analyzing signals on the sphere~\\cite{Antoine:1999,Marinucci:2008,McEwen:2006,McEwen:2007,Starck:2006,Wiaux:2005,Wiaux:2006,Wiaux:2008,Yeo:2008} and have been utilized in various applications~(e.g., in astrophysics~\\cite{Barreiro:2000,McEwen:2005:WMAP,McEwen:2007:cosmology,Pietrobon:2008,Schmitt:2010,Vielva:2004} and geophysics~\\cite{Audet:2011,Simons:2011_1,Simons:2011_2}). Some of the wavelet techniques on the sphere also incorporate directional phenomena in the spatial-scale decomposition of a signal~(e.g., \\cite{Wiaux:2006,Wiaux:2008,Yeo:2008}). As an alternative to spatial-scale decomposition, spatio-spectral~(spatial-spectral) techniques have also been developed and applied for localized spectral analysis, spectral estimation and spatially varying spectral filtering of signals~\\cite{Khalid:2012,Khalid2:2012,Simons:1997,Wieczorek:2005,Wieczorek:2007}. The spectral domain is formed through the spherical harmonic transform which serves as a counterpart of the Fourier transform for signals on the sphere~\\cite{Colton:1998,Driscoll:1994,McEwen:2011,Sakurai:1994}. The localized spherical harmonic transform, composed of spatial windowing followed by spherical harmonic transform, was first devised in \\cite{Simons:1997} for localized spectral analysis. We note that the localized spherical harmonic transform was defined in \\cite{Simons:1997} for azimuthally asymmetric (i.e., directional) window functions, however, it was applied and investigated for azimuthally symmetric functions only. Furthermore, a spectrally truncated azimuthally symmetric window function was used for spatial localization \\cite{Simons:1997}. Due to spectral truncation, the window used for spatial localization may not be concentrated in the region of interest. This issue was resolved in \\cite{Wieczorek:2005}, where azimuthally symmetric eigenfunctions obtained from the Slepian concentration problem on the sphere were used as window functions (the Slepian concentration problem is studied for arbitrary regions on the sphere in \\cite{Simons:2006}). Following \\cite{Simons:1997}, the spatially localized spherical harmonic transform~(SLSHT) for signals on the sphere has been devised in \\cite{Khalid:2012} to obtain the spatio-spectral representation of signals for azimuthally symmetric window functions, where the effect of different window functions on the SLSHT distribution is studied. Subsequently, the SLSHT has been used to perform spatially varying spectral filtering \\cite{Khalid2:2012}, again with azimuthally symmetric window functions. In obtaining the SLSHT distribution for spatio-spectral representation of a signal, the use of an azimuthally symmetric window function provides mathematical simplifications. However, such an approach cannot discriminate localized directional features in the spatio-spectral domain. This motivates the use of asymmetric window functions in the spatio-spectral transformation of a signal using the SLSHT. In order to serve this objective, we employ the definition of the localized spherical harmonic transform in \\cite{Simons:1997} and define the SLSHT and the SLSHT distributions using azimuthally asymmetric window functions for spatial localization. Since the use of an asymmetric window function enables the transform to reveal directional features in the spatio-spectral domain, we call the proposed transform the directional SLSHT. We also provide a harmonic analysis of the proposed transform and present an inversion relation to recover the signal from its directional SLSHT distribution. Since the directional SLSHT distribution of a signal is required to be computed for each spatial position and for each spectral component, and data-sets on the sphere are of considerable size~(e.g., three million samples on the sphere for current data-sets~\\cite{Jarosik:2010} and fifty million samples for forthcoming data-sets~\\cite{planck:bluebook}), the evaluation of the directional SLSHT distribution is computationally challenging. We develop fast algorithms for this purpose. Through experimental results we show the numerical accuracy and efficient computation of the proposed directional SLSHT transform. Furthermore, due to the fact that the proposed directional SLSHT distribution depends on the window function used for spatial localization, we analyze the asymmetric band-limited window function with nominal concentration in an elliptical region around the north pole, which is obtained from the Slepian concentration problem on the sphere. We also illustrate, through an example, the capability of the proposed directional SLSHT to reveal directional features in the spatio-spectral domain. The remainder of the paper is structured as follows. In Section~\\ref{sec:models}, we review mathematical preliminaries related to the signals on the sphere, which are required in the sequel. We present the formulation of the directional SLSHT, its harmonic analysis and signal reconstruction from the SLSHT distribution in Section~\\ref{sec:SLSHT}. Different algorithms for the evaluation of the SLSHT distribution are provided in Section~\\ref{sec:Algos}. In Section~\\ref{sec:results}, we show timing and accuracy results of our algorithms and an illustration of the transform. Concluding remarks are presented in Section~\\ref{sec:conclusions}. %%---------------------------------------------------------------------- ", "conclusions": "\\label{sec:conclusions} We have presented the directional SLSHT to project a signal on the sphere onto its joint spatio-spectral domain as a directional SLSHT distribution. In spirit, the directional SLSHT is composed of SO(3) spatial localization followed by the spherical harmonic transform. Here, we have proposed the use of an azimuthally asymmetric window function to obtain spatial localization, which enables the transform to resolve directional features in the spatio-spectral domain. We have also presented an inversion relation to synthesize the original signal from its directional SLSHT distribution. Since data-sets on the sphere are of considerable size, we have developed a fast algorithm for the efficient computation of the directional SLSHT distribution of a signal. The computational complexity of computing the directional SLSHT is reduced by providing an alternative harmonic formulation of the transform and then exploiting the factoring of rotation approach~\\cite{Risbo:1996} and the fast Fourier transform. The computational complexity of the proposed fast algorithm to evaluate SLSHT distribution of a signal with band-limit $L_f$ using window function with band-limit $L_h$ is $O(L_f^3 L_h^2 + L_f^2 L_h^4)$ as compared to the complexity of direct evaluation, which is $O(L_f^4L_h^3)$. The numerical accuracy and the speed of our fast algorithm has also been studied. The directional SLSHT distribution relies on a window function for spatial localization; we have analyzed the band-limited window function obtained from the Slepian concentration problem on the sphere, with nominal concentration in an elliptical region around the north pole. We provided an illustration which highlighted the capability of the directional SLSHT to reveal directional features in the spatio-spectral domain, which is likely to be of use in many applications. \\begin{appendices}" }, "1207/1207.2787_arXiv.txt": { "abstract": "We performed for the first time stereoscopic triangulation of coronal loops in active regions over the entire range of spacecraft separation angles ($\\alpha_{sep}\\approx 6^\\circ, 43^\\circ, 89^\\circ, 127^\\circ$, and $170^\\circ$). The accuracy of stereoscopic correlation depends mostly on the viewing angle with respect to the solar surface for each spacecraft, which affects the stereoscopic correspondence identification of loops in image pairs. From a simple theoretical model we predict an optimum range of $\\alpha_{sep} \\approx 22^\\circ-125^\\circ$, which is also experimentally confirmed. The best accuracy is generally obtained when an active region passes the central meridian (viewed from Earth), which yields a symmetric view for both STEREO spacecraft and causes minimum horizontal foreshortening. For the extended angular range of $\\alpha_{sep}\\approx 6^\\circ-127^{\\circ}$ we find a mean 3D misalignment angle of $\\mu_{PF} \\approx 21^\\circ-39^\\circ$ of stereoscopically triangulated loops with magnetic potential field models, and $\\mu_{FFF} \\approx 15^\\circ-21^\\circ$ for a force-free field model, which is partly caused by stereoscopic uncertainties $\\mu_{SE} \\approx 9^\\circ$. We predict optimum conditions for solar stereoscopy during the time intervals of 2012--2014, 2016--2017, and 2021--2023. ", "introduction": "Ferdinand Magellan's expedition was the first that completed the circumnavigation of our globe during 1519-1522, after discovering the {\\sl Strait of Magellan} between the Atlantic and Pacific ocean in search for a westward route to the ``Spice Islands'' (Indonesia), and thus gave us a first $360^\\circ$ view of our planet Earth. Five centuries later, NASA has sent two spacecraft of the STEREO mission on circumsolar orbits, which reached in 2011 vantage points on opposite sides of the Sun that give us a first $360^\\circ$ view of our central star. Both discovery missions are of similar importance for geographic and heliographic charting, and the scientific results of both missions rely on geometric triangulation. The twin STEREO/A(head) and B(ehind) spacecraft (Kaiser et al.~2008), launched on 2006 October 26, started to separate at end of January 2007 by a lunar swingby and became injected into a heliocentric orbit, one propagating ``ahead'' and the other ``behind'' the Earth, increasing the spacecraft separation angle (measured from Sun center) progressively by about $45^\\circ$ per year. The two spacecraft reached the largest separation angle of $180^\\circ$ on 2011 February 6. A STEREO SECCHI COR1-A/B intercalibration was executed at $180^\\circ$ separation (Thompson et al.~2011). Thus, we are now in the possession of imaging data from the two STEREO/EUVI instruments (Howard et al.~2008; W\\\"ulser et al.~2004) that cover the whole range from smallest to largest stereoscopic angles and can evaluate the entire angular range over which stereoscopic triangulation is feasible. It was anticipated that small angles in the order of $\\approx 10^\\circ$ should be most favorable, similar to the stereoscopic depth perception by eye, while large stereoscopic angles that are provided in the later phase of the mission would be more suitable for tomographic 3D reconstruction. The first stereoscopic triangulations using the STEREO spacecraft have been performed for coronal loops in active regions, observed on 2007 May 9 with a separation angle of $\\alpha_{sep}=7.3^\\circ$ (Aschwanden et al.~2008) and observed on 2007 June 8 with $\\alpha_{sep}=12^\\circ$ (Feng et al.~2007). Further stereoscopic triangulations have been applied to oscillating loops observed on 2007 June 26 with a stereoscopic angle of $\\alpha_{sep}=15^\\circ$ (Aschwanden 2009), to polar plumes observed on 2007 Apr 7 with $\\alpha_{sep}=3.6^\\circ$ (Feng et al.~2009), to an erupting filament observed on 2007 May 19 with $\\alpha_{sep}=8.5^\\circ$ (Liewer et al.~2009), to an erupting prominence observed on 2007 May 9 with $\\alpha_{sep}=7.3^{\\circ}$ (Bemporad 2009), and to a rotating, erupting, quiescent polar crown prominence observed on 2007 June 5-6 with $\\alpha_{sep}=11.4^\\circ$ (Thompson 2011). Thus, all published stereoscopic triangulations have been performed within a typical (small) stereoscopic angular range of $\\alpha_{sep} \\approx 3^\\circ-15^\\circ$, as it was available during the initial first months of the STEREO mission. The largest stereoscopic angle used for triangualtion of coronal loops was used for active region 10978, observed on 2007 December 11, with a spacecraft separation of $\\alpha_{sep}=42.7^\\circ$ (Aschwanden and Sandman 2010; Sandman and Aschwanden 2011), which produced results with similar accuracy as those obtained from smaller stereoscopic angles. So there exists also an intermediate rangle of aspect angles that can be used for stereoscopic triangulation. However, nothing is known whether stereoscopy is also feasible at large angles, say in the range of $\\alpha_{sep} \\approx 50^{\\circ}-180^\\circ$, and how the accuracy of 3D reconstruction depends on the aspect angle, in which range the stereoscopic correspondence problem is intractable, and whether stereoscopy at a maximum angle near $\\alpha_{sep} \\lapprox 180^{\\circ}$ is equally feasible as for $\\alpha_{sep} \\gapprox 0^\\circ$ for optically thin structures (as it is the case in soft X-ray and EUV wavelengths), due to the $180^\\circ$ symmetry of line-of-sight intersections. In this study we are going to explore stereoscopic triangulation of coronal loops in the entire range of $\\alpha_{sep} \\approx 6^\\circ - 170^\\circ$ and quantify the accuracy and quality of the results as a function of the aspect angle. Observations and data analysis are reported in Section 2, while a discussion of the results is given in Section 3, with conclusions in Section 4. \\begin{figure} \\centerline{\\includegraphics[width=1.0\\textwidth]{f1.eps}} \\caption{Schematic figure of the spacecraft orbits of STEREO/A and B relative to Earth, with the spacecraft separation angles $\\alpha_{sep}=\\alpha_A-\\alpha_B$ indicated approximately at the beginning of the years, ranging from $\\approx 5^\\circ$ in April 2007 to $\\approx 180^\\circ$ in February 2011.} \\end{figure} ", "conclusions": "After the STEREO mission reached for the first time a full $360^\\circ$ view of the Sun this year (2007 Feb 6), the two STEREO A and B spacecraft covered also for the first time the complete range of stereoscopic viewing angles from $\\alpha_{sep} \\gapprox 0^\\circ$ to $\\alpha_{sep} \\lapprox 180^\\circ$. We explored the feasibility of stereoscopic triangulation for coronal loops in the entire angular range by selecting 5 datasets with viewing angles at $\\alpha_{sep} \\approx 6^\\circ, 43^\\circ, 89^\\circ, 127^\\circ$ and $169^\\circ$. Because previous efforts for solar stereoscopy covered only a range of small stereoscopic angles ($\\alpha_{sep} \\lapprox 45^\\circ$), we had to generalize the stereoscopic triangulation code for large angles up to $\\alpha_{sep} \\le 180^\\circ$. We find that stereoscopy of coronal loops is feasible with good accuracy for cases in the range $\\alpha_{sep} \\approx 6^\\circ - 127^\\circ$, a range that is also theoretically predicted by taking into account the triangulation errors due to finite spatial resolution and confusion in the stereoscopic correspondence identification in image pairs, which is hampered by projection effects and foreshortening for viewing angles near the limb. Accurate stereoscopy (within a factor of 2 of the best possible accuracy) is predicted for a spacecraft separation angle range of $\\alpha_{sep} \\approx 22^\\circ - 125^\\circ$. Based on this model we predict that the best periods for stereoscopic 3D reconstruction during a full 16-year STREREO mission cycle occur during 2012-2014, 2016-2017, and 2021-2023, taking the variation in the number of active regions during the solar cycle into account also. Why is the accuracy of stereoscopic 3D reconstruction so important? Solar stereoscopy has the potential to quantify the coronal magnetic field independently of conventional 2D magnetogram and 3D vector magnetograph extrapolation methods, and thus serves as an important arbiter in testing theoretical models of magnetic potential fields, linear force-free field models (LFFF), and nonlinear force-free field models (NLFFF). A benchmark test of a dozen of NLFFF codes has been compared with stereoscopic 3D reconstruction of coronal loops and a mismatch in the 3D misalignment angle of $\\mu \\approx 24^\\circ-44^\\circ$ has been identified (DeRosa et al.~2009), which is attributed partially to the non-force-freeness of the photospheric magnetic field, and partially to insufficient constraints of the boundary conditions of the extrapolation codes. Empirical estimates of the error of stereoscopic triangulation based on the non-parallelity of loops in close proximity has yielded uncertainties of $\\mu_{SE} \\approx 7^\\circ-12^\\circ$. Thus the residual difference in the misalignment is attributed to either the non-potentiality of the magnetic field (in the case of potential field models), or to the non-force-freeness of the photospheric field (for NLFFF models). We calculated also magnetic potential fields here for all stereoscopically triangulated active regions and found mean misalignment angles of $\\mu_{PF}\\approx 21^\\circ-39^\\circ$, which improved to $\\mu_{FFF}\\approx 15^\\circ-21^\\circ$ for a nonlinear force-free model, which testifies the reliability of stereoscopic reconstruction for the first time over a large angular range. The only case where stereoscopy clearly fails is found for an extremely large separation angle of ($\\alpha_{sep}\\approx 170^\\circ$), which is also reflected in the largest deviation of misalignment angles found ($\\mu_{NP} \\approx 39^\\circ$, $\\mu_{FFF} \\approx 21^\\circ$). Based on these positive results of stereoscopic accuracy over an extended angular range from small to large spacecraft separation angles we anticipate that 3D reconstruction of coronal loops by stereoscopic triangulation will continue to play an important role in testing theoretical magnetic field models for the future phases of the STEREO mission, especially since stereoscopy of a single image pair does not require a high cadence and telemetry rate at large distances behind the Sun." }, "1207/1207.0866_arXiv.txt": { "abstract": "Through a series of observations with the Australia Telescope Compact Array we have monitored the variability of ground-state hydroxyl maser emission from G12.889+0.489 in all four Stokes polarisation products. These observations were motivated by the known periodicity in the associated 6.7-GHz methanol maser emission. A total of 27 epochs of observations were made over 16 months. No emission was seen from either the 1612 or 1720 MHz satellite line transitions (to a typical five sigma upper limit of 0.2 Jy). The peak flux densities of the 1665 and 1667 MHz emission were observed to vary at a level of $\\sim$20\\% (with the exception of one epoch which dropped by $\\le$40\\%). There was no distinct flaring activity at any epoch, but there was a weak indication of periodic variability, with a period and phase of minimum emission similar to that of methanol. There is no significant variation in the polarised properties of the hydroxyl, with Stokes $Q$ and $U$ flux densities varying in accord with the Stokes $I$ intensity (linear polarisation, $P$, varying by $\\le$20\\%) and the right and left circularly polarised components varying by $\\le$33\\% at 1665-MHz and $\\le$38\\% at 1667-MHz. These observations are the first monitoring observations of the hydroxyl maser emission from G12.889+0.489. ", "introduction": "A site of hydroxyl maser emission, G12.889+0.489, initially found in a search towards IRAS 18089$-$1732 \\citep{cohen88}, was found to also be a prominent methanol maser at 6.7 GHz \\citep{menten91}. Variability of the methanol was identified by \\citet{caswell95d} and led to the extensive monitoring of \\citet{goedhart04} and the discovery of periodicity \\citep{goedhart09}. 6.7-GHz methanol masers exclusively trace the formation of high-mass stars and are hence closely studied to gain insight into the largely unknown processes of high mass star formation. The discovery of periodicity in the methanol maser emission indicates a periodic variation in the inputs of the maser emission process which has significant implications both for the mechanism of maser emission and the nature of high-mass star formation, indicating periodic processes such as the interaction of winds from binary pre-main sequence high mass stars \\citep[e.g.][]{walt09}. \\citet{goedhart09} found intensity variations in the methanol emission with a period of nearly 30 days that appeared stable, although variations from cycle to cycle in the peak amplitude of the flare were apparent. The flaring features peaked anywhere within an 11-day window, but the phase of the minima was stable. An amplitude variation with the same period is present in features at different velocities and at both 6.7 GHz and 12.2-GHz transitions. Delays of up to six days were found between individual 6.7-GHz features and a one day delay was found between the 12.2-GHz and 6.7-GHz flares at the same velocities. Hydroxyl observations since the discovery of the maser by \\citet{cohen88} have included the positioning observations of \\citet{caswell98} and \\citet{argon00}, and a detailed polarization study by \\citet{szymczak09}. Apparent variations in total intensity exceed 20\\%, but intervals between observations were typically longer than one year and precise comparison is not possible with the significantly different instrumentation. Observations with the Parkes radio telescope (Caswell et al. in prep.) measured the hydroxyl emission in 2004 and 2005, finding peak flux densities of 8 Jy at 1665 MHz and 1.8 Jy at 1667 MHz. The Parkes spectra are in agreement at the two epochs to within 10\\% and similar values were obtained by \\citet{szymczak09} in observations made in 2003. Although the amplitudes have shown some variation, the velocities of the features are consistent over all previous observations. The majority ($\\sim$80\\%) of hydroxyl masers in regions of high-mass star formation have associated methanol maser emission \\citep{caswell98}, with many sites exhibiting a close spatial coincidence, implying a common masing gas and pumping source. Models of maser pumping suggest that both species are pumped by infrared emission from dust surrounding the high mass pre-main sequence object \\citep[e.g.][]{moore88, cragg05, gray07}. Hence there is an expectation that any variability in emission of one species will correlate with variability in the other. Early evidence for hydroxyl maser variability \\citep[e.g.][]{robinson70} was followed by several more focused variability studies \\citep[e.g.][]{sullivan76, clegg91}, but these showed no clear periodicity of hydroxyl masers in star forming regions. More recent investigations of possible common flaring behaviour between the methanol and the ground-state hydroxyl masers has been inconclusive \\citep[e.g.][]{macLeod96}, but correlated variability of excited-state hydroxyl with methanol \\citep{almarzouk12} and formaldehyde with methanol \\citep{araya10} has been found recently. With methanol emission from G12.889+0.489 established to be periodic in nature, we were motivated to initiate monitoring observations of the associated hydroxyl emission. ", "conclusions": "As stated in the introduction, the methanol maser counterpart of G12.889+0.489 has quasi-periodic flaring at both 6.7 GHz and 12.2 GHz \\citep{goedhart09}. There are time delays between the two frequencies and the flaring period occurs anywhere within an 11 day window, but minima are regular (with a period of 29.5 days). This periodic variability was ascribed by the authors as likely being due either to variation in the radio continuum emission (the seed photons for the maser emission) or variation in the infrared emission (pumping the maser). We note that, extrapolating from the work of \\citet{goedhart09}, our range of epochs should have covered flaring activity and minima (three minima should occur in the first session of observations and the second session should be wholly between two adjacent minima). Goedhart et al. show that the methanol exhibits flux density variations of 60 to 70\\%. With the exception of one epoch, the hydroxyl emission we measure varies by only 20\\%. The current data have no strong indication of periodicity, but do suggest at low significance variation with a comparable period to the methanol, with the minimum at MJD 55379 nominally agreeing with the extrapolated minima epoch of the methanol. Further observing epochs will be required to make firm conclusions on the periodicity of variation. The polarisation properties we observe are comparable to the work of \\citet{szymczak09}, who measured the full polarisation properties at 1665 and 1667 MHz with all four Stokes parameters with the Nancay telescope in 2003. They found linear polarization of 87.3\\% and 76.7\\% in two 1665-MHz features at 31.59 (the kinematic and spatial outlier mentioned previously) and 32.76 km\\,s$^{-1}$ respectively, and both features had low circular polarisation at the epoch of the observations (14.6\\% and 8.9\\% respectively). If the persistently strong features, 1665-MHz RHCP at 35.3\\,km\\,s$^{-1}$ and the 1665-MHz LHCP at 33.1\\,km\\,s$^{-1}$, are a Zeeman pair, the implied magnetic field strength is +3.7 mG. This is corroborated with the 1667-MHz emission where the brightest RHCP at 34.8\\,km\\,s$^{-1}$ and the brightest LHCP feature at 33.4\\,km\\,s$^{-1}$, also give an implied magnetic field strength of +3.7 mG. A qualitative explanation for some of the methanol maser sources with long periods (e.g. 9.62+0.20 and 188.95+0.89) is that of a colliding wind binary system \\citep{walt09,walt11}, a system which could periodically alter either the seed flux or pumping mechanism of the maser emission (although for the example sources in the work of van der Walt et al. the seed flux is the more likely in view of their long periods). The hydrogen winds of the binary pre-main sequence stars either heat the circumstellar dust or cause additional ionisation of hydrogen surrounding the forming high mass star. In the case of G12.889+0.489, an upper limit to continuum emission of 1 mJy has been established \\citep{walsh98}, but if there is a very weak region of ionised hydrogen providing the seed photons, the possible offset location of the hydroxyl relative to the methanol (by $\\sim$1500 AU) may place the hydroxyl sufficiently far away so as to be minimally affected, explaining the lower prominence of periodicity in the hydroxyl emission. On the other hand the optical depth at the lower frequency of the hydroxyl emission could be such that the seed radiation does not respond as it does at the methanol frequency. An aspect of this model is that we may expect to see rotation of polarisation angle of features with the passing of shocks associated with the colliding winds (and the flares of maser emission). As mentioned in the results we have one feature at 31.6\\,km\\,s$^{-1}$ with high linear polarisation, but this does not show significant variation in polarisation angle (variation is within the noise). An alternative theory to account for periodicity in masers has been put forward by \\citet{araya10}. This was proposed to explain the periodicity in formaldehyde and methanol maser emission in IRAS 18566+0408 through circumbinary disk accretion. In this model the accreting material heats the dust, increasing the photons pumping the maser. In this variety of model, an offset location could also be sufficient to diminish the effect on the hydroxyl emission. An important aspect of G12.889+0.489 is the shortness of the periodicity (29.5 days), which any model for the system must account for. Additionally, the suggestion that the minima of hydroxyl emission may coincide with the minima of the methanol implies that the periodicity may arise from a quenching or suppression mechanism, rather than flaring. Continued monitoring of this maser source across the various maser transitions is required to provide further insight." }, "1207/1207.0934_arXiv.txt": { "abstract": "s{ We develop a numerical statistical method to study linear cosmological fluctuations in inflationary scenarios with multiple fields, and apply it to an ensemble of six-field inflection point models in string theory. The latter are concrete microphysical realizations of quasi-single-field inflation, in which scalar masses are of order the Hubble parameter and an adiabatic limit is reached before the end of inflation. We find that slow-roll violations, bending trajectories and ``many-field'' effects are commonplace in realizations yielding more than 60 e-folds of inflation, although models in which these effects are substantial are in tension with observational constraints on the tilt of the scalar power spectrum.} ", "introduction": " ", "conclusions": "We have developed a statistical approach to the study of inflationary scenarios with multiple fields: it consists in comparing, in a large ensemble of realizations of a class of inflationary models, the {\\sf exact} dynamics of linear fluctuations to three approximate descriptions: {\\sf slow-roll}, {\\sf single-field}, and {\\sf two-field}, enabling us to characterize to which extent these simpler effective models are able to describe the physics of our class of models. We have applied this method to inflection point inflationary models with multiple fields of masses of order the Hubble parameter, as expected in generic low-energy effective field theories. We have observed that this mass spectrum implies that cosmological perturbations dynamically reach an ``adiabatic limit'' by the end of inflation, thereby alleviating the need for a precise description of reheating. We found that slow-roll violations and strongly bending trajectories are common in realizations yielding more than 60 e-folds of inflation, although the majority of models with an observably acceptable spectral index turns out to be effectively slow roll single-field. We think however that this pessimistic conclusion regarding the likelihood of models displaying large multifield effects and consistent with observations should be tempered, as it might be particular to inflection point inflation, and as we found that multifield effects most often tend to significantly redden the spectrum. We also demonstrated the existence of a critical threshold of turning to generate a given amplitude of multifield effects. Eventually, we pointed out the generic importance on cosmological perturbations of \\textit{many} (beyond 2)-field effects, implying that the prevailing trend to study 2-field models of inflation may well be misleading to unveil the true nature of inflation with multiple fields. More generally, while a great deal of attention has been recently devoted to the study of the primordial non-Gaussianities generated during inflation, we think our study shows that many aspects of the \\textit{linear} dynamics of cosmological perturbations during inflation with multiple fields remain to be explored, and we hope our work and our methodology will pave the way for further developments in this direction." }, "1207/1207.5516_arXiv.txt": { "abstract": "We present the first results of our spectroscopic follow-up of $6.55$-$\\sigma$ significance using a 16hrs long exposure with FORS2 VLT. Based on the absence of flux in bluer broad-band filters, the blue color of the source, and the absence of additional lines, we identify the line as Lyman-$\\alpha$ at $z=6.740\\pm0.003$. The integrated line flux is $f=(0.7\\pm0.1\\pm0.3){\\times}10^{-17}\\lunit$ (the uncertainties are due to random and flux calibration errors, respectively) making it the faintest Lyman-$\\alpha$ flux detected at these redshifts. Given the magnification of $\\mu=3.0\\pm0.2$ the intrinsic (corrected for lensing) flux is $f^{\\rm int}=(0.23\\pm0.03\\pm0.10\\pm0.02){\\times}10^{-17}\\lunit$ (additional uncertainty due to magnification), which is ${\\sim}2-3$ times fainter than other such measurements in $z{\\sim}7$ galaxies. The intrinsic {$\\H$}-band magnitude of the object is $m^{\\rm int}_{H_{\\rm 160W}}=27.57\\pm0.17$, corresponding to $0.5L^*$ for LBGs at these redshifts. The galaxy is one of the two sub-$L^*$ LBG galaxies spectroscopically confirmed at these high redshifts (the other is also a lensed $z=7.045$ galaxy), making it a valuable probe for the neutral hydrogen fraction in the early Universe. ", "introduction": "\\label{sec:intro} The epoch of reionization, which marks the end of the ``Dark Ages'' and the transformation of the universe from opaque to transparent, is poorly understood. It is thought that $z > 6$ faint proto-galaxies were responsible for this transformation, but recent observations of $z\\gtrsim 7$ objects (e.g., \\citealp{robertson10} for a review) complicate that scenario. Finding robust samples of sources, representative of the population contributing a significant amount of energetic photons, is crucial. Wide Field Camera 3 (WFC3) on HST enables a preliminary identification of such galaxies. Substantial progress has been made in detecting $z \\gtrsim 7$ galaxies using the dropout technique \\citep{steidel96}, both in blank fields (HUDF, Candels, e.g., \\citealp{bouwens12, oesch12,finkelstein12,mclure11}), and behind galaxy clusters (e.g., \\citealp{kneib04,egami05,bradley12,richard11,zheng12,zitrin12} and references therein). One of the most obvious limitations of the dropout technique, however, is that unambiguously confirming the object's redshift usually requires spectroscopic follow-up. This is hard to do for typically faint high-z sources, and it is thus an area where gravitational lensing magnification helps greatly, as demonstrated in this paper. In addition to the redshift confirmation, spectroscopy provides information on properties of the interstellar and intergalactic media (ISM and IGM). In particular, Lyman-$\\alpha$ emission from sources close to the reionization era is a valuable diagnostic given that it is easily erased by neutral gas within and around galaxies. Its observed strength in distant galaxies is a gauge of the time when reionization was completed \\citep{robertson10}. Furthermore, we expect Lyman-$\\alpha$ Emitters (LAEs) to be predominantly dust-free galaxies; hence their numbers should increase with redshift until the state of the IGM becomes neutral, at which point their numbers should decline. Significant progress has been made in detecting Lyman-$\\alpha$ emitters (LAEs) in narrow band and spectroscopic surveys at $z\\gtrsim 6$ (e.g., \\citealp{kashikawa06,rhoads12,schenker12,clement12, curtislake12,ono12, stark11, pentericci11}). Most studies see a decline in the LAE population at $z>7$, but not all do \\citep{krug12, tilvi10}. The declining fraction of LAEs within the LBG population \\citep{stark10,kashikawa11,pentericci11} is consistent with this decline being due to changes in the ISM/IGM, specifically to an increased amount of neutral gas. However, current studies only probe the bright end of the luminosity function of LBGs. Furthermore, as noted by \\citet{dijkstra11}, and \\citet{dayal12}, measuring the rest frame Equivalent Width (EW) distribution of LAEs as a function of redshift {\\it and} luminosity is a powerful tool to study reionization. The EW distribution changes with redshift and source luminosity. Simulations suggest that reionization is the key factor driving this trend \\citep{dayal12}, because unlike continuum photons, the Lyman-$\\alpha$ photons that escape the galactic environment are attenuated by the \\ion{H}{1} in the IGM. With a measurement of the EW distribution in LAEs we can therefore help distinguish between effects of ISM dust and neutral IGM and study the epoch of reionization (see also \\citealp{treu12}). The main missing observational ingredient is a measurement of the EW distribution for both luminous and sub-$L^*$ galaxies at the redshifts of reionization and this can only be achieved with spectroscopy. Current spectroscopic observations unfortunately fall short of matching the extremely deep near-IR {HST/WFC3} data for a significant sample of $z\\gtrsim 7$ dropout selected galaxies. Even with state of the art facilities (e.g. the new spectrograph MOSFIRE on Keck, \\citealp{mclean10}) this will be a challenge. While samples of $\\lesssim L^*$ galaxies at $z<6.5$ is steadily increasing (e.g.,\\citealp{richard11, schenker12,labbe10} for spectroscopic and imaging detections), to date, very few ${\\lesssim}L^*$ galaxies at $z \\gtrsim 6.5$ are spectroscopically confirmed. The only examples are a lensed $z=7.045$ galaxy and a marginal detection at $z=6.905$ by \\citet{schenker12}. At ${\\sim}L^*$ a $z=6.944$ galaxy was detected by \\citet{rhoads12}. At $z{\\sim}8$ \\citet{lehnert10} report a marginal detection of an emission line, but independent observations do not detect it \\citep[][in prep.]{bunker12}. Other surveys \\citep{shibuya12, ono12,vanzella11,pentericci11,fontana10} target mostly brighter sources. It is important to increase the sample at $z>6.5$ and compare it to $z<6.5$ because the timescale for changing the number of LAEs and their observed EW distribution close to the reionization epoch is shorter than the interval of cosmic time between $z{\\simeq}6$ and $z{\\simeq}7$ \\citep{dayal12}. A powerful way to detect emission lines from faint sources is to use galaxy clusters as cosmic telescopes (e.g,\\citealp{treu10} for a recent review). Gravitational lensing magnifies solid angles while preserving colors and surface brightness. Thus, sources appear brighter than in the absence of lensing. The advantages of cosmic telescopes are that we can probe deeper (due to magnification), sources are practically always enlarged, and identification is further eased if sources are multiply imaged. Typically, one can gain several magnitudes of magnification, thus enabling the study of intrinsically lower-luminosity galaxies that we would otherwise not be able to detect with even the largest telescopes. Indeed the highest redshift sub-$L^*$ LBG currently spectroscopically confirmed is the $z=7.045$ galaxy lensed by a cluster A1703 \\citep{schenker12}. Observations using galaxy clusters as cosmic telescopes are consistently delivering record holders in the search for the highest redshift galaxies \\citep{kneib04,bradley08,zheng12}. For this reason we have started a large campaign of spectroscopic follow-up of $z > 6.5$ candidates behind the best cosmic telescopes. In this paper we present the first spectroscopic confirmation from this campaign: a $z=6.740\\pm0.003$ galaxy behind the Bullet Cluster. This paper is structured as follows. In Section~\\ref{sec:data} we describe the data acquisition and reduction, in Section~\\ref{sec:results} we present the spectrum of the galaxy and we summarize our conclusions in Section~\\ref{sec:conclusions}. Throughout the paper we assume a $\\Lambda$CDM cosmology with $\\Omega_{\\rm m}=0.3$, $\\Omega_{\\Lambda}=0.7$, and Hubble constant $H_0=70{\\rm\\ kms^{-1}\\:\\mbox{Mpc}^{-1}}$. Coordinates are given for the epoch J2000.0, magnitudes are in the AB system. ", "conclusions": "\\label{sec:conclusions} We have presented deep VLT spectroscopy of a strongly lensed {$\\z$}-band dropout galaxy behind the Bullet cluster. We detected an emission line at $9412\\mbox{\\AA}$ with $>5\\mbox{-}\\sigma$ significance, which we identify as Lyman-$\\alpha$ at $z=6.740\\pm0.003$ at $>99$\\%CL. Correcting for magnification \\citep[by a factor of $\\mu=3.0\\pm0.2$ as discussed by][]{hall12,bradac09}, the intrinsic (unlensed) line flux is $f=(0.23\\pm0.03\\pm0.10\\pm0.02){\\times}10^{-17}\\lunit$ (Table~\\ref{tab:drops}), which is ${\\sim}2-3$ times fainter than the faintest spectroscopic detection of an LAE at $z{\\sim}7$ \\citep{schenker12}. Its intrinsic {$\\H$}-band magnitude is $m^{\\rm{int}}_{H_{\\rm{160W}}}=27.57\\pm0.17$, corresponding to an intrinsic luminosity of $0.5L^*$ (where $L^*$ was calculated from the best fit LBG luminosity function from \\citealp{bouwens11}). The source is undetected in the four IRAC bands, which is not surprising given that we would only be able to detect extremely red galaxies with $m_{H_{\\rm{160W}}}-m_{{\\chone}}=4$ at 3-$\\sigma$ in $\\chone$ for sources this faint. For comparison, $z=6.027$ source behind A383 \\citep{richard11} with an unusually mature stellar population ${\\sim}800 \\mbox{Myr}$ and a multiply-imaged $z=6.2$ object \\citep{zitrin12} with a younger age $\\sim180\\mbox{Myr}$ have much bluer colors $m_{H_{\\rm{160W}}}-m_{\\chone}{\\sim}1.5$. Deeper Spitzer data will be needed to probe the presence of mature stellar populations in the galaxy we present here and other systems at high redshift. While this work presents only a single spectroscopic detection at $z>6.5$, it nonetheless probes a very important region of parameter space. As noted above, measuring the EW distribution of LAEs as a function of redshift {\\it and} luminosity is a very powerful tool to study reionization, because the latter is likely the key factor driving the trend of EW in luminosity \\citep{dayal12}. The main missing observational ingredient is a measurement of the EW distribution for both luminous and sub-$L^*$ galaxies at the redshifts of reionization. Our source is the faintest one (in line flux) detected thus far and is only the second firm spectroscopic detection of a sub-$L^{*}$ source at $z>6.5$. With future observations of dropouts magnified by cosmic telescopes we plan to further increase this sample. Once completed, this survey will help constrain the duration and physical processes occurring at the epoch of reionization." }, "1207/1207.3795_arXiv.txt": { "abstract": "We present \\textit{Hubble Space Telescope} (\\HST) imaging and spectroscopy of the gravitational lens SL2SJ02176-0513, a cusp arc at $z=1.847$. The UV continuum of the lensed galaxy is very blue, which is seemingly at odds with its redder optical colors. The 3D-HST WFC3/G141 near-infrared spectrum of the lens reveals the source of this discrepancy to be extremely strong [\\ion{O}{3}]$\\lambda$5007 and H$\\beta$ emission lines with rest-frame equivalent widths of $2000\\pm100$ and $520\\pm40$~\\AA, respectively. The source has a stellar mass $\\sim10^8~M_\\odot$, $sSFR\\sim100/\\mathrm{Gyr}$, and detection of [\\ion{O}{3}]$\\lambda$4363 yields a metallicity of $\\logOH=7.5\\pm0.2$. We identify local blue compact dwarf analogs to \\lensID, which are among the most metal-poor galaxies in the SDSS. The local analogs resemble the lensed galaxy in many ways, including UV/optical SED, spatial morphology and emission line equivalent widths and ratios. Common to \\lensID\\ and its local counterparts is an upturn at mid-IR wavelengths likely arising from hot dust heated by starbursts. The emission lines of \\lensID\\ are spatially resolved owing to the combination of the lens and the high spatial resolution of \\HST. The lensed galaxy is composed of two clumps with combined size $r_e\\sim$300~pc, and we resolve significant differences in UV color and emission line equivalent width between them. Though it has characteristics occasionally attributed to active galactic nuclei, we conclude that \\lensID\\ is a low-metallicity star-bursting dwarf galaxy. Such galaxies will be found in significant numbers in the full 3D-HST grism survey. ", "introduction": "\\label{s:introduction} Tracing galaxy populations over cosmic time is key to understanding their evolutionary processes. Observations of a particular galaxy type over a range of redshifts yield complementary information, from spatially-resolved, high signal-to-noise studies of individual local examples to more distant statistical samples \\citep[e.g.,][]{overzier:09}. While local low-mass dwarfs ($\\lesssim10^9~M_\\odot$) can be studied relatively easily, more distant examples are usually missing in typical flux-, mass-, or star-formation-rate-limited surveys. The best-studied examples at cosmological distances are usually strongly magnified by gravitational lenses \\citep[e.g.,][]{fosbury:03, yuan:09, ewuyts:12}. The lowest mass galaxies tend to have the lowest metallicities \\citep{tremonti:04}, and the normalization of this mass-metallicity relation evolves such that galaxies at a given stellar mass have lower metallicities at higher redshifts to at least $z\\sim3$ \\citep{erb:06, mannucci:09}. The physical conditions of low-metallicity star-forming galaxies can be very different than their higher-mass counterparts, characterized by a hard radiation field and strong stellar winds \\citep[e.g.,][]{erb:10}. Consequently they may have strong emission lines, which can be used to select them \\citep{atek:11, vanderwel:11} and also must be accounted for when modeling their broad-band photometry \\citep{atek:11, finkelstein:11}. Active galactic nuclei (AGN) are also sources of hard ionizing radiation that can result in similar spectral features, so accurately separating the influence of AGN and metallicity/star formation is crucial for understanding the dominant processes that shape galaxies at $z>1$ \\citep[e.g.,][]{trump:11}. In this Letter, we present \\HST\\ observations of a gravitational lens system discovered by the Strong Lensing Legacy Survey consisting of a bright cusp arc at $z=1.8470$ that is lensed by a massive galaxy at $z=0.6459$ \\citep[SL2SJ02176-0513,][]{tu:09}. The system lies within the CANDELS imaging \\citep{grogin:11, koekemoer:11} and 3D-HST spectroscopic \\citep{brammer:3dhst} surveys of the UKIDSS/UDS field. We use the unique combination of the \\HST\\ datasets and the natural lens to demonstrate that the lensed galaxy is a low-mass, low-metallicity galaxy undergoing an extreme starburst, and that it shares many characteristics of local low-metallicity blue compact dwarf galaxies. We adopt cosmological parameters $h=0.7$, $\\Omega_m=0.3$, and $\\Omega_\\Lambda=0.7$ throughout. ", "conclusions": "\\label{s:discussion_and_summary} We take advantage of the unique combination of a natural gravitational lens and high spatial resolution \\HST\\ imaging and near-IR spectroscopy to extract the detailed properties of \\lensID. The near-IR spectrum is dominated by extremely strong emission lines of H$\\gamma$, H$\\beta$ and [\\ion{O}{3}] at $z=1.847$. From the UV/optical spectrum and photometry, we determine that the source of \\lensID\\ is a young starbursting ($sSFR\\sim100/\\mathrm{Gyr}$) dwarf galaxy ($M\\sim1.3\\times10^{8}~M_\\odot$) with an extremely low gas-phase metallicity ($\\logOH\\sim7.5$). Even with so few metals, \\lensID\\ shows detectable hot dust emission observed at $\\lambda_\\mathrm{obs}=24~\\micronm$. We find the unique properties of \\lensID\\ to be remarkably similar to those of nearby low-metallicity blue compact dwarf galaxies selected to have similar [\\ion{O}{3}] equivalent widths and ([\\ion{O}{3}]/H$\\beta$) line ratios. Such galaxies are frequently considered to be local analogs to galaxies expected to be more common at earlier cosmic times and in \\lensID\\ we have discovered a compelling connection at high redshift (which itself is likely a bright, magnified example of the galaxies discovered by \\citealp{vanderwel:11}). \\lensID\\ has characteristics that are frequently attributed to AGN at similar redshifts: extreme equivalent widths and line ratios of the optical emission lines, high-ionization UV lines, and the presence of a strong IR excess from heated dust. We demonstrate that all of these properties are largely consistent with a hard ionization field produced by a compact, low metallicity starburst (see also examples from \\citealp{fosbury:03, erb:10} and discussion by \\citealp{hunt:10}). Further evidence comes from the fact that the lens resolves two line-emitting components of \\lensID, though extended narrow-line regions excited by nuclear activity have been observed \\citep[e.g.,][]{unger:87}. Both star-formation and AGN reach a peak in their activity at $z\\sim2$ so robust identification and separation of the two contributions is critical for understanding their effect on galaxy evolution. We conclude by noting that we will obtain a substantial sample of (unlensed) galaxies with similar, if not quite as extreme, properties to \\lensID\\ at $1.3 < z < 2.2$ (covering [\\ion{O}{3}]+H$\\beta$) in the full 3D-HST survey. Without the factor of $\\sim$25 lens magnification, such an object would have $m_\\mathrm{F140W}=25.2$, where 3D-HST is sensitive to line equivalent widths $\\gtrsim$1000~\\AA\\ \\citep{brammer:3dhst}. Indeed many such galaxies have recently been found with WFC3 \\citep{atek:11, vanderwel:11}. While individual unlensed objects will not allow such detailed study as performed here, statistical samples of these galaxies will offer insights into the metallicity and star-formation properties of the low-mass building blocks of galaxies observed today." }, "1207/1207.1874_arXiv.txt": { "abstract": "{ The physics of radio emission from cosmic-ray induced air showers is shortly summarized. It will be shown that the radio signal at different distances from the shower axis provides complementary information on the longitudinal shower evolution, in particular the early part, and on the distribution of the electrons in the shower core. This complements the information obtained from surface, fluorescence, and muon detectors and is very useful in getting a comprehensive picture of an air shower. } % ", "introduction": "\\label{intro} There are several approaches that are followed to model radio emission from extensive air showers. These can crudely be separated into two categories, Microscopic and Macroscopic. In a Microscopic approach the tracks of the individual electrons are followed and the emitted radiation of each is summed to yield the total radiation of the extensive air shower. This approach is incorporated in the REAS~\\cite{REAS} and the ZHAireS~\\cite{ZHaireS} codes. In the alternative, Macroscopic, approach the velocity distribution of the electrons is summed locally to yield the macroscopic charge and current density distribution in the shower. From this four-current distributions the (relativistic) Maxwell equations~\\cite{Jac-CE} are used to generate the radiated fields. This approach is followed in the MGMR~\\cite{MGMR,Wer08} and EVA~\\cite{EVA} calculations. Up to a few years ago the two approaches showed large inconsistencies which were shown to be mainly due to a rather subtle term that was missing in the Microscopic calculations. By correcting this, consistency could be reached~\\cite{Hue11} which is a major achievement and indicates that, based on first principles, we basically understand the emission process. It also shows that a dual approach is necessary to be able to trust the results. It should be noted that for coherent radiation the two approaches should give the same results, they differ in the incoherent contribution to the radiation field. The main advantage of a macroscopic approach is that it clearly indicates the aspects of the shower dynamics that are responsible for certain features of the detected radio signal. The macroscopic approach requires a parametrization of the position and time dependence of the four-currents in the shower which can be performed at various levels of sophistication. In the following we use the Macroscopic approach. Since radio emission is basically governed by wave mechanics any features in the frequency spectrum of the pulse are directly related to critical length scales in the distribution of the particles in the shower. ", "conclusions": "" }, "1207/1207.4456_arXiv.txt": { "abstract": "The \\msig~relation has been studied extensively for local galaxies, but to date there have been scarce few direct measurements of stellar velocity dispersions for systems beyond the local universe. We investigate black hole and host galaxy properties of six ``post-starburst quasars'' at $z\\sim0.3$. Spectra of these objects simultaneously display features from the active nucleus including broad emission lines and a host galaxy Balmer absorption series indicative of the post-starburst stellar population. These are the first measurements of \\ssig~in such objects, and we significantly increase the number of directly-measured non-local objects on the \\msig~diagram. The ``post-starburst quasars'' of our sample fall on or above the locally defined \\msig~relation, a result that is consistent with previous \\msig~studies of samples at $z>0.1$. However, they are generally consistent with the \\mlum~relation. Futhermore, their location on the Faber-Jackson relation suggests that some of the bulges may be dynamically peculiar. \\\\ ", "introduction": "The \\msig~relation \\citep{FM00,Gebhardt00,Gultekin09,Woo10} has been extensively studied in both active and inactive galaxies at low redshifts. The existence of this and other black hole - host galaxy scaling relations \\citep{Magorrian98, McLureDunlop02} implies that the growth of black holes is intimately connected with that of the host galaxies. However, the mechanism by which this growth is regulated is still an open matter. One proposed mechanism is through black hole accretion feedback. \\citet{DiMatteo05} show that only a modest amount of the accretion energy is required to shut down star formation in the host galaxy. Their simulations of galaxy mergers reproduce the local \\msig~relation with only $\\sim5$\\% of the AGN luminosity being thermodynamically coupled to the surrounding gas. On the other hand, the need for accretion feedback is not clear. \\citet{Peng07} argues that mergers between galaxies (major and minor) naturally leads to a tight log-linear M$_{\\rm BH}$-M$_{\\rm bulge}$ relation. To investigate the co-evolution of black holes with their hosts, it is necessary to measure the relation at higher redshifts. \\citet{Robertson06} study the theoretical evolution in the \\msig~relation by simulating galactic mergers and including both supernova and black hole feedback. They predict that the slope of the relation remains roughly constant out to redshift $z=6$, while the scaling of the relation slightly decreases due to an evolving Faber-Jackson relation. The effect of this would be that $\\sigma_{*}$ increases for a given M$_{\\rm BH}$ as a function of redshift. Conversely, \\citet{Croton06} predicts that the M$_{\\rm BH}$-M$_{\\rm bulge}$ relation evolves in the opposite sense, \\ie~black hole mass increases with redshift for a given bulge mass. Observational investigations into the evolution of the \\msig~relation face difficulties. The sphere of influence of the black hole cannot be resolved at high redshifts, thus the methods using stellar or gas dynamics \\citep{KR95} to measure M$_{\\rm BH}$ in the local universe are unavailable. Alternatively, one can use the virial method to estimate black hole masses, which invokes AGN broad line widths and the R$_{\\rm BLR}$-L$_{5100}$ relation \\citep{Kaspi00,Bentz09a}. However, in type-1 active galactic nuclei (AGN), the active nucleus often outshines the host, drowning out stellar features in the spectra. Despite this effect, \\citet{Woo06, Woo08} studied a sample of Seyfert galaxies and found that black holes at $z=0.36$ and $z=0.57$ were over-massive for their given stellar velocity dispersions compared to the local relation. \\citet{Canalizo12} avoid the problem of an over-powering AGN continuum by studying dust-reddened quasars between $0.14 < z < 0.37$, which have spectra that show both active nucleus and host galaxy features. They find a result similar to Woo \\etal~in that reddened quasars fall above the local \\msig~relation. The red quasars of Canalizo \\etal~also fall above the \\mlum~relation, supporting the notion that the black holes are over-massive compared to their bulges. Other recent studies have discussed regimes in which the local inactive galaxy relation may not hold. \\citet{Graham11} calibrate the \\msig~relation using elliptical, barred and unbarred galaxies and find that the slope of the \\msig~relation can vary according to morphology. Conversely, \\citet{Beifiori11} and \\citet{Vika11} find no correlation between black hole mass and S\\'{e}rsic index, {\\it n}. At the low-mass end of the \\msig~relation, \\citet{Xiao11} find little difference in the relation as defined by barred and unbarred Seyfert 1 galaxies. Xiao \\etal~also state that inclination angle of a galactic disk components may increase the scatter of the \\msig~relation. This is consistent with the results of \\citet{Bennert11}, who find that \\ssig~can be biased to larger or smaller values by the disk component of the galaxy. Recently, \\citet{McConnell11b, McConnell11a} showed that brightest cluster galaxies have black holes that are significantly more massive than predicted from the \\msig~relation given their velocity dispersions. In this paper we exhibit the first results of our investigation of the \\msig~relation using a sample of AGN whose host galaxies contain luminous post-starburst stellar populations between redshifts $0.2 < z < 0.4$. \\citet[][(in preparation)]{Brotherton12} catalog a subset of quasars with strong Balmer absorption series visible in their spectra. These ``post-starburst quasars'' (PSQs) are a convenient case for the study of the \\msig~relation at higher$-z$, because spectra of these objects simultaneously show the broad emission lines distinctive of an active black hole as well as strong stellar absorption lines indicative of a post-starburst stellar population in the host galaxies. The moniker ``post-starburst quasars'' is not a definitive classification of the objects' luminosities. While some are luminous enough to be considered quasars, others are not. To be consistent with previous work \\citep{Cales11, Brotherton10}, we will continue to refer to the sample as ``post-starburst quasars''. In this paper we present the stellar velocity dispersions ($\\sigma_{*}$) and black hole masses (M$_{\\rm BH}$) of six post-starburst quasars as measured from spectroscopy performed with the Low Resolution Imaging Spectrograph (LRIS) on Keck I. Each of these PSQs were observed with $HST$ ACS F606W as part of a snapshot program (10588, PI Brotherton, M.). These images and a detailed morphological study of the host galaxies are presented in \\citet{Cales11}. In our analysis we rely on photometry reported by Cales \\etal, and we adopt the same cosmology that they use: a flat universe with H$_{\\rm o} = 73$ km s$^{-1}$, $\\Omega_{\\rm M} = 0.27$, and $\\Omega_{\\Lambda} = 0.73$. The paper is organized as follows: we discuss the sample and observations in section 2, describe our spectral fitting code and $\\sigma_{*}$ measurements in section 3, and discuss virial black hole mass calculations in section 4. We present our main results and discussion of possible biases in sections 5 and 6, respectively. ", "conclusions": "We have measured \\mbh~and \\ssig~in six AGN that are hosted by galaxies showing post-starburst stellar populations (``post-starburst quasars''). These objects fall on or above the locally defined \\msig~and \\mlum~relations. The location of the PSQs on the \\msig~diagram is consistent with other studies of objects at $z>0.1$ \\citep{Canalizo12,Woo08,Woo06}. We have used the Balmer region, which includes the Balmer series, Ca lines, the G band, and other stellar features to estimate the stellar velocity dispersions. While we have used host galaxy templates that approximate the stellar population of the host galaxies, it is possible that the young and old populations have different velocity dispersions. The kinematics of the young population are influenced by the mechanism that triggered the starburst. In a theoretical work, \\citet{Bekki05} describe various models of galaxy interactions (major/minor mergers, strong tidal interactions, etc.) that lead to E+A spectral signatures. They find a range of morphologies, velocity dispersions, and starburst mass fractions resulting from the merger event depending on the specific model used. The kinematic properties of the resulting post-starburst population depend on the initial parameters of the galaxy-galaxy interaction that induces the starburst. The morphologies presented by \\citet{Cales11} for our sample indicate that these objects have only minorly disturbed features, if any, that could indicate a recent merger event. Four objects of our sample are described by Cales \\etal~as containing a disk component. However, the overall sample of PSQs display a range of morphologies, and Cales \\etal~found that half of the observed sample appear to be post-merger remnants. The variety of morphologies seen in the parent sample is consistent with the simulations conducted by \\citet{Bekki05}, indicating a range of possible galaxy-galaxy interactions that trigger the starburst event (\\eg~tidal interactions, minor/major mergers). Regardless of the triggering mechanism, the PSQs, at least in this subsample, have \\mbh~and \\ssig~that are generally consistent with previously studied $z>0.1$ galaxies. The origin of the velocity dispersion (pressure supported versus rotational) should influence whether the objects fall on the \\msig~relation, which is best defined for pressure supported bulges. Observationally, \\citet{Norton01} studied the young and old populations independently in E+A galaxies and found them to be largely pressure supported, with the younger post-starburst population having higher \\ssig~than the old population. However, \\citet{Bekki05} had trouble reproducing this result with their models. By contrast, \\citet{Pracy09} found rotation to play a dominant role in their studied E+As. We do not separate the young and old populations in our analysis, so the kinematical differences between the post-starburst population and older population in PSQs remains a question. In principle, we could adopt an additional free parameter in our model to measure the velocity dispersion of the old population separately from the young population. However, this would introduce degeneracies that would be difficult to disentangle, especially with the presence of the active nucleus. We are currently conducting an investigation into the kinematic differences between young and old populations in quiescent E+A galaxies that lack the active nucleus continuum contamination \\citep[][in preparation]{Hiner12b}. This analysis will be similar to that of Norton \\etal, but we will include multiple wavelength regions. By including longer wavelength regions, such as the \\ion{Mg}{1b} and \\ion{Ca}{2} triplet regions, we will be able to further distinguish between the young and old stellar populations. We have also measured the bulge luminosities of the host galaxies. After plotting the \\mlum\\ and Faber-Jackson relations, we find a scenario which rules out \\ssig~bias as the source of the offset from the \\msig~relation for three of the six objects we measured. Two objects (SDSS $0030-1035$ and SDSS $2306-0100$) are significantly offset from the \\msig~relation. SDSS $2306-0100$ is additionally offset from the Faber-Jackson relation, indicating that the measured \\ssig~may be significantly influenced by the face-on disk of the host galaxy. SDSS $0030-1035$ is consistent with the relation as defined by \\citet{Nigoche10}, but it falls significantly below that of \\citet{Desroches07}, so it may also be influenced by the face-on disk. While SDSS $0237-0101$ falls above the Desroches \\etal~Faber-Jackson relation by less than 2$\\sigma$, it is also the object that falls closest to the best fit \\msig~relation. A reduction of $\\sim50$ km s$^{-1}$ in \\ssig~would bring it in line with the Desroches \\etal~fit, but drive the object left of the \\msig~relation (consistent with the other offsets). Peculiar dynamics can clearly influence the measured velocity dispersion, and with a more statistically significant sample, the general trends will become more apparent. The offset from the \\msig~relation is consistent with that found by \\citet{Canalizo12}. Similar to their sample and the samples of \\citet{Woo06} and \\citet{Woo08}, the black holes measured here all have masses $\\ge 10^8$ M$_{\\odot}$. This brings up the question of sample selection at high-$z$ ($z>0.1$). \\citet{Lauer07} describe a potential bias that selects for high-mass black holes at higher-$z$, which is a product of a steep luminosity function and cosmic scatter in both the \\msig~and \\mlum~relations. \\citet{Canalizo12} also show that highly luminous AGN (log(L$_{5100}$/erg s$^{-1}$) $> 43.6$) are offset from the \\msig~relation regardless of their redshifts. All but one of the PSQs in our sample have luminosities in this range. We have presented six objects of a larger class of post-starburst quasars \\citep[][in preparation]{Brotherton12}. This is clearly not statistically significant and we may have by chance chosen objects that all fall above the \\msig~relation. To discuss PSQs as a whole, we require measurements of more objects that define a statistically significant sample. We have obtained data for additional PSQs from the same parent sample that will be exhibited in a future work \\citep[][in preparation]{Hiner12c}. In addition, it will be important to disentangle the kinematics of the younger post-starburst population from the underlying older stellar population. For this purpose, we have obtained spectroscopy of quiescent E+A galaxies. We will probe possible differences in \\ssig~kinematics between the populations using both the Balmer series and longer wavelength regions (\\eg~\\ion{Mg}{1b}, CaT)." }, "1207/1207.6515_arXiv.txt": { "abstract": "Recent results from short--baseline neutrino oscillation experiments and Cosmic Microwave Background anisotropy measurements suggest the presence of additional sterile neutrinos. In this paper we properly combine these data sets to derive bounds on the sterile neutrino masses in the 3+1 and 3+2 frameworks, finding a potentially good agreement between the two datasets. However, when galaxy clustering is included in the analysis a tension between the oscillation and cosmological data is clearly present. ", "introduction": "In recent years, the impressive experimental discoveries in two fields of investigation, namely neutrino physics and cosmic microwave background anisotropies, have revolutionized our knowledge in particle physics and cosmology. Neutrino oscillations experiments have not only firmly established that neutrino are massive and mixed particles (for reviews, see e.g. Refs. \\cite{Giunti:2007ry,Bilenky:2010zza,Xing:2011zza}), but have also provided precise measurements of the three-neutrino mixing parameters (see the recent global fits in Refs. \\cite{Tortola:2012te,Fogli:2012ua}). On the other hand, the measurements of the angular spectrum of the Cosmic Microwave Background (CMB) anisotropies (see e.g. Ref. \\cite{wmap7}) have not only fully confirmed the expectations of the standard cosmological scenario but also provided a precise determination of most of its parameters. Moreover, with the continuous experimental improvements, a clear interplay between neutrino physics and cosmology is emerging. Neutrinos are indeed a fundamental energy component in modern cosmology. A cosmological neutrino background is expected in the standard model and affects both the shape of the CMB and the formation of cosmological structures (see e.g. Ref. \\cite{Lesgourgues:2006nd}). The recent cosmological data have provided a clear evidence (more than $5$ standard deviations) for the existence of the primordial neutrino background and have strongly constrained the absolute neutrino mass scale (see e.g. Ref. \\cite{Hannestad:2007tu}). However, the measurements of CMB anisotropies made by the ACT (Atacama Cosmology Telescope) \\cite{act} and SPT (South Pole Telescope) \\cite{spt} experiments, when combined with the measurements of the Hubble constant $H_0$ and galaxy clustering data, have provided interesting hints for an {\\it extra} relativistic weakly interacting component, coined {\\it dark radiation}. Parameterizing this energy component with the effective number of neutrino species $N_{\\rm eff}$, the recent data bound it to $N_{\\rm eff}=4.08\\pm0.8$ at $95 \\%$ C.L. (see e.g. Ref. \\cite{Hou:2011ec,Archidiacono:2011gq,zahn,Hamann:2011hu}) whereas the standard prediction for only three active neutrino species is $N_{\\rm eff}=3.046$ \\cite{Mangano:2005cc}. While this result should be taken with some grain of salt, since it is derived from a combination of cosmological data and some tension does exist between the data (see e.g. Ref. \\cite{calaratra}) it is anyway interesting since a fourth, or fifth, neutrino species seems also suggested by short--baseline (SBL) oscillation experiments. The appearance and disappearance data of several SBL experiments can be explained by the mixing of the three active neutrinos with one or two additional sterile neutrinos in the so-called 3+1 and 3+2 models (see Refs. \\cite{Kopp:2011qd,Giunti:2011gz,Giunti:2011hn,Giunti:2011cp,Karagiorgi:2012kw,Donini:2012tt}). This work is aimed to determine the masses of the sterile neutrinos in 3+1 and 3+2 models using data from SBL experiments and recent cosmological data and check if the results are mutually compatible. Finally, we combine the bounds from the two different analyses to have a joint probability for the masses of sterile neutrinos. The paper is organized as follows: in Sec.~\\ref{sec:SBL} and in Sec.~\\ref{sec:ii} we present the data sets we make use of, the method we adopt to analyze them and the results we obtain regarding the SBL experiments and in the cosmological context, respectively; in Sec.~\\ref{sec:iiii} the joint analysis method and results are shown; finally we summarize our conclusions in Sec.~\\ref{sec:iiiii}. ", "conclusions": "\\label{sec:iiiii} Measuring the number and the mass of sterile neutrinos is one of the most interesting challenges both in cosmology and in neutrino physics. The existing cosmological data indicate that the energy density of the Universe may contain dark radiation composed of one or two sterile neutrinos, which may correspond to those in 3+1 or 3+2 models which have been invoked for the explanation of short--baseline neutrino oscillation anomalies. We have performed analyses of the cosmological and SBL data in the frameworks of both the 3+1 and 3+2 models. Then we have compared the results obtained with the same Bayesian method, to figure out if the indications of cosmological and SBL data are compatible. At the state of art, cosmological data are sensitive to the sum of neutrino masses, for which they give an upper limit at the scale of about 1 eV. Hence they do not allow us to resolve the degeneracy between the mass of the first and the second sterile neutrino in a 3+2 model, although in the numerical calculation we leave them as independent parameters. Instead, short--baseline neutrino oscillations have a completely different parameterization and in the 3+2 model the degeneracy between the two square mass differences $\\Delta{m}^2_{41}$ and $\\Delta{m}^2_{51}$ is broken. The results of our analysis show that the cosmological and SBL data give compatible results when the cosmological analysis takes into account only CMB data. But if the information on the matter power spectrum coming from galaxies surveys are also considered there is a tension between the sterile neutrino masses needed to have SBL neutrino oscillations and the cosmological upper limit on the sum of the masses. The combined analysis of cosmological and SBL data gives an allowed region for $m_4$ in the 3+1 scheme around 1 eV. In the 3+2 scheme, the cosmological data reduce the allowance of the second massive sterile neutrino given by SBL data, leading to a combined fit which prefers the case of only one massive sterile neutrino at the scale of about 1 eV. In conclusion, our analysis shows that cosmological data are marginally compatible with the existence of one massive sterile neutrino with a mass of about 1 eV, which can explain the anomalies observed in SBL neutrino oscillation experiments. The case of massive sterile neutrinos is less tolerated by cosmological data and in any case the second sterile neutrino must have a mass smaller than about 0.6 eV." }, "1207/1207.0808_arXiv.txt": { "abstract": "Using kilometric arrays of air Cherenkov telescopes at short wavelengths, intensity interferometry may increase the spatial resolution achieved in optical astronomy by an order of magnitude, enabling images of rapidly rotating hot stars with structures in their circumstellar disks and winds, or mapping out patterns of nonradial pulsations across stellar surfaces. Intensity interferometry (once pioneered by Hanbury Brown and Twiss) connects telescopes only electronically, and is practically insensitive to atmospheric turbulence and optical imperfections, permitting observations over long baselines and through large airmasses, also at short optical wavelengths. The required large telescopes ($\\sim$\\,10 m) with very fast detectors ($\\sim$\\,ns) are becoming available as the arrays primarily erected to measure Cherenkov light emitted in air by particle cascades initiated by energetic gamma rays. Planned facilities (e.g., CTA, {\\it{Cherenkov Telescope Array}}) envision many tens of telescopes distributed over a few square km. Digital signal handling enables very many baselines (from tens of meters to over a kilometer) to be simultaneously synthesized between many pairs of telescopes, while stars may be tracked across the sky with electronic time delays, in effect synthesizing an optical interferometer in software. Simulated observations indicate limiting magnitudes around m$_{V}$=\\,8, reaching angular resolutions $\\sim$30$\\,{\\mu}$arcsec in the violet. The signal-to-noise ratio favors high-temperature sources and emission-line structures, and is independent of the optical passband, be it a single spectral line or the broad spectral continuum. Intensity interferometry directly provides the modulus (but not phase) of any spatial frequency component of the source image; for this reason a full image reconstruction requires phase retrieval techniques. This is feasible if sufficient coverage of the interferometric $(u,v)-$plane is available, as was verified through numerical simulations. Laboratory and field experiments are in progress; test telescopes have been erected, intensity interferometry has been achieved in the laboratory, and first full-scale tests of connecting large Cherenkov telescopes have been carried out. This paper reviews this interferometric method in view of the new possibilities offered by arrays of air Cherenkov telescopes, and outlines observational programs that should become realistic already in the rather near future. ", "introduction": "Much of astronomy is driven by imaging with improved spatial resolution and science cases for constantly higher resolution are overwhelming. Our local Universe is teeming with stars but astronomers are still basically incapable of observing stars as such. We do observe the light radiated by them but -- with few exceptions -- are still unable to observe the stars themselves, i.e., resolve their disks or view structures across and outside their surfaces (except for the Sun, of course). One can just speculate what new worlds will be revealed once stars will no longer be seen as mere point sources but as extended and irregular objects with magnetic or thermal spots, flattened or distorted by rapid rotation, and with mass ejections through their circumstellar shells monitored in different spectral features as they flow towards their binary companions. It is not long ago that the satellites of the outer planets passed from being mere point sources to a plethora of different worlds, and one could speculate what meager state extragalactic astronomy would be in, were galaxies observed as point sources only. Tantalizing results from current optical interferometers show how stars are beginning to be seen as a vast diversity of objects, and a great leap forward will be enabled by improving angular resolution by just another order of magnitude. Bright stars have typical diameters of a few milliarcseconds, requiring optical interferometry over hundreds of meters or some kilometer to enable surface imaging. However, amplitude (phase-) interferometers require optical precisions of both their optics, and of the atmosphere above, to within a small fraction of a wavelength, and atmospheric turbulence constrains their operation when baselines exceed some 100 m, especially at shorter visual wavelengths. Using a simple {$\\lambda$}/$r$ criterion for the required optical baseline, a resolution of 1 milliarcsecond (mas) at {$\\lambda$}~500 nm requires a length around 100 meters, while 1 km enables 100 ${\\mu}$as. The potential of very long baseline optical interferometry for imaging stellar surfaces has been realized by several (e.g., Labeyrie 1996; Quirrenbach 2004), and proposed concepts include extended amplitude interferometer arrays in space: {\\it{Stellar Imager}} (Carpenter et al.\\ 2007) and the {\\it{Luciola hypertelescope}} (Labeyrie et al.\\ 2009), or possibly placed at high-altitude locations in Antarctica (Vakili et al.\\ 2005). However, despite their scientific appeal, the complexity and probable expense of these projects make the timescales for their realization somewhat uncertain, prompting searches for alternative approaches. One promising possibility is ground-based intensity interferometry. \\subsection{Intensity interferometry} Intensity interferometry was pioneered by Robert Hanbury Brown and Richard Q.\\ Twiss already long ago (Hanbury Brown 1974) for the original purpose of measuring stellar sizes, and a dedicated instrument was built at Narrabri, Australia. What is observed is the second-order coherence of light (i.e., that of intensity, not of amplitude or phase), by measuring temporal correlations of arrival times between photons recorded in different telescopes. At the time of its design, the understanding of its functioning was a source of considerable confusion (although it was explained in terms of classical optical waves undergoing random phase shifts), and even now it may be challenging to intuitively comprehend. Somewhat later, a more complete semi-classical theory was developed (e.g., Mandel \\& Wolf 1995). Seen in a quantum context, this is a two-photon process, and the intensity interferometer is often seen as the first quantum-optical experiment. It laid the foundation for a series of experiments of photon correlations including also states of light that do not have classical counterparts (such as photon antibunching). A key person in developing the quantum theory of optical coherence was Roy Glauber (1963abc; 2007), acknowledged with the 2005 Nobel prize in physics. The name {\\it{intensity interferometer}} itself is sort of a misnomer: there actually is nothing interfering in the instrument, rather its name was chosen for its analogy to the ordinary amplitude interferometer, which at that time had similar scientific aims in measuring source diameters. Two separate telescopes are simultaneously measuring the random and very rapid intrinsic fluctuations in the light from some particular star. When the telescopes are placed sufficiently close to one another, the fluctuations measured in both telescopes are correlated, but when moving them apart, the fluctuations gradually become decorrelated. How rapidly this occurs for increasing telescope separations gives a measure of the spatial coherence of starlight, and thus the spatial properties of the star. The signal is a measure of the second-order spatial coherence, the square of that visibility which would be observed in any classical amplitude interferometer, and the spatial baselines for obtaining any given resolution are thus the same as would be required in ordinary interferometry. The great observational advantage of intensity interferometry (compared to amplitude interferometry) is that it is practically insensitive to either atmospheric turbulence or to telescope optical imperfections, enabling very long baselines as well as observing at short optical wavelengths, even through large airmasses far away from zenith. Telescopes are connected only electronically (rather than optically), from which it follows that the noise budget relates to the relatively long electronic timescales (nanoseconds, and light-travel distances of centimeters or meters) rather than those of the light wave itself (femtoseconds and nanometers). A realistic time resolution of perhaps 10 nanoseconds corresponds to 3\\,m light-travel distance, and the control of atmospheric path-lengths and telescope imperfections then only needs to correspond to some reasonable fraction of those 3 meters. Since the measured quantity is the {\\it{square}} of the ordinary visibility, it always remains positive (save for measurement noise), only diminishing in magnitude when smeared over time intervals longer than the optical coherence time of starlight (due to finite time resolution in the electronics or imprecise telescope placements along the wavefront). However, for realistic time resolutions (much longer than an optical coherence time of perhaps $\\sim$10$^{-14}$~s), the magnitude of any measured signal is tiny, requiring very good photon statistics for its reliable determination. Large photon fluxes (and thus large telescopes) are therefore required; already the flux collectors used in the original intensity interferometer at Narrabri, were larger than any other optical telescope at that time. Although the signal can be enhanced by improving the electronic time resolution, faster electronics can only be exploited up to a point since there follows a matching requirement on the optomechanical systems. A timing improvement to 100 ps, say, would require mechanical accuracies on mm levels, going beyond what typically is achieved in flux collectors, and beginning to approach the level of fluctuations in path-length differences induced by atmospheric turbulence (Cavazzani et al.\\ 2012; Wijaya \\& Brunner 2011). Details of the original intensity interferometer at Narrabri and its observing program were documented in Hanbury Brown et al.\\ (1967ab), with retrospective overviews by Hanbury Brown (1974; 1985; 1991). The principles are also explained in various textbooks and reference publications, e.g., Glindemann (2011), Goodman (1985), Loudon (2000), Mandel \\& Wolf (1995), and very lucidly in Labeyrie et al.\\ (2006), Saha (2011), and Shih (2011). The original intensity interferometer at Narrabri had two reflecting telescopes of 6.5\\,m diameter, formed by mosaics of 252 hexagonal mirrors, providing images of 12\\,arcmin diameter. In order to maintain a fixed baseline while tracking (and to avoid the need for variable signal delays), the telescopes moved on a railway track of 188\\,m diameter. (The design parameters are said to have been chosen to enable it to spatially resolve the O5 star {$\\zeta$}~Puppis). Its main observing program, completed in 1973, measured angular diameters of 32 stars brighter than about m$_{V}$=\\,2.5 and hotter than T$_{eff}$= 7000\\,K, producing an effective-temperature scale for early-type stars of spectral types between O5 and F8. Following the completion of that program, the design for a second-generation intensity interferometer was worked out (Davis 1975; Hanbury Brown 1979; 1991), envisioned to have 12-m diameter telescopes, movable over 2\\,km. However, the then concurrent developments in astronomical amplitude interferometry, demonstrated already with very small telescopes, led this Australian group in that direction instead and (although a few experiments have been made in radio; e.g., Erukhimov et al.\\ 1970) astronomical intensity interferometry saw no further development. However, following its start in astronomy, intensity interferometry has been vigorously pursued in other fields, both for studying optical light in the laboratory, and in analyzing interactions in high-energy particle physics. For laboratory studies of scattered light, photon correlation spectroscopy can be considered as intensity interferometry in the temporal (not spatial) domain, and is a tool to measure the {\\it{temporal}} coherence of light, and to deduce its spectral broadening (e.g., Becker 2005; Degiorgio \\& Lastovka 1971; Oliver 1978; Saleh 1978). In particle physics, the same basic quantum principles of measuring intensity correlations apply to all bosons, i.e., particles which -- just like photons -- carry an integer value of quantum spin, and therefore share the same type of Bose-Einstein quantum statistics (Alexander 2003; Baym 1998; Boal et al.\\ 1990). In a 1959 bubble-chamber study of charged pion production in proton/antiproton annihilation, the angular distribution of like-charge pion-pairs was found to differ from the unlike-charge ones. In a now classic paper (Goldhaber et al.\\ 1960), this was interpreted as due to Bose-Einstein correlations, although the realization that the effect was equivalent to the astronomical intensity interferometer came only in the 1970's. These studies in particle physics are now generally referred to as `HBT-interferometry' (for Hanbury Brown-Twiss), although also terms such as `femtoscopy' or just `Bose-Einstein correlations' are used. A historical overview of that extensive field is by Padula (2005), while more current activities are reviewed by Bauer et al.\\ (1992); Cs{\\\"o}rg{\\H o} (2006); Heinz et al.\\ (1999); Lisa et al.\\ (2005); Wiedemann \\& Heinz (1999), or in the monograph by Weiner (2000). Thus, intensity interferometry has not been further pursued in astronomy since a long time ago, largely due to its demanding requirements for large and movable optical flux collectors, spread over long baselines, and equipped with fast detectors and high-speed electronics. However, all of these requirements are now rapidly being satisfied through the combination of high-speed digital signal handling with the construction of telescope complexes, erected for a different primary purpose, namely to optically record atmospheric Cherenkov light for the study of the most energetic gamma rays. The purpose of this paper is to review this interferometric method in view of such new possibilities, and to outline observational programs that should become realistic already in the rather near future. Current efforts to develop the several stages required towards its realization are described, including several issues that are specific to this method. Given that no astronomical intensity interferometer currently exists -- and that its functioning is fundamentally different from that of any other astronomical imaging system -- this review aims at connecting the past pioneering efforts by Hanbury Brown et al., with the potential offered by the forthcoming new arrays of air Cherenkov telescopes, and to explain the possibilities (and limitations) also for readers that might not yet be familiar with these techniques. \\begin{figure} \\includegraphics[width=9cm, angle=90]{f1.eps} \\centering \\caption{Angular resolution for existing and future observatories at different wavelengths. Except for X-rays, resolutions were taken as diffraction-limited. HST = Hubble Space Telescope; JWST = James Webb Space Telescope; NSII = Narrabri Stellar Intensity Interferometer; E-ELT = European Extremely Large Telescope; VLTI = Very Large Telescope Interferometer; VLA = Very Large Array; ALMA = Atacama Large Millimeter Array; VLBI = Very Long Baseline Interferometry (here for a baseline equal to the Earth diameter); CTA = Cherenkov Telescope Array. Intensity interferometry with large Cherenkov arrays offers unprecedented angular resolution, challenged only by radio interferometers operating between Earth and antennas in deep space. \\label{fig1}} \\end{figure} \\begin{figure} \\centering \\includegraphics[width=6cm, angle=90]{f2.eps} \\caption{Basic components of an intensity interferometer. Two telescopes observe the same source, and the measured time-variable intensities are electronically cross correlated. \\label{fig2}} \\end{figure} \\clearpage \\begin{table}[b] \\caption{Properties of the three examined configurations of Cherenkov telescope arrays (\\#1 corresponds to the upper rows in Figures \\ref{fig3} and \\ref{magnitudes}; \\#2 to the middle ones). $N$ is the number of telescopes, $A$ is the light collection area of each type of telescope, $b$ is the number of unique baselines available, $B_{min}, B_{max}$ indicates the range of baselines for observations in zenith. The corresponding range of angular diameters in milliarcseconds $(1.22 \\lambda/r)$ for observations at $\\lambda$~400~nm is indicated by $\\theta_{min}, \\theta_{max}$. In the original CTA design study, these three configurations were designated with the letters B, D, and I (Actis et al.\\ 2011).} \\begin{center} \\begin{tabular}{lrrrrr|} Array & $N$ & $A$ [m$^2$] & $b$ & $B_{min}, B_{max}$ [m] & $\\theta_{min}, \\theta_{max}$ [mas]\\\\ \\hline \\hline \\#1 & 42 & 113, 415 & 253 & 32, 759 & 0.13, 3.2 \\\\ \\#2 & 57 & 113 & 487 & 170, 2180 & 0.05, 0.6 \\\\ \\#3 & 77 & 28, 113, 415 & 1606 & 90, 2200 & 0.05, 1.13 \\end{tabular} \\end{center} \\label{configurationstable} \\end{table} \\subsection{Air Cherenkov telescopes} Seemingly ideal flux collectors for intensity interferometry are those air Cherenkov telescopes that are being erected for gamma-ray astronomy. These measure the feeble and brief flashes of Cherenkov light produced in air by cascades of secondary particles initiated by very energetic gamma rays. Time resolution has to be no worse than a few nanoseconds (duration of the Cherenkov light flash); they must be sensitive to short optical wavelengths (Cherenkov light is bluish); they must be large (Cherenkov light is faint), and they must be spread out over hundreds of meters (size of the Cherenkov light-pool onto the ground). The image seen by any one telescope shows the track of the air shower, but multiple telescopes are required for a more precise stereoscopic reconstruction of the shower geometry, and thus the direction to the source. For imaging the air shower, a modest optical imaging quality is sufficient (3--5 arcminutes, say), but possibly diverse path-lengths within the optics must not temporally smear out the Cherenkov pulse more than a few nanoseconds. The success of this concept has prompted the recent construction of several arrays with large flux collectors, including H.E.S.S. in Namibia, MAGIC on La Palma, and VERITAS in Arizona. These telescopes are large (12\\,m diameter for VERITAS, 17\\,m for MAGIC, and for H.E.S.S. even one 28\\,m dish is being completed), distributed over distances of typically 50-200 meters (V{\\\"o}lk \\& Bernl{\\\"o}hr 2009). These telescope parameters are remarkably similar to the requirements for intensity interferometry, and the compatibility is made even greater when realizing that the faintness of the Cherenkov light might preclude its efficient observation during brighter moonlight, a condition that does not inhibit interferometric observations of brighter sources (which can be made over a narrow optical bandwidth). Further, electronic time delays can now be used to compensate for different arrival times of a wavefront to the different telescopes, removing the past requirement of having the telescopes continuously moving during observation. However, the most striking potential comes from planned future facilities which will improve both the gamma-ray flux sensitivity and the angular resolution by having great many, and widely distributed flux collectors. The major international project is CTA, the Cherenkov Telescope Array (2012) which envisions a total of 50--100 telescopes with differently sized apertures between about 5 and 25 meters, distributed over an area of 2--3\\,km{$^2$}. Such an array permits an enormous number of baseline pairs to be synthesized, allowing to probe angular scales between milli- and microarcseconds. The potential of using such arrays for intensity interferometry has indeed been noticed by several (e.g., de Wit et al.\\ 2008; LeBohec \\& Holder 2006; LeBohec et al.\\ 2008a) and, within the CTA project, a working group now has the task to specify how to enable it for also intensity interferometry. If a baseline of 2\\,km could be utilized at {$\\lambda$}\\,350 nm, resolutions would approach 30\\,{$\\mu$}as. This would offer unprecedented spatial resolution in astronomy (Figure \\ref{fig1}), challenged only by radio interferometers operating between Earth and antennas in deep space (Kardashev 2009), or possibly future X-ray interferometers (MAXIM 2012). In this paper, we will examine the methodology for those types of observations, and the astrophysical targets that may be imaged when entering the new microarcsecond parameter domain. ", "conclusions": "" }, "1207/1207.2196.txt": { "abstract": "We present a photometric catalogue of compact groups of galaxies (\\emph{p2MCG}s) automatically extracted from the 2MASS extended source catalogue. A total of $262$ \\emph{p2MCG}s are identified, following the criteria defined by \\cite{Hickson82}, of which 230 survive visual inspection (given occasional galaxy fragmentation and blends in the 2MASS parent catalogue). Only one quarter of these 230 groups were previously known compact groups (CGs). Among the 144 \\emph{p2MCG}s that have all their galaxies with known redshifts, 85 (59\\%) have 4 or more accordant galaxies. This \\emph{v2MCG} sample of velocity-filtered \\emph{p2MCG}s constitutes the largest sample of CGs (with $N\\geq4$) catalogued to date, with both well-defined selection criteria and velocity filtering, and is the first CG sample selected by stellar mass. It is fairly complete up to $K_{\\rm group} \\sim 9$ and radial velocity of $\\sim 6000 \\, \\rm km \\, s^{-1}$. We compared the properties of the 78 \\emph{v2MCG}s with median velocities greater than $3000 \\, \\rm km \\, s^{-1}$ with the properties of other CG samples, as well as those (\\emph{mvCG}s) extracted from the semi-analytical model (SAM) of Guo et al. (2011) run on the high-resolution Millennium-II simulation. This \\emph{mvCG} sample is similar (i.e. with 2/3 of physically dense CGs) to those we had previously extracted on three other SAMs run on the Millennium simulation with 125 times worse spatial and mass resolutions. The space density of \\emph{v2MCG}s within $6000 \\, \\rm km \\, s^{-1}$ is $8.0 \\times 10^{-5} \\,h^3 \\,\\rm Mpc^{-3}$, i.e. 4 times that of the Hickson sample (HCG) up to the same distance and with the same criteria used in this work, but still 40\\% less than that of \\emph{mvCG}s. The \\emph{v2MCG} constitutes the first group catalogue to show a statistically large first-second ranked galaxy magnitude difference (in units of the dispersion of the first-ranked absolute magnitudes) according to Tremaine-Richstone statistics, as expected if the first ranked group members tend to be the products of galaxy mergers, and as confirmed in the \\emph{mvCG}s. The \\emph{v2MCG} is also the first observed sample to show that first-ranked galaxies tend to be centrally located, again consistent with the predictions obtained from \\emph{mvCG}s. We found no significant correlation of group apparent elongation and velocity dispersion in the quartets among the \\emph{v2MCG}s, and the velocity dispersions of apparently round quartets are not significantly larger than those of chain-like ones, in contrast to what has been previously reported in HCGs. By virtue of its automatic selection with the popular Hickson criteria, its size, its selection on stellar mass, and its statistical signs of mergers and centrally located brightest galaxies, the \\emph{v2MCG} catalogue appears to be the laboratory of choice to study physically dense groups of 4 or more galaxies of comparable luminosity. ", "introduction": "Compact Groups (hereafter, CGs) of at least 4 galaxies of comparable luminosity are the densest galaxy associations known at present. The compactness of these groups is so high that the typical projected separations between galaxies are of the order of their own diameters \\citep{HMdOHP92,FK02}, hence their space densities can exceed those of the cores of rich clusters. The combination of their very high number densities and low velocity dispersion makes CGs the ideal site of galaxy mergers (\\citealp{Mamon92}, see also \\citealp{CCS81,Barnes85,Mamon87,BCL93}). Since the discovery of Stephan's Quintet \\citep{Stephan1877} and Seyfert's Sextet \\citep{Seyfert48}, several surveys of CGs have been undertaken: \\cite{Rose77} and \\cite{Hickson82} performed visual identifications of CGs on the POSS I photographic plates. Thereafter, the new catalogues of CGs used automatic searches: from the COSMOS/UKST Southern Galaxy Catalogue \\citep{PIM94,Iovino02}, the DPOSS catalogue \\citep{Iovino03,decarvalho05}, and the Sloan Digital Sky Survey (SDSS) photometric catalogue DR1 \\citep{Lee+04} and DR6 \\citep{McCPES09}. All of the above studies used only 2-dimensional information of the galaxies (i.e., angular positions). Other CG catalogues were obtained by searches in redshift space, e.g.: \\cite{Barton+96} from the the CfA2 catalogue, \\cite{AT00} from the Las Campanas Redshift Survey, \\cite{FK02} from the UZC Galaxy Catalogue, and \\cite{Deng+08} from the SDSS-DR6 spectroscopic catalogue. Since the nearly full spectroscopic followup by \\cite{HMdOHP92} of the original Hickson Compact Groups (\\citealp{Hickson82}, hereafter, HCGs), the velocity-filtered sample of 92 HCGs with at least 3 accordant-redshift members and 69 with at least 4 has been, by far, the most studied to date (e.g. \\citealp{HKA89} for optical photometry; \\citealp{MdOH91} for galaxy morphologies; \\citealp{Moles+94}, \\citealp{delarosa07} and \\citealp{tzanavaris10} for star formation rates; \\citealp{CRdCC98} for nuclear activity; \\citealp{delarosa01} and \\citealp{torres-flores10} for galaxy scaling relations; \\citealp{VerdesMontenegro+01} and \\citealp{borthakur10} for neutral gas content; \\citealp{PBEB96} for hot gas content, etc.). However, the visual inspection performed by \\cite{Hickson82} led to a sample of CGs that is not reproducible, incomplete and not homogeneous \\citep{HKA89,WM89,PIM94,Sulentic97,diaz-mamon10}. In particular, using the $z=0$ outputs of semi-analytical models of galaxy formation run on the Millennium cosmological dark matter simulation \\citep{Springel+05}, \\cite{diaz-mamon10} have shown that the HCG sample is typically less than 10\\% complete at the median distance of the sample. The properties of CGs and their member galaxies must be studied using complete and well-defined observed samples. To achieve this goal, we present a new sample of automatically selected CGs extracted from the largest solid angle catalogue at present, the 2 Micron All Sky Survey. Using 2MASS has two strong advantages: 1) it provides us with a full-sky survey and 2) the $K$-band photometry is only weakly sensitive to both galactic extinction, internal extinction and recent star formation, and is thus a very good tracer of the stellar mass content of galaxies. For these reasons, it is ideal to build a CG sample from a wide $K$-band galaxy survey such as 2MASS (which has the additional benefit of being all-sky) with (nearly) full redshift information available from other sources \\citep{Mamon94}. The layout of this paper is as follows. In Sect.~\\ref{2mass}, we describe the parent catalogue. In Sect.~\\ref{2mcgc}, we present the CG catalogue. We perform a cross-identification between the 2MASS-CGs and other samples of groups in Sect.~\\ref{cross-id}. In Sect.~\\ref{v-filter}, we present a sample of CG after applying a velocity filtering, while we present some general properties of the samples in Sect.~\\ref{props}, and summarise and discuss our results in Sect.~\\ref{conclude}. Throughout this paper we use a Hubble constant $H_0= 100 \\, h \\, {\\rm km \\,s^{-1}\\, Mpc^{-1}}$, and for all cosmology-dependent calculations, we assume a flat cosmological model with a non-vanishing cosmological constant: $\\Omega_{\\rm m}=0.25$ and $\\Omega_\\Lambda=0.75$. ", "conclusions": "\\label{conclude} In this work we have catalogued a new sample of CGs from the 2MASS survey and we have compared them with existing CG samples. Following the criteria defined by \\cite{Hickson82}, we have identified 230 CGs in projection in the $K$-band covering $23\\,844 \\, {\\rm deg}^2$. This catalogue has well defined criteria which produced an homogeneous sample useful to perform statistical analyses on it. 25\\% of them (57 CGs) were previously identified in other catalogues as compact groups, triplets of galaxies or interacting galaxies. A total of $144$ \\emph{p2MCG}s have \\emph{all} their members with redshifts available in the literature, and among them $85$ groups have 4 or more accordant galaxies, which makes this catalogue the largest sample of CGs with 4 or more spectroscopically confirmed members. The percentage of groups with accordant galaxies (59\\%) is slightly lower than that obtained from the HCG sample (67\\%), and very similar to the predicted by \\cite{diaz-mamon10} from the semi-analytical models of \\cite{Bower+06} and \\cite{dLB07}.\\footnote{For the SAM of \\cite{Guo11}, we find that 70\\% of mock CGs found in projected space survive the velocity filtering.} As a side note, we have now built additional mock CG catalogues using the \\cite{Croton+06} SAM in the $K$-band and the \\cite{Guo11} SAM in the $r$ band, where the latter was run on the Millennium-II Simulation, which has 512 times the mass resolution of the Millennium Simulation. For both samples, we found that two-thirds of the mock CGs were physically dense systems of at least 4 galaxies of accordant magnitudes, while the remaining third was caused by chance alignments of galaxies along the line-of-sight, mostly within larger virialised groups, confirming similar conclusions of \\cite{diaz-mamon10}. In comparison with other CG catalogues, the \\emph{v2MCG} catalogue presented in this work is one of the nearest and brightest samples of CGs, although these CGs have larger projected radii and interparticle separations. The \\emph{v2MCG} does not show any significant correlation for quartets between apparent elongation and velocity dispersion nor significantly larger velocity dispersion in round groups relative to chain-like groups, contrary to what \\cite{TMT99} claimed in HCGs. The \\emph{v2MCG} is the only CG sample to display significantly large differences between second- and first-ranked absolute magnitudes (from Tremaine-Richstone statistics) as well as centrally located first-ranked galaxies, both in agreement with mock \\emph{mvCG}s, but in sharp contrast with all other observed velocity-filtered CG samples. Galaxy mergers are an obvious way to decrease $T_1$ and $T_2$ \\citep{Mamon87}, and we cannot think of any other physical mechanism that may cause both $T_1$ and $T_2$ to be significantly smaller than unity in a group catalogue. One major difference of our sample with others is that ours has many more groups with dominant galaxies accounting for over half the total luminosity. While this increases the gap between first and 2nd-ranked magnitudes, we found that our sample also has a small standard deviation of first-ranked absolute magnitudes, which enhances the significance of the Tremaine-Richstone $T_1$ statistic. Why don't we find significant magnitude gaps and luminosity segregation in the other CG samples? It is clear that in his visually search for CGs, \\cite{Hickson82} missed groups with dominant galaxies (\\citealp{PIM94,diaz-mamon10} and Table~\\ref{median_v}). Could those \\emph{v2MCG}s in common with \\emph{vHCG}s have weaker signs of magnitude gaps and luminosity segregation? One expects that if mergers cause the magnitude gap, the masses, i.e. stellar masses, of the galaxies are the crucial variable. Similarly, if luminosity segregation is produced by dynamical friction or by energy equipartition from two-body relaxation, the galaxy (stellar) masses should be the important variable. Therefore, the magnitude gaps and luminosity segregation should be weaker in the $R$ band, where the luminosity is less a measure of stellar mass than in the $K$-band. Unfortunately, we have only 14 groups in common, among which 10 (HCG 7, 10, 23, 25, 40, 58, 86, 88, 93, 99)\\footnote{and with slightly variations: HCG 15, 16, 51, 97.} have exactly the same galaxies. For these 10 groups, we find $T_1=0.68\\pm0.24$ and $T_2=0.87\\pm0.24$ in the $K$ band, while in the $R$ band we find values greater than unity: $T_1 = 1.75\\pm1.69$ and $T_2=1.33\\pm0.39$. So, indeed, the $K$-band luminosities are more sensitive than their $R$-band counterparts to the magnitude gap, but given the bootstrap errors, the large differences in $T_1$ and $T_2$ between $R$ and $K$-based absolute magnitudes are not statistically significant (for $T_2$ the difference is roughly $1\\,\\sigma$, while it is much less for $T_1$). On the other hand, luminosity segregation is not seen in either waveband: worse it is inverted, with the brightest galaxy on average further away from the group centroid than the 2nd-brightest galaxy. The other two CG samples, UZC-CG and LCCG, are based upon Friends-of-Friends (LCCG) or similar (UZC-CG) algorithms, both with velocity linking length of $1000 \\, \\rm km \\, s^{-1}$. Such a velocity link is much more liberal than imposing that the velocities all lie within $1000 \\, \\rm km \\, s^{-1}$ from the median as done in \\cite{HMdOHP92} and here. Indeed, according to Table~\\ref{median_v}, the median velocity dispersion of UZC-CG groups of 4 or more galaxies is $295 \\, \\rm km \\, s^{-1}$, i.e. 25\\% greater than in our sample. This suggests that the UZC-CG sample is more contaminated by chance alignments of galaxies along the line-of-sight (as \\citealp{Mamon86} had suggested for the HCG sample) than is our sample. Moreover, the UZC-CG has a liberal linking length on projected distances of $200 \\, h^{-1} \\, \\rm kpc$, making these groups not so compact (as can be checked by their low mean group surface brightness as seen in Table~\\ref{median_v}). Finally, the UZC-CG is based upon Zwicky's visually estimated magnitudes, which may carry rms errors as large as 0.5 mag, thus washing out in part the effects of the magnitude gap and luminosity segregation. On the other hand, the LCCG sample (again, restricted to groups with at least 4 members) has a very similar median velocity dispersion to our sample (and a similar median mass-to-light ratio). Note that the linking length for projected distances of the LCCG is only $50 \\, h^{-1} \\, \\rm kpc$, i.e. 4 times less than in UZC-CG. The problem with the LCCG is that its parent catalogue (the LCRS, \\citealp{Shectman+96}) is a collection of two samples with $16.0 < R < 17.3$ and $15.0 < R < 17.7$. Thus the magnitude range is very restricted. Hence, it is not a surprise that $\\langle M_2-M_1\\rangle$ is half our value (Table~\\ref{median_v}), leading to $T_1>1$ and $T_2>1$. What does this tell us on the nature of the groups in the different CG samples? Over 25 years ago, \\cite{Mamon86} found $T_1=1.16$ and no signs of luminosity segregation in the largest sample then available of 41 velocity-filtered HCGs with at least 4 members. This was in sharp contrast with the low values of $T_1$ and significant luminosity segregation he was finding in coalescing dense groups \\citep{Mamon87}. This provided him with two arguments (among several) to conclude that most HCGs were caused by chance alignments of galaxies within larger groups \\citep{Mamon86}. As confirmed here with the SAM by \\cite{Guo11}, roughly two-thirds of mock CGs are physically dense (\\citealp{diaz-mamon10} and Table~\\ref{tabSAMs} above). The statistically large magnitude gaps and luminosity segregation in both the observed \\emph{v2MCG}s and the mock \\emph{mvCG}s suggest that \\cite{Mamon86} was misled by the bias of the HCG sample against large gaps into concluding that most of them were not physically dense. So the \\emph{v2MCG} appear to be mostly \\emph{bona fide} physically dense groups. But can we conclude that the other CG samples are dominated by chance alignments? \\cite{diaz-mamon10} attempted to build sample of mock HCGs that include the same biases as they had measured by comparing with the three SAMs that they had built mock CGs from. They found that the same fraction (if not slightly higher) of the mock (biased) HCGs were physically dense. The nature of the groups in other CG samples could be studied in similar ways, using mock CG samples from cosmological galaxy formation simulations, mimicking their selection criteria and observational select effects. Could the lack of HCGs with strongly dominant brightest galaxies prevent the visibility of luminosity segregation? We performed KS tests to compare the distribution of relative positions of 1st and 2nd-ranked group members for subsamples split between those dominated by 1st-ranked members (`Dom') and those with galaxies of more comparable luminosities (`Non-Dom'), making our splits at the median magnitude difference $\\langle M_2-M_1\\rangle$. We performed this analysis for the HCG and v2MCG samples as well as for the C06K and G11 \\emph{mvCG} samples. % \\begin{table} \\begin{center} \\caption{Luminosity segregation split by magnitude gap\\label{lsegsplit}} \\begin{tabular}{llrc} \\hline \\hline Catalogue & subsample & $N$ & $P_{\\rm KS}$ \\\\ \\hline \\emph{v2MCG} & Dom & 39 & $2.9\\times10^{-7}$ \\\\ \\emph{v2MCG} & Non-Dom & 39 & (3.9\\%) \\\\ \\emph{vHCG} & Dom & 17 & 67\\% \\\\ \\emph{vHCG} & Non-Dom & 16 & (6.6\\%) \\\\ \\emph{mvCG}-G11 & Dom & 163 & $1.2\\times 10^{-11}$ \\\\ \\emph{mvCG}-G11 & Non-Dom & 163 & 1.0\\% \\\\ \\emph{mvCG}-C06K & Dom & 223 & $4.0\\times 10^{-9}$ \\\\ \\emph{mvCG}-C06K & Non-Dom & 225 & 0.4\\% \\\\ \\hline \\end{tabular} \\end{center} \\parbox{\\hsize}{ \\small {\\bf Notes.} Dom and Non-Dom subsamples are those with $M_2-M_1$ above and below the median value of the full sample, respectively. $N$ is the number of groups in the subsample. $P_{\\rm KS}$ is the KS probability that a difference in the distributions of normalised distances to the non-weighted group centre is greater than `observed' by chance. Values of $P_{\\rm KS}$ given in parentheses denote reverse luminosity segregation: the 2nd-ranked galaxy is more centrally located than the first-ranked.} \\end{table} Table~\\ref{lsegsplit} shows that indeed luminosity segregation is much stronger for all catalogues in the Dom subsamples, and statistically significant in all of them except the \\emph{vHCG}. Surprisingly, while the Non-Dom subsamples of the two mock CG samples display much weaker, but still statistically significant luminosity segregation, the Non-Dom subsamples of both the \\emph{v2MCG} and \\emph{vHCG} catalogues display \\emph{reversed luminosity segregation}: the 2nd-ranked galaxy is more centrally located than the (slightly more luminous) first-ranked galaxy. We can only see one explanation for this reverse luminosity segregation, if it occurs in wavebands bluer than $K$: late-type galaxies that are second-ranked in stellar mass, hence not centrally located, can end up more luminous (thanks to their efficient star formation) than early-type galaxies of slightly higher stellar mass. However the effect is also present in the $K$-selected \\emph{v2MCG}, and with even greater statistical significance (96.1\\% vs. 93.7\\% confidence for the Non-Dom CGs of the \\emph{v2MCG} and \\emph{vHCG} catalogues, respectively). So, we can only explain this marginal effect as a statistical fluke. Nevertheless, the absence of luminosity segregation in the HCG catalogue could be consistent with their physical reality because the sample is too small to detect the weak luminosity segregation expected from the mocks. Moreover, if the reverse luminosity segregation for Non-Dom groups is real, then the lack of groups with very dominant galaxies in the HCG (caused by the visual selection bias) would cause the Non-Dom groups to cancel the luminosity segregation of the Dom groups. In conclusion, the \\emph{v2MCG} sample has numerous advantages over other CG samples: \\vspace{-0.5\\baselineskip} \\begin{enumerate} \\itemindent 0pt \\item It is the largest available sample of velocity-filtered groups of at least 4 members of comparable luminosity (3 mags, i.e. factor 16). \\item It has an isolation criterion (in contrast with other CG samples except for the HCG). \\item It is automatically extracted (contrary to the HCG). \\item It has a well-defined magnitude limit (which the HCG sample does not). \\item It is deep enough (which some may find surprising given the shallowness of its parent 2MASS catalogue) to have a selection on brightest galaxy magnitude, so as to ensure that all groups can span the maximum allowed magnitude gap of 3. \\item It is selected by stellar mass ($K$ band), which is expected to be a better tracer for magnitude gaps and luminosity segregation (among other things). \\item It is the only sample to show statistical signs of mergers (magnitude gaps) and luminosity segregation, expected in physically dense groups (in contrast with all other CG samples). \\end{enumerate} The last point implies that the \\emph{v2MCG} is the only CG sample for which one is reasonably sure that it is dominated by physically dense groups. For all these reasons, the \\emph{v2MCG} appears to be the sample one ought to study in depth to probe the effects on galaxies of this unique environment of 4 galaxies of comparable luminosity lying close together in real space. As a next step in this project, we are in the process of measuring redshifts for the members for which no spectra are available, and we are continuing our statistical studies of the \\emph{v2MCG}s and their galaxies." }, "1207/1207.1453_arXiv.txt": { "abstract": "The asymmetric molecular emission lines from dense cores reveal slow, inward motion in the clouds' outer regions. This motion is present both before and after the formation of a central star. Motivated by these observations, we revisit the classic problem of steady, spherical accretion of gas onto a gravitating point mass, but now include self-gravity of the gas and impose a finite, subsonic velocity as the outer boundary condition. We find that the accretion rate onto the protostar is lower than values obtained for isolated, collapsing clouds, by a factor that is the Mach number of the outer flow. Moreover, the region of infall surrounding the protostar spreads out more slowly, at a speed close to the subsonic, incoming velocity. Our calculation, while highly idealized, provides insight into two longstanding problems -- the surprisingly low accretion luminosities of even the most deeply embedded stellar sources, and the failure so far to detect spatially extended, supersonic infall within their parent dense cores. Indeed, the observed subsonic contraction in the outer regions of dense cores following star formation appears to rule out a purely hydrodynamic origin for these clouds. ", "introduction": "Astronomers are elucidating, in much greater detail than ever before, the earliest phases of low-mass star formation. It has long been known that at least half of dense cores within molecular clouds do not yet contain young stars \\citep{b86}. It is also known that these starless cores have, on average, less mass than cores with embedded stars \\citep{jma99}. What is now clearer is the chemical and structural progression of cores as they approach the collapse threshold and then cross it. The most centrally concentrated starless cores tend to exhibit the asymmetric emission line profiles signifying inward contraction, with the inferred velocity being several tenths of the sound speed \\citep{lmt01}. The same blueward asymmetry is seen in those cores with embedded stars, where it is especially common \\citep{m97}. Detailed examination of line profiles reveals no qualitative change across the collapse threshold, except for the appearance of broad wings that are attributable to molecular outflows \\citep{ge00}. It is likely that the contraction observed in starless cores is driven by self-gravity, as material in the core's outer region settles toward the center \\citep[e.g.,][]{cb01}. A similar, exterior contraction should also occur in cores that have produced stars, and observations thus far support this conclusion. Here, targeted studies are rare, in part because outflows often contaminate the emission line spectra. As one example, \\citet{m96} utilized a simple radiative transfer model to measure the spatially averaged contraction speed in L1526, a core with an embedded star that drives a weak outflow. Their derived contraction speed of 0.025~km~s$^{-1}$ is indeed close to the inferred speeds in starless cores. The CS and H$_2$CO spectra for the low-mass infall candidates identified by \\citet{m97} have a similar degree of asymmetry as L1526, but have not been modeled to this level of detail. Taken as a whole, these observations suggest that a dense core reaches the point of collapse by accruing matter externally. The core itself is a relatively quiescent structure, since its observed emission lines have nearly thermal widths \\citep{bg98}. The exterior molecular gas, however, has typical speeds of 1 to 2~km~s$^{-1}$, corresponding to Mach numbers of 5 to 10 \\citep{bt07}. How exactly material passes from this environment into the core is unknown, despite increasingly detailed mapping of the transition region \\citep{p10}. In any event, the core is dynamically evolving while it collects mass. Gas in its outer mantle creeps inward, and continues to do so even after the central region collapses to form a star. The standard theory of collapse, first formulated prior to the discovery of dense cores, paints a very different picture. In the still widely used self-similar model of \\citet{s77}, the cloud is a perfectly static entity before the first appearance of a protostar. The impulsive suctioning of gas caused by this event sends out a sharp rarefaction wave, expanding at the speed of sound. It is within this wavefront that infall occurs. Gas leaves the rarefaction wave with zero velocity and attains the sound speed roughly halfway inside it. For a typical dense core radius of 0.1~pc and sound speed of 0.2~km~s$^{-1}$, the rarefaction wave reaches the boundary in 0.5~Myr. This interval, or one comparable, has been traditionally considered to be the lifetime of the protostar phase. Comparing the dynamical picture from theory with observations, it is striking that there are no known cores with embedded stars in which the sonic transition occurs reasonably close to the cloud boundary. On the contrary, the inferred location of this point in the best spectroscopic collapse candidates is only 0.01-0.02~pc from the protostar itself \\citep{c95,ge97,df01,b02}. Again, we must remember that strong molecular outflows can both mask and mimic large-scale motion in these studies. On the other hand, submillimeter continuum mapping of cores with embedded stars also indicates that infall is confined to a small, central region. \\citet{ser02} have shown that more widespread collapse would give rise to a change in the interior density profile that is not observed. The continuing absence of dense cores that contain stars and exhibit true global infall is forcing us to conclude either that all observed protostars happen to be especially young, or else that infall actually spreads more slowly than traditional theory posits. The collapse model of \\citet{s77} also predicts that the mass accretion rate onto the protostar is a constant in time: \\begin{equation} {\\dot M} \\,=\\, m_\\circ\\,{c_s^3 \\over G} \\,\\,. \\end{equation} Here, $c_s$ is the isothermal sound speed in the core and $m_\\circ$ is a nondimensional number with a value of 0.975. Many authors, starting with \\citet{k90}, have noted that the accretion luminosity associated with this rate exceeds those observed for embedded, infrared sources. Subsequent numerical simulations relaxed the assumption of self-similarity, but usually imposed a rigid boundary on the dense core, i.e., a fixed surface where the fluid velocity vanishes \\citep[e.g.,][]{fc93,o99,vb05}. In these calculations, ${\\dot M}$ varies in time, but is even higher, at least before gas begins to drain from the boundary region. The sonic transition occurs only slightly inside this boundary \\citep[see Fig.~1 of][]{fc93}. In response to the persistent and growing disparity between theory and observation, researchers have taken two very different paths. Some have focused on the possibility that, while cloud collapse occurs in the traditional manner, infalling gas is retained in a circumstellar disk, and only released periodically in bursts onto the protostar \\citep[e.g.,][]{dv12}. If the duration of bursts is sufficiently brief, we could be viewing protostars only when $\\dot M$ is relatively low. Whatever the merits of this idea, it does not address the troubling aspects of collapse models generally. Protostellar luminosity and cloud collapse are clearly linked, and a solution that considers both is preferable. Following this second, unified path, \\citet{f04} generalized the model of \\citet{s77} and investigated the self-similar evolution of a cloud that has a specified velocity profile at the start. \\citet{go07} and \\citet{go09} simulated numerically the birth of a dense core as arising from the convergence of a supersonic inflow. Both groups tracked the evolution past the start of core collapse. \\citet{sy09} used perturbation theory to find the contraction in starless cores that are just approaching collapse, and were able to reproduce the asymmetric emission lines \\citep{sy10}. In this brief contribution, we focus on how continuous, external mass addition to the dense core affects protostellar infall itself. We show that $\\dot M$ in this case is equal to $c_s^3/G$ times a factor that is essentially the Mach number of the incoming flow. This key result is easy to understand physically. For a cloud that is marginally stable against self-gravity, the free-fall and sound-crossing times are comparable. It follows that the mean speed of infalling fluid elements in an isolated cloud undergoing collapse is approximately $c_s$. When accretion is imposed from the outside, this mean speed is closer to the externally impressed value. The mass transport rate is lowered by their ratio, which is the incoming Mach number. Since the contraction speed in the outer portion of the dense core is well under $c_s$, the reduction in $\\dot M$ may be substantial. The accretion luminosity is similarly reduced from traditional values, alleviating the discrepancy between theory and observation in this regard. We further show that the region of infall spreads out at a speed that is close to the subsonic injection velocity in the core's exterior. This relatively slow expansion helps to explain why global, supersonic infall has been difficult to observe. In Section~2 below, we describe the physical problem to be solved, present the relevant equations in nondimensional form, and then give our method of solution. Section~3 presents numerical results. Finally, Section~4 compares our findings with those of other studies, and indicates future directions for extending this work. ", "conclusions": "Having solved the problem as posed, our next task is to check that the assumption of steady-state flow is reasonable. The outer portion of the cloud evolves slowly with time, as a structure that is never far from hydrostatic balance. It is the interior region that concerns us. Again, we need to make sure that the crossing time from $r_s$ to the origin is less than $M_\\ast/{\\dot M}$, which in turn is close to the evolutionary time $t$.\\footnote{The time $M_\\ast/{\\dot M}$ would {\\it not} be close to $t$ if $\\dot M$ varied rapidly in time, as it does in simulations of isolated, collapsing clouds. However, we do not expect such rapid variation when a subsonic, external flow is imposed from the start.} Thus, we require \\begin{equation} \\int_0^{r_s} \\!{{dr}\\over u} \\,<\\,t \\,\\,. \\end{equation} But \\hbox{$u\\,>\\,c_s$} in this region, so that the integral is less than $r_s/c_s$. Since the infall region expands subsonically, this ratio is indeed less than $t$. We stress the need for a fully time-dependent calculation to verify our main results. Assuming there is no qualitative change, it is instructive to compare our findings with those of \\citet{go09}. Using direct simulation, these authors built up a dense core from a spherically converging, supersonic flow. The accretion shock bounding the core first expands and then contracts supersonically before the core itself undergoes interior collapse. When the protostar first appears, the velocity profile throughout the dense core is uniform and supersonic, with \\hbox{$u\\,\\approx\\,3\\,c_s$}. A rarefaction wave erupts from the origin, but now expands supersonically. The wave quickly overtakes the accretion shock and effectively destroys it, so that the converging inflow thereafter directly impacts the protostar. \\citet{go11} repeated the simulation using a planar, external flow and found substantially the same results. We observe once more the close relationship between the velocity of $r_s$ and that of the gas introduced externally. To see this connection in yet another light, note that $r_s$ is located roughly where the free-fall velocity onto the star matches the sound speed: \\begin{equation} \\sqrt{{2\\,G\\,M_\\ast}\\over{r_s}} \\,\\approx \\, c_s \\,\\,. \\end{equation} If we let \\hbox{$M_\\ast\\,=\\,{\\dot M}\\,t$} and use equation~(21), we find that \\hbox{$r_s\\,\\approx\\,u_\\infty\\,t$}. The fact that the outer regions of dense cores are contracting subsonically {\\it both before and after} the appearance of a protostar appears to rule out a purely hydrodynamic origin for these objects. The many researchers who have simulated cluster formation in turbulent molecular gas have effectively sidestepped this issue. In these studies, random velocity fluctuations are applied to a representative cube of molecular gas. Once these fluctations thoroughly stir up the gas, self-gravity is switched on and overdense structures promptly collapse \\citep[see, e.g.,][]{bhv99,hmk01,pn02,okm08}. More realistically, a dense core is self-gravitating and gains mass prior to this collapse, and a force in addition to the thermal pressure gradient must support it during this early epoch. This force is already known -- the pressure gradient associated with the interstellar magnetic field. The field may not play a key role in the collapse phase studied here, although it apparently impacts the formation of circumstellar disks \\citep{lks11}. In any case, the magnetic field is surely key in the early buildup of dense cores within their parent cloud filaments, and must also mediate the transition between the supersonic exterior of each core and its supersonic exterior." }, "1207/1207.1723_arXiv.txt": { "abstract": "We study the impact of early dark energy (EDE) cosmologies on galaxy properties by coupling high-resolution numerical simulations with semi-analytic modeling (SAM) of galaxy formation and evolution. EDE models are characterized by a non-vanishing high-redshift contribution of dark energy, producing an earlier growth of structures and a modification of large-scale structure evolution. They can be viewed as typical representatives of non-standard dark energy models in which only the expansion history is modified, and hence the impact on galaxy formation is indirect. We show that in EDE cosmologies the predicted space density of galaxies is enhanced at all scales with respect to the standard \\lcdm scenario, and the corresponding cosmic star formation history and stellar mass density is increased at high-redshift. We compare these results with a set of theoretical predictions obtained with alternative SAMs applied to our reference \\lcdm simulation, yielding a rough measure of the systematic uncertainty of the models. We find that the modifications in galaxy properties induced by EDE cosmologies are of the same order of magnitude as intra-SAM variations for a standard \\lcdm realization (unless rather extreme EDE models are considered), suggesting that is difficult to use such predictions alone to disentangle between different cosmological scenarios. However, when independent information on the underlying properties of host dark matter haloes is included, the SAM predictions on galaxy bias may provide important clues on the expansion history and the equation-of-state evolution. ", "introduction": "\\label{sec:intro} The last decade has seen considerably advances in our understanding of the properties of our Universe as a whole, in particular thanks to the accurate measurement of the most important cosmological parameters (see e.g. \\citealt{Komatsu09} and references herein). One of the most surprising discoveries of this ``precision cosmology'' epoch is that some unknown form of {\\it Dark Energy} (DE, hereafter) accounts for more than $70\\%$ of the energy density of the present-day Universe, and is responsible for its accelerated expansion today (see e.g., \\citealt{Perlmutter99}). The physical nature of DE, together with its origin and time evolution, is one of the most enigmatic puzzles in modern cosmology. In the standard \\lcdm cosmological model, DE is treated as a classic cosmological constant (following Einstein's original conjecture) where a homogeneous and static energy density fills the whole Universe. This simple assumption, however, gives rise to a number of theoretical problems, such as the high degree of ``fine-tuning'' required to accommodate the present day value of $\\Omega_\\Lambda$ (see \\citealt{Weinberg89} for a review). For this reason, many alternative scenarios to explain DE and the accelerated expansion have been proposed, ranging from scalar field models such as quintessence to radical modifications of the laws of gravity. Most observational efforts currently concentrate on constraining effective parametrisations of the DE equation of state (see e.g. \\citealt{Wetterich88, RatraPeebles88}), which is adequate for `non-coupled' dark energy models in which the dark energy can be approximated as uniform and influences structure formation only through a modification of the cosmic expansion rate. An interesting sub-class of these models are scenarios that involve a non-negligible DE contribution at early times during recombination and primordial structure formation, which also implies non-negligible modifications of the cosmic microwave background \\citep{Doran01}, of big-bang nucleosynthesis \\citep{Muller04} and of large-scale structure formation \\citep{Bartelmann06}. The non-linear structure formation predicted in such early dark energy (EDE) models has been studied through high-resolution $N$-body simulation \\citep{Baldi12, GrossiSpringel09}. In particular, the impact on the statistical properties of dark matter (DM) haloes (such as the abundance of high-redshift galaxy clusters) as a function of cosmic time has been compared with the corresponding predictions for the standard \\lcdm cosmology, with the aim of identifying observational tests that are able to disentangle between the different cosmologies based on future surveys (like the EUCLID satellite, \\citealt{Laureijs11}). In contrast, the influence of modified DE models on the properties and the evolution of galaxy populations has not yet been explored in full detail, even though many cosmological tests for constraining the DE equation of state ultimately rely on a precise understanding of how galaxies trace mass. This can in part be understood as a result of our limited present understanding of galaxy formation even in the $\\Lambda$CDM cosmology, which is hampered by several long-standing discrepancies between the predictions of theoretical models and observations, as seen, for example, in the redshift evolution of the galaxy stellar mass function (see e.g.~\\citealt{Fontanot09b}), the properties of the Milky Way satellites (see e.g.~\\citealt{Maccio10, BoylanKolchin12}), the low baryon fraction in galaxy clusters (see e.g.~\\citealt{McCarthy07}), or in the shallow DM profiles associated with different galaxy populations (``cusp-core'' problem, see e.g.~\\citealt{Moore94}). This in turn fuels some interest in exploring both the possibility that these discrepancies may be reduced by assuming an evolving DE, and in finding observational tests based on statistical galaxy properties that could potentially distinguish between different DE scenarios (which one may also hope to be easier to perform compared with tests working purely in the ``dark sector''). The catch of course is that galaxy evolution is a complex process, involving a non-linear blend of many different physical processes acting on the baryonic gas. All models of galaxy formation on cosmological scales (both semi-analytic models, SAMs hereafter, and numerical simulations alike) need to simplify this intrinsic complexity by means of quite coarse, yet physically grounded, analytic approximations (in SAMs) or sub-grid models (in simulations). In SAMs, these analytic approximations are meant to describe the physical processes acting on the baryonic gas (such as gas cooling, star formation and feedback) as a function of the physical properties of model galaxies (like their cold gas and stellar content). This then yields a convenient parametrisation of the physics, which is calibrated against a small subset of low-redshift observations. The ability of modern models of this kind to successfully match a large and diverse set of observations not used in the calibration can be viewed as a powerful confirmation of the viability of the basic paradigm of hierarchical galaxy formation. However, it has also been shown that the SAM approach entails a significant level of degeneracies among the different parameters \\citep[e.g.][]{Henriques09}, and the theoretical uncertainty is exacerbated by the fact that different models tend to adopt different parametrizations for the physical processes acting on the baryonic gas. Even though the comparison of different model predictions shows a reassuring level of consistency in many cases (see e.g. \\citealt{Fontanot09b}), there is hence a certain degree of systematic uncertainty in this approach. Therefore, in order to assess the role of modified DE on galaxy evolution and, viceversa, to use galaxy formation to constrain DE models, it is important to not only quantify the impact of modified cosmologies on galaxy formation within a specific model, but also to check how the size of the changes in the model predictions compares to the intra-model variance for a fixed $\\Lambda$CDM cosmology. This paper is organized as follows. In Section~\\ref{sec:models}, we introduce the cosmological numerical simulations and semi-analytic models we use in our analysis. We then compare the various predictions in Section~\\ref{sec:results}. Finally, we discuss our conclusions in Section~\\ref{sec:final}. ", "conclusions": "\\label{sec:final} In this paper we discuss the expected impact of Early Dark Energy (EDE) cosmologies on the properties of galaxy populations, as predicted by semi-analytic models. We consider the latest version of the \\munich model \\citep{Guo11}, modified to account for time-dependent variations of the equation of state parameter $w$. We also consider earlier versions of the same SAM \\citep{Croton06, DeLuciaBlaizot07} to compare the size of differences in galactic properties induced by changes in the underlying dark energy model with the intra-model variance due to differences in the modeling of the baryonic physics. The extension of our approach to consider other cosmological models with alternative dark energy scenarios \\citep[see e.g.][]{Baldi12} is in principle straightforward, provided high-resolution N-body simulations of such scenarios can be calculated. The latter has very recently become possible even for more complicated theories of gravity, such as DGP or $f(R)$-gravity \\citep[see e.g.][]{Oyaizu08, Schmidt09b, KhouryWyman09, LiZhaoKoyama12}, such that our method can be fruitfully used to connect such scenarios with observations of galaxies. We plan to expand our present results into this direction in forthcoming work. Our results highlight that EDE cosmologies lead to important modifications in the galaxy properties with respect to a standard \\lcdm universe. Nonetheless, they also show that these deviations are of the same order of magnitude as those induced by different assumptions in the modeling of physical properties and/or by parameter recalibrations in a homologous set of $\\Lambda$CDM-based SAMs \\citep[see also][]{Guo12}. Stronger effects are seen at higher redshift ($z>4$), but these redshifts correspond to a cosmic epoch where direct observational constraints on galaxies are extremely difficult to obtain and where we expect SAM predictions to be comparatively uncertain (given that their usual calibration samples are selected at lower redshifts). This behaviour is mainly due to our still limited understanding of the physical properties acting in the baryonic sector, and to the high level of degeneracy between the analytic approximations employed in different SAMs. We thus conclude that predicted galaxy properties alone provide only weak constraints on disentangling a standard \\lcdm cosmology from an Early Dark Energy model. These results are consistent with studies that analysed the response of SAMs against variations of cosmological parameters (corresponding to the best constraints from different data releases of the WMAP satellite) in the \\lcdm Universe \\citep[see e.g][]{Wang08, Guo12}. We have also shown that some of these degeneracies are broken if additional information on the properties and distribution of the underlying DM field became available. The combination of the stronger evolutionary patterns expected for the ``dark sector'' and the definition of galactic populations with suitably stable properties in SAMs indeed leads to the identification of reliable cosmological tests. In particular, we show the potential of a measurement of galaxy bias on scales larger than a few tenths of Mpc as a cosmological discriminant. Future observational efforts aimed at constraining $w$ and its redshift evolution, like the EUCLID mission \\citep{Laureijs11}, are indeed devised in order to provide, at the same time, constraints on the DM distribution (via weak lensing) and on the properties of large galaxy populations (via spectroscopy and photometry). We stress however, that this mission is designed to primarly focus on the $z\\sim1$ diagnostics, where the predicted effects of EDE cosmologies on galaxy bias are smaller than at higher redshifts. While we confirm that the combination of these different probes is indeed one of the most promising methods to achieve constraints on the dark energy component of our Universe, we also would like to emphasise the power of the particular simulation approach taken here, which is able to accurately predict galaxy bias in non-standard cosmologies as a function of scale, epoch, and galaxy type. The combination of high-resolution dark matter simulations and semi-analytic galaxy formation models used in this paper should also be highly useful to explore and understand generic effects of galaxy formation on cosmological constraints, for instance those derived from BAO or redshift space distortion measurements \\citep{Angulo12}. This would ultimately lead to an optimal and accurate exploitation of the next generation of galaxy surveys." }, "1207/1207.1515_arXiv.txt": { "abstract": "We present the result of our low-luminosity quasar survey in the redshift range of 4.5 $\\lesssim$ {\\it z} $\\lesssim$ 5.5 in the COSMOS field. Using the COSMOS photometric catalog, we selected 15 quasar candidates with 22 \\verb|<| {\\it i\\arcmin} \\verb|<| 24 at {\\it z} $\\sim$\\ 5, that are $\\sim$ 3 mag fainter than the SDSS quasars in the same redshift range. We obtained optical spectra for 14 of the 15 candidates using FOCAS on the Subaru Telescope and did not identify any low-luminosity type-1 quasars at $z\\sim5$ while a low-luminosity type-2 quasar at $z\\sim5.07$ was discovered. In order to constrain the faint end of the quasar luminosity function at $z\\sim5$, we calculated the 1$\\sigma$ confidence upper limits of the space density of type-1 quasars. As a result, the 1$\\sigma$ confidence upper limits on the quasar space density are $\\Phi<$ 1.33 $\\times$ 10$^{-7}$ Mpc$^{-3}$ mag$^{-1}$ for $-24.52 < M_{1450} < -23.52$ and $\\Phi<$ 2.88 $\\times$ 10$^{-7}$ Mpc$^{-3}$ mag$^{-1}$ for $-23.52 < M_{1450} < -22.52$. The inferred 1$\\sigma$ confidence upper limits of the space density are then used to provide constrains on the faint-end slope and the break absolute magnitude of the quasar luminosity function at $z\\sim5$. We find that the quasar space density decreases gradually as a function of redshift at low luminosity ($M_{1450}\\sim -23$), being similar to the trend found for quasars with high luminosity ($M_{1450}<-26$). This result is consistent with the so-called downsizing evolution of quasars seen at lower redshifts. ", "introduction": "The evolution of supermassive black holes (SMBHs) is now regarded as one of the most important unresolved issues in the modern astronomy, after the discovery of the galaxy-SMBH ``co-evolution\" inferred from, e.g., a tight relationship between the mass of SMBHs and their host bulges \\citep[e.g.,][]{2003ApJ...589L..21M, 2004ApJ...604L..89H, 2009ApJ...698..198G}. Measuring the whole shape of the quasar luminosity function (QLF) is particularly important to understand how the SMBHs grew, since it is highly dependent on some key parameters of SMBHs such as the growth timescale of SMBHs \\citep[e.g.,][]{ 2003PASJ...55..133E}. The QLF at $z\\lesssim3$ has been well quantified over a wide luminosity range \\citep[e.g.,][]{2009MNRAS.399.1755C} and is best represented by a double power law \\citep[e.g.,][]{1988MNRAS.235..935B, 1995ApJ...438..623P}. Recently, the faint end of the QLF at $z\\sim4$ has been measured \\citep{2010ApJ...710.1498G, 2011ApJ...728L..25I, 2011ApJ...728L..26G} and is also best represented by a double power law. More interestingly, recent studies on the optical QLF show that the space density of low-luminosity active galactic nuclei (AGNs) peaks at a lower redshift than that of more luminous AGNs \\citep{2009MNRAS.399.1755C, 2011ApJ...728L..25I}. This result can be interpreted as AGN (or SMBH) downsizing evolution, since the brighter AGNs tend to harbor the more massive SMBHs if the dispersion of the Eddington ratio of quasars is not very large \\citep[see, e.g.,][]{2011ApJ...733...60T}. The AGN downsizing has been also reported by X-ray surveys (\\citealt{2003ApJ...598..886U}; \\citealt{2005A&A...441..417H}; see also \\citealt{2009ApJ...693....8B} and \\citealt{2011ApJ...741...91C}). However, the physical origin of the AGN downsizing is totally unclear, that makes high-$z$ low-luminosity quasar surveys more important \\citep[see][for theoretical works on the AGN downsizing evolution]{2012MNRAS.419.2797F}. Recently, some low-luminosity quasar surveys have been performed at $z>5$ \\citep{2006AJ....132..823C,2005ApJ...634L...9M}. \\cite{2006AJ....132..823C} identified three quasars at $z>5$ with $z'<22$ and included a quasar at $z$ = 5.85 with $z'$ = 20.68, in the AGES survey which covers 8.5 $\\rm deg^2$. Jiang et al. (2008) also identified five new quasars at $z>5.8$ with $20 < z' < 21$ in the Sloan Digital Sky Survey (SDSS) deep stripe which covers 260 $\\rm deg^2$. The space density of quasars at $z\\sim6$ which is calculated by the result of Cool et al. (2006) is about six times larger than the result of Jiang et al. (2008). This large discrepancy may be caused by the small survey area of Cool et al. (2006). \\cite{2005ApJ...634L...9M} identified a very faint quasar at $z=5.70$ with $z' =$ 23.0 in the total quasar survey area of 2.5 $\\rm deg^2$. \\cite{2005ApJ...634L...9M} mentioned that the surface density of quasars at similar redshifts is roughly consistent with previous extrapolations of the faint end of the QLF. In this way, some low-luminosity quasars have been discovered although the faint end of QLF is not determined exactly, due to the lack of low-luminosity quasars. At $z>6$, a number of luminous quasars have been found up to $z\\sim6.5$ by the SDSS \\citep[e.g.,][]{2006AJ....131.1203F, 2006MNRAS.371..769G, 2008AJ....135.1057J, 2009AJ....138..305J} and the Canada-France High-$z$ Quasar Survey (CFHQS; \\citealt{2007AJ....134.2435W}; \\citealt{2009AJ....137.3541W}; \\citealt{2010AJ....139..906W}). Recently, a luminous quasar at $z=7.085$ has been found \\citep{2011Natur.474..616M} through the data obtained by the United Kingdom Infrared Telescope Infrared Deep Sky Survey \\citep[UKIDSS;][]{2007MNRAS.379.1599L}. Although some low-luminosity quasars at $z>5$ have been discovered as mentioned above, the faint-end slope of the $z>5$ QLF is still very poorly constrained. Consequently it is not understood how low-luminosity quasars evolve at high redshifts, or if the AGN downsizing evolution is also seen in the earlier universe. Since the number density of low-luminosity quasars is expected to be much higher than that of high-luminosity quasars, the whole picture of SMBH evolution cannot be understood without firm measurements of the number density of low-luminosity quasars at such high redshifts. Motivated by these issues, we have searched for low-luminosity quasars at $z\\sim5$ in the COSMOS field \\citep{2007ApJS..172....1S}. In Section 2, we describe the data and method that were used for the photometric selection of quasar candidates. In Section 3, we report the results of the follow-up spectroscopic observations. In Section 4, we describe how we estimate the photometric completeness to derive the QLF. In Section 5, we present the upper limits of the QLF at $z\\sim5$ and briefly discuss it. Throughout this paper we use a $\\Lambda$ cosmology with $\\Omega_m$ = 0.3,\\ $\\Omega_{\\Lambda}$ = 0.7, and the Hubble constant of $H_0$ = 70 km s$^{-1}$ Mpc$^{-1}$.\\\\ ", "conclusions": "In order to examine the faint end of the QLF at $z\\sim5$, we select photometric candidates of quasars at $z\\sim5$ in the COSMOS field. The main results of this study are: \\begin{itemize} \\item Although we discover the type-2 quasar at $z\\sim5.07$, no type-1 quasars at $z\\sim5$ are identified through the follow-up spectroscopic observation. \\item The upper limits on the type-1 quasar space density are $\\Phi<$ 1.33 $\\times$ 10$^{-7}$ Mpc$^{-3}$ mag$^{-1}$ for $-24.52 < M_{1450} < -23.52$ and $\\Phi<$ 2.88 $\\times$ 10$^{-7}$ Mpc$^{-3}$ mag$^{-1}$ for $-23.52 < M_{1450} < -22.52$. \\item The quasar space density and its error when we include a type-2 quasar at $z\\sim5.07$ are $\\Phi=$ $0.87_{-0.72}^{+2.01}$ $\\times$ 10$^{-7}$ Mpc$^{-3}$ mag$^{-1}$ for $-23.52 < M_{1450} < -22.52$. \\item The derived upper limits on the quasar space density are consistent with the QLF inferred by the previous works at $z\\sim5$. \\item The characteristic absolute magnitude of the QLF shows a significant redshift evolution between $z\\sim4$ ($M^{*}_{1450}>-26$) and $z\\sim5$ ($M^{*}_{1450}<-26$). \\item A continuous decrease of the space density of low-luminosity ($-24\\lesssim M_{1450}\\lesssim -23$) quasars is inferred, that is roughly consistent with the picture of the AGN downsizing evolution. \\end{itemize}" }, "1207/1207.4610_arXiv.txt": { "abstract": "We search for ``solar twins'' in the Geneva-Copenhagen Survey (GCS) using high resolution optical spectroscopy. We initially select Sun--like stars from the GCS by absolute magnitude, $(\\mathrm{b}-\\mathrm{y})$ colour and metallicity close to the solar values. Our aim is to find the stars which are spectroscopically very close to the Sun using line depth ratios and the median equivalent widths and depths of selected lines with a range of excitation potentials. We present the ten best stars fulfilling combined photometric and spectroscopic criteria, of which six are new twins. We use our full sample of Sun--like stars to examine the calibration of the metallicity and temperature scale in the GCS. Our results give rise to the conclusion that the GCS may be offset from the solar temperature and metallicity for sun-like stars by 100\\,K and 0.1\\,dex, respectively. ", "introduction": "``Solar twins'' (or ``solar analogues'') are stars which are very close matches to the spectroscopic and photometric properties of the Sun \\citep{b6}, and while there is currently no ``perfect twin'', a few stars are known which are very close matches to the Sun. For over a decade, the star considered most similar to the Sun has been 18\\,Sco/HD\\,146233 \\citep{b2,b1}, and recent asteroseismological and interferometric measurements have confirmed its radius and mass to be solar within a few percent \\citep{b33}. Based on its spectroscopic properties, HD\\,98618 was considered the next best solar twin by \\citet{b3}, while \\citet{b13} have found HIP\\,100963 to be as good a twin as 18\\,Sco (although both stars have a higher Li abundance than the Sun by 0.5 and 0.8\\,dex respectively). Currently, HD\\,56948 is the star considered closest to the Sun \\citep{b11} -- its lithium abundance is very similar to the Sun, and together with another close twin (HIP\\,73815), it shows that the solar lithium abundance is not atypical. Some tens of solar twins (or solar analogues) have been published in recent years \\citep{b13,b32}; and a solar twin has even been identified in the open cluster M67 \\citep{b5}. To date all solar twins show small but interesting differences with respect to the Sun: such as the lithium abundance being high \\citep{b3} or the stars being variable and showing chromospheric activity (e.g. 18\\,Sco, \\citealt{b4}). Solar twins are useful because, obviously enough, one cannot point the same instruments/telescopes at the Sun as used on faint objects in the night sky. This gives rise to a major difficulty with calibrating the stellar metallicity and temperature scale to the same scale as the Sun, in which these parameters are measured (notwithstanding that the absolute solar metallicity has been under considerable discussion in the last decade, \\citealt{b20}). The largest extant sample of solar type stars (F, G and K dwarfs), the Geneva-Copenhagen-Survey (hereafter GCS, \\citealt{b9}), has a photometric metallicity and temperature scale which is tied to the stellar colours. Recently, the metallicity and temperature calibrations have been called into question by \\citet{b24}, who used the InfraRed Flux Method on 423 stars, and the properties of 10 solar twins, to argue that the GCS scale may need shifting by about 100\\,K in temperature and 0.1\\,dex in metallicity. We address this topic in this paper, presenting a new method for testing the scales in the GCS catalogue using solar analogues, and finding a similar shift. Interest has revived in solar twins since the advent of exoplanet detection. Interestingly, the first discovered exoplanet host, 51 Peg \\citep{b7}, is a very good solar twin \\citep{b6}. Exoplanets have been found around other solar twins \\citep{b28,b29}, but to date no systematic searches have been undertaken. One of our aims is to provide a list of nearby solar twins by searching systematically for them in the GCS. Recently \\citet{b32} and \\citet{b15} have suggested that whether a star hosts exoplanets or not might be revealed in the detailed chemical composition obtained from high resolution spectra. Another motivation for this study is to provide new targets to probe this correlation further. In this paper we present the results of our own quest for solar twins, based on photometric selection from the GCS catalogue, combined with high resolution spectroscopic data. Most previous surveys looking for solar twins have focused on objects of the Northern Hemisphere, such as at the Observatoire des Haute Provence \\citep{b1}, Keck \\citep{b3,b8} or the McDonald Observatory in Texas \\citep{b15}. We explore the relatively understudied Southern Hemisphere from the Max-Planck-Gesellschaft (MPG)/European Southern Observatory (ESO) 2.2m at La Silla Observatory, giving us the opportunity to extend our search to targets not considered previously, and we turn up with six new twins. In this paper we provide a description of our photometric candidate selection process in section 2; in section 3 we outline the observations and data reduction; in section 4 we present our methods for finding solar twins and the results. In section 5 we use our full spectroscopic sample of about a hundred Sun--like stars to test the temperature and metallicity scale in the GCS, by differential comparison to the solar spectrum, and we draw our conclusions in section 6. \\section[]{Candidate selection} Our candidate solar-twins were selected from the first release of the Geneva-Copenhagen-Survey (GCS-I) \\citep{b9} of Stromgren colours, absolute magnitudes, metallicity and temperature estimates for $\\sim14\\,000$ nearby F to K type stars. We selected stars bracketting the solar $(\\mathrm{b}-\\mathrm{y})$ colour, absolute visual magnitude $\\mathrm{M}_\\mathrm{V}$ and metallicity, for which we adopt the solar values of $(\\mathrm{b}-\\mathrm{y})_\\odot = 0.403$ \\citep{b19}, $\\mathrm{M}_\\mathrm{V} = 4.83$ \\citep{b22} and [Fe/H] $=0.0$ (by definition). Our ranges are : $0.371 < (\\mathrm{b}-\\mathrm{y}) < 0.435$, in absolute magnitude $4.63 < \\mathrm{M}_\\mathrm{V} < 5.03$ and in metallicity (GCS-I scale) $-0.15 < [\\mathrm{Fe/H}] < 0.15$. These criteria resulted in 338 stars, of which 80 were chosen for the proposal as accessible from La Silla, and not already available in the Fiber-fed Extended Range Optical Spectrograph (FEROS) archive. In the end, 70 of these 80 targets were observed (all in service mode). In Fig.\\,\\ref{proposal} we show the colour-magnitude diagram (CMD) of the main sequence and turnoff stars from the GCS-I catalogue. The dots show our initial photometric twin candidates, chosen for spectroscopic follow--up. An interesting effect appears if we show the 10 very good spectroscopically matched solar twins found by \\citet{b1} (triangles). Their twins tend to lie redward (cooler) and tend to be brighter than the Sun. This highlights the point (concluded by those authors) that purely spectroscopic matching might still yield stars with systematic differences to the Sun in other properties -- thus, in this paper, we examine the effects of using both spectroscopic and photometric criteria to find solar twins. \\begin{figure} \\includegraphics[width=90mm]{proposalplotnew.ps} \\caption{Stars from the Nordstr\\\"om et al. (2004) catalogue (GCS-I) (dots), in the metallicity range of $-0.15 < [\\mathrm{Fe/H}] < 0.15$, together with our initial twin candidates (circles) and the candidates from \\citet{b1}(triangles). The Sun is marked in the middle of the box. The solar twin candidates of the \\citet{b1} sample (triangles) tend to lie redward (cooler) and tend to be intrinsically brighter than the Sun.} \\label{proposal} \\end{figure} While this project was being undertaken, two revisions of the Geneva-Copenhagen Survey (GCS-II and GCS-III) \\citep{b27, b21} were published. Both revisions addressed the metallicity, age and temperature scales of the stars, and utilised the updated Hipparcos parallaxes \\citep{b26} when they became available. This resulted in increased temperatures, which went up by an average of 80\\,K at solar values and in decreased metallicities, which went down by an average of 0.05\\,dex. With the revised absolute magnitudes and metallicities, some of our target stars moved slightly out of the original selection window while others moved in. These new arrivals have been included in our sample, as we searched the FEROS archive at ESO and found spectra for 28 of these candidates. We also included stars in our sample which have been classified as 'Sun-like' stars by others \\citep{b12,b13,b14,b15,b1}. This yielded another 47 stars, bringing our total sample to 145 objects. They range in apparent $V$ magnitudes from 3.5 to 9, with the majority lying around $V = 8$. \\begin{figure} \\includegraphics[width=90mm]{wholesample.ps} \\caption{As Fig.\\,\\ref{proposal}, showing our target sample as open circles (from GCS-I and GCS-III) and the targets taken from other papers (filled circles). The Sun is shown in the middle of the box, which represents our original selection window.} \\label{newsample} \\end{figure} Fig.\\,\\ref{newsample} shows our final spectroscopic sample in the $\\mathrm{M}_{\\mathrm{V}}$ vs. $(\\mathrm{b}-\\mathrm{y})$ plane. Our basic candidate stars, initially selected from the GCS-I, but with revised GCS-III data and stars which entered the selection window as a result of the revised GCS data, are shown as open circles. Finally, Sun-like stars from other studies in the literature are shown as filled circles. The Sun's location is shown in the middle of our selection box. Some stars selected from the literature lie very far in absolute magnitude and colour from our selection window. This is because often the authors adopted a broad definition of Sun--like stars, sometimes extending to F and K dwarfs \\citep{b25} or metal-rich stars \\citep{b12}. We decided to keep all those stars in our analysis, to have a chance to test broad trends with temperature and metallicity in the spectroscopic selection of our best twin candidates. ", "conclusions": "In this paper we use high resolution optical spectroscopy to search for solar twins in the Geneva Copenhagen Survey, by applying various methods adopted from recent literature. We have shown that there is no unique way to search for a solar twin and that it is necessary to include photometric as well as spectroscopic selection criteria to really determine which stars are solar twins. We confirm HD\\,146233 (18\\,Sco) as the best twin in our list, being selected by all 4 spectroscopic methods used; HD\\,126525 and HD\\,138573 are second best, being selected in 3 out of 4 methods; 6 out of our 10 twins are new additions to the literature; HD\\,117860, HD\\,97356, HD\\,142415, HD\\,163441, HD\\,173071 and HD\\,126525. We use our entire sample to probe for offsets in the temperature and metallicity scale in the GCS for Sun--like stars, introducing a new method (degeneracy lines method) which disentangles the differing metallicity and temperature degeneracies in the measured indices for our stars. We estimate that, for Sun--like stars, the GCS-III scale is offset by $(-0.12\\pm0.02)$\\,dex and $(-97\\pm35)$\\,K respectively -- i.e. we find it is a little too metal poor and cool. This result is in good agreement with similar offsets claimed in recent literature, based on solar twins: \\citet{b23} find the GCS values to be $\\Delta$T = 48\\,K too cool and $\\Delta$[Fe/H] = 0.09 too metal poor. \\citet{b24} finds offsets of about $-$100\\,K and $-$0.1\\,dex, respectively. Our new method has been successfully tested both internally in GCS (i.e.\\ recovering the metallicity and temperature of random reference stars, which replaced the Sun/Ceres for the sake of the test) and on theoretical spectra. We are currently applying a similar degeneracy lines method to the very high quality High Accuracy Radial velocity Planet Searcher (HARPS) archive spectra, using both neutral and ionised species for more elements than just Fe. Early results confirm the offsets found here and will be discussed in a forthcoming paper (Datson et al., in preparation). Despite the agreement of our results with other studies for an offset in the temperature and metallicity scales of GCS for Sun--like stars, we point out that recent measurements of the temperatures of Sun-like stars via interferometry with the Center for High Angular Resolution Astronomy (CHARA) array instrument \\citep{b37} show, for over a dozen stars with the most secure angular diameters ($>$1~mas) {\\it excellent agreement} with GCS temperatures (Holmberg, private comm.). This contrasts with the conclusions drawn by us and other authors, based on solar twins and Sun-like stars. We leave the discussion of this intriguing result to future work. The offsets we find in the GCS would imply the solar $(\\mathrm{b}-\\mathrm{y})$ colour to be $(\\mathrm{b}-\\mathrm{y}) = 0.414\\pm0.007$, also determined via our degeneracy--lines approach. This colour is redder than the $(\\mathrm{b}-\\mathrm{y})=0.403\\pm0.013$ found earlier by our group (Holmberg et al. 2006) but very close to the recent result of Melendez et al.\\ (2010), based on solar twins. One of our best solar twins, HD\\,126525, has a temperature and metallicity, that are so offset from solar (5585 K and $-0.19$ dex, see Table~6), that even the proposed corrections to the GCS scale still leave it in tension with the solar values. We cannot rule out that the twin selection methods used here are still affected by systematics; a more detailed study of the individual spectra and of metallicity-temperature degeneracy issues is currently underway. Three of our twins are known to host an exoplanet: HD\\,142415, HD\\,147513 \\citep{b30} and HD\\,126525 \\citep{b31}. There have been no confirmed detections of exoplanets for the other seven twins so far, which will hopefully be included as targets for future planet searches." }, "1207/1207.6086_arXiv.txt": { "abstract": "We investigate the launching of jets and outflows from magnetically diffusive accretion disks. Using the PLUTO code we solve the time-dependent resistive MHD equations taking into account the disk and jet evolution simultaneously. The main question we address is {\\em which kind of disks do launch jets and which kind of disks do not?} In particular, we study how the magnitude and distribution of the (turbulent) magnetic diffusivity affect mass loading and jet acceleration. We have applied a turbulent magnetic diffusivity based on $\\alpha$-prescription, but have also investigate examples where the scale height of diffusivity is larger than that of the disk gas pressure. We further investigate how the ejection efficiency is governed by the magnetic field strength. Our simulations last for up to 5000 dynamical time scales corresponding to 900 orbital periods of the inner disk. As a general result we observe a continuous and robust outflow launched from the inner part of the disk, expanding into a collimated jet of super fast magneto-sonic speed. For long time scales the disk internal dynamics changes, as due to outflow ejection and disk accretion the disk mass decreases. For magneto-centrifugally driven jets we find that for i) less diffusive disks, ii) a stronger magnetic field, iii) a low poloidal diffusivity, or a iv) lower numerical diffusivity (resolution), the mass loading of the outflow is increased - resulting in more powerful jets with high mass flux. For weak magnetization the (weak) outflow is driven by the magnetic pressure gradient. We consider in detail the advection and diffusion of magnetic flux within the disk and we find that the disk and outflow magnetization may substantially change in time. % This may have severe impact on the launching and formation process - an initially highly magnetized disk may evolve into a disk of weak magnetization which cannot drive strong outflows. We further investigate the jet asymptotic velocity and the jet rotational velocity in respect of the different launching scenarios. We find a lower degree of jet collimation than previous studies, most probably due to our revised outflow boundary condition. ", "introduction": "Jets as highly collimated beams of high velocity material and outflows of comparatively lower degree of collimation and lower speed are an ubiquitous phenomenon among astrophysical objects. Jets are powerful signs of activity and are observed over a wide range of luminosity and spatial scale. Among the jet sources are young stellar objects (YSO), micro-quasars, active galactic nuclei (AGN), and most probably also gamma ray bursts \\citep{Fanaroff1974, Abell1979, Mundt1983, Rhoads1997, Mirabel1994Na}. The common models of launching, acceleration, and collimation work in the framework of magnetohydrodynamic (MHD) forces (see e.g.~\\citealt{BP1982, Pudritz1983, Uchida1985}), although the details of the process are not fully understood. Jets and outflows from YSO and AGN affect their environment, and, thus, the formation process of the objects they are launching them. Numerous studies investigate effects of such feedback mechanisms in star formation and galaxy formation (see e.g. \\citealt{Banerjee2007,Carroll2009, Gaibler2011}). However, a quantitative investigation of how much of mass, momentum, or energy from the infall is actually recycled into a high speed outflow needs to resolve the innermost jet-launching region and to model the physical process of launching directly. This is the major aim of the present paper. According to the current understanding, accretion and ejection are related to each other. One efficient way to remove angular momentum from a disk is to connect it to a magnetized outflow. This has been motivated by the observed correlation between signatures of accretion and ejection in jet sources (see e.g. \\citealt{Cabrit1990, Hartigan1995}). The overall idea is that the energy and angular momentum are extracted from the disk by an efficient magnetic torque relying on a global, i.e. large-scale magnetic field threading the disk. If the inclination of the field lines is sufficiently small, magneto-centrifugal forces can accelerate the matter along the field line. Beyond the Alfv\\'en point also Lorentz forces contribute to the acceleration. The collimation of the outflow is thought to be achieved by magnetic tension due to a toroidal component of the magnetic field. Still, we have to keep in mind the fact the the toroidal field pressure gradient acts de-collimating, and an existing external gas over-pressure may contribute to jet collimation. Before going into further details, we like to make clear that with jet {\\em formation} we denote the process of accelerating and collimating an already existing slow disk wind or stellar wind in to a jet beam. With jet {\\em launching} we denote the process which conveys material from radial accretion into a vertical ejection, thereby lifting it from the disk plane into the corona, and thus establishing a disk wind. A vast literature exists on magnetohydrodynamics jet modeling. We may distinguish i) between steady-state models and time-dependent numerical simulations, or ii) between simulations considering the jet formation only from a fixed-in-time disk surface and simulations considering also the launching process, thus taking into account disk and jet evolution together. Steady-state modeling have mostly followed the self-similar Blandford \\& Payne approach (e.g. \\citep{Sauty1994AA, Contopoulos1994C}), but also fully two-dimensional models were proposed \\citep{Pelletier1992, Li1993}, some of them taking even into account the central stellar dipole \\citep{Fendt1995, Paatz1996}. Further, some numerical solutions have been proposed by .\\citep{konigl2010M, Salmeron2011S, Wardle1993} in a weakly ionized protostellar accretion discs that are threaded by a large-scale magnetic field as a wind-driving accretion disk. They have studied the effects of different regimes for ambipolar diffusion or Hall and Ohm diffusivity dominance in these disk. (Self-similar) steady-state models have also been applied to the jet launching domain \\citep{Ferreira1995, Li1995,Casse2000}, connecting the collimating outflow with the accretion disk structure. In addition to the steady-state approaches, the magneto-centrifugal jet formation mechanism has been subject of a number of time-dependent numerical studies. In particular, \\citet{Ustyugova1995} and \\citet{Ouyed1997} have demonstrated for the first time the feasibility of the MHD self-collimation property of jets. We note, however, that it was already 1985 when the first jet formation simulations were published, in that case for much shorter simulation time scale \\citep{Shibata1985}. Among these works, some studies investigated artificial collimation \\citep{Ustyugova1999}, a more consistent disk boundary condition \\citep{Krasnopolsky1999}, the effect of magnetic diffusivity on collimation \\citep{Fendt2002}, or the impact of the disk magnetization profile on collimation \\citep{Fendt2006}. In the aforementioned studies, the jet-launching accretion disk is taken into account as a boundary condition, {\\em prescribing} a certain mass flux or magnetic flux profile in the outflow. This is a reasonable setup in order to investigate jet formation, i.e. the acceleration and collimation process of a jet. However, such simulations cannot tell the efficiency of mass loading or angular momentum loss from flux from disk to jet, or cannot help answering the question which kind of disks do launch jets and under which circumstances. It is therefore essential to extend the jet formation setup and include the launching process in the simulations. Numerical simulations of the MHD jet launching from accretion disks have been presented first by \\citet{Kudoh1998} and \\citet{Casse2002}, treating the ejection of a collimated outflow out of an evolving-in time resistive accretion disk. \\citet{Zanni2007} further developed this approach with emphasis on how resistive affect the dynamical evolution. An additional central stellar wind was considered by \\citet{Meliani2006}. Further studies were concerned about the effects of the absolute field strength or the field geometry, in particular investigating field strengths around and below equipartition \\citep{Kuwabara2005, Tzeferacos2009, Murphy2010}. These latter were {\\em long-term} simulations for several 100s of (inner) disk orbital periods, providing sufficient time evolution to also reach a (quasi) steady state for the fast jet flow. Finally, we like to emphasize the fact that jets and outflows are observed as {\\em bipolar} streams. Jet and counter jet appear typically asymmetric with only very few exceptions. One exception is the protostellar jet HH\\,212 showing an almost perfectly symmetric bipolar structure \\citep{Zinnecker1998}. So far, only very few numerical simulations investigating the bipolar launching of disk jets have been performed. Among them are the works of \\citet{Rekowski2003} or \\citet{Rekowski2004} which even included a disk dynamo action. Recent publications consider asymmetric ejections of stellar wind components from an offset multi-pole stellar magnetosphere \\citep{Long2008M,Lovelace2010, Long2012}. It is therefore interesting to investigate the evolution of both hemispheres of a {\\em global} jet-disk system in order to see whether and how a global asymmetry in the large-scale outflow can be governed by the disk evolution. This has not been done so far. In a series of two papers, we will address both the detailed physics of MHD jet launching (paper I) and the bipolarity aspects of jets (paper II). In the present paper (paper I) we investigate the details of the launching, acceleration, and collimation of MHD jets from resistive magnetized accretion disks. We investigate how the mass and angular momentum fluxes depend on the internal disk physics applying long-lasting global MHD simulations of the disk-jet system. We hereby investigate different resistivity profiles of the disk. This paper is organized as follows. Section 2 is dedicated to MHD equations and to describe the numerical setup, the initial and boundary conditions of our simulations. The general evolution of jet launching and the physical processes involved are presented in section 3 with the help of a reference simulation. Section 4 is then devoted to a parameter study comparing jets from different setups. In paper II we will present the bipolar jets simulations, discussing their symmetry properties and how symmetry can be broken by the intrinsic disk evolution. ", "conclusions": "We have presented results of MHD simulations investigating the launching of jets and outflows from a magnetically diffusive disk in Keplerian rotation. The time evolution of the accretion disk structure is self-consistently taken into account. The simulations are performed in axisymmetry applying the MHD code PLUTO. The main goal of our simulations was to study how magnetic diffusivity (its magnitude and distribution) and magnetization affect the disk and outflow properties, such as mass and angular momentum fluxes, jet collimation, or jet radius. Our grids extend to $(96 \\times 288)$ inner disk radii with a resolution of $(0.064 \\times 0.066)$, respectively $(50 \\times 180)$ inner disk radii with a higher resolution of $(0.025 \\times 0.025)$. An internal boundary (sink) is placed close to the origin absorbing the accreted mass and angular momentum. We have prescribed a magnetic diffusivity in the disk based on an $\\alpha$-prescription. One of our parameters was the scale height of the magnetic diffusivity with the option to have it higher than the thermal scale height. This can be motivated considering that it is the turbulent disk material which is loaded into the outflow, and that the turbulence pattern is swept along with the disk wind until it decays. We have investigated disks carrying a magnetic flux corresponding to an initial plasma-beta ranging from 10 to 5000 at the inner disk radius. As a general result we observe a continuous and robust outflow launched from the inner part of the disk, expanding into a collimated jet and is accelerated to super fast magneto-sonic speed. The key results of our simulations can be summarized as follows. (1) Concerning the {\\em acceleration of the outflows}, our simulations confirm that the magneto-centrifugal acceleration mechanism is most efficient in the low plasma-$\\beta$ regime, while for weak magnetic fields the toroidal magnetic pressure gradient drives the ejected material. We also confirm that the magnetocentrifugal mechanism also depends on the mass load in the outflow, as this mechanism works more efficiently for outflows with low mass fluxes. However, compared to the magnetic pressure driven outflows, jets in the magnetocentrifugal acceleration regime have usually higher mass fluxes. In our simulations with very high plasma-$\\beta$ we detect a highly unsteady behaviour (2) Efficient magneto-centrifugal driving which can accelerate jets to high kinetic energy relies on a strong coupling between magnetic field and the rotating disk, thus on a low diffusivity. We find that it is the poloidal diffusivity $\\eta_{\\rm p}$ which mainly affects the driving of the outflow. However, besides the coupling needed for acceleration, also the {\\em launching of material} depends on diffusivity. With increasing $\\eta_{\\rm p}$, the mass fluxes (both the accretion rate and the ejection rate) decrease. Subsequently, the higher ejection rates result in a lower asymptotic outflow velocities. (3) We measure typical {\\em outflow velocities} are in the range of $0.3 - 0.8$ times the inner disk rotational velocity with the tendency that the mass fluxes obtained in magneto-centrifugally driven outflow are substantially higher. Here, we confirm the clear (inverse) correlation between jet velocity and mass load, as it is well known from the literature. Note that we have apply a mass-flux-weighted jet velocity which we find more applicable to the observations. For the bulk mass flux we find lower velocities compared to other papers, which are mostly dealing with the maximum speed obtained in the simulation. Our maximum velocities are similar. The relatively high speed of the outflows with low Poynting flux which are driven by poloidal pressure gradient is unclear. We find that the toroidal diffusivity affects the {\\em outflow rotation} - a small toroidal diffusivity implies a larger jet rotational velocity. This has not shown before in simulations. (4) We do not find a clear correlation between the {\\em outflow collimation} and the magnetic field strength. Also weakly magnetized outflows, which are driven by the magnetic pressure gradient, and which we find to be quite unsteady, show a high degree of collimation. The question of collimation for the weakly magnetized outflows is not answered. We find that outflows within the magneto-centrifugal-driving regime the flows ejected from a weakly diffusive disks are only weakly collimated. Similarly, their jet radius (here defined as mass flux weighted radius) is larger in case of a lower poloidal magnetic diffusivity. Follwoing the magnetic flux surface along the bulk mass flux from the asymtotic regime to the launching area, we can defined the {\\em launching area} of the outflow. We find a size of the launching area from which the bulk of the mass flux originates in the range between 3 and 8 inner disk radii or about 0.4 AU. (5) Depending on the strength of magnetic diffusivity, the disk-jet structure may evolve into a steady-state. We found that the cases with the strong field with $\\beta_{\\rm i} \\sim 10$ and poloidal diffusivity $\\eta_{\\rm p, i} \\geq 0.03$ will reach a quasi steady state, confirming the literature. (6) The magnetic flux profile along the disk is subject to advection and diffusion. We find that the magnetization (or plasma-$\\beta$) of disk and outflow may therefore substantially change during the time evolution. We have observed that the initial disk magnetization may change by a factor of 100. This may have severe impact on the launching process and the formation of the outflow in the sense that a rather highly magnetized disk may evolve into a weakly magnetized disk which cannot drive strong outflows. This issue has not been discussed before in the literature. (7) For very long time scales the accretion disk changes its internal dynamics, as due to outflow ejection and disk accretion the disk mass decreases. As a consequence, the accretion and ejection rates slightly decrease. In order to compensate for this effect, we have applied a large outer disk radius providing a large mass reservoir for the inner jet-launching disk. (8) For our simulations, we found that 10-50\\% of the accreting plasma can be diverted into the outflow. For i) less diffusive disks, ii) a strong magnetic field, iii) a low poloidal diffusivity, or a iv) lower numerical diffusivity (resolution), the mass loading into the outflow is increased - resulting in more massive jets. We interpret as physical reason the more efficient extraction of angular momentum from the disk, due to the stronger matter-field coupling. Note that we do not consider in our simulations viscosity or the wealth of thermal effects which play an essential role for launching \\citep{Casse2000}. (9) We found that jets launched in a setup with smaller diffusive scale height are more perturbed. The same effect is seen in outflows launched from disks with weaker diffusivity. (10) We finally remark that it is essential to do long-term simulations covering thousands of rotational periods in order to find a steady state situation of the accretion-ejection dynamics. Our simulations run for 5000 dynamical time steps, corresponding to about 900 revolutions at the inner disk radius as adopted in our reference run. In summary, we confirm the hypothesis that efficient magneto-centrifugal jet driving requires a strong magnetic flux (i.e. a low plasma beta), together with a large enough magnetic torque in order to produce a powerful jet. In addition, the magnitude of (turbulent) magnetic diffusivity plays the major role in the ejection efficiency, while the anisotropy in the diffusivity mainly affects the jet rotation. Both results imply that the structure of the asymptotic jet is indeed governed by the properties of the accretion disk, here parametrized by the magnetization and magnetic diffusivity. The mass ejection-to-accretion ratio along with the momentum and energy transfer rates from inflow to outflow are essential properties for any feedback mechanism in star formation or galaxy formation scenarios and could only be derived from simulations resolving the inner region of the jet-launching accretion disk." }, "1207/1207.1459_arXiv.txt": { "abstract": "Since its launch in 2008 the Fermi Large Area Telescope provides regular monitoring of a large sample of gamma-ray sources on time scales from hours to years. Together with observations at other wavelengths it is now possible to study variability and correlation properties in a much more systematic and detailed way than ever before. The paper describe some of the statistical methods and tools that have been, or can be, used to characterize variability and to study the relation between multiwavelength light curves. Effects and limitations due to time sampling, measurement noise, non-stationarity etc are illustrated and discussed. ", "introduction": "The combination of gamma-ray and radio observations together with measurements in other wavelength bands has created extraordinary opportunities to study multiwavelength properties of Active Galactic Nuclei, in particular for blazars. Among the aims are to better answer questions like, how are the components of the spectral energy distribution (SED) related? Do the components originate from one or more spatial regions? Does Compton seed photons originate locally or from a source external to the jet? Are hadronic cascades an important contribution to the gamma-ray emission? There are theoretical predictions of potentially observable effects that can be searched for, such as spectral softening associated with particle cooling during flare decay~\\cite{dotson2011}. From an observational point of view however, the first aim is to characterize source variability and correlations between different spectral bands. Simultaneously the wealth of new data also allow us to search for new and previously unobserved phenomena. The increasing amount of multiwavelength data on blazars allow us to attack some of the unresolved questions in a statistically more comprehensive way instead of studies based on a restricted number of individual cases. The importance of this is particularly clear when one considers the complex and apparently unsystematic multiwavelength behaviour revealed by earlier investigations. Available data from different spectral bands differ in terms of Signal-to-Noise, time sampling, observation length etc. The interpretation of observed variability and correlations is complicated not only by measurment noise and sampling but also by the stochastic variability itself, which means that observed variability properties may change with time and that even unrelated time series can exhibit chance correlations in time limited observations. All these effects needs to be considered in the analysis and interpretation of the observations. \\subsection{Data Properties} Fermi Large Area Telescope (LAT) has a large field of view, covering about 20 $\\%$ of the sky at any particular time. Except for some shorter pointed observations it is operated in a sky surveying mode, mapping the full sky every 3 hours. This provides a regular monitoring of all the gamma-bright blazars on the sky. In total, about 1000 blazars were detected in the data of the first two years of operation and included in the second LAT AGN catalog~\\cite{2LAC}. The signal-to-noise is such that a few tens of sources are typically detected in time bins of one to a few days, while hundreds of AGNs are significantly measured in time bins of one to a few weeks. The regular sampling of the Fermi LAT observations is a great advantage in the data analysis, not just for the gamma-ray analysis itself but also for the analysis of observations at other wavelengths. Radio and optical observations are often of high signal-to-noise but the time sampling is in most cases less regular. It is then an advantage that the sparse data can be compared with a more densely covered light curve. One limitation with Fermi light curves is however, that it is often not possible to choose a single bin width that both resolve variability at high flux levels and at the same time detect the source at low flux levels. This commonly results in upper limits which are hard to handle in the high level analysis. A way to remedy this is to use a non-constant bin width chosen to give approximately the same signal-to-noise for each bin. A procedure to create such adaptive binned light curves for Fermi data is described in~\\cite{lott2012}. An alternative approach based on Bayesian Blocks has also been developed~\\cite{scargle2012}. While Fermi has been operating for 3.5 years many blazars have been followed in optical and radio for tens of years. This provides valuable information on variability on longer time scales than is accessible by Fermi. This has a bearing on for example duty cycles and on the question of non-stationarity of the variability. \\subsection{Analysis tools} The analysis tools that have been applied to study blazar variability include, excess variance, flare profile fitting, flux distribution (duty cycles), power density spectra, auto correlation function, structure function, and wavelets. Multiwavelength correlations can be investigated by e.g. direct light curve comparisons, flux - flux plots, the cross correlation function and the cross spectrum. An overview of these methods with direct relevance to blazar variability analysis is given in~\\cite{scargle2011}. Here we will instead focus on some of the practical aspects of cross correlation analysis. ", "conclusions": "" }, "1207/1207.1942_arXiv.txt": { "abstract": "We construct magnetized stars composed of a fluid stably stratified by entropy gradients in the framework of general relativity, assuming ideal magnetohydrodynamics and employing a barotropic equation of state. We first revisit basic equations for describing stably-stratified stationary axisymmetric stars containing both poloidal and toroidal magnetic fields. As sample models, the magnetized stars considered by Ioka and Sasaki~\\cite{Ioka2004}, inside which the magnetic fields are confined, are modified to the ones stably stratified. The magnetized stars newly constructed in this study are believed to be more stable than the existing relativistic models because they have both poloidal and toroidal magnetic fields with comparable strength, and magnetic buoyancy instabilities near the surface of the star, which can be stabilized by the stratification, are suppressed. ", "introduction": "Recent observations established that soft-gamma repeaters (SGRs) and anomalous x-ray pulsars (AXPs) are the so-called magnetars, i.e., highly magnetized neutron stars whose surface field strength is as large as $\\sim 10^{14}-10^{15}$~G~\\cite{duncan,paczynsk,thompsona,thompsonb,woods}. The presence of the magnetars has activated studies on equilibrium configurations of magnetized stars, which have a long history. Since the pioneering work by Chandrasekhar and Fermi \\cite{chandra}, an enormous number of studies has been done for exploring structures of magnetized stars: Prendergast~\\cite{prendergast} and Woltjer \\cite{woltjer} calculated equilibrium configurations of magnetized stars having mixed poloidal-toroidal fields, where the magnetic fields are treated as first-order perturbations around a spherical star (see, also, Ref.~\\cite{roxburgh}). Monaghan \\cite{monaghan} studied magnetized stars containing purely poloidal magnetic fields. Ioka~\\cite{ioka} developed the works by Prendergast and Woltjer into those of second-order perturbations to study magnetic effects on the stellar structures. He also employed the results obtained to explain magnetar activities. Miketinac obtained magnetized stars containing purely toroidal fields~\\cite{miketinac} and purely poloidal fields \\cite{miketinac2} by solving exact master equations numerically. Tomimura and Eriguchi developed a numerical method for obtaining magnetized stars with mixed poloidal-toroidal fields using a non-perturbative technique \\cite{tomimura} (see, also, Refs.~\\cite{yoshida0,yoshida,lander}). Duez and Mathis variationally considered the lowest-energy equilibrium states for a fixed magnetic helicity and constructed equilibria of magnetized stars having mixed poloidal-toroidal fields by a perturbation technique~\\cite{duez}. General relativistic models of magnetized neutron stars have been also explored. Bocquet et al.~\\cite{bocquet} and Cardall et al.~\\cite{cardall} obtained relativistic neutron star models with purely poloidal magnetic fields. Using a perturbative technique, Konno et al.~\\cite{konno} calculated similar models. Kiuchi and Yoshida \\cite{kiuchi} computed magnetized stars with purely toroidal fields. Ioka and Sasaki~\\cite{Ioka2004}, Colaiuda et al.~\\cite{colaiuda}, and Ciolfi et al.~\\cite{ciolfia,ciolfib} derived relativistic stellar models having both toroidal and poloidal magnetic fields with a perturbative technique. Although progress has been achieved in this field, further studies are required because all the magnetized star models are constructed by some special magnetic-field configurations which may not be realistic. In particular, it is not clear at all whether their models are stable. The stability of the magnetized star is an important issue, because only stable equilibrium models are viable. Stability analyses of magnetized stars have been performed by many works, since the pioneering work by Tayler~\\cite{tayler73}, who showed that stars having purely toroidal magnetic fields are unstable. Wright~\\cite{wright73} subsequently showed that there is the same type of the instability, the so-called pinch-type instability, for stars containing purely poloidal magnetic fields. He also suggested the possibility that stars having mixed poloidal-toroidal magnetic fields may be stable if the strength of both components is comparable (see, also, Refs.~\\cite{markey,tayler80}). Assche et al.~\\cite{assche} proved that the pinch-type instability in general arises in magnetized stars with purely poloidal fields, and Wright~\\cite{wright73} and Markey and Tayler~\\cite{markey} studied this instability for particular magnetic-field configurations. Flowers and Ruderman~\\cite{flowers} found that another type of instability occurs in purely poloidal magnetic-field configurations. All those classical stability analyses have been done by a method of an energy principle in the framework of Newtonian dynamics (see, also, Refs.~\\cite{pitts,goossens}). Another approach is a local analysis, with which Acheson \\cite{acheson1978} investigated the stability of rotating magnetized stars containing purely toroidal fields in detail in the framework of Newtonian dynamics (see, also, Refs.~\\cite{acheson1979,spruit}) and derived detailed stability conditions for purely toroidal magnetic fields buried inside rotating stars with dissipation. Note that although it is an approximate approach, the local analysis can take account of realistic effects on the stability like rotation, heat conduction, and resistivity, which cannot be included in a method of the energy principle. Bonanno and Urpin analyzed the axisymmetric stability~\\cite{bonanno} and the non-axisymmetric stability~\\cite{bonanno2} of cylindrical equilibrium configurations possessing mixed poloidal-toroidal fields, while ignoring compressibility and stratification of the fluid. Recently the stability problem of the magnetized star has been approached from another direction. By following the time evolution of small random initial magnetic fields around a spherical star in the framework of Newtonian resistive magnetohydrodynamics, Braithwaite and Spruit~\\cite{Braithwaite2004,Braithwaite2006} obtained stable configurations of a magnetized star that are formed as a self-organization phenomenon. The resulting stable magnetic fields have both poloidal and toroidal components with comparable strength and support the conjecture for stability conditions of the magnetized star given by the classical studies mentioned before. By using the numerical magnetohydrodynamics simulation, Braithwaite \\cite{Braithwaite2009} studied stability conditions for the magnetized stars and obtained a stability condition for his models given in terms of the ratio of the poloidal magnetic energy to the total magnetic energy which is of order unity. Duez et al. showed that magnetized stars constructed in Ref.~\\cite{duez} exhibit no instability for several Alfv\\'{e}n time scales in their numerical simulations \\cite{duez10}. This fact reconfirms the results given by Braithwaite. Lander and Jones explored the stability of magnetized stars by numerically solving the time evolution of linear perturbations around the stars in their series of papers \\cite{lander11,lander11b,lander12}. For the purely toroidal/poloidal field cases, their results are consistent with those of the classical stability analysis, i.e., the pinch-type instability is observed near the symmetry and the magnetic axes for the purely toroidal and purely poloidal field cases, respectively. They also assessed the stability of various magnetized stars with mixed poloidal-toroidal fields and found that all their models considered suffer from the pinch-type instability even for the cases in which the poloidal and toroidal components have comparable strength \\cite{lander12}. It is obvious that the results by Lander and Jones are incompatible with those by Braithwaite and his collaborators~\\cite{Braithwaite2009,duez10}. Lander and Jones discussed the possibility that some physics missing in their study would suppress the instability they found. We infer that in particular, {\\em stratification of the fluid} will be a key ingredient, which is taken into account in the analyses of Refs.~\\cite{Braithwaite2009,duez10} but not in the analyses of Ref.~\\cite{lander12}. We will return to this point later. General relativistic magnetized stars have been also analyzed recently. By numerical-relativity simulations, Kiuchi et al.~\\cite{kiuchia,kiuchib} investigated the stability of the magnetized stars with purely toroidal magnetic fields obtained by Kiuchi and Yoshida~\\cite{kiuchi}. They showed that the stars with some specific distributions of magnetic fields are stable against axisymmetric perturbations but all the models considered are unstable against non-axisymmetric perturbations due to the strong magnetic buoyancy instability near the surface of the stars. The initial behavior of the instability observed in Refs.~\\cite{kiuchia,kiuchib} is consistent with that expected by the Newtonian linear analyses by Acheson~\\cite{acheson1978}. Lasky et al.~\\cite{lasky} and Ciolfi et al.~\\cite{ciolfi} showed by numerical-relativity simulations that the purely poloidal magnetic fields, obtained by Bocquet et al.~\\cite{bocquet}, are unstable due to the pinch-type instability near the magnetic axis as predicted by the Newtonian linear analyses~\\cite{wright73,markey}. All these recent general relativistic magnetohydrodynamics simulations have contributed a lot to the progress of the stability analyses of general relativistic magnetized stars. However, the numerical simulations have been performed for equilibrium stars composed of {\\em non-stratified} fluids. The assumption of non-stratification is often used and quite reasonable as a first approximation for exploring cold neutron stars, even though the cold neutron star is expected to be highly stably stratified by the composition gradient (see, e.g., Refs.~\\cite{finn,reisenegger,lai}). If the effects of magnetic fields are taken into account for the neutron star models in their stability analysis, however, the situation changes drastically. As discussed by many authors, e.g., Reisenegger \\cite{reisenegger2008} and Kiuchi et al. \\cite{kiuchib}, it has long been known that the magnetic buoyancy makes the magnetized star unstable and that the stable stratification is necessary to remove the magnetic buoyancy instability (the so-called Parker instability~\\cite{parker}). It should be emphasized that in the assumption that the stellar matter is stably stratified, Braithwaite and Spruit~\\cite{Braithwaite2004,Braithwaite2006} obtained stable magnetized stars of simple magnetic configurations in their numerical simulations. Therefore, we infer that a stable stratification is one of the key ingredients for stable configurations of the magnetized stars. Note that Lander and Jones~\\cite{lander12} indicated the possibility that for stars containing mixed poloidal-toroidal magnetic fields, some weak instability associated with poloidal magnetic fields may not be removed by stable stratification (see, also, Ref.~\\cite{bonanno}). However, at the moment, no definite conclusion has been obtained. A reason that non-stratified magnetized stars are employed for the stability analyses in general relativity is that no model of stratified magnetized stars have been constructed, although in the framework of Newtonian dynamics, magnetized stars with stable stratification due to composition gradients have been obtained~\\cite{mastrano,lander12b,glampedakis}. In this paper, thus, we study stably stratified and magnetized stars in the framework of general relativity aiming at giving a prescription for constructing them. First of all, we describe a general formulation to obtain stationary axisymmetric magnetized stars composed of both toroidal and poloidal magnetic fields with stratification due to entropy gradients assuming ideal magnetohydrodynamics and employing a barotropic equation of state. As sample models, the magnetized stars considered by Ioka and Sasaki~\\cite{Ioka2004}, which contain poloidal and toroidal fields of comparable strength only inside the stars, are modified to be the ones stably stratified. Note that to date, no general relativistic magnetized stars containing mixed poloidal-toroidal fields have been constructed with a non-perturbative approach because of difficulties in the treatment of non-circular spacetimes (see, e.g., Ref.~\\cite{erik}). Finally, we describe the reason that the magnetized stars obtained in the present study are more stable than the existing relativistic models. ", "conclusions": "Checking the stability is an important issue to be explored after obtaining an equilibrium state of a magnetized star, because unstable solutions lose their physical meaning in the sense that they are not realized in nature. For the present models, as observed in Sec. V, the gravitational mass and the total electromagnetic energy increase with the mode number $i$ of the eigensolution, if we consider the equilibrium sequences for a given baryon mass and a magnetic helicity. This result is quite important and interesting due to the following reason. If the total baryon mass and magnetic helicity are conserved during the formation process of neutron stars, as discussed before, it is likely that the final state of the magnetic fields characterized by the arbitrary functions of the flux function (\\ref{Def_f_Om}) -- (\\ref{Def_f_Gam}) and the surface boundary condition (\\ref{surface_bc}) is the lowest-order eigensolution characterized by the smallest eigenvalue $L=L_1$ because it has the lowest gravitational mass and total electromagnetic energy among the equilibrium solutions characterized by the fixed baryon mass and magnetic helicity. We therefore conjecture that all the high-order solutions are unstable because there is an equilibrium state with energy lower than theirs and that the solutions associated with the lowest eigenvalue $L=L_1$ can be stable if there is a stable solution among the present models. Another important fact is that for the magnetic-field profile characterized by $L=L_1$, their magnetic energy is most equally divided into the poloidal and the toroidal magnetic energies among the eigensolutions. This property is consistent with the conjecture for stable magnetic configurations given by the linear analyses; stable magnetized stars contain both poloidal and toroidal components with comparable magnetic energies. Most of the magnetized-star models constructed in the framework of general relativity so far were non-stratified, and therefore, marginally stable against convection. Those non-stratified models are in general highly unstable against the magnetic buoyancy in the vicinity of the stellar surface, in the presence of magnetic fields. For strongly magnetized stars like magnetars, as shown by Kiuchi et al.~\\cite{kiuchib}, the magnetic buoyancy instability induces a convective motion near the surface of the star and fully destroys initially coherent magnetic fields inside the star. To stabilize this magnetic-buoyancy instability, the stratification with the strength sufficient to overcome the magnetic buoyancy are necessary as a stabilizing agent. This stabilization effect prevails in non-rotating diffusion-less stars with purely toroidal magnetic fields as argued by Acheson \\cite{acheson1978}. When $N^2 \\gg \\omega_A^2>0$, his dispersion relation (see Equation (3.20) of Ref.~\\cite{acheson1978}) has four solutions, $\\omega \\approx \\pm k_\\theta N(k^2_r+k^2_\\theta)^{-{1\\over 2}}$ and $\\omega \\approx \\pm m \\omega_A$, where $\\omega$ and $\\omega_A$ mean the oscillation and the Alfv{\\'e}n frequencies, respectively, and $k_r$, $k_\\theta$, and $m$ denote the vertical, horizontal, and azimuthal wave numbers, respectively. These four solutions are composed of propagating waves; the former is an internal gravity wave and the latter an Alfv{\\'e}n wave. We therefore confirm that there is no magnetic-buoyancy instability as long as $N^2 \\gg \\omega_A^2>0$. In Figure~\\ref{f04}, we plot squares of the Brunt--V{\\\"a}is{\\\"a}l{\\\"a} frequency $N^2$ and the Alfv{\\'e}n frequency $\\omega_A^2$ for a $\\Gamma=2.1$ model characterized by $M/R=0.2$ and $L=L_1$ as functions of the dimensionless radius $r/R$. Here, the Alfv{\\'e}n frequency is evaluated on the equatorial plane and defined by $\\omega_A \\equiv\\sqrt{ B^\\mu B_\\mu/((4\\pi\\rho h + B^\\mu B_\\mu)r^2)}$, and the strength of the magnetic fields is determined by the condition $E_{\\rm EM}/|W|=2.5\\times 10^{-2}$, which corresponds to very strong magnetic fields $\\approx 10^{16}$~G for a typical neutron star. This figure shows that the Brunt--V{\\\"a}is{\\\"a}l{\\\"a} frequency, $N$, is much larger than the Alfv{\\'e}n frequency $\\omega_A$ in the region of $r > 0.5 R$ for the models with $L=L_1$. Thus, we can predict that this model is stable against the magnetic buoyancy vicinity of the stellar surface. It should be noted that the physical origin of the anti-buoyancy force in our models is different from that in real neutron stars, as mentioned before, because the former comes from the entropy gradient and the later mainly from the composition gradient (see, e.g., Ref.~\\cite{reisenegger}). Magnetic fields, however, do not care the origin of the anti-buoyancy force, but do whether stars are stably stratified or not. As pointed out by Acheson \\cite{acheson1979}, the magnetic hoop stress caused by the strong toroidal magnetic fields governs dynamics of the perturbation near the center of the star. This fact may be confirmed from Figure~\\ref{f04}, which shows $N \\ll\\omega_A$ near the center of the star. Thus, the stratification is not helpful near the center of the star in the presence of the magnetic instability. For the central part of the star, as mentioned before, the presence of the poloidal magnetic fields having comparable strength with the toroidal ones will suppress the pinch-type instabilities. By solving linear perturbation equations around magnetized star models in the framework of Newtonian dynamics, Lander and Jones showed that this suppression of the magnetic instability indeed occurs, although the presence of the poloidal magnetic fields leads another instability associated with themselves \\cite{lander12} (see, also, Refs.~\\cite{bonanno,bonanno2}). Thus, the results of their numerical simulation lessen the possibility that the pinch-type magnetic instabilities are completely removed for the stars containing the mixed poloidal-toroidal magnetic fields with comparable strength. Although Lander and Jones's results obviously conflict with those by Braithwaite and his collaborators \\cite{Braithwaite2004,Braithwaite2006,Braithwaite2009,duez10}, we have so far had no definite conclusion to this controversy. One important difference between their analyses is in the treatment of the resistivity of the matter; Lander and Jones employed the ideal magnetohydrodynamics approximation, whereas Braithwaite and his collaborators took into account the resistivity. This might be a key to the solution of this problem. Toward a definite answer to this stability problem, we need further studies on the stability of the magnetized stars. As such a study, we plan to perform a stability analysis of the present magnetized star models by general relativistic magnetohydrodynamics simulations like those done by Kiuchi et al.~\\cite{kiuchia,kiuchib}. The present stably stratified models characterized by $L=L_1$ satisfy the following conditions; they have both poloidal and toroidal magnetic fields with comparable strength to suppress the hoop-stress instability inside the star, and in addition, the fluid constructing the magnetized star is stably stratified with strength sufficient to overcome the magnetic buoyancy near the stellar surface. Thus, it is obvious that some major magnetic instabilities will be reduced in the present magnetized star models. This will make the stability problem more tractable. As discussed in Kiuchi et al.~\\cite{kiuchi,kiuchia}, it is reasonable to expect that the toroidal component of the magnetic fields is much larger than the poloidal component inside the neutron star at least soon after its birth. The reason for this is that the winding of poloidal magnetic fields caused by a rapid and differential rotation during the core collapse would create large toroidal fields. Most of general relativistic magnetized star models obtained in numerical computations so far can, however, have toroidal magnetic fields much weaker than poloidal ones \\cite{colaiuda,ciolfia,ciolfib}. Their minimum ratio of the poloidal magnetic field energy to the total magnetic energy is $\\approx 0.92$, which is much larger than those of the present models. (For magnetized star models in the framework of the Newtonian dynamics, see, e.g., Ref.~\\cite{haskell}). Their magnetic fields are composed of the mixed poloidal--toroidal twisted torus magnetic fields inside the star and nearly dipolar magnetic fields outside the star. They may look quite plausible for the magnetosphere of neutron stars. However, weak toroidal magnetic-field strength seems to be unlikely for the strongly magnetized neutron stars. Although the magnetic field vanishes outside the star for the present models, which is quite unrealistic, the present models would give a reasonable inside structure of the strongly magnetized neutron stars because of their large toroidal magnetic-field strength. For obtaining more realistic models composed of a stably stratified fluid, basic equations given in Equations (\\ref{Def_Psi}) -- (\\ref{current}) may be employed as far as ideal magnetohydrodynamics and the barotropic equations of state are employed." }, "1207/1207.7080_arXiv.txt": { "abstract": "We measure the integrated contributions of dusty AGB stars and other luminous red mid-IR sources to the mid-IR luminosities of 6 galaxies (M81, NGC~2403, NGC~300, M33 and the Magellanic Clouds). We find the dusty AGB stars whose mid-IR fluxes are dominated by dust rather than photospheric emission contribute from 0.6\\% (M81) to 5.6\\% (SMC) of the $3.6~\\micron$ flux and 1.0\\% (M81) to 10.1\\% (SMC) of the $4.5~\\micron$ flux. We find a trend of decreasing AGB contribution with increasing galaxy metallicity, luminosity and mass and decreasing SSFR. However, these galaxy properties are strongly correlated in our sample and the simplest explanation of the trend is galaxy metallicity. Bright, red sources other than dusty AGB stars represent a smaller fraction of the luminosity, $\\sim$1.2\\% at 3.6~$\\micron$, however their dust is likely cooler and their contributions are likely larger at longer wavelengths. Excluding the SMC, the contribution from these red sources correlates with the specific star formation rate as we would expect for massive stars. In total, after correcting for dust emission at other wavelengths, the dust around AGB stars radiates 0.1-0.8\\% of the bolometric luminosities of the galaxies. Thus, hot dust emission from AGB and other luminous dusty stars represent a small fraction of the total luminosities of the galaxies but a significant fraction of their mid-IR emissions. ", "introduction": "Understanding the emissions from AGB stars is important for stellar evolution, models of galaxy spectral energy distributions (SEDs) and inferences about stellar populations and their evolution. AGB stars are H and He shell burning stars that evolve from ~0.8-8 $M_{\\odot}$ Main Sequence stars. During the AGB phase, which lasts 1-13 Myr \\citep{vass1993}, the star can undergo significant mass-loss, particularly in the Thermal-Pulsing AGB phase (TP-AGB, $\\sim1.7M_{\\odot} \\le M \\le 6M_{\\odot}$). During this last 0.2-2 Myr of the AGB phase, the AGB star experiences thermal pulses from a series of shell flashes that drive high mass loss rate winds conducive to the formation of dust. The short lifetime of this phase makes these stars relatively rare, so they are generally absent from the star clusters used to calibrate stellar evolution models. As a result, this is one of the most uncertain phases of stellar evolution. Their large range of mass loss rates combined with dust formation mean that AGB stars can be important over a broad range of wavelengths from the optical to the mid-IR. For distant galaxies, the properties of stellar populations must be inferred from spectra or the SEDs of ensembles of stars. Since AGB stars represent a non-trivial fraction of the luminosity of $\\sim$Gyr old stellar populations, any uncertainties in the AGB phase propagate into uncertainties in overall galaxy properties because the stellar population synthesis (SPS) model must assume some treatment of the (TP) AGB stars. \\citet{maraston2006} shows how different calibrations of the TP-AGB stars in the SPS models can cause large, systematic differences in the fits to SEDs and the resulting estimates of a galaxy's mass and age. \\citet{conroy2010} and \\citet{conroy2010b} show in a more general way how TP-AGB stars can effect the determination of a galaxy's properties. Much progress has been made recently in moving beyond the Milky Way and Magellanic Cloud star clusters to calibrate the AGB phase of stellar evolution (e.g. \\citealt{kriek2010}; \\citealt{meidt2012b}; \\citealt{zibetti2012}). In particular, \\citet{melbourne2012} used NIR observations to build on the optical study of 12 metal poor galaxies by \\citet{girardi2010} to constrain the lifetimes of AGB stars based on the number and luminosities of the AGB stars. While they found that the numbers of observed AGB stars agreed with the updated \\citep{girardi2010} SPS models given their uncertainties, the predicted luminosities are too large. They go on to demonstrate how more accurate lifetimes and luminosities are needed to be able to accurately describe galaxies at high redshift. \\citet{boyer2009} examined AGB stars in eight Local Group dwarf irregular galaxies in the Spitzer IRAC bands. They examined the number of AGB stars, their mass-loss rates and how this mass returning to the ISM could effect the current SFR. They also suggest that optical AGB searches will miss 30\\%-40\\% of AGB stars due to dust obscuration. In this paper we focus on the contribution of hot dust around luminous stars, particularly AGB stars, to the mid-IR emission of 6 nearby galaxies of varying properties. In addition to AGB stars, these sources include dusty young star clusters and other luminous stars that have experienced dense, episodic winds or eruptive outbursts. We use archival Spitzer IRAC 3.6 and $4.5~\\micron$ observations of M33, M81, NGC~300, NGC~2403, the SMC and the LMC to survey these populations. We are focusing on AGB stars whose mid-IR fluxes are dominated by dust rather than photospheric emission, and we will refer to them as DAGB (dusty AGB) stars in order to distinguish them from the remainder of the AGB population whose mid-IR emission is predominately photospheric. Section 2 describes the data and its analysis. We identified these stars as those with red [3.6]$-$[4.5] colors indicating the presence of hot circumstellar dust. The same color criteria also identify other, more luminous stars whose mid-IR fluxes are dominated by hot dust, and we will examine these as a separate population of ``red'' stars. Section 3 presents the data analysis, including the source classification and background corrections. In Section 4 we determine the fraction of the [3.6] and [4.5] flux contributed by the DAGB and red stars and examine the trends in the DAGB and red star flux fractions with galaxy properties. Our final summary is in Section 5. ", "conclusions": "We surveyed the luminous dusty stars in 6 nearby galaxies. Using 3.6 and 4.5~$\\micron$ data we identified DAGB and other luminous red sources with significant emission from hot circumstellar dust. We determine their contribution to the flux at these bands and whether there is any correlation between their luminosity contribution and galaxy metallicity, stellar mass, luminosity or SSFR. Such studies can help to constrain one of the uncertainties considered in \\citet{conroy2010}, namely the fraction of stellar emission obscured by stellar dust. By selecting the stars in the mid-IR we automatically include heavily obscured sources that would be missed in optical studies. Additionally, by looking at whole galaxies, we have a relatively large sample of these relatively rare stars and should obtain reliable population statistics. The contributions of obscured DAGB stars range from 0.9\\% (M81) to 5.6\\% (SMC) at $3.6~\\micron$ and from 1.0\\% (M81) to 10.1\\% (SMC) at 4.5~$\\micron$. We see a trend of higher mid-IR with metallicity, SSFR, stellar mass and luminosity, but this is most likely a correlation with metallicity since metallicity is correlated with the other properties and none of the other properties should be correlated with DAGB stars. It is a legitimate concern, however, that without complete star formation histories, we cannot be certain the SFR does not significantly effect the correlation. The dependence of AGB dust production and the resulting IR emission on initial metallicity is an area of active study, with somewhat conflicting results. It is a difficult problem, combining stellar evolution, chemistry, wind formation and dust properties. For example, \\citet{ventura2012} examined models of dust production by AGB stars at LMC and SMC metallicities. They find that for the more massive stars that undergo hot bottom burning ($8M_{\\odot} > M > 3.5M_{\\odot}$), a higher metallicity results in more dust production. For lower mass stars that undergo third dredge up ($3.5M_{\\odot} > M > 1.5M_{\\odot}$), they find that the dust production is nearly independent of metallicity. The more massive stars become oxygen rich at the surface, in which case dust formation is limited by the availability of silicon, which is controlled by the metallicity. For the low mass stars, third dredge up leads to an enhanced surface abundance of carbon and thus dust formation that depends little on initial metallicity. \\citet{ventura2012} in the end argue that theoretical predictions of dust production around lower mass AGB stars are not robust due to uncertainties in the amount of mixing and the extent of the third dredge up. Similar results are found by \\citet{marigo2008}, \\citet{bowen1991} and \\citet{matsuura2007}. \\citet{groenewegen2009} compare their sample of LMC and SMC AGB stars to models for mass loss, taking different dust grain composition into account, and compare to Galactic estimates to find that there is no strong dependence of mass-loss rate on metallicity to factors of order 2-4. This is broadly consistent with \\citet{wachter2008}, who find the mass loss rate of the Milky Way AGB stars to be a factor of 2 higher than that of the SMC. While higher metallicity may result in more dust formation, this dust is dominated by silicates which have a lower opacity than graphitic dust. These issues make it difficult to translate any trends into observational predictions. The \\citet{maraston2005} models of the AGB contribution are based on the fuel consumption theorem \\citep{renzini1981} and assume mass loss is not significantly effected by metallicity. Instead, it is the surface abundances that are modified, so that the number ratio of oxygen-rich to carbon-rich AGB stars depends on metallicity. A metal-poor population is expected to have more carbon-rich stars. A metal-poor AGB star has a lower abundance of oxygen in its envelope, so less carbon has to be dredged up before the oxygen has all been bound in CO and the residual carbon can go on to form dust (\\citealt{iben1983}; \\citealt{maraston2006}). In this scenario, a metal-poor population could result in a higher mid-IR luminosity. \\citet{bird2011} examined the energy output of AGB stars using fuel consumption theory, but find insufficient calibration data to determine an accurate metallicity dependence. The earlier Padova models (\\citealt{marigo2007}; \\citealt{marigo2008}) have more carbon-rich stars at low metallicity due to the higher efficiency of third dredge up, while the more recent \\citet{girardi2010} models have increased mass loss rates and shorter AGB phase lifetimes at low metallicity. While the effect of metallicity is not fully explored by these studies, the increased dust production and mass loss rate suggested at low metallicity would result in a higher mid-IR contribution, as we find in our sample. Previous observational constraints on the metallicity dependence of the dust production also have mixed conclusions. \\citet{meidt2012b} found that lower metallicity clusters in M~100 have more dusty AGB stars, arguing that since the dust optical depth depends upon whether an AGB star is O-rich or C-rich, the AGB contribution to the cluster depends upon metallicity. They also find larger mass-loss rates for younger clusters. Since the younger clusters also have a higher metallicity, they cannot address the overall effect of metallicity on the AGB contribution. Looking at the AGB population of eight local group dwarf irregular galaxies, \\citet{boyer2009} find that all galaxies in their sample have obscured AGB stars. For AGB stars with optical counterparts, they find that the more metal-rich galaxies have a larger population of red, dust enshrouded AGB stars. However, requiring optical counterparts biases the sample against dustier stars. An alternate explanation for the increased number of dusty stars in the more metal-rich galaxies is the age of the populations, because the higher metallicity galaxies in their sample also have more recent ($<$ 3 Gyr) star formation. They find no clear correlation between metallicity and optical completeness, as would be expected if metallicity were the dominating factor in the degree of obscuration. \\citet{meidt2012a} examined the contribution from red and AGB stars to the luminosity of M81. They find hot dust and polycyclic aromatic hydrocarbons (PAHs) contributes 5.0$\\pm0.9\\%$ while red and AGB stars contribution only 0.10$\\pm0.02\\%$ of the $3.6~\\micron$ luminosity, rather than our significantly higher estimate of 0.58$\\pm0.02\\%$ for the DAGB stars. \\citet{meidt2012a} used a very different method to estimate the percentage of the light due to AGB and intermediate age stars. Essentially, they divide the 3.6 and 4.5~$\\micron$ emission into that from stars and that due to ``contaminating'' sources that contribute to the mid-IR flux while representing little stellar mass. This includes, dusty stars of all luminosities and PAH emission. They do this by using independent component analysis (ICA) to separate out statistically significant source distributions and maximize $[3.6]-[4.5]$ color differences. Using the $8.0~\\micron$ images, they were able to obtain a rough estimate of the contribution of stellar and non-stellar (dust) sources. They note a selection bias against spatially coincident dust and clusters dominated by intermediate age stars and also a possible bias from mismatched [3.6] and [8.0] PSFs which would confuse the division between stellar and non-stellar emission. They caution the distinction between 3.6~$\\micron$ emission due to evolved stars and PAH emission are good for ``rough estimates'' of the components and that a more detailed analysis of the ``contaminates'' is possible but beyond their scope. However, \\citet{meidt2012b} explores evolved stars in M100 in more detail. For comparison, \\citet{melbourne2012} estimate the near-IR (HST WFC3/IR F160W) H-band contributions of AGB stars in three of our galaxies, M81, NGC~300 and NGC~2403. The estimates are only for a single WFC3 fields rather than global estimates. They find a NIR AGB luminosity contribution of 8\\% for M81, 15\\% for NGC~300 and 17\\% for NGC~2403. Since the $1.6~\\micron$ emission is dominated by the photosphere of the AGB stars rather than the reprocessed emission from dust around stars, these fractions need not match ours. Like \\citet{kriek2010}, they find that the stellar population models over-predict the AGB NIR luminosity and that the $1.6~\\micron$ AGB luminosity fraction increases with the mass fraction of intermediate age stars ( $<$ 2 Gyr and $<$ 0.3 Gyr) with no obvious metallicity trends. The brighter, red sources contribute less, typically accounting for $\\sim$1.2\\% of the $3.6~\\micron$ and $\\sim$1.8\\% of the $4.5~\\micron$ luminosity. These sources are a mixture of dusty young massive stars and star clusters, and we expect their luminosity fraction could be strongly correlated with the SSFR. The correlation may exist if there is also strong suppression of the red stars at the low metallicity of the SMC (\\citealt{bonanos2009}; \\citealt{bonanos2010}). We lack a large enough distribution of galaxy properties to address this question. These sources are probably more important at longer wavelengths. The 3.6~$\\micron$ and 4.5~$\\micron$ emissions are due to fairly hot dust close to the stars, while many of these sources are episodic dust sources rather than having long lived winds (see discussion in \\citet{kochanek2012}). During most of the period where they have a significant dust optical depth, the dust is at larger distances and cooler temperatures radiating at longer wavelengths. Simple estimates suggest that their contribution will peak $\\sim$24~$\\micron$, similar to $\\eta$ Carinae's present day SED (\\citealt{robinson1973}; \\citealt{humphreys1994}). Dusty young star clusters will also tend to be dominated by cooler dust \\citep{whelan2011}. Overall, our findings are in agreement with other recent studies, finding that while (D)AGB stars have a significant effect on a galaxy's SED in the IR, the contribution is not as large as originally thought (\\citealt{kelson2010}; \\citealt{boyer2011}; \\citealt{meidt2012b}). \\citet{boyer2011} find that AGB stars contribute about 20\\% of the 3.6~$\\micron$ flux in the SMC, compared to our $\\sim$6\\%. Our mid-IR identification of sources, likely selects a smaller, dustier sample of stars causing this difference. We also examined how the dusty stars we study effect the SED of a galaxy. By comparing to galaxy templates from \\citet{assef2010}, we found that the summed 3.6 and 4.5~$\\micron$ emission represents $\\sim$ 5\\% of a galaxy's bolometric luminosity, and that the DAGB star contribution ranges from $\\sim$ 0.04\\% (M81) to $\\sim$ 0.37\\% (SMC) of the bolometric luminosity. By modeling the AGB star contribution with DUSTY \\citep{ivezic1999} we found that the DAGB luminosity constitutes $\\sim$ 50\\% of the hot dust emission from these stars. This means that while dusty stars have a small effect on the overall galaxy SED, they can significantly effect the mid-IR colors. Note that here we have focused on the contribution of dust emission by AGB stars to the overall luminosity of the galaxy, while many of these other studies are focused on the contribution from the stellar photospheres of the AGB stars. These less or non-dusty AGB stars likely contribute a significant amount of mid-IR flux without distorting the SED enough to produce a red mid-IR color. The uncertainty in the treatment of the AGB phase in the SPS models can have a significant impact on the estimation of galaxy properties (\\citealt{maraston2006}; \\citealt{conroy2009}; \\citealt{conroy2010}; \\citealt{conroy2010b}). For example, \\citet{conroy2010} find that lower metallicity populations prefer cooler AGB stars which could in part reflect increased dust production at lower metallicity. They also find that both TP-AGB stars and dust cause significant uncertainties in the properties of star forming galaxies. In studies like \\citet{conroy2010}, it would be straight forward to include a reasonable phenomenological model of AGB or other stars with dusty winds. For a period $T_{wind}$ stars are assigned a dust optical depth $\\tau$ with the dust radiating at $T_{dust}\\sim1000$ K. The SED can easily be calculated using DUSTY \\citep{ivezic1999} or similar dust radiation transfer models. Extending the UV through near-IR wavelengths of \\citet{conroy2010} and \\citet{conroy2010b} to the mid-IR would then constrain the hot dust contribution and its effects on the overall SED because it will be difficult for other variables to mimic the mid-IR dust emission. Focusing on the 4.5 and 5.8 $\\micron$ bands would minimize the need to worry about the strong PAH emissions at 8 $\\micron$ and the weaker emissions at 3.6 $\\micron$. Although the overall amount of flux from dusty stars may be small, the effect on the mid-IR can be substantial and should be explored further. Our results are primarily limited by the low resolution of Spitzer. The James Webb Space Telescope will allow this type of study to be performed on many more galaxies and over a broader wavelength range, providing better constraints on how hot dust around stars, and AGB stars in particular, effect a galaxy's SED and how their contribution correlates with galaxy properties such as SFH and metallicity." }, "1207/1207.1938_arXiv.txt": { "abstract": "Deeper understanding of the properties of dark energy via SNIa surveys, and to a large extent other methods as well, will require unprecedented photometric precision. Laboratory and solar photometry and radiometry regularly achieve precisions on the order of parts in ten thousand, but photometric calibration for non-solar astronomy presently remains stuck at the percent or greater level. We discuss our project to erase this discrepancy, and our steps toward achieving laboratory-level photometric precision for surveys late this decade. In particular, we show near-field observations of the balloon-borne light source we are presently testing, in addition to previous work with a calibrated laser source presently in low-Earth orbit. Our technique is additionally applicable to microwave astronomy. Observation of gravitational waves in the polarized CMB will similarly require unprecedented polarimetric and radiometric precision, and we briefly discuss our plans for a calibrated microwave source above the atmosphere as well. ", "introduction": "The presence of artificial flux standards above the Earth's atmosphere may provide significant reduction of photometric uncertainties for measurements that depend on such calibration. The combined effect of atmospheric and instrumental extinction in the visible and near-infrared is presently the source of the largest uncertainty on the properties and amount of dark energy (see Fig.~\\ref{fig:Alexplots}). Man-made visible and NIR light sources can be measured to a precision of up to 100 times better than standard stellar sources, as presently shown by both laboratory and solar irradiance measurements, allowing for vast reduction of astronomical and cosmological uncertainties due to photometry. Separately, in the microwave spectrum, polarization patterns in the cosmic microwave background (CMB) encode a vast amount of other information (beyond dark energy) about the early universe, ranging from the amplitude of primordial gravitational waves and the energy scale of cosmic inflation, to the gravitational lensing potential integrated over cosmic history. The interpretation of these signals is entirely contingent on an accurate calibration of polarized instrumental sensitivity. As yet, no well-calibrated, polarized, celestial microwave sources exist, and in the coming generation of microwave telescopes, this deficit may become a limiting factor in our ability to measure and understand the polarized microwave sky. Current calibration solutions rely on relatively nearby ground-based sources, requiring refocusing and special low-sensitivity detectors to handle the near-field source and additional atmospheric loading. The cleanest and simplest calibration solution would solve both of these issues, by lofting a source into the far field, at distances of order tens of kilometres on large microwave telescopes. \\begin{figure}[!t] \\begin{center} \\includegraphics[angle=0,width=6.5cm]{albert_j_1a.eps} \\hfill \\includegraphics[angle=0,width=6.5cm]{albert_j_1b.eps} \\caption{ Constraints on the dark energy equation of state parameters w$_0$ ($\\equiv p/\\rho$ of the universe at the present day), and w$_a \\equiv$ dw/d$a$ (also at present, where $a$ is the scale factor of the universe), for all current and upcoming SNIa dark energy projects (plus the Planck CMB mission) with the new calibration provided by our calibration program, vs.~with all current means of calibration. The left plot shows the uncertainties if one artificially constraints the universe to be flat, and the right plot shows the uncertainties with flatness constraint removed. Dashed lines show 68\\% and 95\\% confidence intervals for each case. The calibration will improve the dark energy ``figure of merit''~\\citep{alb06} by a factor of 2.4, with additional improvement beyond 2014. } \\label{fig:Alexplots} \\end{center} \\end{figure} Prior to the second half of the 20th century, the only sources of light above the Earth's atmosphere were natural in origin: stars, and reflected light from planets, moons, comets, etc. Natural sources have of course served extremely well in astronomy: through understanding the physical processes governing stellar evolution, we are now able to fairly precisely understand the spectra of stars used as calibration sources [\\textit{e.g.} \\citet{boh00}]. Nevertheless, in all stars the vast bulk of material, and the thermonuclear processes that themselves provide the light, lie beyond our sight below the surface of the star. Superb models of stellar structure are available, but uncertainties of many types always remain. Since the launching of the first high-altitude balloons and satellites, a separate class of potential light sources in space and in near-space has become available. Observable light from most satellites is primarily due to direct solar reflection, or reflection from Earth's albedo. While providing a convenient method of observing satellites, this light is typically unsuitable for use as a calibrated light source due to large uncertainties in the reflectivity (and, to a lesser extent, the precise orientation and reflective area) of satellites' surfaces.\\footnote{Reflected solar light has, however, been successfully used as an absolute infrared calibration source by the Midcourse Space Experiment (MSX), using 2 cm diameter black-coated spheres ejected from the MSX satellite, whose infrared emission was monitored by the instruments aboard MSX~\\citep{pri04}. This technique proved highly effective for the MSX infrared calibration; however, the technique is not easily applicable to measuring extinction of visible light in the atmosphere.} However, balloon- and satellite-based calibration sources for ground-based telescopes are by no means technically prohibitive. As an example, a standard household 25-watt tungsten filament lightbulb (which typically have a temperature of the order of 3000 K and usually produce approximately 1 watt of visible light between 390 and 780 nm) which radiates light equally in all directions from a 700 km low Earth orbit has an equivalent brightness to a 12.5-magnitude star (in the AB system, although for this approximate value the system makes little difference). In general, the apparent magnitude of an orbiting lamp at a typical incandescent temperature which radiates isotropically is approximately given by \\begin{equation} m \\approx -5.0 \\log_{10} \\left ( \\frac{ \\left ( \\ln \\left ( \\frac {P}{\\rm{1\\;watt}} \\right ) \\right )^3 }{h} \\right ) + 5.9, \\end{equation} where $P$ is the power of the lamp in watts, and $h$ is the height of the orbit in kilometers. The systematic uncertainty on the radiance of an optimally-designed above-atmosphere lamp, where cost is no object, would be dominated by the precision of radiometric monitoring technology, and be in the range of approximately 200 parts per million if the best presently available radiometric technology is used. An alternative to an isotropic or near-isotropic lamp would be a laser source, with beam pointed at the observer (with a small moveable mirror, for example). Divergences of laser beams are typically on the order of a milliradian (which can be reduced to microradians with a beam expander) so much less output power than a lamp would be required for a laser beam to mimic the brightness of a typical star. The apparent magnitude of an orbiting laser with Gaussian beam divergence pointed directly at a ground-based telescope, is given by \\begin{equation} m \\approx -2.5 \\log_{10}\\left ( \\frac{P}{h^2 d^2} \\right ) - 20.1, \\end{equation} where $P$ is the laser power in milliwatts, $h$ is the height of the orbit in kilometers, and $d$ is the RMS divergence of the laser beam in milliradians, under the assumption that the aperture of the telescope is small compared with the RMS width of the beam at the ground, $hd$. The RMS divergence would be the combination of the divergence at the source, and the divergence due to the atmosphere. In clear conditions, total atmospheric divergence in a vertical path is at the level of approximately 5 microradians~\\citep{tat61}, and this of course only acts on the last fraction of the laser path that is within the atmosphere, so as long as the source divergence is significantly larger than this, atmospheric divergence would be negligible. The uncertainty on the apparent magnitude of an orbiting laser stemming from uncertainties in the radiometrically-monitored laser power would likely be limited by the precision of current radiometer technology. Modern electrical substitution radiometers can achieve a precision of approximately 100 parts per million when aperture uncertainties can be neglected, as in the case of laser radiometry~\\citep{kop05}. Uncertainties on the magnitude due to uncertainty in the pointing and beam profile would potentially be limited by the size of the array of outboard telescopes for monitoring the laser spot, and by calibration differences between the individual telecopes in the array and with the main central telescope. The latter could clearly be minimized by a ground system for ensuring the relative calibration of the outboard telescopes and main telescope are all consistent. The uncertainties considered above assume that the exposure time is long compared with the coherence time of the atmosphere. With short exposures --- or in the case of a laser that either quickly sweeps past, or is pulsed --- atmospheric scintillation can play a major role in uncertainty in apparent magnitude of an above-atmosphere source. A typical timescale for a CW laser with 1 milliradian divergence in low Earth orbit to sweep past is tens of milliseconds, which is of the same order as characteristic timescales of atmospheric scintillation, and the typical timescale of single laser pulses is nanoseconds, much shorter than scintillation timescales, thus one cannot assume that such effects can be time-averaged over. In idealized conditions, for small apertures $D < \\sim 5$ cm and sub-millisecond integration times, the relative standard deviation in intensity $\\sigma_I \\equiv \\Delta I / \\langle I \\rangle$, where $\\Delta I$ is the root-mean-square value of $I$, is given by the square root of \\begin{equation} \\sigma_I^2 = 19.12 \\lambda^{-7/6} \\int_0^\\infty C_n^2(h)h^{5/6}dh, \\end{equation} where $\\lambda$ is optical wavelength (in meters), $C_n^2(h)$ is known as the refractive-index structure coefficient, and $h$ is altitude (in meters) \\citep{tat61}. Large apertures $D > \\sim 50$ cm have a relative standard deviation in intensity given by the square root of \\begin{equation} \\sigma_I^2 = 29.48 D^{-7/3} \\int_0^\\infty C_n^2(h)h^{2}dh \\end{equation} \\citep{tat61}. The values and functional form of $C_n^2(h)$ are entirely dependent on the particular atmospheric conditions at the time of observation, however a relatively typical profile is given by the Hufnagel-Valley form: \\begin{eqnarray} \\!\\!\\! C_n^2(h) \\! & \\!\\!\\! = \\!\\!\\! & \\! 5.94 \\times 10^{-53}(v/27)^2 h^{10}e^{-h/1000} + \\nonumber\\\\ \\! & & \\! 2.7 \\times 10^{-16} e^{-h/1500} + Ae^{-h/100}, \\end{eqnarray} where $A$ and $v$ are free parameters \\citep{huf74}. Commonly-used values for the $A$ and $v$ parameters, which represent the strength of turbulence near ground level and the high-altitude wind speed respectively, are $A = 1.7 \\times 10^{-14}$ m$^{-2/3}$ and $v = 21$ m/s \\citep{rog96}. Using these particular values, for a small aperture, the relative standard deviation $\\sigma_I$ would be expected to be 0.466 for 532 nm light, which is not far off experimental scintillation values for a clear night at a typical location [\\textit{e.g.}~\\citet{jak78}]. For a single small camera, this is an extremely large uncertainty. Other than by increasing integration time (which is not possible with a pulsed laser) or by significantly increasing the camera aperture, the only way to reduce this uncertainty is to increase the number of cameras. With $N$ cameras performing an observation, which are spaced further apart than the coherence length of atmospheric turbulence (typically 5 to 50 cm), the uncertainty from scintillation can be reduced by a factor $\\sqrt{N}$ (for large $N$). The considerations above are necessarily both speculative and approximate. However, at present there is an actual laser in low Earth orbit, visible with both equipment and with the naked eye, and analysis of ground-based observational data of the laser spot can be used for comparisons with the above, as well as for development of and predictions for potential future satellite-based photometric calibration sources of ground telescopes. \\begin{figure*}[p] \\vspace*{-8mm} \\begin{center} \\includegraphics[angle=0,width=6.5cm]{albert_j_2a.eps} \\hfill \\includegraphics[angle=0,width=6.5cm]{albert_j_2b.eps} \\caption{ Seven-camera observations of CALIPSO overpasses, taken (left) on Apr.~17, 2007 near Carefree, Arizona, and (right) on May 1, 2007 near the Great Salt Lake, Utah. } \\label{fig:Carefreeobs} \\vspace*{9mm} \\includegraphics[angle=0,width=6.3cm]{albert_j_3a.eps} \\hfill \\includegraphics[angle=0,width=6.7cm]{albert_j_3b.eps} \\vspace*{1mm} \\includegraphics[angle=0,width=6.3cm]{albert_j_3c.eps} \\hfill \\includegraphics[angle=0,width=6.7cm]{albert_j_3d.eps} \\caption{ Camera time-integrated irradiance data, and the resulting fitted time-integrated irradiance maps, for the (upper left) Apr.~17, (upper right) May 1, (lower left) May 30, and (lower right) Jun.~1 observations. The numbers at each camera refer to the measured value of time-integrated irradiance by that camera, with its associated uncertainty (68\\% CL), and the expectation value from the fitted function at the location of that camera. The contours on each plot are spaced at 1 $\\mu$J/m$^2$ intervals. The upper-right inset on each plot extends the $x$ and $y$ axes in order to see the three 2-D Gaussians that comprise the fitted function (as described in the text), and the lower-left inset shows a different 3-D view (from the side, rather than from above) of the fitted function and the data points. } \\label{fig:irradmaps} \\end{center} \\end{figure*} ", "conclusions": "\\label{conclusion} We have performed initial tests and measurements of space- and near-space-based precision photometric calibration sources. Our initial tests are highly promising, and we believe this is a viable means to reduce the dominant uncertainty in measurements of dark energy this decade. Our technique can additionally be applied to microwave astrophysics and to other regions of the spectrum, to impact cosmological and other measurements in those areas as well. Improved precision in photometric calibration will be nearly as critical for astronomy as increased aperture and etendue telescopes in upcoming decades. The future of precision photometry is extremely promising, and laboratory-based standards in near-space and in space allow one to forsee many-fold improvement in photometric calibration as a near-term prospect." }, "1207/1207.3083_arXiv.txt": { "abstract": "Recently, Widrow and collaborators announced the discovery of vertical density waves in the Milky Way disk. Here we investigate a scenario where these waves were induced by the Sagittarius dwarf galaxy as it plunged through the Galaxy. Using numerical simulations, we find that the the Sagittarius impact produces North-South asymmetries and vertical wave-like behavior that qualitatively agrees with what is observed. The extent to which vertical modes can radially penetrate into the disc, as well as their amplitudes, depend on the mass of the perturbing satellite. We show that the mean height of the disc is expected to vary more rapidly in the radial than in the azimuthal direction. If the observed vertical density asymmetry is indeed caused by vertical oscillations, we predict radial and azimuthal variations of the mean vertical velocity, correlating with the spatial structure. These variations can have amplitudes as large as 8 km s$^{-1}$. ", "introduction": "Minor mergers can significantly perturb the overall structure of their host galactic disc \\citep[][]{quinn93,alvaro08}. As they merge, relatively small satellite galaxies can induce the formation of spiral arms and ring-like structures as well as radial migration, significantly flare or warp the disc, and influence the growth of a central bar \\citep[][]{tutu,kaza08,alvaro08,younger,m09,quillen09,p11,bird,g12b}. \\citet[][hereafter P11]{p11} presented simulations of the response of the Milky Way (MW) disc to tidal interaction with the Sagittarius dwarf galaxy (Sgr). They showed that many of the global morphological features observed in the Galactic disc can be simultaneously explained by this interaction. An example is the kinematically cold structure known as the Monoceros ring \\citep{newberg02,juric08}, which naturally emerges in these simulations, although its origin is still a matter of debate \\citep[see][]{conn,linew,mono}. \\citet[][hereafter G12]{g12a} showed that perturbations observed in the phase-space distribution of old disc stars in the Solar Neighbourhood (SN) can be reproduced qualitatively with these simulations. They interpreted such perturbations as signatures of \\emph{radial} density waves excited on the plane of the disc by Sgr. Perturbations in the \\emph{vertical} direction were not explored in their work. However, using a much larger photometric and spectroscopic data set, \\citet[][hereafter W12]{w12} recently identified a North-South asymmetry in both the spatial density and the velocity distribution of SN stars. The asymmetry has the appearance of a coherent, wave-like perturbation, intrinsic to the disc. W12 speculate that this perturbation could have been excited by the passage of a satellite galaxy through the Galactic disc. In this Letter, we explore the possibility of Sgr being the perturber associated with {\\it both} the vertical and radial modes. ", "conclusions": "\\label{sec:disc} In this work, we have explored a scenario in which Sgr is the perturber behind the North-South asymmetry recently observed in the number density and mean vertical velocity of Solar Neighbourhood stars (W12). For this purpose, we have searched for both local and global signatures of vertical density waves in two simulations modelling the response of the MW to the infall of Sgr. Distributions of stellar particles as a function of height in Solar Neighbourhood-like volumes extracted from our more massive Sgr's progenitor simulation present clear indications of a perturbation in the vertical direction of the disc. This asymmetry becomes more evident when comparing to a model of the underlying smooth vertical distribution of particles. Within the $|{\\rm Z}|$-range allowed by our finite mass resolution, the phase-space distribution of certain SN-like volumes can qualitatively reproduce the North-South asymmetry observed by W12. Remarkably, the same phase-space distributions can simultaneously reproduce the signatures of radial density waves observed by G12. By creating maps of the mean height of the disc within $R \\leq 20$ kpc, we have shown that the vertical perturbations observed in local volumes are signatures of a global mode perturbing the entire disc. As in the case of the radial density modes (G12), the amplitude and the extent to which vertical modes can radially penetrate into the disc depends on the mass of the perturbing satellite. Interestingly, we have shown that the mean height of the disc is expected to vary much more rapidly in the radial than in the azimuthal direction. Furthermore, the mean height of overdense spiral features vary azimuthally, moving from below to above the midplane of the disc. Signatures of vertical modes should also be observable in maps of the Galactic disc's mean vertical velocity since, not surprisingly, they present a clear oscillatory behavior. In contrast to radial modes that can be excited by a number of different external and internal mechanisms, vertical modes {\\it must be excited} by some agent external to the disc. Altough we have shown that perturbations induced by Sgr could be enough to account for several features observed in the Solar Neighborhood, it is likely that MCs are playing a role in shapping the vertical structure of MW disc. We plan to characterize the coupling of these two perturbations in a follow-up study. Contrasting the results presented in this work against currently available samples of Galactic disc stars could help to understand the origin of the observed vertical and in-plane perturbations." }, "1207/1207.1086_arXiv.txt": { "abstract": "We show that a single imperfect fluid can be used as a source to obtain a mass-varying black hole in an expanding universe. This approach generalizes the well-known McVittie spacetime, by allowing the mass to vary thanks to a novel mechanism based on the presence of a temperature gradient. This fully dynamical solution, which does not require phantom fields or fine-tuning, is a step forward in a new direction in the study of systems whose local gravitational attraction is coupled to the expansion of the universe. We present a simple but instructive example for the mass function and briefly discuss the structure of the apparent horizons and the past singularity. ", "introduction": " ", "conclusions": "" }, "1207/1207.4340_arXiv.txt": { "abstract": "{} {In this paper, we study the properties of solar granulation in a facular region from the photosphere up to the lower chromosphere. Our aim is to investigate the dependence of granular structure on magnetic field strength.} {We use observations obtained at the German Vacuum Tower Telescope (Observatorio del Teide, Tenerife) using two different instruments: Triple Etalon SOlar Spectrometer (TESOS), in the \\BaII\\ 4554 \\AA\\ line to measure velocity and intensity variations along the photosphere; and, simultaneously, Tenerife Infrared Polarimeter (TIP-II), in the \\FeI\\ 1.56 $\\mu$m lines to the measure Stokes parameters and the magnetic field strength at the lower photosphere.} {We obtain that the convective velocities of granules in the facular area decrease with magnetic field while the convective velocities of intergranular lanes increase with the field strength. Similar to the quiet areas, there is a contrast and velocity sign reversal taking place in the middle photosphere. The reversal heights depend on the magnetic field strength and are, on average, about 100 km higher than in the quiet regions. The correlation between convective velocity and intensity decreases with magnetic field at the bottom photosphere, but increases in the upper photosphere. The contrast of intergranular lanes observed close to the disc center is almost independent of the magnetic field strength. } {The strong magnetic field of facular area seems to stabilize the convection and to promote more effective energy transfer in the upper layers of the solar atmosphere, since the convective elements reach larger heights.} ", "introduction": "Solar faculae are areas on the solar surface surrounding active regions and appearing bright toward the limb. Understanding their properties is important from several points of view. On the one hand, facular contrast influences total solar irradiance variations. On the other hand, faculae are believed to be composed of conglomerates of magnetic elements. Therefore, in faculae we can observe the interaction of convection with relatively strong magnetic fields, which is interesting in itself from a physical point of view and also to constrain state of art numerical models of magneto-convection. It is well established that the origin of facular brightness excess is related to the presence of magnetic field: the brightness excess is caused by the modified radiative transfer through magnetic atmospheres viewed from different angles \\citep[e.g.][]{Keller+etal2004, Carlsson+etal2004, Vogler2005, Okunev+Kneer2005, Steiner2005}. The variation of facular brightness near the solar limb can provide information about the structure and properties of the magnetic elements composing facular regions. Measurements of centre-to limb variations of the contrast done by different authors mostly agree that the contrast increases till $\\mu=0.2-0.4$ \\citep[][]{Muller1977, Auffret+Muller1991, Topka+etal1997, Ortiz+etal2002, Okunev+Kneer2004, Hirzberger+Wiehr2005}. A controversy exists for the extreme limb, where it is not clear if the contrast further increases \\citep[][]{Boer1997, Sutterlin1999} or falls off toward the limb. The variety of results can be attributed to the difference in spatial resolution of observations, facular size, its magnetic field strength, as well as selection criteria. At high spatial resolution (0.\\arcsec1--0.\\arcsec2) and at disk centre, facular regions appear as conglomerates of bright points, small pores and small-scale granular structures \\citep{Lites+etal2004, Berger2007, Narayan+etal2010}. The contrast of faculae at disk centre is slightly positive or close to zero, though some negative values are also detected \\citep[see][and references therein]{Title1992, Topka+etal1997}. Several classes of models claim to reproduce facular brightness properties. The ``hot wall'' model, initially proposed by \\citet{Spruit1976}, consists of a magnetic flux tube, evacuated due to the presence of magnetic field, similar to Wilson depression in sunspots. Viewed from an angle, the line-of-sight penetrates deeper into the tube because of its lower density and, thus, hotter layers become visible. This idea was followed in works by e.g., \\citet{Topka+etal1997, Steiner2005, Okunev+Kneer2005}, who provide evidences that the ``hot wall'' model can closely reproduce observed facular properties, such as continuum contrast, its centre-to-limb variation, as well as the dependence of contrast on the magnetic field. Much more complex models of facular regions based on ``realistic'' MHD simulations \\citep{Keller+etal2004, Carlsson+etal2004, Vogler2005} suggest that faculae are seen bright on the limb because hot granular walls become visible through transparent magnetic field concentrations. The three-dimensional structure of granules becomes apparent at faculae near the solar limb in high resolution observations and is well reproduced in the simulations. Two factors are crucial for the facula formation: the shape of the granule limb-ward of the flux concentration and the size of the magnetic field concentration. If the flux concentration is too small, the opacity reduction is not sufficient to have the continuum intensity formed exclusively in the hot granule. A quantitative disagreement still remains in the values of the peak brightness in faculae, being substantially larger in simulations compared to observations \\citep{Keller+etal2004}. \\citet{Berger2007} question the dependence of granular brightness on magnetic field strength obtained in earlier works \\citep{Topka+etal1997, Ortiz+etal2002}. \\citet{Berger2007} analyze extremely high spatial resolution data (close to 0.\\arcsec1) obtained at the Swedish Solar Telescope on La Palma. They claim that if, instead of analyzing the brightness for binned magnetogram signals \\citep[as is done in][]{Topka+etal1997}, one analyzes magnetic flux density for facular points segmented from the data set, there is no dependence of the brightness on the magnetic flux density. \\citet{Berger2007} propose that what we see as faculae are granular walls, and not interiors of the magnetic flux tubes. Since granules all have similar properties, there is no dependence of their brightness on the magnetic field, and the magnetic field only plays an indirect role, making the atmosphere transparent in front of the granules. The strong magnetic field of facular regions modifies the properties of convection. Simulations of magnetic flux emergence through the convection zone show that during the emergence granules become larger, more elongated and smoother \\citep{Cheung+etal2008}. Observationally, in already formed facular region granulation presents a lot of fine structuring at high spatial resolution: isolated bright points, strings of bright points and dark micro-pores, ribbons, or more circular flower structures \\citep{Title1992, Berger2004, Narayan+etal2010}. Granules become smaller than in quiet areas \\citep{Muller1977, Schmidt1988, Title1992} and intergranular lanes are characterized by the presence of micro-pores. Granular velocities in facular and plage areas are found to be generally lower than in the nearby quiet areas \\citep{Nesis+Mattig1989, Title1992, Narayan+etal2010}. Velocities measured in plage and facular areas are found to depend on the magnetic field strength. From medium-to-high (430 km) resolution observations of the disk centre plage in \\NiI\\ 6768 \\AA\\ line, \\citet{Title1992} obtain an increase of downflow velocities with magnetic field strength up to 600 G, and a decrease for stronger fields. A similar result is reported by \\citet{Montagne1996, Berger2004, Morinaga2008, Narayan+etal2010}. The correlation between velocity and intensity fields, typical for granulation, is found to be partially destroyed by the magnetic field \\citep{Rimmele2004, Narayan+etal2010}. One of the possible reasons of the de-correlation, proposed by \\citet{Narayan+etal2010} is that small scale convection in plage areas does not overshoot to the same height as field-free convection and leaves only weak traces at the line-forming region. It is clear that further observational studies of magneto-convection in strongly magnetized plage and facular areas are needed to constrain the relationship between magnetic field and granulation, as well as the properties of magnetic elements. Here we report on such an observational study based on state-of-the-art simultaneous observations done at the German Vacuum Tower Telescope at the Observatorio del Teide (Tenerife) with two instruments: TIP-II \\citep{Collados2007} and TESOS \\citep{Tritschler+etal2002}. We analyze correlation between granular/intergranular intensities, velocities, and magnetic field; as well as the results of the heights of sign reversals of contrast and velocity as a function of the magnetic field. \\begin{figure*} \\centering \\includegraphics[width=18cm]{kostik_f1.eps} \\caption{An overview of the observations. The panels from left to right are: \\BaII\\ 4554 \\AA\\ continuum intensity in the units of its spatially average value; mask applied to locate granules and intergranular lanes; \\BaII\\ 4554 \\AA\\ line centre intensity in the units of its spatially average value; magnetic field strength from the inversion of \\FeI\\ IR lines. Contours (same in all panels) mark the locations of granular areas with magnetic field above 1.2 kG. } \\label{fig:fov} \\end{figure*} ", "conclusions": " \\begin{enumerate} \\item At the bottom photosphere, the convective velocities of granules decrease with magnetic field strength, while the convective velocities of intergranules increase with magnetic field strength. \\item Similar to the quiet regions, we detect that in facular regions the contrast of granulation and the velocities of convective elements reverse their sign with height. The height where such reversal takes place depends on the strength of convective elements at the bottom photosphere, but also on the magnetic field strength. Convective elements above stronger magnetic features reach higher in the atmosphere without breaking. \\item The correlation coefficient between velocity and intensity at the bottom-most level decreases with the magnetic field. In the upper atmosphere, above 500 km, the correlation gets larger with increasing the magnetic field. \\item The contrast of convective elements decreases with increasing magnetic field strength. But if intergranular lanes are considered separately, their contrast is nearly independent of the field strength. \\end{enumerate}" }, "1207/1207.2750_arXiv.txt": { "abstract": "We investigate the effect of the environment on the Faber Jackson (FJ) relation, using a sample of 384 nearby elliptical galaxies and estimating objectively their environment on the typical scale of galaxy clusters. We show that the intrinsic scatter of the FJ is significantly reduced when ellipticals in high density environments are compared to ellipticals in low density ones. This result, which holds on a limited range of overdensities, is likely to provide an important observational link between scaling relations and formation mechanisms in galaxies. ", "introduction": "The Faber Jackson (FJ) is the first scaling relation discovered for elliptical galaxies. Already described by Morgan \\& Mayall (1957), although in a qualitative way, it was given its first quantified form 20 years later by Faber \\& Jackson (1976) who, on the basis of a handful of nearby early type galaxies, were able to proof the existence of a power law relation linking luminosity ($L_{\\rm B}$) to central velocity dispersion ($\\sigma_0$). Soon thereafter Kormendy (1977) found a second scaling relation which holds for elliptical galaxies, relating the effective surface brightness ($\\mu_{\\rm e}$) to the effective radius ($R_{\\rm e}$). The Kormendy relation was refined ten years later for ellipticals by Hamabe \\& Kromendy (1987) and in that same year Dressler et al (1987) and Djorgovsky \\& Davis (1987) discovered a more general relation (the Fundamental Plane, FP) linking Log $R_{\\rm e}$, Log $\\sigma_0$ and $\\mu_{\\rm e}$. The scaling relations are powerful tools that can be used to derive galaxy distances and, even more important, constitue an invaluable observational bench mark for theoretical models. It is especially for this latter reason that they have been the subject of much interest since their discovery. Understanding origin and nature of the scaling relations is a fundamental quest for any successful theory of galaxy formation which is expected to be able to predict the observed slope, scatter, possible variation (as a function of luminosity, wavelength, environment) and evolution (with z). The much narrower scatter displayed by the FP with respect to the FJ and Kormendy relations made the former one to become rapidly more attractive than the latter two which were easily interpreted to be partial representations (projections) of the FP onto a lower dimensional space \\cite{dr1987,dj1987,fa1987,dz1991}. Very recently the FJ appears to have captured again attention on the theoretical point of view, as Sanders (2010) has claimed it to be more fundamental and universal than the FP within the context of MOND (modified Newtonian dynamics). This finding is expected to motivate renewed interest in the FJ, which so far has not been largely investigated. There is no much work which has been carried out on the FJ relation if one excepts studies which have provided evidence for a decrease of its stepness at low luminosity \\cite{to1981,da1983,he1992,fr2005,ma2005,be2006,ds2007,la2007,vo2007,ko2012} and studies devoted to investigate the effect of luminosity, mass and redshift on it \\cite{fr2005,be2006,ds2007,ni2010,ni2011}. At variance with the FP for which the effect of the environment has been largely investigated, although with rather conflicting results \\cite{de1992,ma1996,ma1999,de2001,tr2001,be2003,ev2002,go2003,re2004,re2005,de2005,do2008,la2010}, so far only 4 studies exist \\cite{zi2005,fr2005,fr2009,fz2009} which have looked for possible effects induced by the environment on the FJ relation of rather distant (z $\\in $ [0.2 -- 0.7]) early-type galaxies, without finding however any strong evidence for them. According to the standard cosmological paradigm, structures in the present day Universe have formed through a hiearchical scenario process predicting rather different assembling time scales and evolutionary paths for galaxies in high and low density regions (Baugh et al. 1996; Kauffmann \\& Charlot 1998; Somerville \\& Primack 1999; Kauffmann et al. 2004). Environment is thus expected to play a relevant role in shaping galaxy properties and is likely to leave its imprint in the scaling relations as well. This is the reason which has motivated the above mentioned studies (mostly concentrated on the FP) and the present work devoted to investigate the effect of the environment on the FJ relation, using a sample of 384 nearby ellipticals and estimating their environment on the typical scale of galaxy clusters. The structure of the paper is the following: in $\\S$2 we present the sample, in section $\\S$3 we derive the FJ relation for the whole sample and for its bright and faint components and test the robustness of our results accounting both for errors on $\\sigma_0$ and $m_{\\rm B}$, in $\\S$4 we illustrate the method that we have used to estimate the environment, in $\\S$5 we show that the scatter of the FJ gets largely reduced in high density environments and increased in low density ones and that this difference is neither induced by errors on $\\sigma_0$ nor by luminosity difference between the samples, in $\\S$6 we show that the scatter of the FJ relation increases with decreasing density in overdense environments and decreases with increasing density in underdense environments, in $\\S7$ we draw the conclusions. A Hubble constant of $H_0$ = 70 km s$^{-1}$ Mpc$^{-1}$ is adopted throughout. ", "conclusions": "Using a sample of 384 nearby elliptical galaxies and objectively estimating their environment on the basis of the number of neighbours within the typical galaxy cluster and group scale we have provided evidence for an effect relating the intrinsic scatter of the FJ relation to the environment. We have shown that the scatter of the FJ is reduced to almost half of its value when ellipticals in highest overdensities are compared to ellipticals in less-density environments, that the effect is not induced by luminosity differences between the samples and that it holds for overdensities ranging between 3.5 and 5 the median value of the number of neighbours distribution. Besides indicating a rather simple and quite natural way to reduce the large scatter affecting the FJ relation, our result, if confirmed on larger samples, is very likely to open an interesting perspective for models of galaxy formation." }, "1207/1207.5493_arXiv.txt": { "abstract": "Traditional approaches to the study of the dynamics of spacetime curvature in a very real sense hide the intricacies of the nonlinear regime. Whether it be huge formulae, or mountains of numerical data, standard methods of presentation make little use of our remarkable skill, as humans, at pattern recognition. Here we introduce a new approach to the visualization of spacetime curvature. We examine the flows associated with the gradient fields of scalar invariants derived from the spacetime. These flows reveal a remarkably rich structure, and offer fresh insights, even for well known analytical solutions to Einstein's equations. The intent, however, is to go beyond idealized analytical solutions and eventually consider physically realistic situations. This requires a careful analysis of exactly which invariants that can actually be used in this approach. The present analysis serves as an overview and as an introduction to this program. ", "introduction": "To quote from a recent monograph \\cite{gp}: \\textit{``It appears that it is much easier to find a new solution of Einstein's equations than it is to understand it\".} The present work involves the novel use of both computer algebra (for background calculations fundamental to the approach) and numerical routines (for numerical integration and visualization) with the objective being the development of a fresh view of ``curvature\", and an consequent ``understanding\" (for example, in the case of Einstein's theory) of a given spacetime. In this endeavor we are not alone. In \\cite{dan}, a visualization of spacetime curvature, restricted to a study of the projected electric and magnetic components of the Weyl tensor, is being developed. In \\cite{rez} a complementary approach is also under development. Here, we examine something akin to a highly complex dynamical systems approach: we look at the ``flows\" associated with the gradient fields of scalar invariants of the background space. Preliminary results show that these flows reveal a remarkably rich structure and fresh insights even for well known geometries. The subject of invariants, even at dimension 4, is not closed in a mathematical sense. The approach considered here then necessarily imposes restrictions on the invariants used. We demand that any ``useable\" invariant does not single out any particular observer (which requires the definition of ``observer independent\" invariants, as explained below). The objects fundamental to \\cite{dan} are also used here but in quite a different way: we use no projection. The approach we use does not, in a fundamental sense, use any particular theory. However, we also demand that the invariants used be connected directly to the underlying physics (and not offer just some general geometrical interpretation) and so we must adopt a particular theory to form this connection. We formulate the presentation given here within Einstein's theory of gravity. ", "conclusions": "This paper serves as an introduction to the study of gradient flows of scalar invariants as a means to visualize curvature. No particular observer, or theory, is fundamental to this approach, but the intention is to apply the techniques to spacetimes associated with Einstein's theory of gravity. This allows for a physical understanding of the scalars in use. For example, a physical understanding of the Ricci invariants follows directly from Einstein's equations. Because the approach used here does not single out particular observes, we have shown that this restriction limits the number of invariants that can be used to construct gradient flows. Only the 4 Ricci invariants and the 4 Weyl invariants (in terms of its electric and magnetic components) can be used as building blocks to study gradient flows in the most general case. We have shown that higher order (mixed) invariants explicitly exhibit vectors and scalars associated with the timelike 4-vector used to split the Weyl tensor. It has been shown that gradient flows of invariants polynomial in the Riemann tensor are necessarily orthogonal to Killing flows, should they exist. (This allows for a rigorous definition and generalization of the classical notions of $\\mathcal{R}$ and $\\mathcal{T}$ regions of spacetime.) From the point of view of dynamical systems, gradient flows are simple in the sense that critical points are the only possible limit sets of the flow. However, the classification of critical points must be made in a coordinate independent fashion and so a covariant classification scheme was developed. This, along with the Poincar\\'{e} index, classify the ``topology\" of the flow essential to the visualization. For purely electric Ricci - flat spacetimes one can construct strict Newtonian analogues, based on the topology of flows; for the first electric (tidal) invariant in spacetime, and for the tidal invariant in Newtonian theory. This is discussed at length elsewhere \\cite{curzon}. We have reviewed examples, restricted here to spherical symmetry. A rather complete analysis of the Kruskal - Szekeres vacuum was given, interpreting the associated bifurcate 2-sphere as the isolated critical point of the solution. More generally, isolated critical points of the Weyl invariant within spherical symmetry were distinguished as locally isotropic or anisotropic, with explicit examples given for each. \\bigskip" }, "1207/1207.5988_arXiv.txt": { "abstract": "We report on a search for particle dark matter with the XENON100 experiment, operated at the Laboratori Nazionali del Gran Sasso (LNGS) for 13~months during 2011 and 2012. XENON100 features an ultra-low electromagnetic background of $(5.3 \\pm 0.6) \\times 10^{-3}$\\,events/(keV$_{\\n{ee}}\\times$kg$\\times$day) in the energy region of interest. A blind analysis of 224.6\\,live days $\\times$ 34\\,kg exposure has yielded no evidence for dark matter interactions. The two candidate events observed in the pre-defined nuclear recoil energy range of 6.6-30.5\\,keV$_{\\n{nr}}$ are consistent with the background expectation of $(1.0 \\pm 0.2)$\\,events. A Profile Likelihood analysis using a 6.6-43.3\\,keV$_{\\n{nr}}$ energy range sets the most stringent limit on the spin-independent elastic WIMP-nucleon scattering cross section for WIMP masses above 8\\,GeV/$c^2$, with a minimum of $2 \\times 10^{-45}$\\,cm$^2$ at 55\\,GeV/$c^2$ and 90\\% confidence level. ", "introduction": " ", "conclusions": "" }, "1207/1207.5807_arXiv.txt": { "abstract": "We perform binary stellar evolutionary calculations following the simultaneous evolution of both stars in the system to study a potential progenitor system for the Type IIb supernova 2011dh. Pre-explosion photometry as well as light-curve modeling have provided constraints on the physical properties of the progenitor system. Here we present a close binary system that is compatible with such constraints. The system is formed by stars of solar composition with 16~$\\mathrm{M_{\\odot}}$~+~10~$\\mathrm{M_{\\odot}}$ on a circular orbit with an initial period of 125~days. The primary star ends its evolution as a yellow supergiant with a mass of $\\approx 4 \\, \\mathrm{M_{\\odot}}$, a final hydrogen content of $\\approx 3-5 \\times 10^{-3}$~$\\mathrm{M_{\\odot}}$ and with an effective temperature and luminosity in agreement with the HST pre-explosion observations of SN~2011dh. These results are nearly insensitive to the adopted accretion efficiency factor $\\beta$. At the time of explosion, the companion star has an effective temperature of 22 to 40 thousand Kelvin, depending on the value of $\\beta$, and lies near the zero age main sequence. Considering the uncertainties in the HST pre-SN photometry the secondary star is only marginally detectable in the bluest observed band. Close binary systems, as opposed to single stars, provide a natural frame to explain the properties of SN~2011dh. ", "introduction": "\\label{sec:intro} Core-collapse supernovae (CCSNe) are the explosive end of massive stars with $M_{\\mathrm{ZAMS}} \\gtrsim 8 \\mathrm{M_{\\odot}}$. There is a diversity in the spectroscopic and photometric properties of CCSNe which are mainly related to the ability of the progenitor to retain its outermost layers. Type II SNe, with clear H lines in their spectra, represents the case where a thick H envelope is kept before the explosion. Type Ib SNe, with no H lines but with clear He lines, have lost their H envelope but not the He layers. Finally Type Ic SNe, with no H and He lines in the spectra, represent a more extreme case where not only the the H but also the He envelopes are likely lost before the explosion. There are also transitional objects between these different types. One example is that of Type IIb SNe, which show H lines at early times but then the spectrum is transformed into that of typical SNe~Ib \\citep[see][for a classification scheme]{1997ARA&A..35..309F}. Type IIb, Ib and Ic objects are collectively called striped-envelope SNe \\citep{1996ApJ...462..462C}. Progenitor models of SNe~IIb comprising a helium star surrounded by a very thin hydrogen-rich envelope (of $\\lesssim 1 \\, \\mathrm{M_{\\odot}}$) have been successful to explain the observed light curves (LC) and spectral features \\citep{1994ApJ...420..341S,1994ApJ...429..300W,1998ApJ...496..454B}. However, it is not clear which is the mechanism responsible for the removal of the outer envelope before the explosion. One possibility is strong winds that occur in massive stars with $M_{\\mathrm{ZAMS}} \\gtrsim 25 \\mathrm{M_{\\odot}}$. Alternatively, in close binary systems (CBS) stars are expected to exchange mass providing an efficient mechanism to allow for the removal of outer layers. Currently, the binary channel is favored particularly for the case of SNe~IIb \\citep{2008MNRAS.384.1109E,2011A&A_528A_131C,2011MNRAS.412.1522S}. Additional support for the binary scenario in SNe~IIb comes from the detection of a hot companion for the famous SN~IIb 1993J~\\citep{2004Natur_427_129M}. This was initially suggested by pre-explosion photometry \\citep{1994AJ....107..662A}. The LC and evolutionary models of SN~1993J were also in favor of the binary channel \\citep{1993Natur_364_507N,1993Natur_364_509P,1994ApJ...429..300W}. Some evidence for a companion was also reported for another SN~IIb 2001ig \\citep{2006MNRAS.369L..32R}. The Type IIb SN~2011dh was recently discovered in the nearby galaxy M51 attracting the attention of many observers because of its proximity and brightness. It was discovered almost immediately after explosion \\citep{2011ApJ_742L_18A}. It showed early radio and X-ray emission \\citep{2012ApJ...752...78S}. Using pre-explosion images obtained from the HST archive \\citet{2011ApJ_739L_37M} and \\citet{2011ApJ_741L_28V} detected a source at the location of SN~2011dh. They derived similar values of luminosity and effective temperature for pre-SN source. The object was consistent with a yellow supergiant (YSG) star with a radius $R \\approx 270 \\, \\mathrm{R_{\\odot}}$ and without any clear evidence of a companion star contributing to the observed spectral energy distribution (SED). At present there is a controversy in the literature as to whether the YSG is the actual progenitor of SN~2011dh. Some authors have suggested that the exploding star should be more compact \\citep{2011ApJ_742L_18A,2012ApJ...752...78S,2011ApJ_741L_28V} based on (1) a simple comparison between the early light curve (LC) of SN~2011dh and SN~1993J; (2) a discrepancy between the temperature derived from an early-time spectrum and that predicted by an analytic expressions for an extended progenitor; and (3) the large shock velocity derived from radio observations. Recently, we have performed a detailed hydrodynamical modeling of SN~2011dh using stellar evolutionary progenitors \\citep{2012ApJ...757...31B}. These models indicate that observations are compatible with a helium star progenitor of a mass near $4$~$\\mathrm{M_{\\odot}}$ surrounded by a thin hydrogen-rich envelope ($\\approx$ $0.1$ $\\mathrm{M_{\\odot}}$) with a radius of $\\approx 200 \\,\\mathrm{R_{\\odot}}$ (similar to that of the detected YSG star) that underwent an explosion with an energy of $8\\times10^{50}$~erg that synthesized $0.063$~$\\mathrm{M_{\\odot}}$ of \\Ni. Such large radius values are needed to reproduce the early light curve of SN~2011dh without contradicting the temperatures derived from the spectra. In addition, our hydrodynamical modeling rules out progenitors with He core masses larger than 8~$\\mathrm{M_{\\odot}}$, which corresponds to $M_{\\mathrm{ZAMS}} \\gtrsim 25 \\mathrm{M_{\\odot}}$. It is very difficult for a single star to reach these pre-SN conditions. The existence of a strong wind capable of removing most of the envelope requires a massive star of $\\approx$ 25 $\\mathrm{M_{\\odot}}$ or more \\citep{2003ApJ...591..288H,2009A&A...502..611G}, which is in contradiction with the LC models. Moreover, in order to retain a thin hydrogen-rich layer, the mass loss rate would have to be on a very narrow interval. These facts strongly suggests that the progenitor of SN~2011dh should be a component of a binary system. However, a recent work by \\citet{2012A&A_538L_8G} proposed that single YSG stars such as the one detected at the location of SN~2011dh are plausible SN progenitors. This is based on stellar evolution calculations of stars with main sequence masses of 12--15 $\\mathrm{M_{\\odot}}$ under the assumption of an increased mass-loss rate several times above the standard values. However, no physical explanation is given for such an increase. Also note that a recent paper by \\citet{2011A&A...526A.156M} found a good agreement between modern determinations of mass-loss for RSGs and the standard mass-loss prescription \\citep{1988A&AS_72_259D}. Although, other mass-loss formulation as proposed by \\citet{2005A&A...438..273V} point towards higher mass-loss rates but this prescription seems to be applicable only to dusty stars and gives a overestimates of the mass-loss rates for Galactic RSGs. The aim of this work is to show the plausibility that the progenitor of SN~2011dh was part of a close binary system (CBS) with properties compatible with the pre-SN observations and the results of LC modeling. The observational properties of the remaining companion star are discussed in anticipation of future detections. Although we do not perform a complete exploration of the parameter space (stellar masses, initial orbital period, and mass-transfer efficiency $\\beta$), we show that our results are robust if we consider moderate changes of the initial conditions. The remainder of this paper is organized as follows. In Section~\\ref{sec:code} we present a brief description of our binary stellar evolution code paying special attention to the characteristics that enabled us to compute pre-SN models. In Section~\\ref{sec:results} we present the main results of this paper regarding the adopted binary configuration (\\S~\\ref{subsec:configuration}), evolutionary calculations (\\S~\\ref{subsec:evolu}) as well as the spectra of the components at the moment of the explosion (\\S~\\ref{subsec:sed}). In Section~\\ref{sec:disc} we present a discussion of our results and finally, in Section~\\ref{sec:conclu} we provide some concluding remarks. ", "conclusions": "\\label{sec:conclu} With the aim of providing a description of the progenitor of SN~2011dh, we have studied the evolution of close binary systems of solar composition stars with masses of 16~$\\mathrm{M_{\\odot}}$~+~10~$\\mathrm{M_{\\odot}}$. We considered an initial period of 125~days and different efficiencies ($\\beta$) of the mass transfer process. We followed the simultaneous evolution of the donor and accreting stars from the zero age main sequence up to the oxygen core exhaustion of the donor. We found that the donor star, independently of $\\beta$, ends its evolution with effective temperature and luminosity consistent with the YSG object detected in the HST pre-SN photometry. The exploding star has a mass $M\\approx$~4\\,$\\mathrm{M_{\\odot}}$, a radius $R\\approx$~250\\,$\\mathrm{R_{\\odot}}$ and an outermost layer containing $3-5 \\times 10^{-3}$~$\\mathrm{M_{\\odot}}$ of hydrogen. This is generally consistent with the type IIb classification and the results of LC modeling of SN~2011dh by \\citet{2012ApJ...757...31B}. These results are a natural consequence of the close binary evolution and require no external adjustment of any physical condition. Regarding the accretion efficiency, $\\beta$, we found that (1) the evolution of the donor star is almost independent of $\\beta$ while the secondary strongly depends on it, and (2) the evolution of the orbital period, the MTR and the total hydrogen content are almost independent of the value of $\\beta$. Our calculations indicate that the donor star is lossing mass at the moment of the explosion with rates that differ markedly from constant. Inferences on the mass of the donor star at the time of the explosion should take into account the appropriate mass loss in this phase. We also found some indication that the total hydrogen content may be a function of the initial orbital period with larger period producing a larger the hydrogen content. A more detailed study of this point is left for future work. Note that the structure of the donor star at the moment of the explosion is consistent with an extended SN IIb but with very little H mass ($< 0.1$ $\\mathrm{M_{\\odot}}$). We analyzed the effect of the secondary star on the observed HST pre-explosion photometry. For all the values of $\\beta$, at the moment of the explosion of the donor, the secondary star is still near the ZAMS. This is a direct consequence of our assumption that the object has a mass appreciably lower than that of the donor. The effective temperature of the companion is far higher than that of the donor with values within of 22 to 40 thousand Kelvin. Thus, the largest contribution to the flux of the system from the secondary is in the bluest observed band, F336W, producing a marginal detection of 0.6--2$\\sigma$ level depending of the value of $\\beta$. Unfortunately, the available HST pre-SN observations are not very suitable to constrain the properties of the secondary. The ultimate proof of the binary nature of SN~2011dh must come from the possible detection of a very hot star once the SN light fades enough. This situation would be similar to what occurred with SN~1993J but with different properties of the companion. In any case, we should remark that detecting the companion star of SN~2011dh would provide valuable information on the efficiency of the mass transfer process and evolution of massive CBSs in general." }, "1207/1207.5955_arXiv.txt": { "abstract": "We revisit potential impacts of nuclear burning on the onset of the neutrino-driven explosions of core-collapse supernovae. By changing the neutrino luminosity and its decay time to obtain parametric explosions in one-(1D) and two-dimensional (2D) models with or without a 13-isotope $\\alpha$ network, we study how the inclusion of nuclear burning could affect the postbounce dynamics for four progenitor models; three for $15.0 \\Msun$ stars, one for an $11.2 \\Msun$ star. We find that the energy supply due to nuclear burning of infalling material behind the shock can energize the shock expansion especially for models that produce only marginal explosions in the absence of nuclear burning. These models are energized by nuclear energy deposition when the shock front passes through the silicon-rich layer and/or later it touches the oxygen-rich layer. Depending on the neutrino luminosity and its decay time, a diagnostic energy of explosion increases up to a few times $10^{50}$ erg for models with nuclear burning compared to the corresponding models without. We point out that these features are most remarkable for the Limongi-Chieffi progenitor in both 1D and 2D, because the progenitor model possesses a massive oxygen layer with its inner-edge radius being smallest among the employed progenitors, so that the shock can touch the rich fuel on a shorter timescale after bounce. The energy difference is generally smaller ($\\sim 0.1-0.2 \\times 10^{51}$ erg) in 2D than in 1D (at most $\\sim 0.6 \\times 10^{51}$ erg). This is because neutrino-driven convection and the shock instability in 2D models enhance the neutrino heating efficiency, which makes the contribution of nuclear burning relatively smaller compared to 1D models. Considering uncertainties in progenitor models, our results indicate that nuclear burning should remain as one of the important ingredients to foster the onset of neutrino-driven explosions. ", "introduction": "Ever since the dawn of modern core-collapse supernova (CCSN) theory, the neutrino-heating mechanism \\citep{colgate}, in which a supernova shock is revived by neutrino energy deposition to trigger explosions \\citep{wils85,bethe85}, has been the leading candidate for the explosion mechanism for more than four decades. However, the simplest, spherically-symmetric (1D) form of this mechanism fails, except for super-AGB stars at the low-mass end \\citep{bernhard12a}, to explode canonical massive stars \\citep{Rampp00,lieb01,thom03,Sumiyoshi05}. Pushed by accumulating supernova observations of the blast morphology \\citep[e.g.,][and references therein]{wang08,tanaka12}, a number of multi-dimensional (multi-D) hydrodynamic simulations have been reported so far, which gives us a confidence that hydrodynamic motions associated with convection \\citep[e.g.,][]{Herant92,Burrows95,Janka96,frye02,fryer04a} and the Standing-Accretion-Shock-Instability \\citep[SASI, e.g., ][]{Blondin03,Scheck04,scheck06,Ohnishi06,ohnishi07,ott_multi,Murphy08,Foglizzo06,thierry07,endeve12,thierry12,iwakami1,iwakami2,rodrigo09_2,rodrigo09,rodrigo10,hanke,rodrigo13} can help the onset of neutrino-driven explosions \\citep[see collective references in][]{thomas12,kotake12}. In fact, neutrino-driven explosions have been obtained in first-principle two-(2D) and three-(3D) dimensional simulations in which spectral neutrino transport is solved by various approximations \\citep[e.g.,][]{kotake12b}. The Garching group \\citep{Buras06a,Buras06b,Marek09, hanke13, bernhard11,bernhard12a,bernhard12b,bernhard13} included one of the best available neutrino transfer approximations by the ray-by-ray variable Eddington factor method. The Oak Ridge group \\citep{bruenn09,bruenn13} included a ray-by-ray multi-group flux-limited diffusion transport with the best available weak interactions. The Nippon group\\footnote{\"Nippon\" stands for Japan in Japanese, and from now on we like to call our team as so whose members come all around Japan.} \\citep{Suwa10,Suwa11,suwa12,Takiwaki11,takiwaki13} employed a ray-by-ray isotropic diffusion source approximation \\citep{idsa} with a reduced set of weak interactions\\footnote{See \\citet{sumiyoshi12} for collective references about detailed neutrino transport schemes.}. This success, however, is accompanying further new question. The explosion energies obtained in some 2D models are underpowered by up to a factor of 10 compared to the canonical supernova kinetic energy ($\\sim 10^{51}$ erg, see table 1 in \\citet{Kotake11} for a current summary). What on earth is missing furthermore ? 3D hydrodynamics has been pointed out to boost the onset of neutrino-driven explosions compared to 2D \\citep{nordhaus10}, although it is still under considerable debate \\citep{hanke,hanke13,Takiwaki11,couch12}. Very recently, general relativity has been reported to help the onset of multi-D neutrino-driven explosions by \\citet{bernhard11,bernhard12b} in 2D simulations with detailed neutrino transport and by \\citet{kuroda12,kuroda13} in 3D simulations but with approximate neutrino transport. Impacts of nuclear equations of state (EOSes) have been investigated in multi-D simulations by \\citet{Marek09,marek09b,suwa12} and \\citet{couch13}. However, there may still remain further room to study more detailed nuclear physical impacts in these first principle multi-D simulations, such as the density dependence of symmetry energy and the skewness of compressibility \\citep{steiner,lattimer12} and influences of light nuclei \\citep[e.g.,][]{sumi08,arcones,nakamura} and of inelastic neutrino-nucleus scattering \\citep[e.g.,][]{haxton,ohnishi07,langanke08} on enhancing the neutrino heating rates in the gain region. More recently, impacts of improved neutrino interactions based on the 1D full Boltzmann simulations have been elaborately investigated \\citep{lentz11, lentz12}. The neutrino-driven mechanism would be assisted by other candidate mechanisms such as the acoustic mechanism \\citep[e.g.,][]{burr06} or the magnetohydrodynamic mechanism (e.g., \\citet{kota04a,kota04b,taki04,taki09,burr07,fogli_B,martin11,taki_kota}, see also \\citet{kota06} for collective references therein). Other possibilities include QCD phase transitions in the core of the protoneutron star \\citep[e.g.,][]{takahara88,sage09} viscous heating by the magnetorotational instability \\citep{thomp05}, or energy dissipation via Alfv\\'en waves \\citep{suzu08}. Joining in these efforts to look for some possible ingredients to foster explosions, we pay attention to the roles of nuclear burning in this study. To the best of our knowledge, \\citet{janka01a} were the first to clearly point out that an additional energy released by nuclear burning of infalling material behind the shock could make a significant contribution to affect the explosion energy (see their Eq.(5)). The mass in the silicon (Si) layer, depending sensitively on the progenitor masses and structures, is in the range of $\\sim 0.3 - 0.6 M_{\\sun}$ \\citep{WW95,woos02,limongi}. Since the release of nuclear energy in Si burning is $\\approx 10^{18}~{\\rm erg \\, g}^{-1}$, a few $10^{50}$ erg are expected to be deposited by the explosive nuclear burning. It should be noted that nuclear burning has been included in the full-scale simulations by the Garching group \\citep{rampp02,Buras06a, Buras06b,Marek09}, in which composition changes of silicon, oxygen (and similarly neon and magnesium), and carbon and their nuclear energy release are computed by a ``flashing\" treatment \\citep[see Appendix \\ref{app-flash}, and also Appendix B.2 in ][]{rampp02}. In a series of multi-D simulations in which neutrino transport is treated by a more approximative way to follow a long-term postbounce evolution in the context of the neutrino-driven mechanism, nuclear burning is included by a small network calculation \\citep[e.g.,][]{kifo,scheck06,annop,hammer,arcones_2011,ugliano}. However in these literatures, impacts of nuclear burning on the supernova dynamics have not been unambiguously investigated so far. In conference proceedings, the Oak Ridge group reported 2D explosion models based on their radiation-hydrodynamic simulations \\citep{bruenn06,mezza07} for 11.2$\\Msun$ and 15.0$\\Msun$ stars, only when an alpha network calculation was included, but not when they applied the flashing treatment. They pointed out that oxygen burning assists the (weak) shock to move farther out due to the additional pressure support in the vicinity of the weak shock. These situations motivate us to revisit the impacts of nuclear burning on assisting the shock propagation by performing hydrodynamic simulations including a network calculation. In the present work, we take the following strategy to clearly see the roles of nuclear burning. Firstly we try, in the spirit of \\citet{burogoshy} and \\citet{Janka01}, to find a critical condition in 1D, in which nuclear burning affects the criteria of explosion. Instead of performing full-scale radiation-hydrodynamic simulations which are computationally expensive, we employ a light-bulb scheme to trigger explosions \\citep[e.g.,][]{Janka96} for the sake of our systematic survey. Previously the role of nuclear burning seems to be considered as negligible using a very limited set of progenitor models but we will show that for a previously untested progenitor model, nuclear burning can really push the weak shock farther out to help explosions. This paper opens with the description of numerical setup including information about our hydrodynamic code with nuclear network and about initial models (Section 2). Results are given in section 3. After giving a detailed explanation in 1D models how nuclear burning could affect the postbounce dynamics (section \\ref{sec-1d}, \\ref{sec-prog}), we move on to discuss our 2D models to study how nuclear burning interacts with multi-D hydrodynamics (section \\ref{sec-2d}). We summarize our results and discuss their implications in Section 4. ", "conclusions": "We revisited the potential impacts of nuclear burning on the onset of neutrino-driven explosions of core-collapse supernovae. By changing the neutrino luminosity and its decay time to obtain parametric explosions in 1D and 2D models with or without a 13-isotope $\\alpha$ network, we studied how the inclusion of nuclear burning could affect the postbounce dynamics for four progenitor models; three for $15.0 \\Msun$ stars of \\citet{limongi}, \\citet{WW95}, and \\citet{woos02}, and one for an $11.2 \\Msun$ star of \\citet{woos02} Our results showed that the energy gain due to nuclear burning of infalling material behind the shock can energize the shock expansion especially for models that produce only marginal explosions in the absence of nuclear burning. These models enjoy the assistance from nuclear burning typically in the following two ways, whether the shock front passes through the silicon-rich layer, or later it touches to the oxygen-rich layer. Depending on the neutrino luminosity and its decay time, the diagnostic energy of explosion was found to increase up to a few times $10^{50}$ erg for models with nuclear burning compared to the corresponding models without. The energy difference becomes generally smaller in 2D than in 1D, because neutrino-driven convection and the SASI in 2D models enhance the neutrino heating efficiency, making the contribution of nuclear burning relatively smaller compared to 1D models. It was pointed out that these features are most remarkable for the LC15 progenitor, which possesses a massive oxygen layer with its inner-edge radius being smallest among the employed progenitors, which makes the timescale shorter for the shock to encounter the rich fuel. Considering reduction of the critical luminosity and increase of the diagnostic energy by nuclear burning, and also uncertainties in the structure of progenitors, our results indicate that nuclear burning should still remain as one of the important ingredients to foster the onset of neutrino-driven explosions." }, "1207/1207.5678_arXiv.txt": { "abstract": "We present the implementation of a fast estimator for the full dark matter bispectrum of a three-dimensional particle distribution relying on a separable modal expansion of the bispectrum. The computational cost of accurate bispectrum estimation is negligible relative to simulation evolution, so the isotropic bispectrum can be used as a standard diagnostic whenever the power spectrum is evaluated. As an application we measure the evolution of gravitational and primordial dark matter bispectra in $N$-body simulations with Gaussian and non-Gaussian initial conditions of the local, equilateral, orthogonal and flattened shape. The results are compared to theoretical models using a 3D visualisation, 3D shape correlations and the cumulative bispectrum signal-to-noise, all of which can be evaluated extremely quickly. Our measured bispectra are determined by $\\mathcal{O}(50)$ coefficients, which can be used as fitting formulae in the nonlinear regime and for non-Gaussian initial conditions. In the nonlinear regime with $k<2h\\,\\mathrm{Mpc}^{-1}$, we find an excellent correlation between the measured dark matter bispectrum and a simple model based on a `constant' bispectrum plus a (nonlinear) tree-level gravitational bispectrum. In the same range for non-Gaussian simulations, we find an excellent correlation between the measured additional bispectrum and a constant model plus a (nonlinear) tree-level primordial bispectrum. We demonstrate that the constant contribution to the non-Gaussian bispectrum can be understood as a time-shift of the constant mode in the gravitational bispectrum, which is motivated by the one-halo model. The final amplitude of this extra non-Gaussian constant contribution is directly related to the initial amplitude of the constant mode in the primordial bispectrum. We also comment on the effects of regular grid and glass initial conditions on the bispectrum. ", "introduction": "Observations of the cosmic microwave background (CMB) \\cite{wmap7,shellard1006,Cooray2, FRS2} and large scale structure (LSS) \\cite{citeulike:2833069,citeulike:9218385} are currently consistent with Gaussian primordial cosmological perturbations. Much effort has been undertaken in recent times to constrain a possible non-Gaussian contribution. Detection of a significant primordial bispectrum or trispectrum would have major implications for the mechanism of inflation, possibly ruling out the simplest paradigm of canonical single-field slow-roll inflation. Given this possibility, the importance of developing methods that discriminate between the plethora of inflationary models is clear. Such general methods have been developed in the case of the CMB in \\cite{shellard0812,shellard0912,shellard1006,RSF10,FRS2}. This work exploited the use of a separable expansion of the underlying bispectrum or trispectrum in order to greatly reduce the computational cost involved in analysing general shapes. The CMB bispectrum is currently the most powerful and direct probe of primordial non-Gaussianity, and higher resolution temperature and polarisation data will soon be available from Planck. Nevertheless, interest in the possibility of using observables of large scale structure to test primordial non-Gaussianity has undergone a recent resurgence, due in part to the potentially three dimensional nature of the dataset as a result of its redshift dependence. In particular, the scale-dependent bias induced by non-Gaussian local initial conditions on the halo power spectrum has been shown in \\cite{dalal08} to offer an additional and powerful probe. Much work has been undertaken to improve analytic predictions for the impact of non-Gaussian initial conditions on the matter and galaxy power spectrum and bispectrum (see for example \\cite{sefusatti09,sefusatti1003,scoccimarro-couchman01,verde1111,Scocc2012,seljak-stochasticity1104,Sef2011,Scocc0906,Desj1105,kendrick-stochasticity,10035020} and references therein). In \\cite{Verde2010, citeulike:9218385} constraints on local-type non-Gaussianity were found at levels competitive with those obtained using CMB data. Despite these advances it is notable that until recently only local type non-Gaussianity had been studied using large-scale structure, owing to the difficulty in generating generic initial conditions. This situation greatly improved due to work by Wagner and Verde \\cite{wagner-verde1006,wagner-verde1102} and work by Scoccimarro et al.~\\cite{scocci1108}. However, both of these approaches become inefficient for non-separable bispectra. In \\cite{shellard1008} the possibility of using the separable expansion method, exploited to great effect in the case of the CMB \\cite{shellard1006}, was explored. This approach was verified in \\cite{shellard1108} to allow for a far more efficient generation of non-Gaussian initial conditions than had previously been possible in the literature. Furthermore, the separable decomposition methodology allows for the study of non-separable shapes. This has allowed for a robust and truly general approach towards the generation of primordial non-Gaussian initial conditions for use in $N$-body simulations. In this paper we exploit this approach to set up and run $N$-body simulations with non-Gaussian initial conditions. In particular, we study initial conditions given by local, equilateral, orthogonal and flattened bispectra, respectively. We choose the flattened (or trans-Planckian) model as an explicit example of an inherently non-separable bispectrum. We exploit the modal method to efficiently and accurately reconstruct the full 3D matter bispectrum at each redshift of interest for each of these non-Gaussian models. This allows us to correlate the results of the $N$-body simulations with analytic estimates, allowing us to identify clearly the regime at which nonlinear corrections become important. This applies first to accurately determining the gravitational bispectrum well into the nonlinear regime, which is necessary in order to differentiate the more subtle impact of primordial non-Gaussianity. The correlation measure proves useful in distinguishing the different primordial shapes. We also thoroughly test the impact of starting $N$-body simulations using either glass or regular grid initial conditions. We emphasise that our purpose here is to study the detailed nature of the underlying matter bispectrum derived from $N$-body simulations. These methods can be equally applied to efficiently extract the galaxy bispectrum from huge survey data sets or to predict the halo bispectrum from simulations in a realistic observational context, but this will be the subject of future work. The paper is organised as follows. In Section II we review the generation of non-Gaussianity due to nonlinear gravitational evolution, together with the physically-motivated models we will use to describe our results. In Section III we review primordial non-Gaussianity, introduce the specific models studied in this paper and discuss their impact on the matter bispectrum, proposing a new time-shift model. In Section IV we discuss the bispectrum estimation methodology based on a separable expansion which allows extremely efficient estimation of the full $N$-body bispectrum, as well as the cumulative signal-to-noise and its correlation to theoretically-predicted bispectra. We discuss the simulation setup, impact of glass initial conditions, validations and convergence tests in Section V. In Section VI we present our results for the gravitational bispectrum and the simple fits thereof, while primordial bispectra in non-Gaussian simulations and their fits are discussed in section VII. Finally, in Section VIII we present our conclusions. Readers familiar with the modal methodology and interested primarily in the measured bispectra and testing of various fitting formulae may wish to start with Section VI and follow the references to the earlier sections given there. ", "conclusions": "We have presented an implementation of a bispectrum estimator for $N$-body simulations using a separable modal expansion of the bispectrum as described in \\cite{shellard1008}. While a brute force estimation of the full bispectrum is computationally expensive, requiring $\\mathcal{O}(N^6)$ operations for $N$ particles per dimension, we find that the gravitational and the most prominent primordial bispectra can be approximated by only $n_\\mathrm{max}=\\mathcal{O}(50)$ separable basis functions (for the range of scales relevant for $N$-body simulations). The bispectrum projection on the corresponding subspace of all possible bispectra is estimated with $\\mathcal{O}(n_\\mathrm{max}N^3)$ operations, which is faster than brute force estimation by a factor of $\\mathcal{O}(N^3/n_\\mathrm{max})\\sim \\mathcal{O}(10^7)$ for typical simulations. Thus the computational cost for accurate 3D bispectrum estimation is almost negligible compared to the cost for running the $N$-body simulations (e.g.~we can estimate the full bispectrum of a $1024^3$ grid up to $k_\\mathrm{max}=\\tfrac{N}{4}\\tfrac{2\\pi}{L}$ in one hour on only $6$ cores). This allows us to estimate the bispectrum as a standard simulation diagnostic whenever the power spectrum is measured. Expressing the bispectrum using its $n_\\mathrm{max}$ separable components yields a radical compression of data which simplifies further analysis like comparisons between theory and simulations. The separable bispectrum estimator therefore provides a very useful additional statistic characterising the formation of structures in $N$-body simulations, with high sensitivity to different shapes of primordial non-Gaussianity corresponding to different models of inflation. We have performed many $N$-body simulations with Gaussian initial conditions as well as non-Gaussian initial conditions of the local, equilateral, orthogonal and (non-separable) flattened shape, exploiting the separable bispectrum expansion for efficient generation of initial conditions as described for primordial fields in \\cite{shellard1008,shellard1108}. On large scales the measured gravitational and primordial bispectra agree with leading order perturbation theory and with measurements for Gaussian initial conditions by \\cite{verde1111} in the mildly nonlinear regime, demonstrating the unbiasedness of the estimator and the initial conditions. In the nonlinear regime, the gravitational bispectrum becomes dominated by a large `constant' signal receiving elongated and equilateral contributions not captured by tree level perturbation theory. However, it remains suppressed in the squeezed limit, where primordial bispectrum signals can peak (see \\fig{tet3dplots_grav_loc_eq_orth}). Our measured $N$-body bispectra for Gaussian and non-Gaussian simulations can be expressed by $50$ components $\\beta^R_n$, which we provide in Table \\ref{tab:betas-nbody} for the key models. They can be used as fitting formulae for the gravitational and primordial bispectrum in the nonlinear regime. Less accurate but simpler fitting formulae are obtained by modeling the bispectrum as a combination of partially loop-corrected perturbative bispectra and a simple `constant' bispectrum, which is constant on slices $\\sum k_i= \\mathrm{const.}$ and which is obtained as an approximation to the $1$-halo bispectrum. While the former contribution dominates on large scales and early times, the latter constant contribution dominates in the nonlinear regime. Interpreting the effect of primordial non-Gaussianity on the constant bispectrum contribution as a time-shift with respect to Gaussian simulations allows us to model the time dependence of the constant bispectrum contribution for non-Gaussian initial conditions. For Gaussian initial conditions, the simple fits achieve a shape correlation of at least $99.8\\%$ with the measured gravitational bispectrum for $z\\leq 20$ and $k_\\mathrm{max}=\\{0.5,2\\}h/\\mathrm{Mpc}$. For local, equilateral and flattened non-Gaussian initial conditions the primordial bispectrum is fitted with a shape correlation of at least $97.9\\%$ at $z=0$ and at least $94.4\\%$ for $z\\leq 20$ (with correlation typically $\\gtrsim 98\\%$ for most shapes and redshifts, see Table \\ref{tab:msfit_NG_table}). The impact of the orthogonal shape seems to be somewhat harder to model, because it does not have a constant component initially, but our simple fit still achieves correlations of at least $91\\%$ for $z\\leq 20$. Throughout this work we have visualised the measured bispectra in three-dimensional tetrapyd plots \\cite{shellard0912}, which show the bispectrum shape and amplitude at different length scales and generalise commonly used plots of one- or two-dimensional slices through the tetrapyd volume. For a more quantitative analysis, particularly to test analytic predictions and fitting formulae, we have made extensive use of full three-dimensional shape correlations, the cumulative signal-to-noise of the bispectra and their projection $f_\\mathrm{NL}$ on theoretical shapes. These quantities have been evaluated extremely efficiently using the bispectrum components obtained from the separable estimator. We find that regular grid initial conditions produce an initial spurious bispectrum due to the anisotropy of the regular grid, which can be avoided by using glass initial conditions. However the difference between regular grid and glass initial conditions decreases with time as gravitational perturbations grow such that both initial conditions yield similar results at late times. Effects of order $f_\\mathrm{NL}^2$ were shown to affect the bispectrum measurements by less than $5\\%$ in our large scale simulations. Clearly, further work is required to study the effects of general primordial non-Gaussianity on observable quantities like the bispectrum of galaxies, particularly in the nonlinear regime. However the present work represents, we believe, an important step forward in the understanding of structure formation in the presence of primordial non-Gaussianity and the search for primordial non-Gaussianity in large scale structures." }, "1207/1207.4190.txt": { "abstract": "The growth of supermassive black holes appears to be driven by {\\changed galaxy mergers, violent merger-free processes and/or `secular' processes}. In order to quantify the effects of secular evolution on black hole growth, we {\\changed study} a sample of active galactic nuclei (AGN) in galaxies {\\changed with a calm formation history free of significant mergers}, a population that heretofore has been difficult to locate. Here we present an initial sample of 13 AGN in massive ($M_\\ast \\gtrsim 10^{10}~\\mmsun $) bulgeless galaxies --- which lack the classical bulges believed inevitably to result from mergers --- selected from the Sloan Digital Sky Survey using visual classifications from Galaxy Zoo. Parametric morphological fitting confirms the host galaxies lack classical bulges; any contributions from pseudobulges are very small (typically $< 5$\\%). %This is the largest such sample yet assembled. % BDS: I am taking the above sentence out because \"largest such sample\" is too unclear, and I don't know how to make it clear in the abstract. It's not the largest sample of bulgeless AGN hosts. It's the largest sample of massive AGN host galaxies at z < 0.1 with (probably) no classical bulges and (definitely) no big pseudobulges, where the black hole mass plays no direct role in the selection of the sample. For example, Jiang et al. (2011a,b) have a larger number of galaxies with bulge-to-total < 5%, but a) their sample extends to higher redshifts, b) the mass distribution of their ~bulgeless host galaxies is lower than that of our sample (though not by a lot), and c) they select first on low BH mass, so the vast majority of their sample has optical broad lines that indicate MBH < 1e6 Msun. Based on luminosity/Eddington arguments it seems likely that most of our growing BHs are ~1e6 Msun or larger (with many an order of magnitude larger), and *that* is what hasn't been seen before in a bulgeless galaxy sample. I feel more comfortable with the later statement that we're demonstrating BH growth beyond 1e6 Msun is possible without mergers. We compute black hole masses for the two broad-line objects in the sample ($4.2 \\times 10^6$ and $1.2 \\times 10^7~\\mmsun $) and place lower limits on black hole masses for the remaining sample (typically $M_{\\mathrm{BH}} \\gtrsim 10^6~\\mmsun $), showing that significant black hole growth must be possible in the absence of mergers {\\changed or violent disk instabilities}. The black hole masses are systematically higher than expected from established bulge-black hole relations. However, if the mean Eddington ratio of the systems with measured black hole masses ($L/L_{\\rm{Edd}} \\approx 0.065$) is typical, 10 of 13 sources are consistent with the correlation between black hole mass and \\emph{total} stellar mass. That pure disk galaxies and their central black holes may be consistent with a relation derived from elliptical and bulge-dominated galaxies with very different formation histories implies the details of stellar galaxy evolution and dynamics may not be fundamental to the co-evolution of galaxies and black holes. ", "introduction": "% % %%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%% Constraining the contribution of mergers to the evolution of the galaxy population is one of the fundamental challenges in modern galaxy formation theory. Galaxies have long been believed to form hierarchically, building up to their observed sizes through a series of mergers \\citep{white78,kauffmann93}. The merger history of each galaxy thus contributes significantly to the galaxy's stellar and gas dynamics, and is also thought to drive the co-evolution of a galaxy with its central supermassive black hole \\citep[SMBH;][]{sanders88,dimatteo05,croton06,hopkins06b,hopkins08a}. Given such a {\\changed fundamental} effect, distinguishing between the effect of mergers and {\\changed that of other evolutionary pathways, such as the slow, internal processes collectively known as `secular' evolution}, is difficult. In this paper we present a sample of massive galaxies {\\changed chosen to have had no} significant merger in their history, discuss their properties, {\\changed and demonstrate for the first time with such a large sample that substantial black hole growth (to $M_{\\rm BH} \\gtrsim 10^{6-7}~\\rm{M}_{\\odot}$) is possible without the advent of a significant merger.} A galaxy's morphology contains signatures of its evolutionary history. In particular, the assembly of massive disk galaxies through mergers inevitably produces a central bulge component \\citep[e.g.,][]{toomre77,walker96,hopkins11c,martig12}, {\\changed and some merger-free processes such as violent disk instabilities can also form a bulge \\citep[e.g.,][]{noguchi99,d_elmegreen04}.} The bulge is dynamically hot, rising vertically above the disk, and has a steeper density profile than an exponential disk \\citep{devaucouleurs}. A galaxy lacking a central bulge thus {\\changed must have} a formation history free of {\\changed violent formation processes. This implies a lack of} significant mergers, with a strong limit on the mass ratio between the main galaxy and any accreting satellite galaxies \\citep[$\\sim 1:10$;][though \\citeauthor{brook12} \\citeyear{brook12} suggest the ratio may be {\\changed as high as $1:4$}]{walker96,hopkins11c}. Such bulgeless galaxies, with a purely secular formation history, might be expected to be rare in a hierarchical scenario. The presence amongst the galaxy population of large bulgeless galaxies thus presents a serious challenge to this picture \\citep{kormendy10}, as they cannot have undergone a significant merger yet have assembled stellar masses of $M_\\ast \\gtrsim 10^{10} \\rm{M}_{\\odot}$. Additionally, the well-established {\\changed correlations between} galaxies and their central SMBHs {\\changed \\citep[e.g.,][]{magorrian98,kormendy01,ferrarese00,tremaine02,marconi03,haringrix04} have} led to the prevalence of major-merger-driven theories for black hole-galaxy co-evolution \\citep{sanders88,dimatteo05,croton06,hopkins08a}. {\\changed However,} a growing body of recent work suggests minor mergers{\\changed , cold accretion} and secular processes may be a more typical means of growing a galaxy and its central black hole, both locally \\citep[e.g.,][]{greene10b,jiang11b} and at higher redshift \\citep[e.g.,][]{simmons11,cisternas11,schawinski11,schawinski12,kocevski12}. Owing in part to the compounded rarity of both massive, bulgeless galaxies and active galactic nuclei (AGN), the extent to which a SMBH can grow in the absence of merger processes remains difficult to characterise. {\\changed Galaxies lacking classical bulges but hosting AGN have previously been found; these typically have lower stellar masses compared to the general galaxy population, and/or host black holes with relatively low black hole masses (e.g., NGC 4395, \\citeauthor{filippenko03} \\citeyear{filippenko03}; NGC 3621, \\citeauthor{satyapal07} \\citeyear{satyapal07}; NGC 4178, \\citeauthor{satyapal09} \\citeyear{satyapal09}, \\citeauthor{secrest12} \\citeyear{secrest12}; NGC 3367 and NGC 4536, \\citeauthor{mcalpine11} \\citeyear{mcalpine11}; NGC 4561, \\citeauthor{arayasalvo12} \\citeyear{arayasalvo12}). In some cases, these properties are at least in part a direct result of sample selection, as in studies based on samples of low-mass black holes \\citep{gh04, gh07a, greene08,greene10b,jiang11a,jiang11b}.} This paper uses morphological classifications from the Galaxy Zoo\\footnote{www.galaxyzoo.org} project \\citep{lintott08,lintott11} to construct a sample of bulgeless galaxies that host actively growing black holes. Selecting galaxies that lack classical bulges (as opposed to galaxies with a more varied history of both secular and merger-driven evolution) enables the isolated study of black hole growth in the absence of mergers. {\\changed This selection includes optical detection of an AGN but no restriction on its black hole mass.} These galaxies provide a strong challenge to models of galaxy formation, requiring substantial and ongoing secular growth of a central black hole. We aim to use this rare population to assess whether these galaxies fall on the same galaxy-{\\changed black hole} relations seen in galaxies with more merger-driven histories. By comparing upper limits on bulge masses to black hole masses from broad emission lines {\\changed and} lower limits on black hole masses using Eddington limits, we assess the sizes to which black holes can grow over their lifetimes due to secular processes alone. Section \\ref{sec:finding} describes the methods used to select bulgeless galaxies with growing black holes {\\changed from Galaxy Zoo and the Sloan Digital Sky Survey \\citep[SDSS;][]{york00}}. Section \\ref{sec:sample} presents the sample of host galaxies, with Section \\ref{sec:mbh} detailing how the black hole masses and lower limits are calculated. In Section \\ref{sec:discussion} we discuss how bulgeless AGN host galaxies inform our understanding of the co-evolution of black holes and galaxies. Throughout this paper, we assume $H_0 = 71~\\mathrm{km/s/Mpc}$, $\\Omega_M = 0.27$ and $\\Omega_\\Lambda = 0.73$, consistent with the most recent WMAP cosmology \\citep{komatsu11}. %%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%% % % ", "conclusions": "% % %%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%% By using classifications of visual morphologies of Sloan Digital Sky Survey galaxies, drawn from the Galaxy Zoo project, we have selected a large set of bulgeless face-on spiral galaxies. A conservative initial selection identifies 13 of these galaxies which are unambiguously systems with growing black holes. Parameterized fitting of these galaxies provides stringent limits on bulge or pseudobulge mass, with the latter typically contributing $\\sim3$\\% by mass on average. Two of the galaxies in the sample have broad-line AGN, and thus measurements of their black hole mass are possible. For the rest, infrared observations from the \\emph{WISE} mission allow us to place a lower limit on the black hole mass. The black hole masses are substantial, reaching $\\sim10^7~\\mmsun$, and lie above those predicted by the local bulge-black hole mass relation, even when all the pseudobulge component is included as an upper limit to the mass of the classical bulge in each case. One of the two black holes with measured masses is fully consistent with the relation between black hole and total stellar mass (the other is just outside the scatter). If the Eddington limits of the black holes with measured masses are typical of the full sample, $80\\%$ of the systems for which only lower limits are available have black hole masses consistent with predictions based on \\emph{total} galaxy stellar mass. Firm conclusions require further observations, but it is not inconsistent with the idea that black hole mass is more closely related to the overall gravitational potential of the galaxy and its dark matter halo (which is dominated by the halo but traced by the total stellar mass) than to the dynamically hot bulge component. In any case, the presence of massive, growing super-massive black holes in bulgeless galaxies indicates that secular evolution is an important part of the evolution of the galaxy population. Either significant black hole growth is possible even in the absence of significant {\\changed bulge-building mechanisms}, or a dynamical means to keep galaxies bulgeless despite these mechanisms must be found. Future work will include an analysis of the more than 10,000 candidate bulgeless galaxies from which this sample was drawn in order to constrain the properties of this intriguing population; in particular, a search for bulgeless systems in mergers will distinguish between the two scenarios left open. Observational follow-up of the small sample identified here, particularly in order to constrain more tightly the bulge properties and black hole masses is also urgently necessary. Although extending the work to higher redshift will be challenging, Galaxy Zoo classifications for large \\emph{Hubble Space Telescope} studies hold the promise of identifying a similar set of galaxies out to a redshift of approximately one. Even at low redshift, however, the systems identified here present a stringent test of simulations of galaxy formation. %%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%% % %" }, "1207/1207.6398_arXiv.txt": { "abstract": "We present a kinematic analysis of the globular cluster (GC) system of the giant elliptical galaxy NGC 4365 and find several distinct kinematic substructures. This analysis is carried out using radial velocities for 269 GCs, obtained with the DEIMOS instrument on the Keck II telescope as part of the SAGES Legacy Unifying Globulars and Galaxies Survey (SLUGGS). % We find that each of the three (formerly identified) GC colour subpopulations reveal distinct rotation properties. The rotation of the green GC subpopulation is consistent with the bulk of NGC 4365's stellar light, which `rolls' about the photometric major axis. The blue and red GC subpopulations show `normal' rotation about the minor axis. We also find that the red GC subpopulation is rotationally dominated beyond 2.5 arcmin ($\\sim17$ kpc) and that the root mean squared velocity of the green subpopulation declines sharply with radius suggesting a possible bias towards radial orbits relative to the other GC subpopulations. Additionally, we find a population of low velocity GCs that form a linear structure running from the SW to the NE across NGC 4365 which aligns with the recently reported stellar stream towards NGC 4342. These low velocity GCs have $g'-i'$ colours consistent with the overall NGC 4365 GC system but have velocities consistent with the systemic velocity of NGC 4342. We discuss the possible formation scenarios for the three GC subpopulations as well as the possible origin of the low velocity GC population. % \\vspace{0.3in} ", "introduction": "\\begin{figure*}\\centering \\includegraphics[width=0.85\\textwidth]{footprints.jpg} \\caption{The positions of the six DEIMOS slitmasks plotted on the Subaru Suprime-Cam (S-Cam) $i'$ filter image. The scale is shown in the bottom left corner of the $35 \\times 27$ arcmin image. It is centred on $\\alpha$=12:24:26.824; $\\delta$=+07:19:03.52 (J2000.0). At a distance of $23.1\\pm0.8$ Mpc \\citep{FCS5} $1\\,\\mathrm{arcmin} =6.72$ kpc.} \\label{fig:footprints} \\end{figure*} Globular clusters (GCs), as some of the oldest and densest stellar systems in the universe, present intriguing questions about their own formation while also being very useful tracers of the formation of their host galaxy \\citep{Br06}. They probably formed during violent star formation episodes in their host galaxy (or host galaxy progenitors) and therefore we can use their spatial distribution and kinematic properties to unravel the formation history of galaxies. GCs are more numerous and more easily identifiable in elliptical galaxies (where bright star forming clumps and dust obscuration is minimal) than in spiral galaxies. % NGC 4365 is a giant elliptical galaxy (E3), $M_B=-21.3$ $M_V=-22.7$ mag \\citep{VCS6, Ko09}, with various intriguing and unusual properties as highlighted below. We use these unusual properties of NGC 4365 to constrain galaxy and GC formation scenarios by placing the extra information contained in this system in the overall framework of `normal' giant elliptical galaxies. The GC systems of most giant elliptical galaxies are bimodal in optical colours \\citep{Br06,VCS9}. There is a growing body of evidence that the relative age difference between GC subpopulations is small \\citep{St05} and the distinct colours of the commonly observed subpopulations are due to bimodality in GC metallicity \\citep{St07,Wo10,AB11} that are likely explained by two formation mechanisms, sites or epochs. Possible formation scenarios for bimodal metallicity distributions in GC systems are the major merger \\citep{Ze93}, multiphase collapse \\citep*{F97} and accretion scenarios \\citep{C98}. An alternative view is that the bimodal GC colour distributions can be observed from a unimodal GC metallicity distribution, because of strong nonlinearity in the relationship between colour and metallicity \\citep{Y06,Ca07,Bl10} and in this case the kinematic properties of each colour subpopulation are not expected to be significantly different. The GC system of NGC 4365 has three GC subpopulations, that includes an additional subpopulation of `green' GCs between the commonly observed blue and red GC subpopulations. This additional subpopulation, thought at the time to be younger than the `normal' two subpopulations, was discovered by \\citet{Pu02} using a combination of near-IR and optical photometry. \\citet{Br05} suggested that the three subpopulations might be found as trimodality in optical colours and \\citet{La05} found that the green subpopulation was restricted to small galactocentric radii. % Specifically for NGC 4365, two independent methods have confirmed the GCs for all three subpopulations (blue, green and red) to be older than $\\sim10$ Gyr. \\citet{Br05} observed a spectroscopic sample of 22 GCs to measure their ages and metallicities using Lick indices from the Low-Resolution Imaging Spectrograph (LRIS) on the Keck telescope. They found that NGC 4365 has three old GC subpopulations with metal poor (blue), intermediate metallicity (green) and metal rich (red) GCs. \\citet{Ch11} presented a near-IR photometry and optical analysis of 99 GCs with much higher precision than previously possible, and determined that there is no observable offset in the mean age between the GCs of NGC 4365 and GC populations of other giant ellipticals. This indicates that there is no significant population of young GCs in NGC 4365. In addition to its almost unique GC system, NGC 4365 has very unusual stellar kinematic properties. The kinematically distinct core (KDC) at its centre is relatively uncommon, seen in $\\sim10$ per cent of early type (E and S0) galaxies in the ATLAS$^{\\mathrm{3D}}$ volume limited survey \\citep{Kr11}. Also very uncommon is the starlight kinematics outside the KDC which `rolls' about the major axis rather than rotating about the minor axis \\citep{Su95}. Less than 5 per cent of galaxies, in the ATLAS$^{3D}$ volume limited survey, show this kinematic behaviour. \\citet{Da01} used the SAURON integral field spectrograph to map its kinematic and metallicity structure in two dimensions, finding stars both inside and outside the KDC to be older than 10 Gyr. NGC 4365 is one of only two early type galaxies, in the ATLAS$^{3D}$sample of 260, that have both a KDC and a near $90^\\circ$ misalignment between the photometric and kinematic major axes (i.e.\\ rolling rather than rotating stars). Lastly \\citet{Bo12}, see also Mihos et al.\\ (2012, in prep), presented evidence of a stream of stellar light extending $\\sim200$ kpc southwest and $\\sim 100$ kpc northeast from NGC 4365. This strong indicator of a very recent $\\sim10:1$ merger (occurring with in the last few Gyr) presents a possible cause for the unique properties of NGC 4365. \\citet{Blom12} presented an optical photometry analysis of $\\sim 4000$ NGC 4365 GCs, using the most spatially extended sample of GCs analysed to date (observed with Subaru/Suprime-Cam and supplemented with archival imaging from Hubble Space Telescope/Advanced Camera for Surveys) and found different spatial distributions, sizes and mass distributions for the three subpopulations of NGC 4365 GC system. They also concluded that separate formation scenarios are required to explain the existence of the three separate subpopulations of GC in NGC 4365. Kinematic analysis of the three GC subpopulations in NGC 4365 is key to disentangling the possible formation scenarios of GC systems. Several recent investigations have shown that the kinematic features of the two standard GC subpopulations (seen in other giant elliptical galaxies) are different \\citep{Le10a,Ar11,Fos11}. When combined with galaxy formation models this information constrains the possible formation scenarios of GCs. For example, \\citet{Fos11} compare kinematics of the GC subpopulations in NGC 4494 with simulations such as those described in \\citet{Be05} and conclude that the galaxy has undergone a recent major merger of similar disk galaxies. The addition of the kinematic properties of a third GC subpopulation limits the possible formation scenarios even further. In Section 2 we first describe the preparation, observation and reduction of the spectroscopic sample of GCs and then investigate the colour/metallicity and line-of-sight velocity distributions of the GCs. The separation of the GCs into subpopulations and analysis of the kinematic features of the three subpopulations are presented in Section 3. We then discuss the results and summarise our conclusions regarding the possible formation scenarios for NGC 4365 in Sections 4 and 5. The Appendices contain the kinematic fits for all the different methods of subpopulation separation we investigated, and an alternative assumption of position angle for kinematic fitting. ", "conclusions": "NGC 4365 is a giant elliptical galaxy with rare stellar kinematic properties and also, very unusually, three globular cluster (GC) subpopulations. Recently, a stellar stream running across NGC 4365 and its nearby neighbour, NGC 4342, was discovered. With high velocity resolution Keck/DEIMOS spectra we analyse the kinematics of a large number and spatially extended sample of GCs around NGC 4365. These data allow us to separate the GC system into its three colour subpopulations and analyse the kinematics of each subpopulation in detail. We find distinct rotation properties for each GC subpopulation. This indicates that the three GC subpopulations are distinct and that the formation history of NGC 4365 is complex. We conclude that it is unlikely that the third (green) GC subpopulation is related to the existence of the stellar stream and that it is also unlikely to have originated from an accreted galaxy. We also find a further group of low velocity GCs (covering the full range in colour), which might be related to the stellar stream extending across NGC 4365 towards NGC 4342 ($\\sim 35$ arcmin away at a similar systemic velocity to these GCs). Future analysis of spectroscopic observations of GCs in the stream will further illuminate the complex formation history of NGC 4365. %" }, "1207/1207.1747_arXiv.txt": { "abstract": "The gaseous molecular disk that orbits the main sequence A-type star 49 Ceti has been known since 1995, but the stellar age and the origin of the observed carbon monoxide molecules have been unknown. We now identify 49 Ceti as a member of the 40 Myr old Argus Association and present a colliding comet model to explain the high CO concentrations seen at 49 Ceti and the 30 Myr old A-type star HD 21997. The model suggests that massive -- 400 Earth mass -- analogs of the Sun's Kuiper Belt are in orbit about some A-type stars, that these large masses are composed primarily of comet-like objects, and that these objects are rich in CO and perhaps also CO$_2$. We identify additional early-type members of the Argus Association and the Tucana/Horologium and Columba Associations; some of these stars display excess mid-infrared emission as measured with the Widefield Infrared Survey Explorer (WISE). ", "introduction": "Extensive observations have shown that by the time a typical T Tauri star is $\\sim$5 Myr old its protostellar disk will retain insufficient carbon monoxide gas to be detectable with a radio telescope. Stars with ages $<$5 Myr are virtually non-existent within $\\sim$100 pc of Earth (Zuckerman \\& Song 2004; Torres et al 2008). Thus, stars this close to Earth with detectable gas-rich circumstellar disks are few and far between. CO has been detected from three classical T Tauri stars within 100 pc of Earth: TW Hya, V4046 Sgr and MP Mus (Kastner et al 2010; Rodriguez et al 2010; and references therein). Such stars, with likely ages in the range 7-12 Myr, represent the oldest known examples of remnant gaseous protoplanetary disks around solar-type stars. Two dusty A-type stars near Earth with gas-rich disks -- 49 Cet and HD 21997 -- are quite different from classical T Tauri stars. The circumstellar gas at 49 Cet (= HR 451 \\& HIP 7345) was discovered in 1995 and has warranted the attention of single and interferometric radio telescopes (Zuckerman et al 1995; Dent et al 2005; Hughes et al 2008). The dust distribution has been investigated with high spatial resolution mid-infrared imaging (Wahhaj et al 2007). The gas around HD 21997 (= HR 1082 \\& HIP 16449) was discovered much more recently (Moor et al 2011). 49 Cet and HD 21997 are, respectively, 59 and 72 pc from Earth (van Leeuwen 2007). HD 21997 is classified, reasonably securely, as a member the 30 Myr old Columba Association (Moor et al 2006; Torres et al 2008). The age of 49 Cet has been much more up in the air. The earliest estimates of this age were typically 8-10 Myr and 8 Myr is the age adopted in the interferometric study by Hughes et al (2008). With such a young age, no older than that of the three classical T Tauri stars mentioned above, it was reasonable to interpret 49 Cet as a star in transition between a protoplanetary and a debris disk. However, largely because they did not associate the Galactic space motions (UVW) of 49 Cet with any known young (8-100 Myr old) stellar moving group, Rhee et al (2007) tentatively assigned an age of \"20?\" Myr to the star. Based on the discussion that follows below, we believe that all previous age estimates for 49 Cet are too young and that the star is actually a member of the 40 Myr old Argus Association (Torres et al 2008; Zuckerman et al 2011). Thus among main sequence stars, 49 Cet and HD 21997 contain the longest-lived substantial reservoirs of circumstellar gas currently known in astronomy. Moor et al (2011) outline how difficult it is to explain how so much molecular gas can be present in the circumstellar disks of such old stars. In the present paper we present a model that appears capable of accounting for the properties of such stars. ", "conclusions": "We present a massive (400 Earth mass) comet-cloud model to explain the large quantities of carbon monoxide gas seen at the 30-40 Myr old, A-type, stars HD 21997 and 49 Ceti. Because CO is rapidly photodissociated in the stellar and interstellar radiation fields, it must be produced rapidly. We calculate that this production rate is an order of magnitude faster than the rate of production of dust particles via a model of collisional cascade in the steady state. The implications of this ratio are: (1) young, pristine comets are likely to be richer in CO and perhaps CO$_2$ than typically observed solar system comets, and (2) the destruction of the youthful comets around HD 21997 and 49 Cet are not likely to have been in a steady state situation over periods of time exceeding millions of years. If dynamical activity in debris disks is ramping up at the age of these stars, for example through the build-up of large bodies as in the model of the young Kuiper Belt proposed by Schlichting \\& Sari (2011) or the model of debris disks around A-type stars by Kenyon \\& Bromley (2010), then the production of small dust particles may lag well behind the rate of CO outgassing. The primary driving force for CO outgassing is likely to be collisions among the 100 trillion comets that orbit 49 Cet and HD 21997. Such a model for these systems is consistent with origin in a protoplanetary disk of initial mass of order a few tenths of a solar mass and suggests that comets form very early in the history of such disks before there is sufficient time to convert much CO to CH$_4$. If comet collisions provide the observed CO, then the H$_2$/CO ratio is unconstrained and may not be large. If so, then collisional excitation of the CO rotational levels may be dominated by something other than H$_2$, perhaps electrons. Acke et al (2012) have proposed a rapid comet-destruction model for the dusty debris disk that orbits Fomalhaut that, at first glance, might seem similar to our model for HD 21997 and 49 Ceti. But, in fact, the two models are very different. As we note in Appendix C, the Acke et al dust loss rate is 20 times faster than the rate we calculate via a conventional collisional cascade. They do not specify a physical mechanism that can convert comets so rapidly into dust. Their rapid dust loss rate at Fomalhaut is compelled by their quoted very short lifetime (1700 yr) for dust grains near the blowout size. But a detailed model of the similar debris disk at Vega by M\\\"uller et al (2010 and as outlined in our Appendix C) casts doubt on the validity of such a short lifetime. Additional observations and modeling of the Fomalhaut debris disk are warranted. In particular, a deep search for CO at Fomalhaut would be worthwhile; because it is the presence of so much CO that compels the rapid comet destruction rate in our model, while combination of observations of CO and of dust serve to constrain the model. At the risk of being obvious, dusty young stars will be excellent targets for ALMA mapping of dust and CO gas, should the latter be detectable. It should be possible to map the shape and mass of youthful Kuiper Belt analogs and to clarify the composition of young comets. \\vskip 0.2in We thank Meredith Hughes, Alexander Krivov, and Hilke Schlichting for taking time to clarify some aspects of their papers, David Jewitt and Michael Jura for helpful comments, and the referee for useful suggestions. This research was funded in part by NASA grants to UCLA and the University of Georgia. \\appendix" }, "1207/1207.4953_arXiv.txt": { "abstract": "We present new low-resolution \\HI\\ spectral line imaging, obtained with the {\\it Karl G. Jansky Very Large Array} ({\\it JVLA}), of the star-forming Magellanic irregular galaxy UGCA\\,105. This nearby (D = 3.39$\\pm$0.25 Mpc), low mass (M$_{\\rm HI}=$ 4.3$\\pm$0.5\\,$\\times$\\,10$^{8}$ \\msun) system harbors a large neutral gas disk (\\HI\\ radius $\\sim$ 7.2 kpc at the N$_{\\rm HI}$ = 10$^{20}$ cm$^{-2}$ level) that is roughly twice as large as the stellar disk at the B-band R$_{\\rm 25}$ isophote. We explore the neutral gas dynamics of this system, fitting tilted ring models in order to extract a well-sampled rotation curve. The rotation velocity rises in the inner disk, flattens at 72$\\pm$3 \\kms, and remains flat to the last measured point of the disk ($\\sim$7.5 kpc). The dynamical mass of UGCA\\,105 at this outermost point, (9$\\pm$2)\\,$\\times$\\,10$^{9}$ \\msun, is $\\sim$10 times as large as the luminous baryonic components (neutral atomic gas and stars). The proximity and favorable inclination (55\\degr) of UGCA\\,105 make it a promising target for high resolution studies of both star formation and rotational dynamics in a nearby low-mass galaxy. ", "introduction": "\\label{S1} Dwarf galaxies offer an opportunity to study various processes that bear on galaxy evolution. Nearby systems allow an exploration of the interplay between ongoing star formation and the multi-phase interstellar medium (ISM). Further, nearby gas-rich systems are amenable to detailed studies of galactic rotational dynamics in the absence of differential shear. Many dwarfs display solid-body rotation that is well-suited to precision rotation curve work \\citep[e.g.,][]{oh08}. Most nearby systems appear to be dark-matter dominated \\citep{mateo98}, making them important laboratories for studying both the luminous and the dark mass components in galaxies. UGCA\\,105 (see Table~\\ref{t1} for representative qualities) is a Magellanic-type irregular galaxy with ongoing star formation [recent star formation as traced by the \\halpha\\ emission line has been studied by \\citet{kennicutt08}, \\citet{lee09}, and \\citet{karachentsev10}; while the total \\halpha\\ luminosities differ slightly between these works, each finds a significant \\halpha-based ongoing star formation rate of $\\sim$0.06-0.07 \\msun\\,yr$^{-1}$]. Its relative proximity makes it well-suited for detailed studies of the ISM. Using the magnitudes of the brightest stars, \\citet{tikhonov92} and \\citet{karachentsev97} estimated distances of 3.2-3.3 Mpc. Subsequent observations with the {\\it Hubble Space Telescope} ({\\it HST}) provided a distance based on the magnitude of the tip of the red giant branch (TRGB; M$_{I}$ = $-$4.05$\\pm$0.02, with little dependence on metallicity; see {Rizzi \\etal\\ 2007}\\nocite{rizzi07} and references therein) of 3.15$\\pm$0.32 Mpc \\citep{karachentsev02}. Subsequent analyses by \\citet{jacobs09} and by the authors of the Extragalactic Distance Database ({Tully \\etal\\ 2009}\\nocite{tully09}; B. Jacobs, private communication) revise this slightly upward to 3.39$\\pm$0.25 Mpc. We adopt this distance measurement throughout the present work. At 3.39 Mpc, 1\\arcsec\\ corresponds to 16.4 pc. As Figure~\\ref{figcap1} shows, the stellar disk has an irregular morphology and harbors numerous high surface brightness \\HII\\ regions and widespread diffuse \\halpha\\ emission. Despite these characteristics, UGCA\\,105 has by comparison remained poorly studied in the literature. The low Galactic latitude of the system (13.7\\degr) and the significant foreground extinction values (E(B$-$V)=0.351 mag, or 1.51 mag of extinction in the B-band; see discussion in footnotes to Table~\\ref{t1}) may have conspired to keep this system out of many mainstream local galaxy surveys. As we show in this work, the stellar and gaseous components of UGCA\\,105 contain rich morphological and kinematic structure. This work presents the first detailed study of the neutral gas dynamics of this nearby dwarf galaxy. ", "conclusions": "\\label{S4} We have presented new low resolution {\\it JVLA} imaging of the nearby low mass galaxy UGCA\\,105. This system has remained comparatively under-studied in the astrophysical literature. The system is actively forming stars (as traced by high surface brightness \\halpha\\ emission; see {Kennicutt \\etal\\ 2008}\\nocite{kennicutt08}) and is located in sufficient proximity (3.39$\\pm$0.25 Mpc; {Jacobs \\etal\\ 2009}\\nocite{jacobs09}, {Tully \\etal\\ 2009}\\nocite{tully09}) to allow detailed dynamical analysis. In this work we present the first spatially resolved study of the neutral gas dynamics of UGCA\\,105. At 54\\arcsec\\ (890 pc) resolution, the neutral gas morphology and kinematics are measured with high fidelity with these new {\\it JVLA} data. The \\HI\\ gas spans $\\sim$175 \\kms; sampled over $\\sim$54 channels each separated by 3.3 \\kms, the system displays a classical ``butterfly diagram'' and double-horned integrated line profile. We recover 98\\% of the single dish flux \\citep{springob05}, with no correction for \\HI\\ self absorption applied. We derive a systemic velocity of 90.8\\,$\\pm$2.0 \\kms; UGCA\\,105 has an unusual velocity given its location well outside the Local Group. The integrated \\HI\\ column density distribution of UGCA\\,105 contains high surface density gas (N$_{\\rm HI}$ $>$ 10$^{21}$ cm$^{-2}$) throughout the extent of the luminous stellar disk. While we see evidence for regions of low column density (i.e., \\HI\\ holes or shells) and also evidence for spatial agreement between regions of ongoing star formation (as traced by \\halpha\\ emission) and high column density gas, we defer a detailed treatment until higher spatial resolution \\HI\\ imaging is available. We do note that the inclination corrected surface mass density profile of UGCA\\,105 falls in the regions closest to the dynamical center; this could be interpreted as marginal evidence for a molecular region in the inner disk. These new {\\it JVLA} data offer an opportunity to study the bulk neutral gas dynamics of UGCA\\,105. The system displays well-ordered rotation throughout the neutral gas disk and the intensity weighted isovelocity contours are parallel throughout the entire inner disk (cospatial with the high surface brightness stellar disk). We use standard GIPSY tilted ring analysis in order to fit the observed velocity field of UGCA\\,105. Regardless of how the parameters are fixed, we find a robust and well constrained rotation curve at all galactocentric radii. The profile rises steeply in the innermost $\\sim$1.5 kpc, rises more slowly in the region from $\\sim$1.5 kpc to $\\sim$5 kpc, and then remains flat at 72 \\kms\\ out to the last measured point (7.5 kpc). The total dynamical mass of UGCA\\,105, derived from this rotation curve, is M$_{\\rm dyn}$ $=$ (9$\\pm$2)\\,$\\times$\\,10$^{9}$ \\msun. This dynamical mass is larger than the sum of the luminous components [M$_{\\rm gas}$ $=$ (5.9$\\pm$0.7)\\,$\\times$\\,10$^{8}$ inclusive of a 35\\% correction for Helium and molecular material; M$_{\\star}$ $=$ (1.8\\,$\\pm$\\,0.8)\\,$\\times$\\,10$^{8}$ \\msun] by a factor of $\\sim$10; UGCA\\,105 is a typical, dark matter dominated dwarf galaxy (see, e.g., {Oh \\etal\\ 2008}\\nocite{oh08}, {Oh \\etal\\ 2011}\\nocite{oh11}, and references therein). The proximity and favorable inclination (55\\degr) of UGCA\\,105 make it a promising target for high resolution studies of both star formation and rotational dynamics in a nearby low-mass galaxy. In particular, B configuration {\\it JVLA} observations would achieve a synthesized physical resolution element of order 100 pc; the high \\HI\\ surface brightness would guarantee high signal to noise measurements at this spatial resolution. The resulting datasets would facilitate detailed analyses of the interplay of neutral gas and recent star formation on $\\sim$100 pc scales." }, "1207/1207.6217_arXiv.txt": { "abstract": "{\\em Secondary} contributions to the anisotropy of the Cosmic Microwave Background (CMB), such as the integrated Sachs-Wolfe (ISW) effect, the thermal Sunyaev-Zel'dovich effect (tSZ), and the effect of gravitational lensing, have distinctive non-Gaussian signatures, and full descriptions therefore require information beyond that contained in their power spectra. The Minkowski Functionals (MF) are well-known as tools for quantifying any departure from Gaussianity and are affected by noise and other sources of confusion in a different way from the usual methods based on higher-order moments or polyspectra, thus providing complementary tools for CMB analysis and cross-validation of results. In this paper we use the recently introduced skew-spectra associated with the MFs to probe the topology of CMB maps to probe the secondary non-Gaussianity as a function of beam-smoothing in order to separate various contributions. We devise estimators for these spectra in the presence of a realistic observational masks and present expressions for their covariance as a function of instrumental noise. Specific results are derived for the mixed ISW-lensing and tSZ-lensing bispectra as well as contamination due to point sources for noise levels that correspond to the Planck ($143$ GHz channel) and EPIC ($150$ GHz channel) experiments. The cumulative signal to noise ration $S/N$ for one-point generalized skewness-parameters can reach an order of ${\\cal O}(10)$ for Planck and two orders of magnitude higher for EPIC, i.e. ${\\cal O}(10^3)$. We also find that these three spectra skew-spectra are correlated, having correlation coefficients $r \\sim 0.5-1.0$; higher $l$ modes are more strongly correlated. Though the values of $S/N$ increase with decreasing noise, the triplets of skew-spectra that determine the MFs become more correlated; the $S/N$ ratios of lensing-induced skew-spectra are smaller compared to that of a frequency-cleaned tSZ map. ", "introduction": "All-sky multi-frequency Cosmic Microwave Background (CMB) missions, such as the completed WMAP\\footnote {http://map.gsfc.nasa.gov/}, ongoing Planck\\footnote {http://www.rssd.esa.int/index.php?project=Planck} and future (proposed) Experimental Probe of Inflationary Cosmology (EPIC) survey \\citep{Bock08,Bock09,Bau09} or ESAs Cosmic Origin Explorer (COrE, \\cite{Core12}) are major sources of information about the properties of the primordial density fluctuations that seeded the process of galaxy formation in the Universe as well as other key aspects of cosmological theory, including the global isotropy \\citep{Copi07,Hoft09,HanLew09} and topology of the Universe \\citep{Lum03,Rouk04}. The study of non-Gaussianity in the CMB fluctuations can provide valuable and detailed information regarding the physics of the early Universe of the inflationary epoch In the standard slow-roll paradigm, the scalar field responsible for inflation fluctuates with a minimal amount of self interaction which ensures that any non-Gaussianity generated during the inflation through self-interaction is expected to be small \\citep{Salopek90,Salopek91,Falk93,Gangui94,Acq03,Mal03}; see \\cite{Bartolo06} for a review. Variants of the simple inflationary model such as multiple scalar fields, features in the inflationary potential, non-adiabatic fluctuations, non-standard kinetic terms, warm inflation, or deviations from Bunch-Davies vacuum can however all lead to higher level of primordial non-Gaussianity \\citep{Chen10}. However, the detection of departure from Gaussianity in the CMB can be due to either primary or secondary effects (or both), as well as the mode-coupling effects of secondaries and gravitational lensing along the observer's light cone. Secondary anisotropies resulting from the formation of structure are known to dominate at smaller angular scales, are highly non-Gaussian in nature \\citep{Cooray01b,CoorayHu,VS02} and are arguably as interesting as their primary counterpart. One of the prominent contributions to the secondary non-Gaussianity is due to the mode-coupling of weak gravitational lensing and sources of secondary contributions such as the thermal Sunyaev-Zel'dovich effect \\citep{GoldbergSpergel99a,GoldbergSpergel99b, CoorayHu}. Although weak lensing of the CMB produces its own characteristic signature in the angular power spectrum, its detection has proved to be difficult using the CMB power spectrum alone. Non-Gaussianity imprinted by lensing into the primordial CMB remains below the detection level of current experiments, although with Planck the situation is likely to improve. Nevertheless, cross-correlating CMB data with external tracers means lensing signals can be probed at the level of the mixed bispectrum. After the first unsuccessful attempt to cross-correlate WMAP against SDSS, recent efforts by \\cite{SmZaDo00} have found a clear signal of weak lensing of the CMB, by cross-correlating WMAP against NVSS. Their work also underlines the link between three-point statistical estimators and the estimators for weak lensing effects on CMB. The understanding of secondaries are not only important in their own right, but also from the perspective of their impact on estimation of cosmological parameters \\citep{Smidt10}. The study of non-Gaussianity is usually primarily focused on the bispectrum, as this saturates the Cram\\'er-Rao bound \\citep{Babich,KSH11} and is therefore in a sense optimal, however in practice it is difficult to probe the entire configuration dependence in harmonic space contained within the bispectrum using noisy data \\citep{MuHe10}. The cumulant correlators are multi-point correlators collapsed to encode two-point statistics. These were introduced in the context of analyzing galaxy clustering by \\cite{Szapudi}, and were later found to be useful for analyzing projected surveys such as the APM galaxy survey\\citep{Mun}. Being two-point statistics they can be analyzed in multipole space by defining an associated power-spectrum. Recent studies by \\cite{Cooray3} and \\cite{Cooray8} have demonstrated their wider applicability including, e.g., in 21cm studies. In more recent studies the skew- and kurt-spectra were found to be useful for analysing temperature \\citep{MuHe10} as well as polarization maps \\citep{Mu11c} from CMB experiments and in weak lensing studies \\citep{Mu11b,Mu11d}. In addition to studies involving lower order multi-spectra, MFs have been extensively developed as a statistical tool for non-Gaussianity in a cosmological setting for both 2-dimensional (projected) and 3-dimensional (redshift) surveys. Analytic results are known certain properties of the MFs of a Gaussian random field making them suitable for identifying non-Gaussianity. Examples of such studies include CMB data \\citep{Schmalzing98,Novikov00,HikageM08,Natoli10}, weak lensing (\\cite{Matsubara01,Sato01,Taruya02,MuWaSmCo12}), large-scale structure \\citep{Gott86,Coles88,Gott89,Melott89,Gott90,Moore92,Gott92,Canavezes98,SSS98, Schmalzing00,Kerscher01,Hikage02,Park05,Hikage06,Hikage08}, 21cm \\citep{Gleser06}, frequency cleaned Sunyaev-Zel'dovich (SZ) maps \\citep{MuSmJoCo12} and N-body simulations \\citep{Schmalzing00,Kerscher01}. The MFs are spatially-defined topological statistics and, by definition, contain statistical information of all orders in the moments. This makes them complementary to the poly-spectra methods that are defined in Fourier space. It is also possible that the two approaches will be sensitive to different aspects of non-Gaussianity and systematic effects although in the weakly non-Gaussian limit it has been shown that the MFs reduce to a weighted probe of the bispectrum (\\cite{Hikage06}). The skew-spectrum is a weighted statistic that can be tuned to a particular form of non-Gaussianity, such as that which may arise either during inflation at an early stage or from structure formation at a later time. The skew-spectrum retains more information about the specific form of non-Gaussianity than the (one-point) skewness parameter alone. This allows not only the exploration of primary and secondary non-Gaussianity but also the residuals from galactic foreground and unresolved point sources. The skew-spectrum is directly related to the lowest order cumulant correlator and is also known as the two-to-one spectra in the literature \\citep{Cooray01}. In a series of recent publications the concept of skew-spectra was generalized to analyse the morphological properties of cosmological data sets or in particular the MFs by \\cite{MSC10,MuWaSmCo12,MuSmJoCo12,PratMun12}. The first of these three spectra, in the context of secondary-lensing correlation studies, was introduced by \\cite{MuVaCoHe11} and was subsequently used to analyse data release from WMAP by \\cite{Cala10}. The primary aim of this paper is to consider the entire set of generalised skew-spectra resulting from the mode-coupling of secondary anisotropies and lensing of the CMB and the contribution thereof to non-Gaussian morphology of the CMB maps. We will be considering three different secondary-lensing correlation bispectra. The secondaries that we consider are the Integrated Sachs-Wolfe effect (ISW) that dominates at large angular scales \\citep{Cooray02} and the thermal Sunyaev-Zel'dovich (tSZ) effect that dominates at smaller angular scales \\cite{Birk99}. In addition we consider a foreground, namely the contribution from unresolved point sources. We will consider two experimental setups, the ongoing Planck satellite and the the proposed EPIC satellite mission discussed above. The layout of the paper is as follows. In \\textsection\\ref{sec:ani} we briefly outline the bispectrum corresponding to lensing-secondary mode-coupling. Next, in \\textsection\\ref{sec:mf}, we review the formalism underlying the Minkowski Functionals and in \\textsection\\ref{sec:skew_spec} we introduce the generalised skew-spectra associated with the MFs. In \\textsection\\ref{sec:estim} we present the estimators for these spectra and their covariances. Finally, in \\textsection\\ref{sec:concl} we discuss our results and comment on future implementation. Throughout we will use the parameters of a WMAP cosmology \\citep{Lar11}. ", "conclusions": "\\label{sec:concl} Non-Gaussianity is in itself a poorly defined concept, in that there is no unique approach that can be adopted to describe or parametrize an arbirary form of non-Gaussianity in a complete manner. In order to quantify non-Gaussianity as fully as possible it is therefore essential to deploy a battery of complementary approaches each of which exploits different statistical characteristics. Each such technique will have a unique response to real world issues such as the sky-coverage (observational mask), beam and instrumental noise. Any robust detection therefore will have to involve a simultaneous cross-validation of results obtained from independent methods. The most common characterizations of non-Gaussianity involve studying the bispectrum, which represents the lowest-order departure from Gaussianity; higher order non-Gaussianity can be studied using its higher order analogues i.e. the multi-spectra. By contrast the topological estimators (MFs) that we have studied here carry information to all orders, though in a collapsed (one-point) form. Analytical results for MFs for a Gaussian field are well understood, and form the basis of non-Gaussianity studies \\citep{Tomita86}. There have been several previous studies on extraction of the MFs from the CMB data that rely either on simplification of radiative transfer using the Sachs-Wolfe limit \\citep{HKM06} or using a perturbative approach based on a series expansion of the MFs that can be studied order by order \\citep{Mat03,mat94}. The MFs have also been studied using elaborate computer-intensive non-Gaussian simulations \\citep{Komatsu03,Spergel07}. Most of these studies were done using a specific model of non-Gaussianity, namely the {\\em local} model of primordial non-Gaussianity which is parametrized by the well known parameter $f_{\\rm NL}$. The main motivation behind the present study has been to to extend such methods to secondary non-Gaussianities which have not been studied before in the context of morphological statistics analytically. The increase in sensitivity of CMB experiments and near all-sky coverage along with wide frequency range means the study of non-Gaussianities will be feasible in the very near future. Moreover, in the currently favoured adiabatic CDM models it is expected that the contribution from primordial non-Gaussianity is negligible and the main contribution to non-Gaussianities comes from secondaries. The secondary non-Gaussianity signal are associated with large scale structure contributions and through various mode coupling effects such as gravitational lensing \\citep{GoldbergSpergel99a,GoldbergSpergel99b, CoorayHu}. Our primary aim in this work has been to study how well we can probe the secondary signals from mode coupling using morphological descriptors. One of the main difficulties faced by one-point estimators $S^{(i)}(\\fw)$, which also affects the MF-based estimators $V_k(\\nu)$, is their inability to differentiate among various sources of non-Gaussianity. The triplets of skew spectra $S^{(i)}_l(\\fw)$ that we have introduced can be used to separate out contributions from various secondaries as well as to probe and constrain any foreground residuals left from the component separation step of the data analysis chain. Generalizing \\cite{MuHe10} we have introduce a set of triplets of skew-spectra which can be extracted from any realistic data. These skew-spectra do not compress the available information from a bispectrum to a single number, and their shape can help to distinguish among various sources of non-Gaussianity. Exploiting the perturbative expansion of the MFs, we showed that at the leading order of non-Gaussianity the MFs are completely specified by the knowledge of the bispectrum. Our results are most naturally defined in the harmonic domain. Comparison of MFs extracted using harmonic approach can be cross-compared with more traditional approach in the real space as an useful consistency check. The methods based on the skew-spectra that we have presented are simple to implement once the derivative fields $[\\nabla \\Theta\\cdot\\nabla\\Theta]$ or $[\\nabla^2\\Theta]$ are constructed. We have shown that this can be implemented in a model-independent way. Our method is based on a Pseudo-${\\cal C}_l$ approach \\citep{Hiv} and can handle arbitrary sky coverage and inhomogeneous noise distributions. The Pseudo-${\\cal C}_{\\ell}$ approach is well understood in the context of power spectrum studies, and its variance or scatter can be computed analytically. We provide generic analytical results for the computation of scatter around individual estimates. We also provide detailed predictions on how they are cross-correlated. In our method, it is possible indeed to go beyond the lowest level in non-Gaussianity to include the contribution from trispectrum. The main contributions in frequency-cleaned CMB maps will be from lensing of the CMB, though it is expected that such corrections will be sub dominant at least in the context of CMB data analysis. \\begin{table*} \\begin{center} \\begin{tabular}{|c |c|c| c| c} \\hline (Planck,EPIC) & SZ & ISW & PS &\\\\ \\hline $\\left ({S/N} \\right )$ & $(5.0,1137.4)$ & $(1.0, 216.0)$ & $(0.5,209.0)$ \\\\ \\hline $\\left ({S/N} \\right )$ & $(24.0,1354.9)$ & $(62.2, 420.3)$ & $(4.3,552.0)$\\\\ \\hline $\\left ({S/N} \\right )$ & $(19.7,1328.8)$ & $(31.8,246.5)$ & $(1.7, 421.0)$ \\\\ \\hline \\end{tabular} \\caption{The cumulative signal to noise (S/N) for Planck and EPIC surveys, are shown for the three one-point skew-spectra. Parameters used to compute the skew-spectra and the associated scatter for the two different experiments, ongoing Planck \\citep{Planck08} and EPIC \\citep{Bau09}.\\label{table:signal_noise1}} \\end{center} \\end{table*} \\begin{table*} \\begin{center} \\begin{tabular}{|c |c|c| c| c} \\hline (Planck,EPIC) & SZ & ISW & PS &\\\\ \\hline $\\left ({S/N} \\right )$ & $(3.8,503.2)$ & $(0.3,39.3)$ & $(0.4,53.4)$ \\\\ \\hline $\\left ({S/N} \\right )$ & $(5.8,1299.0)$ & $(6.4\\times 10^{-2},70.9)$ & $(0.6,529.1)$ \\\\ \\hline $\\left ({S/N} \\right )$ & $(1.2,625.8)$ & $(1.4\\times 10^{-2},21.1)$ & $(0.1,178.2)$ \\\\ \\hline \\end{tabular} \\caption{The cumulative signal to noise (S/N) for Planck and EPIC surveys for the one-point cumulants $S^{(i)}(\\oh)$, defined in Eq.(\\ref{eq:s2n_def}), are shown for each skew-spectra. Parameters used to compute the skew-spectra and the associated scatter for the two different experiments, ongoing Planck \\citep{Planck08} and EPIC \\citep{Bau09}. We have assumed $f_{\\rm sky}=1$ in our calculations. \\label{table:signal_noise2}} \\end{center} \\end{table*} \\begin{figure} \\begin{center} {\\epsfxsize=4.8 cm \\epsfysize=5.5 cm {\\epsfbox[37 439 294 714]{planck_cross.eps}}} \\end{center} \\caption{Same as the previous Figure but for Planck noise. The expression for scatter Eq.(\\ref{eq:var1})-Eq.(\\ref{eq:var3})and cross-correlation Eq.(\\ref{eq:cross1})-Eq.(\\ref{eq:cross3}) has two contributions. In each of these expressions there are terms which depend on the bispectrum and there are terms which can constructed from power spectrum alone. For Planck noise we found that the expressions for the scatter as well as the cross-correlation are entirely dominated by the terms which depend only on the power spectrum thus making the coefficient $r_{ij}$ independent of the type of underlying bispectrum.} \\label{fig:cross_planck} \\end{figure} \\begin{figure} \\begin{center} {\\epsfxsize=15 cm \\epsfysize=4.8 cm {\\epsfbox[37 531 589 714]{epic_cross.eps}}} \\end{center} \\caption{The information content of the skew-spectra are not completely independent. The level of cross-correlation among various estimator is encoded in the coefficient of cross-correlations $r_{ij}$ defined in Eq.(\\ref{eq:cross_def}). The cross-correlation coefficient $r_{01}$, $r_{02}$ and $r_{12}$ are plotted as a function of harmonics $l$. The noise level correspond to that of EPIC (see Table \\ref{tab:exp}). The correlation structure refelects the underlying spectra $\\beta_l$ as well as the level of noise. The cross-correlation is similar for SZ and PS.} \\label{fig:cross_epic} \\end{figure} We conclude by pointing out that the MFs do not probe the full bispectrum, but involve only weighted sums of modes and are thus equivalent to the three {\\it generalised} skewness parameters we have used. We have also defined three generalized skew-spectra associated with each of these skewness parameters. In this sense, the study of these skew-spectra can replace the study of MFs. The skew-spectra we have introduced can all be probed for arbitrary mask and noise. Ubiased estimators can also be constructed which can work in the presence of partial sky coverage and inhomogeneous noise. Their variance can also be computed {\\em analytically}; thereby avoiding the use of non-Gaussian Monte-Carlo simulations completely. Finally, the MFs can be constructed from the knowledge of generalized skew-spectra and can be compared with the results from real space analysis. The triplets of generalized skew-spectra can be used to separate individual components of NGs using their shape information. From our analytical results of cross-correlation, we find that in the absence of noise, e.g. experiments such as EPIC, the skew-spectra are highly correlated, more so for higher $l$ values. The correlation coefficients are typically in the range $r = 0.5-1$ for a Planck type experiment. The cumulative signal-to-noise ratio, in a Planck type experiment, for bispectrum corresponding to the ISW and SZ and lensing cross-correlation reaches ${\\cal O}(10)$. An improvement of about two orders of magnitude can be expected with experiments such as EPIC. Throughout we have ignored the presence of primordial non-Gaussianity which is expected to be subdominant. Nevertheless, it can be incorporated. Individual skew-spectra from different underlying bispectrum can essentially be combined to construct the total skew-spectra which means that our results can straightfowardly be generalised to incorporate specific models of primordial non-Gaussianity. The generic results derived here are also applicable to other areas in cosmology and have indeed been explored recently. Examples include the analysis of galactic redshift surveys \\citep{PratMun12}, weak lensing surveys \\citep{MuWaSmCo12} and the frequency cleaned SZ maps \\citep{MuSmJoCo12}. The results presented here can be extended beyond the analysis of temperature maps, e.g. to analyze polarisation maps, by extending the spin-$0$ calculations to spin-$2$. Such results can furnish useful probes for tje characterization of morphology of reionization in three dimensions. A few comments are in order about the comparison of our estimators with the so-called optimal estimators. The motivation to construct an optimal estimator is to improve the signal-to-noise of detection which is important in case of weak signals such as the primordial non-Gaussianity. The main motivation in this paper has been to reconstruct the topological properties of the CMB map going beyond Gaussianity, in the harmonic domain; in particular due to the contributions from secondary lensing cross-correlation which will be detected with high signal to noise with the proposed CMB surveys such as EPIC. In addition to the primary and secondary non-Gaussianity, cosmic defects such as textures or cosmic strings \\citep{ABR99,Cr07,RS10} also leave non-Gaussian footprints in CMB maps which can be detected by the change in topological nature of the maps. The estimators we have presented here may have relevance in such investigations. A detailed study will be presented elsewhere. At the level of the bispectrum the effect of lensing can only be studied through its cross-correlation with other secondaries. However weak lensing is also independently responsible for a next order correction to MFs through its effect on the trispectrum; the signal-to-noise is expected to be low. The signal-to-noise of the skew-spectra for secondary-lensing cross-correlation bispectrum is comparable to that of the skew-spectra of frequency-cleaned SZ maps \\citep{MuSmJoCo12}. However the secondary skew-spectra are much higher compared to skew-spectra associated with primary skew-spectra unless we assume a rather high value for the $f_{\\rm NL}$." }, "1207/1207.4024_arXiv.txt": { "abstract": "In this work, we consider a network of cosmic strings to explain possible deviation from $\\Lambda$CDM behaviour. We use different observational data to constrain the model and show that a small but non zero contribution from the string network is allowed by the observational data which can result in a reasonable departure from $\\Lambda$CDM evolution. But by calculating the Bayesian Evidence, we show that the present data still strongly favour the concordance $\\Lambda$CDM model irrespective of the choice of the prior. ", "introduction": "One of the greatest mysteries in cosmology today is the nature of dark energy in the Universe. According to the concordance cosmology, this makes up around $70\\%$ of the total energy density of the universe and causes the universe to accelerate around the present epoch \\cite{review}. Yet, there is no strong theoretical as well as observational indication to pin point the nature of this dark component. In fact we still do not know whether the mysterious late time acceleration is due to the presence of dark energy or due to any modification of gravitational laws on large scales. Nevertheless, majority of the observational data \\cite{obs, suzuki, komatsu, Eisenstein} support the concordance $\\Lambda$CDM model as a possible explanation for the late time acceleration. In this model, the presence of the cosmological constant ($\\Lambda$) poses two serious problems. First one is related to the fine tuning of the value of $\\Lambda$ to a ridiculously small number ($\\sim 10^{-120} M_{pl}^{4}$ in natural units) necessary for the present day cosmic acceleration. The other one is related to its dominance over the matter component of the universe precisely at the present epoch. Any small departure from these two conditions will result in a cosmological evolution that is different from the observable Universe. These conceptual problems dilute the superiority, the $\\Lambda$CDM model enjoys over other dark energy models as far as observational results are concerned. Interestingly, current observational data also allow deviation from the $\\Lambda$CDM behaviour. Dark energy models with $w_{de} = \\frac{p_{de}}{\\rho_{de}} \\neq -1$ are also allowed by current observations \\cite{obs, suzuki, komatsu, Eisenstein}. Although the allowed deviation from $w_{de} = -1$ is not large, but it is still detectable at the precision levels of current and future observational set ups. Hence efforts are on to construct models that can explain this deviation. Almost all the approaches for this purpose, assume that $\\Lambda$ is exactly zero in our universe, and an evolving dark energy is solely responsible for the late time acceleration of the universe. This includes possibilities likes quintessence \\cite{quint}, k-essence \\cite{kess}, an arbitrary barotropic fluid like Chaplygin gas or its various generalization and many more \\cite{gcg}. All these models are generally phenomenological without any strong theoretical motivations and they also burden us with the challenge of explaining several other model parameters. The other option is to assume the existence of a small but non-zero $\\Lambda$ as in $\\Lambda$CDM model but also to consider some extra component to explain the observed deviation from $\\Lambda$CDM behaviour \\cite{aas}. This may be motivated by the fact that string-landscape model may explain the existence of a small but non zero $\\Lambda$ in our present universe \\cite{SL}. But it is important to have well motivated candidate that can explain the deviation from $\\Lambda$CDM; otherwise this approach will have the same shortcomings which are present in standard dark energy models mentioned above. In this study, we propose one such model. There are models where dark energy is assumed to be a perfect fluid having equation of state $w = \\frac{p}{\\rho} <0$. But one major problem for such models is due to its sound velocity $c_{s}^2 = w < 0$ which causes instabilities on the small scales. It was then proposed to include rigidity in the fluid to make it an elastic solid \\cite{spergel}. In this case, if the rigidity is sufficiently large, $c_{s}^{2} > 0$ even if $w < 0$. ", "conclusions": "Cosmological observations have now confirmed the late time acceleration of the Universe. The challenge is to explain the source of this acceleration. Although a concordance $\\Lambda$CDM model is consistent with different observational results, a small deviation from this $\\Lambda$CDM behaviour is also allowed. Moreover recent results also show that there may be some inconsistencies between different observational results if one assumes $\\Lambda$CDM as the correct cosmological model \\cite{tension}. This motivates people to consider scenarios slightly different from an exact $\\Lambda$CDM behaviour. In most of these scenarios, one assumes the $\\Lambda$ term to be exactly zero and introduces dynamical scalar fields that can give rise to late time acceleration of the universe. All these models are phenomenological in nature and lacks theoretical understanding ( See \\cite{SS} for recent attempt to build a quintessence model in string theory). They not only inherit the problems present in the concordance $\\Lambda$CDM model, but also burden us with a number of extra parameters which also have to be fine tuned. We should also stress the fact that except the recently discovered Higgs field in Large Hadron Collider (LHC) \\cite{higgs}, no other scalar field has been detected so far that can explain this late time acceleration. This motivates us to look for alternative approaches that may explain the deviation from the $\\Lambda$CDM evolution. Formation of cosmic string network is possible during phase transitions in the early Universe. Although, on cosmological scales, we have not detected these strings so far, but they are more commonly observed in condensed matter systems in the laboratory. Keeping this is mind, we propose a cosmological model where a small contribution from cosmic strings network is present in the energy budget of the Universe together with the cosmological constant and other usual matter components. One can consider this network of strings as a perfect gas \\cite{KT} and express the overall equation of state $w_{s}$ as a function of its average velocity $v_{s}$. Numerical simulations suggest that this average velocity can not be exactly zero but is constrained to be $v_{s}^2 \\sim 0.17$. We show that such a model is completely consistent with current observational data. We put constraint on the model parameters like $\\Omega_{s0}$ and $\\Omega_{m0}$. Subsequently calculating Bayesian Evidence, we show that statistically data still prefer strongly a model with simple cosmological constant and this result does not change much with the choice of prior. To conclude, if the deviation from the $\\Lambda$CDM evolution is indeed confirmed, cosmological constant with a small contribution from cosmic string network may explain such deviation. At the theoretical level, we believe this set up has similar appeal as in scalar field models as both can arise in reasonable particle physics set up though neither has been detected so far. But given the fact that cosmic strings formation has been observed in condensed matter system, and numerical simulations on the dynamics of string networks on cosmological scales have revealed a lot about their properties, if detected, they may be a good alternative to scalar field dark energy models. \\vspace{5mm}" }, "1207/1207.0027_arXiv.txt": { "abstract": "Spatial extension is an important characteristic for correctly associating $\\gamma$-ray-emitting sources with their counterparts at other wavelengths and for obtaining an unbiased model of their spectra. We present a new method for quantifying the spatial extension of sources detected by the Large Area Telescope (LAT), the primary science instrument on the {\\em \\fermi Gamma-ray Space Telescope} (\\fermi). We perform a series of Monte Carlo simulations to validate this tool and calculate the LAT threshold for detecting the spatial extension of sources. We then test all sources in the second \\fermi-LAT catalog (2FGL) for extension. We report the detection of seven new spatially extended sources. ", "introduction": "A number of astrophysical source classes including supernova remnants (SNRs), pulsar wind nebulae (PWNe), molecular clouds, normal galaxies, and galaxy clusters are expected to be spatially resolvable by the Large Area Telescope (LAT), the primary instrument on the {\\em \\fermi Gamma-ray Space Telescope} (\\fermi). Additionally, dark matter satellites are also hypothesized to be spatially extended. See \\cite{atwood_fermi} for pre-launch predictions. The LAT has detected seven SNRs which are significantly extended at \\gev energies: W51C, W30, IC~443, W28, W44, RX\\,J1713.7$-$3946, and the Cygnus Loop \\citep{w51c,w30_lat,ic443,w28,w44,rx_j1713_lat,cygnus_loop_lat}. In addition, three extended PWN have been detected by the LAT: MSH\\,15$-$52, Vela~X, and HESS\\,J1825$-$137 \\citep{msh1552,velax,fermi_hess_j1825}. Two nearby galaxies, the Large and Small Magellanic Clouds, and the lobes of one radio galaxy, Centaurus A, were spatially resolved at \\gev energies \\citep{lmc,smc,cen_a_lat}. A number of additional sources detected at \\gev energies are positionally coincident with sources that exhibit large enough extension at other wavelengths to be spatially resolvable by the LAT at \\gev energies. In particular, there are 59 \\gev sources in the second Fermi Source Catalog (2FGL) that might be associated with extended SNRs \\citep[2FGL,][]{second_cat}. Previous analyses of extended LAT sources were performed as dedicated studies of individual sources so we expect that a systematic scan of all LAT-detected sources could uncover additional spatially extended sources. The current generation of air Cherenkov detectors have made it apparent that many sources can be spatially resolved at even higher energies. Most prominent was a survey of the Galactic plane using the High Energy Stereoscopic System (H.E.S.S) which reported 14 spatially extended sources with extensions varying from $\\sim0\\fdg1$ to $\\sim0\\fdg25$ \\citep{hess_plane_survey}. Within our Galaxy very few sources detected at \\tev energies, most notably the $\\gamma$-ray binaries LS\\,5039 \\citep{HESSLS5039}, LS I+61$-$303 \\citep{MAGICLSI, VERITASLSI}, HESS\\,J0632+057 \\citep{HESS0632}, and the Crab nebula \\citep{crab_weekes}, have no detectable extension. High-energy $\\gamma$-rays from \\tev sources are produced by the decay of $\\pi^0$s produced by hadronic interactions with interstellar matter and by relativistic electrons due to Inverse Compton (IC) scattering and bremsstrahlung radiation. It is plausible that the \\gev and \\tev emission from these sources originates from the same population of high-energy particles and so at least some of these sources should be detectable at \\gev energies. Studying these \\tev sources at \\gev energies would help to determine the emission mechanisms producing these high energy photons. The LAT is a pair conversion telescope that has been surveying the $\\gamma$-ray sky since 2008 August. The LAT has broad energy coverage (20 \\mev to $>300$ \\gev), wide field of view ($\\sim 2.4$ sr), and large effective area ($\\sim 8000\\ \\cm^2$ at $>1$ \\gev) Additional information about the performance of the LAT can be found in \\cite{atwood_LAT_mission}. Using 2 years of all-sky survey data, the LAT Collaboration published 2FGL \\citep[2FGL,][]{second_cat}. The possible counterparts of many of these sources can be spatially resolved when observed at other frequencies. But detecting the spatial extension of these sources at \\gev energies is difficult because the size of the point-spread function (PSF) of the LAT is comparable to the typical size of many of these sources. The capability to spatially resolve \\gev $\\gamma$-ray sources is important for several reasons. Finding a coherent source extension across different energy bands can help to associate a LAT source to an otherwise confused counterpart. Furthermore, $\\gamma$-ray emission from dark matter annihilation has been predicted to be detectable by the LAT. Some of the dark matter substructure in our Galaxy could be spatially resolvable by the LAT \\citep{pre_luanch_dark_matter_fermi}. Characterization of spatial extension could help to identify this substructure. Also, due to the strong energy dependence of the LAT PSF, the spatial and spectral characterization of a source cannot be decoupled. An inaccurate spatial model will bias the spectral model of the source and vice versa. Specifically, modeling a spatially extended source as point-like will systematically soften measured spectra. Furthermore, correctly modeling source extension is important for understanding an entire region of the sky. For example, an imperfect model of the spatially extended LMC introduced significant residuals in the surrounding region \\citep{first_cat,second_cat}. Such residuals can bias the significance and measured spectra of neighboring sources in the densely populated Galactic plane. For these reasons, in Section~\\ref{analysis_methods_section} we present a new systematic method for analyzing spatially extended LAT sources. In Section~\\ref{validate_ts}, we demonstrate that this method can be used to test the statistical significance of the extension of a LAT source and we assess the expected level of bias introduced by assuming an incorrect spatial model. In Section~\\ref{extension_sensitivity}, we calculate the LAT detection threshold to resolve the extension of a source. In Section~\\ref{dual_localization_method}, we study the ability of the LAT to distinguish between a single extended source and unresolved closely-spaced point-like sources In Section~\\ref{test_2lac_sources}, we further demonstrate that our detection method does not misidentify point-like sources as being extended by testing the extension of active Galactic nuclei (AGN) believed to be unresolvable. In Section~\\ref{validate_known}, we systematically reanalyze the twelve extended sources included in the 2FGL catalog and in Section~\\ref{systematic_errors_on_extension} we describe a way to estimate systematic errors on the measured extension of a source. In Section~\\ref{extended_source_search_method}, we describe a search for new spatially extended LAT sources. Finally, in Section~\\ref{new_ext_srcs_section} we present the detection of the extension of nine spatially extended sources that were reported in the 2FGL catalog but treated as point-like in the analysis. Two of these sources have been previously analyzed in dedicated publications. ", "conclusions": "Twelve extended sources were included in the 2FGL catalog and two additional extended sources were studied in dedicated publications. Using 2 years of LAT data and a new analysis method, we presented the detection of seven additional extended sources. We also reanalyzed the spatial extents of the twelve extended sources in the 2FGL catalog and the two additional sources. The 21 extended LAT sources are located primarily along the Galactic plane and their locations are shown in Figure~\\ref{allsky_extended_sources}. Most of the LAT-detected extended sources are expected to be of Galactic origin as the distances of extragalactic sources (with the exception of the local group Galaxies) are typically too large to be able to resolve them at $\\gamma$-ray energies. For the LAT extended sources also seen at \\tev energies, Figure~\\ref{gev_vs_tev_plot} shows that there is a good correlation between the sizes of the sources at \\gev and \\tev energies. Even so, the sizes of PWNe are expected to vary across the \\gev and \\tev energy range and the size of HESS\\,J1825$-$137 is significantly larger at \\gev than \\tev energies \\citep{fermi_hess_j1825}. It is interesting to compare the sizes of other PWN candidates at \\gev and \\tev energies, but definitively measuring a difference in size would require a more in-depth analysis of the LAT data using the same elliptical Gaussian spatial model. Figure~\\ref{gev_vs_tev_histogram} compares the sizes of the 21 extended LAT sources to the 42 extended H.E.S.S. sources.\\footnote{The \\tev extension of the 42 extended H.E.S.S. sources comes from the H.E.S.S. Source Catalog \\url{http://www.mpi-hd.mpg.de/hfm/HESS/pages/home/sources/}.} Because of the large field of view and all-sky coverage, the LAT can more easily measure larger sources. On the other hand, the better angular resolution of air Cherenkov detectors allows them to measure a population of extended sources below the resolution limit of the LAT (currently about $\\sim0\\fdg2$). \\fermi has a 5 year nominal mission lifetime with a goal of 10 years of operation. As Figure~\\ref{time_sensitivity} shows, the low background of the LAT at high energies allows its sensitivity to these smaller sources to improve by a factor greater than the square root of the relative exposures. With increasing exposure, the LAT will likely begin to detect and resolve some of these smaller \\tev sources. Figure~\\ref{compare_index_2FGL} compares the spectral indices of LAT detected extended sources and of all sources in the 2FGL catalog. This, and Tables~\\ref{known_extended_sources} and~\\ref{new_ext_srcs_table}, show that the LAT observes a population of hard extended sources at energies above 10 \\gev. Figure~\\ref{hess_seds} shows that the spectra of four of these sources (2FGL\\,J1615.0$-$5051, 2FGL\\,J1615.2$-$5138, 2FGL\\,J1632.4$-$4753c, and 2FGL\\,J1837.3$-$0700c) at \\gev energies connects to the spectra of their H.E.S.S. counterparts at \\tev energies. This is also true of Vela Jr., HESS\\,J1825$-$137 \\citep{fermi_hess_j1825}, and RX\\,J1713.7$-$3946 \\citep{rx_j1713_lat}. It is likely that the \\gev and \\tev emission from these sources originates from the same population of high-energy particles. Many of the \\tev-detected extended sources now seen at \\gev energies are currently unidentified and further multiwavelength follow-up observations will be necessary to understand these particle accelerators. Extending the spectra of these \\tev sources towards lower energies with LAT observations may help to determine the origin and nature of the high-energy emission. The \\textit{Fermi} LAT Collaboration acknowledges generous ongoing support from a number of agencies and institutes that have supported both the development and the operation of the LAT as well as scientific data analysis. These include the National Aeronautics and Space Administration and the Department of Energy in the United States, the Commissariat \\`a l'Energie Atomique and the Centre National de la Recherche Scientifique / Institut National de Physique Nucl\\'eaire et de Physique des Particules in France, the Agenzia Spaziale Italiana and the Istituto Nazionale di Fisica Nucleare in Italy, the Ministry of Education, Culture, Sports, Science and Technology (MEXT), High Energy Accelerator Research Organization (KEK) and Japan Aerospace Exploration Agency (JAXA) in Japan, and the K.~A.~Wallenberg Foundation, the Swedish Research Council and the Swedish National Space Board in Sweden. Additional support for science analysis during the operations phase is gratefully acknowledged from the Istituto Nazionale di Astrofisica in Italy and the Centre National d'\\'Etudes Spatiales in France. This research has made use of pywcsgrid2, an open-source plotting package for Python\\footnote{\\url{http://leejjoon.github.com/pywcsgrid2/}}. The authors acknowledge the use of HEALPix\\footnote{\\url{http://healpix.jpl.nasa.gov/}} \\citep{healpix}." }, "1207/1207.2214_arXiv.txt": { "abstract": "We develop empirical methods for modeling the galaxy population and populating cosmological N-body simulations with mock galaxies according to the observed properties of galaxies in survey data. We use these techniques to produce a new set of mock catalogs for the DEEP2 Galaxy Redshift Survey based on the output of the high-resolution Bolshoi simulation, as well as two other simulations with different cosmological parameters, all of which we release for public use. The mock-catalog creation technique uses subhalo abundance matching to assign galaxy luminosities to simulated dark-matter halos. It then adds color information to the resulting mock galaxies in a manner that depends on the local galaxy density, in order to reproduce the measured color-environment relation in the data. In the course of constructing the catalogs, we test various models for including scatter in the relation between halo mass and galaxy luminosity, within the abundance-matching framework. We find that there is no constant-scatter model that can simultaneously reproduce both the luminosity function and the autocorrelation function of DEEP2. This result has implications for galaxy-formation theory, and it restricts the range of contexts in which the mocks can be usefully applied. Nevertheless, careful comparisons show that our new mocks accurately reproduce a wide range of the other properties of the DEEP2 catalog, suggesting that they can be used to gain a detailed understanding of various selection effects in DEEP2. ", "introduction": "\\label{sec:intro} The accurate interpretation of galaxy survey data requires a careful consideration of a wide range of biases and incompleteness that can arise from the selection of the surveyed galaxies. Galaxy selection most commonly occurs as a function of galaxy flux, color, and (particularly for spectroscopic surveys) local density on the sky. Since these three properites are well known to be strongly correlated with one another \\citep[\\eg][]{Hogg04, Cooper06a}, a selection algorithm that chooses galaxies based on one of them will also produce selection biases in the distributions of the others. Estimating the scale of secondary selection effects like these requires a detailed understanding of the interplay between different galaxy properties; this becomes more important and more complex when considering surveys that cover a wide range in redshift, since the galaxy population may evolve substantially over the lookback time of the survey. The use of mock (\\ie, simulated) galaxy catalogs to understand and account for selection effects has therefore been an important part of the analysis in nearly all modern galaxy surveys, particularly those probing redshifts of order unity, such as DEEP2 \\citep{Davis03, DEEP2}, VVDS \\citep{VVDS}, and zCOSMOS \\citep{zCOSMOS}. In this study, we develop a set of empirically based techniques for constructing mock galaxy catalogs for high-redshift surveys. We also produce a new set of catalogs for the DEEP2 survey and make them available to the public, although our techniques should be generally applicable to any spectroscopic galaxy survey. One of our primary goals in constructing the mock catalogs presented here is to produce a simulated catalog appropriate for testing and optimizing algorithms to detect groups and clusters of galaxies in redshift space in DEEP2. We also wish to use these mocks to calibrate the halo-mass selection function that corresponds to the observational selection of galaxy groups (\\ie, systems with two or more observed members in DEEP2). At present, the detailed astrophysics of galaxy formation is not sufficiently well understood to directly simulate the galaxy population from first principles. However, there is strong support, through many different lines of evidence, for a basic picture of galaxy formation which holds that (1) all galaxies are embedded in larger dark matter halos that comprise most of their mass, (2) above some minimum mass threshold, all dark matter halos (whether primary halos or subhalos within a primary) host a single galaxy at their centers, and (3) there is a tight relation between the optical luminosity of a galaxy and the mass of its dark matter host. This paradigm suggests an obvious approach to constructing mock galaxy catalogs by assigning a galaxy properties to each halo or subhalo in a dark-matter-only N-body simulation, using some prescription for mapping between galaxy properties and (sub)halo properties. Techniques for converting periodic N-body simulation volumes into the cone-shaped geometries typical of surveys, and for including the effects of cosmic structure evolution at high redshift, are now well developed \\citep[\\eg,][]{YWC04, Blaizot05, KW07}. What remains is to specify a halo-galaxy connection that reproduces the observed galaxy population with high accuracy. One would ideally like to have a physically well-motivated method for assigning galaxies to dark-matter halos, so numerous authors have constructed mocks using semi-analytic models of galaxy formation applied to the underlying N-body simulations \\citep[\\eg,][]{Eke04, KW07, Henriques12}. At present, however, most such models fail to accurately reproduce the color distribution of galaxies at redshifts near unity. An additional difficulty arises because, until quite recently, N-body simulations of cosmologically interesting volumes have not had sufficient particle numbers to resolve the relatively low-mass halos and subhalos that host the faintest galaxies in modern surveys. \\citet{YWC04} (hereafter YWC) addressed these problems by taking an empirically based halo-model approach, using a conditional luminosity function to construct mock catalogs for the DEEP2 redshift survey, populating massive halos with one central galaxy and a number of satellites depending on the halo mass, in such a way as to match the measured DEEP2 luminosity and autocorrelation functions. Since the satellites had to be randomly assigned to dark matter particles in these halos, rather than to well-defined subhalos, this approach likely mis-estimated the spatial profiles and velocity distributions of galaxy groups and clusters. Given the simulations available at the time, however, even this approach could only account for the relatively bright galaxies probed by DEEP2 at $z\\ga 0.7$; the fainter galaxy population sampled at low $z$ was not included. The YWC mocks also did not include information on galaxy properties besides luminosity, although \\citet{Gerke07a} were able to add color information to these mocks according to the measured relation between color and local galaxy density. Their approach to color assignment was inspired by the ADDGALS mock-catalog creation algorithm that was used to produce mock catalogs for the Sloan Digital Sky Survey \\citep{SDSS} in \\eg, \\citet{Koester07a}, and which will be described in detail by Wechsler et al. (in preparation). Even with the addition of colors, these mocks were not sufficient to model the faint, low-redshift population probed by DEEP2 in one of its three observational fields. In addition, the YWC mocks were based on N-body simulations whose underlying cosmological parameters are at variance with the best fit to current data. For these reasons, we are motivated to improve upon the YWC efforts and construct new mock catalogs for the DEEP2 survey. Since the construction of the YWC mocks, improvements in both software and hardware have allowed for a substantial increase in the mass resolution of N-body simulations of a given volume. Most recently, the Bolshoi simulation \\citep{Bolshoi} simulated the growth of cosmic structure in a cubical volume 250 comoving h$^{-1}$ Mpc on a side, resolving halos and subhalos down to masses smaller than those of the Magellanic Clouds \\citep{Busha10b} and using a set of cosmological parameters that is in good agreement with current constraints from a wide variety of data. The combination of a large volume and excellent mass resolution permit the construction of DEEP2 mock catalogs over the full redshift and luminosity range probed by the survey, except for a handful of very faint low-$z$ dwarfs. Moreover, because Bolshoi resolves subhalos down to the required mass range for DEEP2, we will be able to assign galaxies to halos and subhalos individually, rather than taking a halo-model based approach, which should better reflect the phase-space distributions of galaxies in groups and clusters. Lacking a sufficiently accurate astrophysical model of galaxy formation, we carry on with a purely empirical approach to mock-catalog construction. A popular technique for assigning galaxy luminosities, known as subhalo abundance matching, has been demonstrated as a feasible means of reproducing both the galaxy luminosity function and the autocorrelation function at a variety of redshifts, under certain cosmological assumptions \\citep[e.g.][]{ Kravtsov04, Tasitsiomi04, ValeOstriker04, Conroy06a, Reddick12}; we adopt this approach here and explore its applicability in more detail. To add color information to the resulting mock catalogs, we expand on the environment-dependent approach of \\citet{Gerke07a}, making substantial improvements to accurately model galaxies near the DEEP2 flux limit, as well as various luminosity and color-dependent sources of incompleteness in the survey. We also repeat the mock-making procedure for two other high-resolution N-body simulations with different cosmological parameters from Bolshoi, to facilitate tests for cosmological dependence in the DEEP2 selection function. We proceed as follows. The next section introduces the DEEP2 survey and the N-body simulations. Section~\\ref{sec:algorithm} details our methods for assigning galaxy properties to halos and subhalos, while Section~\\ref{sec:observation} describes the techniques we use to replicate various DEEP2 selection effects in the resulting catalogs. Section ~\\ref{sec:compare} is likely the most useful portion of the paper for users of these mock catalogs, since it makes detailed comparisons between the mock catalogs and the DEEP2 dataset. In the Appendix, we describe the contents of the public DEEP2 mock catalogs and give a brief example of their use for computing the DEEP2 halo-mass selection function. Throughout this paper, unless otherwise specified, distances are quoted in comoving $h^{-1}$ Mpc and absolute magnitude values are given as $M-5\\log h$. ", "conclusions": "\\label{sec:conclusions} We have constructed three sets of mock catalogs for the DEEP2 Galaxy Redshift Survey, derived from three different $N$-body simulations with different background cosmologies. In doing this, we have striven to reproduce, as accurately as possible, the DEEP2 HOD, as well as all of the important galaxy properties that might have an impact on DEEP2 galaxy selection. Meeting these goals required us to develop new empirical techniques to assign colors to our mock galaxies based purely on the color-environment relation in the DEEP2 data, since physically motivated galaxy formation models tend not to reproduce this relation accurately at high redshift. With the release of this paper, we also release the mock catalogs for use by the general public. Instructions for downloading the mock catalogs and detailed information on their contents can be found in the Appendix, along with a specific example showing how these mocks might be useful for estimating DEEP2 selection effects as a function of halo mass. Our mock-making technique assigns mock galaxies to the halos and subhalos in the dark-matter-only simulations by directly imposing the empirical properties of DEEP2 on the simulations. We start by stacking simulation snapshots from different timesteps to construct lightcones with the geometry of DEEP2 fields, which accurately reproduces the redshift evolution of cosmic structure. We then use the subhalo abundance matching technique to assign galaxy luminosities to the halos and subhalos in each lightcone. This technique posits a tight, monotonic relation between the mass of a halo and the luminosity of the galaxy it hosts (with the caveat that the masses of subhalos should be measured before they are accreted into their parent halos). To assign galaxy colors in such a way as to reproduce their strong dependence on luminosity, redshift, and local environment, we bin the DEEP2 and mock galaxies in luminosity and redshift and further subdivide these bins into quintiles of local $n$th-nearest-neighbor overdensity. We then draw DEEP2 galaxies at random from these bins and assign them to the mock galaxies in the corresponding bin, taking care to account for incompleteness in the DEEP2 sample at faint magnitudes. Having produced mock galaxies with redshifts, luminosities, and colors, we can invert the $k$-correction technique of W06 to assign apparent $R$-band magnitudes to the mock galaxies. These allow us to apply the DEEP2 apparent magnitude limit to our mock sample. We can also apply a redshift-dependent selection that closely approximates the effect of the DEEP2 selection in color-color space. Given these mock DEEP2 targets, we can apply the DEEP2 spectroscopic target-selection algorithm to reproduce the spatial selection effects inherent to multiplexed slitmask spectroscopy. Finally, to account for observations that fail to yield a redshift, we exclude a fraction of the mock galaxies according to the color and magnitude-dependent redshift-failure rate of DEEP2 galaxies. This mockmaking technique accounts for all major observational effects that pertain to DEEP2 galaxies. The resulting mocks accurately reproduce a wide array of properties of the DEEP2 catalog, including the luminosity function; the color, magnitude, and redshift distributions; the color-environment relation and its evolution; and the inferred DEEP2 HOD for massive halos, provided that the background cosmology of the simulation is in agreement with current constraints. This accuracy gives us confidence that our DEEP2 mock catalogs will allow us to infer the selection of DEEP2 galaxies with reasonable accuracy as a function of dark-matter-halo mass, as well as the relations between halo mass and galaxy properties. In the process of constructing the mock catalogs, we obtained two results that are of general scientific interest beyond the area of mock-catalog creation. First, when assigning galaxy luminosities using abundance matching techniques, we tried various different assumptions about the scatter in the mass-luminosity relation and compared the resulting projected two-point correlation functions $w_p(r_p)$ to the DEEP2 measurement. We find that no model with fixed log-normal scatter (in luminosity at fixed mass) gives simulated $w_p(r_p)$ consistent with the clustering measured in DEEP2. This is broadly consistent with the results of \\citet{Wetzel10}, who require a large value of the scatter to achieve even approximate agreement with the DEEP2 clustering results. For reasons discussed in Section~\\ref{sec:scatter}, we do not believe that these issues will create problems when using these mocks to optimize group-finding algorithms, although they should be used with care for purposes that involve understanding the selection of low-mass halos. In addition, because we have constructed mock catalogs from three separate N-body simulations with different cosmological parameters, we can probe the cosmology dependence of our mock-making techniques. We find that the HODs in the different mock catalogs vary strongly with cosmology. This arises directly from the technique for luminosity assignment. Because abundance matching maps the observed galaxy luminosity function to the simulated halo mass function at fixed number density, if the mass function $dn/dM$ in a given simuation has a lower normalization than the true mass function in the universe, then the HOD resulting from the galaxy assignment must have a higher normalization and lower mass cutoff to compensate. For cosmologies with nearly identical mass functions (e.g., points lying along the same degenerate curve in $\\Omega_M$--$\\sigma_8$ space), however, the impact on the HOD will be relatively small. Thus, it is important when constructing mock catalogs from $N$-body models to use a simulation whose assumed cosmology is in agreement with all existing constraints, to ensure the best possible accuracy in the resulting HOD. For most applications, we therefore recommend the mock catalogs we have constructed from the Bolshoi simulation, since it is consistent with the best cosmological constraints available. The mocks constructed from the L120 and L160 simulations are likely to be generally useful only for testing the impact of varying assumptions about cosmology on any conclusions drawn from the mocks. Given that the Bolshoi mocks should accurately reproduce the DEEP2 halo selection function, an immediately interesting use to which we can put them is testing and calibrating an algorithm for detecting groups and clusters of galaxies. We have used them for this purpose in \\citet{Gerke12}." }, "1207/1207.0950_arXiv.txt": { "abstract": "We demonstrate that the space formed by the star-formation rate (SFR), gas-phase metallicity ($Z$), and stellar mass (\\Ms), can be reduced to a plane, as first proposed by Lara-L\\'opez et al. We study three different approaches to find the best representation of this 3D space, using a principal component analysis, a regression fit, and binning of the data. The PCA shows that this 3D space can be adequately represented in only 2 dimensions, i.e., a plane. We find that the plane that minimises the $\\chi^2$ for all variables, and hence provides the best representation of the data, corresponds to a regression fit to the stellar mass as a function of SFR and $Z$, \\Ms=$f$(Z,SFR). We find that the distribution resulting from the median values in bins for our data gives the highest $\\chi^2$. We also show that the empirical calibrations to the oxygen abundance used to derive the Fundamental Metallicity Relation (Nagao et al.) have important limitations, which contribute to the apparent inconsistencies. The main problem is that these empirical calibrations do not consider the ionization degree of the gas. Furthermore, the use of the $N2$ index to estimate oxygen abundances cannot be applied for \\abox~$\\gtrsim8.8$ because of the saturation of the [\\ion{N}{ii}]\\,$\\lambda$6584 line in the high-metallicity regime. Finally we provide an update of the Fundamental Plane derived by Lara-L\\'opez et al. ", "introduction": "Stellar mass (\\Ms), metallicity ($Z$), and star-formation rate (SFR) are key galaxy properties. \\Ms\\ reflects the amount of gas locked up in stars over a galaxy's history. SFR indicates the current rate at which gas is being converted into stars. $Z$ reflects the gas reprocessed by stars over the course of stellar evolution, and any exchange of gas between the galaxy and the environment. The relationships between these three properties are fundamental in understanding galaxy evolution. Indeed, models of galaxy formation within the $\\Lambda$-CDM scenario already include chemical hydrodynamic simulations \\citep[e.g.,][]{DeLucia04,Tissera05,DeRossi06,Dave07,Martinez10}. In the last few years it has been found that \\Ms, $Z$, and SFR are strongly interrelated. Analyzing galaxy measurements from the Sloan Digital Sky Survey (SDSS), \\citet{Ellison08} found that the mass-metallicity (\\Ms$-Z$) relation for star-forming (SF) galaxies depends on the SFR. Subsequently, \\citet{Lara10a} reported the existence of a Fundamental Plane (FP) between these three parameters. These authors confirmed that the \\Ms$-Z$ and \\Ms$-$SFR relations are just particular cases of a more general relationship. \\citet{Lara10a} fitted a plane and derived an expression for the stellar mass as a function of the gas metallicity and SFR (the Fundamental Plane, FP). In a parallel and independent study, using the same SDSS data, but different Z and SFR estimations, \\citet{Mannucci10} found a similar fundamental relationship, but instead expressed $Z$ as a combination of \\Ms\\ and SFR with a substantially different quantitative relationship. They refer to this correlation as the Fundamental Metallicity Relation (FMR). In a recent study, \\citet{Yates12} used models and SDSS data to analyze the dependences of different combinations between SFR, $Z$ and \\Ms. They found qualitative differences in the dependencies of those variables depending on the choice of approach in measuring the metallicity and SFR. A fundamental requirement in all these analyses is obtaining a reliable estimation of the galaxy metallicity. The most robust method to derive the metallicity in SF galaxies is via the estimate of metal abundances and abundance ratios, in particular through the determination of the gas-phase oxygen abundance. % This is typically achieved through the analysis of emission-line spectra of \\HII\\ regions within the galaxies. A proper determination of the oxygen abundance relies on the detection of the [\\ion{O}{iii}]\\,$\\lambda$4363 auroral line \\citep[the \\Te\\ method, e.g.,][]{LSE09}, but this emission line is usually not observed because of its faintness. Consequently, it is common to invoke the so-called strong-line methods. These techniques assume that the oxygen abundance of an \\HII\\ region can be derived using only a few bright emission lines. Empirical calibrations based on photoionization models, however, systematically over-predict by 0.2-0.6~dex the oxygen abundances derived using the \\Te\\ method and those calibrations which are based on it \\citep[see][]{Yin+07,KE08,Bresolin09,LSE10b,Moustakas+10, LSDK+12}. However, the absolute metallicity scale is still uncertain, temperature fluctuations and gradients can render the \\Te\\ method incorrect by up to 0.4 dex \\citep{Peimbert07}. Here we explore three different approaches to the representation of the three-dimensional distribution of \\Ms, SFR, and $Z$ for galaxies. We detail our sample selection in \\S\\,\\ref{SampleSelection} and review some issues with metallicity estimators in \\S\\,\\ref{Metissue}. The analysis is presented in \\S\\,\\ref{FPsection}, and we explore the implications for relationships between SFR and $Z$ in \\S\\,\\ref{Z-SFR}. We present a discussion of the outcome of our analysis, and summarise our results, in \\S\\,\\ref{conc}. ", "conclusions": "\\label{conc} The use of a reliable metallicity estimator is crucial when analyzing the SFR dependence of the \\Ms$-Z$ relation as shown in \\citet{Yates12}. For this reason, we recommend that the estimator of N06 be used with caution, and limited to the range (12+log(O/H)$<$8.8) where the saturation of the N2 parameter is not a problem. The emission-line galaxy spectra from SDSS are high quality, and measurements for many emission lines are available, making it possible to determine the gas-phase metallicity more robustly by applying techniques which consider the ionization degree of the gas. Examples of these methods are \\citet{McGaugh91,KD02,KK04} and \\citet{Tremonti04} (which are based on photoionization models) and \\citet{P01a,P01b,PT05} and \\citet{PVT10} (which rely on datasets for which $Z$ is known using the \\Te\\ method). We analyzed the 3D distribution of \\Ms, Z, and SFR using three different approaches: $(i)$ fitting a plane using PCA, $(ii$) fitting a plane through regression \\citep{Lara10a}, and $(iii)$ binning in SFR and \\Ms\\ to obtain the median metallicity of each bin \\citep{Mannucci10}. For the five methods used, we estimated the $\\chi_{red}^2$ as a measure of goodness of fit (Table \\ref{TableChi}). We find that the best representation of the data is the plane defined by regression on \\Ms. We compare the \\citet{Mannucci10} surface (the FMR) and the \\citet{Lara10a} Fundamental Plane (FP), and demonstrate that the best representation of the data corresponds to a plane. While PCA does not provide the best representation, it does demonstrate that the 3D distribution can be adequately represented in two dimensions. The \\Ms$-Z$, \\Ms$-$SFR, and $Z-$SFR relationships are then projections of this plane. We also highlight that the plane found by the regression fit on \\Ms\\ is not developed as a new technique for stellar mass estimation. Rather, this approach is primarily aimed at identifying the most concise representation of \\Ms, $Z$ and SFR in order to facilitate more detailed exploration of the interplay between these properties of galaxies. Nevertheless, in cases when more robust techniques to estimate the stellar mass \\citep[e.g.,][]{Taylor11} are not available, the use of the FP could be used to estimate the stellar mass, being careful to take into account the metallicity and SFR uncertainties. Our analysis of the $Z-$SFR relation with the two approaches toward binning the data highlights a crucial need for caution in interpretation when exploring distributions represented as median values. The inappropriate interpretation of such results will lead to apparently contradictory conclusions, depending on the binning order used. Furthermore, presenting medians in bins as a three-dimensional distribution will lead to differing representations of the data, depending on the binning order chosen. The SFR of a galaxy relates to the amount of gas currently being converted into stars, and correlates with the current mass in stars, while metallicity is a measure of the number of times that the gas has been reprocessed by stars, and also correlates with the current mass in stars in a galaxy. The fact that we can represent \\Ms\\ as a linear combination of SFR and metallicity suggests that the stellar mass of a galaxy can be thought as the rate at which a galaxy is currently forming stars (SFR), plus a measure of the star formation history, here represented by the metallicity ($Z$), corresponding to the amount of reprocessing of the gas by past stellar generations. The SF history and current SFR of a galaxy are closely linked to \\Ms. There is now an abundance of high quality spectroscopic measurements from large surveys such as the SDSS and the Galaxy And Mass Assembly (GAMA) survey \\citep[][]{Driver11}. These resources provide the means to calculate robust metallicity estimators for significant numbers of galaxies. Managing the measurements from this growing data volume brings its own challenges, as well as the opportunity to explore scaling relations and broad population properties for statistically robust samples that can be divided into well-defined subsets in many different ways. Approaching these challenges and opportunities in the most robust way possible, by using the most accurate measurements available, will ensure that the most reliable scientific understanding of galaxy evolution can be produced." }, "1207/1207.5572_arXiv.txt": { "abstract": "We calculate the low red-shift Taylor expansion for the luminosity distance for an observer at the center of a spherically symmetric matter inhomogeneity with a non vanishing cosmological constant. We then test the accuracy of the formulas comparing them to the numerical calculation for different cases for both the luminosity distance and the radial coordinate. The formulas can be used as a starting point to understand the general non linear effects of a local inhomogeneity in presence of a cosmological constant, without making any special assumption about the inhomogeneity profile. ", "introduction": "Modern cosmological observations such as the luminosity distance \\cite{Perlmutter:1999np,Riess:1998cb,Tonry:2003zg,Knop:2003iy,Barris:2003dq,Riess:2004nr} and the WMAP measurements \\cite{WMAP2003,Spergel:2006hy} of cosmic microwave background radiation (CMBR) have provided a strong evidence for the presence of dark energy. One of the main assumptions of the standard cosmological model used in fitting these observational data is spatial homogeneity of the Universe. We cannot nevertheless exclude the presence of a local inhomogeneity around us which could affect our interpretation of cosmological data \\cite{Romano:2010nc,Sinclair:2010sb,Romano:2011mx}. So far most of the efforts in estimating these effects have consisted in using some ansatz for the profile of the inhomogeneity and then calculate numerically the effects on cosmological observables. Such an approach has the limitation of depending on the particular functional form chosen to model the local inhomogeneity, and of relying completely on numerical calculations. In order to provide a more general study of this effects we approach the problem analytically and we derive a low-redshift formula for the luminosity distance relation for an observer at the center of a matter inhomogeneity in presence of a cosmological constant modeled by a LTB solution. The paper is organized as follows We first calculate the low red-shift expansion of the null radial geodesics for a central observer and then use it to obtain the luminosity distance. The calculation is based on using the analytical solution, and the geodesic equation expressed in the same coordinates of the analytical solution. The formula obtained is then compared to the numerical calculation of the luminosity distance to test its accuracy. In the appendix we give details of the derivation and the simplified formulae in the limit in which the inhomogeneity can be treated perturbatively. ", "conclusions": "We have derived the analytical low red-shift expansion of the luminosity distance for a central observer at the center of a spherically symmetric matter inhomogeneity in presence of a cosmological constant. We have first solved the null radial geodesic equation and calculated the local red-shift for $r(z)$ and $\\eta(z)$, and we have then used these to calculate the expansion of the luminosity distance. The formulae obtained take a simpler form in the case in which $K_0=0$, while in general are rather long and complicated, but can be reduced to a more tractable form in the limit in which the deviation form homogeneity can be treated perturbatively. The formulas we have derived can be used to understand the physical effects of local inhomogeneities in presence of a cosmological constant. It has the advantage, contrary to previous numerical studies, of not depending on any functional ansatz for the profile of the local inhomogeneity. This makes it particularly useful to study possible low red-shift inhomogeneities in a model independent way in the regime in which perturbation theory cannot be applied. \\appendix" }, "1207/1207.0482_arXiv.txt": { "abstract": "This paper corrects and completes a previous study of the shape of the extinction curve in the visible and the value of $R_V$. A continuous visible/infrared extinction law proportional to $1/\\lambda^p$ with $p$ close to 1 ($\\pm0.4)$ is indistinguishable from a perfectly linear law ($p=1$) in the visible within observational precision, but the shape of the curve in the infrared can be substantially modified. Values of {\\it p} slightly larger than 1 would account for the increase of extinction (compared to the p = 1 law) reported for $\\lambda > 1 \\mu$m and deeply affect the value of $R_V$. In the absence of gray extinction $R_V$ must be 4.04 if $p=1$. It becomes 3.14 for $p=1.25$, 3.00 for $p=1.30$, and 2.76 for $p=1.40$. Values of {\\it p} near 1.3 are also attributed to extinction by atmospheric aerosols, which indicates that both phenomena may be governed by similar particle size distributions. A power extinction law may harmonize visible and infrared data into a single, continuous, and universal, interstellar extinction law. ", "introduction": "\\begin{figure}[h] \\resizebox{1.\\columnwidth}{!}{\\includegraphics{fig1.eps}} \\caption{The atmospheric aerosol extinction law according to an occultation spectrum of Sirius (dots) observed by the GOMOS satellite (Fig.~1 in Paper~1), after correction for ozone absorption and Rayleigh scattering by nitrogen (N$_2$). The observations are matched closely by either a linear extinction law (solid line $\\propto e^{-0.7\\lambda^{-1}}$) or a quasi-linear one (red dashes $\\propto e^{-0.45\\lambda^{-1.3}}$). } \\label{fig1} \\end{figure} ``The Extinction Curve in the Visible and the Value of $R_V$,'' \\citep[][hereafter Paper~1]{rv} addressed several controversial issues concerning interstellar extinction in the visible and near-infrared ($\\lambda<1.2\\,\\mu\\rm m$), in particular its wavelength dependence, variability with line of sight, and the value of $R_V=A_V/${\\it E(B--V)}. To within the precision of the observations the interstellar extinction law in the visible region is well defined and depends solely on the amount of interstellar matter along the line of sight, as given by {\\it E(B--V)}. Normalized extinction curves in different directions are closely reproduced by a linear extinction law over the 0.43--1.2~$\\mu\\rm m$ wavelength interval. In the absence of gray extinction a linear extinction law implies $R_V\\sim4$, a value \\cite{tu12} found in Carina. As outlined in Paper~1 atmospheric aerosol extinction, much like extinction by interstellar grains, closely follows a linear extinction law (Fig.~1 in Paper~1) suggesting a close similarity in the particle size distribution for the two extinction processes. Power laws for extinction proportional to $1/\\lambda^p$ result from the specific size distribution of the particles responsible for the extinction without regard for the composition of the particles. Atmospheric aerosol extinction does not necessarily follow a simple linear law ($p=1$). According to a series of studies in the 1950s to 1970s a power law with $p$ of order 1.3, related to a power law size distribution of particles with exponent 4.3, would best represent extinction by aerosols \\citep{angstrom61, shaw73,allen73}. Fig.~\\ref{fig1} demonstrates that both laws, a linear extinction law with $p=1$ and a quasi-linear one with $p=1.3$, are difficult to distinguish over the visible wavelength range alone. They are more readily separated if the spectral domain of observation is extended to the infrared. There, a quasi-linear law with $p>1$ differs markedly from a linear law (\\S\\ref{pec}) and displays a similar degree of flattening to what is observed in interstellar extinction observations (\\S\\ref{ecir}). {\\it{p}}-values larger than 1 would also deeply modify $R_V$ and justify standard $R_V$-values close to $3 $ \\citep[see][]{tu76}. Power extinction laws ($\\propto 1/\\lambda^p$, $p \\neq 1$) are thus less restrictive and may provide an alternative to the linear law, maybe in even better agreement with observation. Our purpose is to investigate their properties and implications. The study is limited to visible and infrared wavelengths extending to the $L$ band (3.5~$\\mu\\rm m$), beyond which thermal emission from circumstellar dust may contaminate observed extinction curves. ", "conclusions": "In Paper~1 and here several issues are addressed of importance to a comprehensive understanding of interstellar extinction. What is the shape of the visible/near-infrared extinction curve? Does the extinction law vary with line of sight? And what is the value of $R_V$? Such questions arise from published results for the visible/near-infrared extinction curve that are incompatible with each other. It is not possible for the value of $R_V$ (and the extinction law) to be independent of the line of sight, as suggested by different data sets, yet to vary with direction viewed, as implied in other studies \\citep[][and Paper~1]{tu76,tu89,tu94}. Infrared and visible data suggest that the extinction in any direction, in the infrared/visible wavelength range, depends solely on the quantity of interstellar matter along the sight line, quantified by {\\it E(B--V)}. That is consistent with a common, direction-independent, extinction law governing both wavelength domains. A similar law holds for atmospheric aerosols. In the visible spectral region different power laws with exponents close to 1 describe equally well the observed normalized interstellar extinction curve and permit a large range of $R_V$ values. The recent study by \\citet{tu12} demonstrates that uncertainties in direct estimates (from photometry in the visible) of $R_V$ can be large. Provided that extinction in the visible and infrared results from the same continuous size distribution of interstellar grains, values of {\\it p} and $R_V$ are more easily established from infrared data. However, definitive conclusions about the exact wavelength dependence of the infrared interstellar extinction curve are difficult to reach. Extinction in the infrared is small and ground-based observations are limited to a few photometric bands. Observational uncertainties are consequently large and conspicuous differences exist between published infrared extinction curves. Additional studies of interstellar extinction, presumably from space and from spectroscopic observations, covering both the visible and the infrared, would be most beneficial." }, "1207/1207.3354_arXiv.txt": { "abstract": "We present results from the first fully general relativistic, magnetohydrodynamic (GRMHD) simulations of an equal-mass black hole binary (BHBH) in a magnetized, circumbinary accretion disk. We simulate both the pre and post-decoupling phases of a BHBH-disk system and both ``cooling'' and ``no-cooling'' gas flows. Prior to decoupling, the competition between the binary tidal torques and the effective viscous torques due to MHD turbulence depletes the disk interior to the binary orbit. However, it also induces a two-stream accretion flow and mildly relativistic polar outflows from the BHs. Following decoupling, but before gas fills the low-density ``hollow'' surrounding the remnant, the accretion rate is reduced, while there is a prompt electromagnetic (EM) luminosity enhancement following merger due to shock heating and accretion onto the spinning BH remnant. This investigation, though preliminary, previews more detailed GRMHD simulations we plan to perform in anticipation of future, simultaneous detections of gravitational and EM radiation from a merging BHBH-disk system. ", "introduction": " ", "conclusions": "" }, "1207/1207.6092_arXiv.txt": { "abstract": "{} {Diffuse interstellar bands (DIBs) measured in stellar spectra contain information on the amount of interstellar (IS) matter that is distributed along the line-of-sight, and similarly to other absorbing species may be used to locate IS clouds. Here we present a new database of 5780.5 and 6283.8 \\AA\\ DIB measurements. Those two DIBs have the advantage that they are strong and also broad enough to be detectable in cool-star spectra. We also study their correlation with the reddening.} {The database is based on high-resolution, high-quality spectra of early-type nearby stars located in the southern hemisphere at an average distance of 300 pc. Equivalent widths of the two DIBs were determined by means of a realistic continuum fitting and synthetic atmospheric transmissions. For all stars that possess a precise measurement of their color excess, we compare the DIBs and the extinction.} {We find average linear relationships of the DIBS and the color excess based on $\\simeq$ 120 and 130 objects that agree well with those of a previous survey of $\\simeq$ 130 northern hemisphere stars closer than 550pc. Because our target sky coverage is complementary, this similarity shows that there is no significant spatial dependence of the average relationship in the solar neighborhood within $\\simeq$600 pc. A noticeably different result is our higher degree of correlation of the two DIBs with the extinction, especially for the 5780\\AA\\ DIB. We demonstrate that it is simply due to the lower temperature and intrinsic luminosity of our targets. Using cooler target stars reduces the number of {\\it outliers}, especially for nearby stars, confirming that the radiation field of UV bright stars has a significant influence on the DIB strength. We illustrate the potential use of 3D maps of the ISM for characterizing the DIB sites. There is some evidence that interstellar cavity boundaries are DIB-deficient, although definite conclusions will have to wait for maps with a higher resolution. Finally, we have used the cleanest data to compute updated DIB shapes.\\thanks{available from the CDS, Strasbourg}} {} ", "introduction": "Identifying the diffuse interstellar bands (DIBs), the $\\simeq$ 400 weak absorption features seen in the spectra of reddened stars (\\cite{hobbs2008}; 2009) remains one of the longest-standing spectroscopic problems in astrophysics, and today they are still widely studied with the aim of conclusively identifying their carriers in the interstellar medium (ISM). If large molecules are preferentially designated (\\cite{legdhend85}), nothing has been decisively established as yet on their actual structures and sizes, on their participation from the gaseous or solid phase, on the role of the charge state balance, etc.. (for summaries see \\cite{herbig95}; \\cite{JD94}; \\cite{salama96}; \\cite{fulara00}; \\cite{snow11}; \\cite{fried11} (hereafter FR11) and references therein). Clues to their origin have been investigated in different ways. Laboratory spectra and models of rotational excitation of molecules have been compared to observations with the aim of finding convincing matches (see the recent review by \\cite{sarre2006}). Polarization studies aim at constraining the carrier phase, with most results pointing against the solid state (\\cite{smith77}; and recently \\cite{cox2006}; 2011). At the same time, the DIBs are studied toward specific targets or for large statistical samples of stars with the aim of establishing correlations with other interstellar tracers and identifying {\\it families} among the DIBS that are characterized by similar correlations. Sightline categories have been identified, in particular the so-called $\\sigma$- and $\\zeta$-type sightlines which have specific associations of DIB properties and UV extinction laws (\\cite{kre94}). $\\zeta$-type sightlines correspond to UV-shielded cloud cores while $\\sigma$-type sightlines probe cloud external regions that are partially ionized by the UV radiation field. The influence of the radiation field has been recently studied by Vos et al. (2011) based on stars from the upper Scorpius subgroup of the Sco OB2 association. These authors established significant differences between the properties of the two groups in terms of DIB-DIB, DIB-extinction, and DIB-gas relationships. These relationships have also been found based on sky-distributed datasets: links between the strengths of some DIBs and HI, H$_{2}$, KI, NaI or C$_{2}$ columns, between some of the DIBs and the dust column traced by the color excess, or between DIBS themselves have been found by several authors (e.g. \\cite{krelo87}; \\cite{thor03}; \\cite{welty06}; \\cite{ fried11}; \\cite{mccall2010}). The influence of the UV radiation field of the target stars has again found to be significant. In addition to the potential destruction of the carriers by energetic photons, the combination of the charge state and the carrier size is thought to play a major role in a number of DIBs (\\cite{gala04}; \\cite{salama96}). Instead of focusing on DIB properties and their differential behaviors that may give clues to their carriers, it is a more practical aspect that has motivated the present statistical work. Although correlations between the DIB strengths and the color excesses or the gas columns may be weak, DIBS are carrying some information on the quantity of matter along the path to a target star. In the same way that neutral sodium has been used to build local ISM maps (\\cite{welsh10}; \\cite{vergely10}), despite the absence of a strong correlation between neutral sodium and H columns, DIB absorption strengths toward stars at increasing distances can be used to determine IS cloud locations. Forthcoming extensive stellar datasets and hopefully parallaxes from the ESA Gaia mission will open the way to constructing precise 3D maps of the galactic dust by means of the inversion method (\\cite{vergely1}). These maps will in turn be useful for the interpretation of the stellar surveys, by helping to break the degeneracy between the star temperature and the reddening. Our aim is to build new derivation methods of the DIBS in all stellar spectra, to help constraining the 3D mapping, and in parallel to study the relationships between the DIB strengths and other IS parameters in more detail and to cover all various ISM cloud types and locations in the Galaxy. Here we complement previous statistical studies of two strong DIBS by an analysis of southern hemisphere high-quality, high-spectral resolution data, and in particular to compare the results with the recent results of an extensive northern hemisphere survey presented by FR11, and of a survey dedicated to the Upper Scorpius area (\\cite{Vos11}). While the FR11 survey aimed at detailed studies of DIB-to-DIB pairwise correlations, as part of a search for DIB carrier identification, our data were recorded with the aim of building a three-dimensional map of the nearby ISM. This is why the choice of the targets is drastically different (see section 4). Because we are using nearby targets with low reddening, we focused on two of the strongest bands in the red domain. Both have been studied by several authors, and in particular by FR11. They are those that probably are the most readily extractable from survey data such as the Gaia-ESO spectra, especially for cool stars, because the DIBS in this case must be broader than the stellar lines. In section 2 we describe the spectroscopic dataset used throughout the paper. In section 3, we explain the cleaning and fitting techniques that led to the extraction of the DIBs. We also describe a new method to estimate the error induced by fitting a continuum to both sides of the DIB. In section 4 we describe the correlative studies based on this new dataset, and the combination with the previous measurements of FR11. In section 5 we discuss our results and some potential improvements, in particular, attempts to identify the ISM environments that lead to substantial deviations from the average trends. ", "conclusions": "" }, "1207/1207.6601_arXiv.txt": { "abstract": "A fraction of multiple planet candidate systems discovered from transits by the Kepler mission contain pairs of planet candidates that are in orbital resonance or are spaced slightly too far apart to be in resonance. We focus here on the four planet system, KOI 730, that has planet periods satisfying the ratios 8:6:4:3. By numerically integrating four planets initially in this resonant configuration in proximity to an initially exterior cold planetesimal disk, we find that of the order of a Mars mass of planet-orbit-crossing planetesimals is sufficient to pull this system out of resonance. Approximately one Earth mass of planet-orbit-crossing planetesimals increases the interplanetary spacings sufficiently to resemble the multiple planet candidate Kepler systems that lie just outside of resonance. This suggests that the closely spaced multiple planet Kepler systems, host only low mass debris disks or their debris disks have been extremely stable. We find that the planetary inclinations increase as a function of the mass in planetesimals that have crossed the orbits of the planets. If systems are left at zero inclination and in resonant chains after depletion of the gas disk then we would expect a correlation between distance to resonance and mutual planetary inclinations. This may make it possible to differentiate between dynamical mechanisms that account for the fraction of multiple planet systems just outside of resonance. ", "introduction": "The latest tally of multiple planet candidate systems discovered by the Kepler mission \\citep{batalha12} includes 361 multiple-planet systems \\citep{fabrycky12}. A statistical analysis, focused on the probability that binary stars are the most likely contaminant, finds that most of the multiple planet candidates are real planetary systems \\citep{liss12}. The large number of recently discovered multiple planet systems represents a significant (by an order of magnitude) increase in the number of known multiple planet systems compared to those discovered from radial velocity surveys \\citep{wright11}. Both transit and radial velocity discovered multiple planet systems contain pairs of planets that are in first order mean motion resonance \\citep{wright09,liss11,wright11}. In the Kepler transit systems, there are statistically significant excesses of candidate planet pairs both in resonance and spaced slightly too far apart to be in resonance (as delineated by \\citealt{veras12}), particularly near the 2:1 mean motion resonance \\citep{liss11,fabrycky12}. Planet migration due to tidal interaction with a gas disk is a possible mechanism through which convergent migration induces resonance capture, leaving planets in resonance (e.g. \\citealt{lee02,ferraz03,kley04,lee04,papa05,thommes08,morbi07,libert11,rein12}). Gravitational interactions between planets leading to planet-planet scattering events (e.g., \\citealt{god09,fabrycky10,moore12,rein12}), turbulence in the disk \\citep{pierens11}, and tidal interactions between planets \\citep{papa11,lithwick12} have been proposed as mechanisms for pulling initially resonant planetary systems out of resonance. Interactions with planetesimals, for example as part of the `Nice' model for the early solar system evolution, can also cause planets to diverge away from resonance \\citep{nice} (also see \\citealt{thommes08}). Simulations of planets and planetary embryos embedded in a gas disk allow planets to become trapped in resonance \\citep{kley04,morbi07,rein12}. After the gas disk dissipates, newly formed planets may be left in a chain of mean-motion resonances \\citep{morbi07,matsumura10,moeckel12}. By a chain of mean motion resonances, we mean that each consecutive pair of planets is in a $j+1:j$, first order, mean motion resonance (though the integer $j$ can differ for each pair). Resonant chains have been chosen as initial conditions for studies of planetary system evolution after the depletion of the gas disk (e.g., \\citealt{nice,thommes08,batygin10}). When pairs of planets are in or near mean motion resonances, there may be a librating or nearly fixed Laplace angle involving three or more bodies. Three-body resonances (e.g., those comprised of zero-th order terms; \\citealt{quillen11}) may be present even when pairs of planets are not in first order mean motion resonances. A system in a resonant chain of first order resonances does not necessarily exhibit a librating Laplace angle involving three or more bodies. The four planet candidate system, KOI-730, contains planets with periods that satisfy the ratios 8:6:4:3 to approximately 1 part in 1000 or better \\citep{liss11}. These ratios imply that the outer two and inner two planet pairs are in (or near) a 4:3 mean motion resonance and that the middle pair of planets is in (or near) a 3:2 resonance. In this study we explore the evolution of a similar model four planet system in proximity to a planetesimal disk. We ask: how much mass in orbit-crossing planetesimals is sufficient to pull this system out of resonance? For a pair of planets and a first order mean motion resonance, \\citet{fabrycky12} define a parameter, $\\zeta_{1,1}$, that measures proximity to a first order mean motion resonance, \\begin{equation} \\zeta_{1,1} \\equiv 2 \\left( {1 \\over {\\mathcal P} - 1} - {\\rm Round}\\left({1 \\over {\\mathcal P } - 1} \\right) \\right), \\label{eqn:zeta} \\end{equation} where ${\\mathcal P} = P_i/P_j$ (greater than 1) is the observed period ratio of planets $i,j$ (and as defined in the appendix by \\citealt{fabrycky12}). The `Round' function rounds to the nearest integer. The parameter $\\zeta_{1,1} = 0$ at a $j+1:j$ first order resonance and the function goes from -1 to 1 at second order resonances ($j+2:j$). The parameter used by \\citet{liss11}, $\\zeta_1 = 1.5\\zeta_{1,1}$ and varies from -1.5 to 1.5. The probability density distribution of $\\zeta_{1}$ values generated from all pairs of Kepler planet transit candidates, for pairs residing in a single system, exhibits a peak at about $\\zeta_{1} \\approx - 0.2$, (see Figure 11 by \\citealt{liss11} and Figure 5 by \\citealt{fabrycky12}). For KOI-730, $\\zeta_1 = -0.0123, -0.0186, -0.0063$ for the inner pair, middle pair and outer pair of planets, respectively \\citep{fabrycky12}, consequently the system is likely to be in or very near a resonant chain. By integrating a system modeled after the KOI-730 system that is initially in a chain of resonances, we ask: how much mass in orbit-crossing planetesimals would increase the interplanetary spacing so that $\\zeta_1 \\sim -0.2$? We also note that due relatively low mass of the planets and their proximity to the star, their resonant widths are likely smaller than $\\zeta_1 = 0.1$, as first order mean motion resonance width scales approximately with mass to the 2/3 power (e.g. \\citet{wisdom80}). We first describe how we find resonant chain configurations for the KOI-730 system. We use these configurations to construct initial conditions for N-body integrations that include a planetesimal disk. A summary and discussion follows. \\begin{table} \\caption{\\large Masses and Periods for the planet candidates in the KOI 730 system and initial orbital elements \\label{tab:tab1}} \\begin{tabular}{@{}lccccc} \\hline Mass \t\t\t& Period\t& $a$\t\t& $e$\t\t& $\\omega$\t& $M$\t\t\\\\ ($M_\\oplus$) \t& (days)\t&\t\t\t&\t\t\t&\t\t\t&\t\t\t\\\\ 2.5\t\t\t\t& 7.3840\t& 1.0\t\t& 0.055589\t& -2.808830\t& 1.216544\t\\\\ 3.7\t\t\t\t& 9.8487\t& 1.211221\t& 0.071128\t& -0.739956\t& 2.841707\t\\\\ 8.6\t\t\t\t& 14.7884\t& 1.586171\t& 0.050110\t& 1.453201\t& -0.924252\t\\\\ 6.2\t\t\t\t& 19.7213\t& 1.920630\t& 0.043428\t& -1.990222\t& -0.172501\t\\\\ \\hline \\end{tabular} { Planet masses for the KOI 730 system are computed from radii based on transit durations taken from http://archive.stsci.edu/kepler/planet\\_candidates.html \\citep{batalha12}, using equation \\ref{eqn:mass}. The observed periods are given in days and are taken from the same website. The rightmost four columns give the orbital elements we used as initial conditions (see section 2). Here $\\omega$ is the argument of pericenter and $M$ the mean anomali. Initial inclinations and the longitudes of the ascending node were set to zero. These orbital elements were taken from the integration shown in Figure \\ref{fig:imran}. } \\end{table} \\begin{table} \\caption{\\large N-body Simulations \\label{tab:tab2}} \\begin{tabular}{@{}lcc} \\hline Simulation\t& Planetesimal Mass\t\t& Orbit-crossing mass\t\\\\ & ($M_\\oplus$)\t\t\t& ($M_\\oplus$)\t\t\t\\\\ Z\t\t\t& 0\t\t\t\t\t\t& 0\t\t\t\t\t\t\\\\ M\t\t\t& $10^{-4}$\t\t\t\t& 0.04\t\t\t\t\t\\\\ E3\t\t\t& $3.3 \\times 10^{-4}$ & 0.12\t\t\t\t\t\\\\ E\t\t\t& $10^{-3}$\t\t\t\t& 0.46\t\t\t\t\t\\\\ N5\t\t\t& $3.3 \\times 10^{-3}$\t& 1.7\t\t\t\t\t\\\\ N\t\t\t& $1.67 \\times10^{-2}$\t& 16.6\t\t\t\t\t\\\\ \\hline \\end{tabular} {\\\\ Each simulated disk contains 1024 equal mass planetesimals. The planetesimal masses in units of Earth mass are listed in the second column. The third column shows the total mass in planetesimals (in Earth masses) that crossed the orbits of any of the planets at the end of the simulation. The names of the simulations are related to the masses of the planetesimal disks. The Z simulation has a zero mass disk. The M, E and N simulations have disks with Mars, Earth and Neptune masses, respectively. The E3 and N5 simulations have disks with a third Earth and a fifth Neptune mass, respectively. } \\end{table} ", "conclusions": "Using an N-body integrator with the addition of drag causing convergent migration and eccentricity damping, we constructed a resonant chain for four planets with period ratios 8:6:4:3, and used planet masses estimated for the KOI-730 multiple planet system. Orbital elements from this integration were then used as initial conditions for N-body simulations with the four planets which included an external planetesimal disk. Interactions with the planetesimal disk allowed the planets to migrate, primarily diverging rather than converging, as expected. We find that one Earth mass of orbit-crossing planetesimals is sufficient to pull a system similar to the KOI-730 system out of its chain of mean motion resonances. As planet radii estimated from transit data depend on stellar radii, it is possible that we have over or underestimated the masses of the planets in the KOI-730 system. The distance migrated by a planet should scale with the mass in planetesimals that it interacts with. However, mean motion resonant widths are larger when planet masses are larger. We find that it is more difficult to form a resonant chain via resonant capture and such chains are less stable when the planets are more massive. Consequently, the amount of material required to pull a system out of resonance may not scale linearly with planet mass. Nevertheless, we expect that if the true planet masses are larger than adopted here, a larger mass in orbit-crossing planetesimals would be required to pull this system out of resonance and vice-versa if the planets are less massive. After the system is out of resonance, we find that it remains stable. Our simulations do not exhibit planet/planet orbit crossing events. The KOI-730 system, comprised of sub-Neptunian mass planets, can be compared to the HR8799 system, comprised of hyper-Jovian mass planets in resonance. When the HR8799 system is pulled out of resonance, the system is extremely unstable \\citep{god09,moore12}. The KOI-730 system is likely in (or very close to) a resonant chain. Because interactions with planetesimals and tidal interactions between planets primarily cause planetary orbits to slowly diverge \\citep{papa11,lithwick12} and so pull systems out of resonance, we can be fairly confident that processes following depletion of a gas disk did not put this system in resonance. Because the system is currently near or in resonance, only a small mass in planetesimals could have ever crossed the orbits of these planets. We infer that this system either lacks a debris disk or contains one that is so diffuse or stable that less than an Earth mass of debris has ever crossed the orbits of these planets. We also find that an Earth mass of orbit-crossing planetesimals can cause the planets to migrate far enough that the system lies sufficiently outside of resonance to resemble the Kepler systems with $\\zeta_1 \\sim -0.2$ where the peak of the distribution lies \\citep{fabrycky12,liss11}. However we expect that different planetary systems would have different quantities of orbit-crossing planetesimals. Consequently we would not expect that a narrow peak in the $\\zeta$ distribution would arise in a distribution of planetary systems. Fine tuning in the quantity of planetesimals may be required to account for the peak seen in the $\\zeta$ distribution. The tidal damping scenario for pulling pairs of planets away from resonance \\citep{lithwick12,papa11} may more naturally account for a peak in $\\zeta$. We find that the more an initially flat planetary system interacts with planetesimals, the higher the mean planet inclinations. If somewhere between a few earth masses to about a Neptune mass of planetesimals cross the orbits of the planets, the planet inclinations can increase to a few degrees. This is sufficiently high that they would not be all simultaneously be detected in transit. The tidal circularization scenario would not give a relation between migration distance and inclination. However over short distances, as planets migrate and they cross vertical resonances, we would expect a correlation between migration distance (and so distance out of resonance) and inclination. Study of the relation between the inclination and period distributions of the Kepler systems may differentiate between roles of tidal forces and planetesimals. It is possible that the transiting multiple planet systems discovered by the Kepler mission are more compact or lower mass than radial velocity discovered planetary systems. Can we differentiate between the tidal interaction mechanism for pulling systems out of resonance and that caused by interactions with planetesimals? If planetesimals interact with planets, then it is likely that planet inclinations can increase as planets cross vertical resonances. It may be possible to differentiate between these two mechanisms based on inclination distributions as tidal interactions likely do not increase inclinations but planet/planetesimal interactions can. Note that \\citet{libert11} have shown that resonant capture for higher planet mass systems can also induce inclination variations. However, additional mechanisms, such as turbulence associated with a gas disk or secular perturbations from distant planets could also affect the inclination distributions. Future studies can probe the role of planetesimals, and migration associated with scattering them, in accounting for the inclination distributions of the Kepler planetary systems. We have focused here on interactions with a low eccentricity planetesimal disk. When it encounters a planet, a high eccentricity planetesimal is less strongly gravitationally focused than a low eccentricity one. Consequently, high eccentricity objects are less effective at scattering a planet or inducing migration. Our limit on the total mass in orbital crossing planetesimals can be considered a lower limit as we began with a low eccentricity disk. Future studies can explore the possibility that compact Kepler systems harbor massive outer planetary systems and high eccentricity cometary populations. In summary, we believe that closely spaced, low inclination multiple planet Kepler systems likely have either low mass or extremely stable debris disks. There appears to be a relation between the inclination and amount of migration for a planet. The inclination distributions may make it possible to differentiate between dynamical scenarios for pulling planets out of resonance. Due to the improbability of a Late Heavy Bombardment like scenario for KOI-730, we believe these inclinations are probably caused by crossing vertical resonances. \\vskip 0.3 truein Acknowledgments. This work was in part supported by NSF through award AST-0907841. We thank Hal Levison for helpful discussion." }, "1207/1207.0626_arXiv.txt": { "abstract": "{A mechanism producing the transition from an Euclidean to a Minkowski manifold is described. A global Robertson-Walker symmetry is assumed from the large scale data of the visible universe. Allowing for the strain of the manifold as an additional field in the Lagrangian, we interpret the symmetry as a consequence of a global texture defect. The additional term gives rise to a boundary dividing the manifold into an Euclidean plus a Lorentzian region. It is also shown that the presence in the early epoch of homogeneous matter/energy fields preserves the horizon and the signature change across it. The horizon has properties much similar to the ones of the Big Bang of the Standard Model, including the need for a phase transition of the scalar field producing particles and fields as we know them now.} ", "introduction": "The commonly assumed background of special relativity is a Minkowski space-time, i.e. a flat four-dimensional manifold equipped with a Lorentzian signature. When generalizing the theory of relativity (GR) by the introduction of curvature in the actual space-time, the Minkowski manifold is typical of all tangent spaces to the curved manifold. Minkowski space-time as such does in a sense not exist, but is the asymptotic form of any real space-time when all kinds of matter/energy are taken out. This simple view, however, hides a puzzling feature. A Minkowski manifold is not the most general undifferentiated flat four-dimensional manifold, because of the light cones, i.e. of the Lorentzian signature. The structure of the light cones picks out a bunch of directions stemming from any given event in the manifold (the time-like worldlines) which cannot be confused with the rest. Where does this symmetry diminution come from? Indeed the most general four-dimensional manifold should be Euclidean: perfect isotropy and homogeneity. This is the question for which I shall try to find an answer in the present work. The problem of the signature of space-time has already been considered in the past from different viewpoints and with different motivations. A role of Euclidean signature geometries in four dimensions has been introduced, for instance, in quantum cosmology where Euclidean path integrals over geometries have been generalized from flat quantum field theory to gravity and the universe \\cite{hartle} \\cite{hawking}. In a quantum approach the transition from Euclidean to Lorentzian signature is achieved by analytic continuation in the complex domain or formally just by means of a Wick's rotation (introduction of imaginary time $t\\rightarrow it$). The existence of a Euclidean patch with the geometry of a four-dimensional half sphere was intended, in Hawking's approach, to remove the singularity in the remote past of the universe. How a real classical interpretation of the results emerges out of the quantum methods is however a non-trivial matter. The approach I will adopt in this paper is entirely classical \\emph{ab initio}. Even from a classical viewpoint the possibility of having different signatures in different regions of space-time has been discussed in the literature, especially in the 90's of the past century \\cite{ellis1}\\cite{ellis2}. The discussion was soon concentrated on the general constraints posed by a possible signature flip somewhere in the manifold and on different viable approaches \\cite{embacher}\\cite{hellaby}. The continuity conditions at the border between the Euclidean and the Lorentzian domains were analyzed together with the conservation laws on the boundary \\cite{dray}\\cite{sumeruk}. Though establishing the compatibility conditions for the co-existence of Euclidean and Lorentzian domains, the actual presence of Euclidean regions is then left to further theories, without being considered as a logical implication of the symmetry breaking manifested by the presence of the light cones. The purpose of the present paper is precisely to provide a consistent classical framework wherein the matching between Euclidean and Lorentzian signature is obtained \"naturally\". The question is: can we envisage a mechanism by which, pouring matter/energy (or something else) in a Euclidean manifold, the Lorentzian signature appears? In fact a way to induce a peculiar symmetry (to reduce the total symmetry) in a continuum is to introduce a defect in the sense used when describing material continua\\cite{volterra}% \\cite{punti}. This is what I meant by \"something else\" in the line above. However it is not clear how a defect can convert a ($++++$) into a ($+---$) signature. In order to find the way out I shall follow a path outlined in \\cite{cqg} and verified by various cosmological tests \\cite{test}; for it I shall use the name Strained State Cosmology (SSC) or Strained State Theory (SST). It consists in assuming that space-time behaves as a continuous deformable medium in four dimensions, extending the usual general relativistic approach by the introduction in the Lagrangian density of empty space-time of an \"elastic potential\" contribution built from the strain tensor; the final Lagrangian looks very similar to the ones used in the so called \"massive gravity\" but does indeed not coincide with them. The strain tensor of the manifold is intended as being half the difference between the actual metric tensor and the metric tensor of a reference undeformed manifold. According to the considerations I made above the reference manifold should be an Euclidean one. Mentioning two metric tensors gives the impression that I am presenting a bimetric theory. This however is not the case. Only one of the manifolds is actually existing: the one corresponding to our space-time (the natural manifold). The reference manifold is not present anywhere; it is not even a background. It simply is part of a logical description in which the universe is thought to behave as a deformed continuum. For this reason the mentioned metric tensor of the Euclidean manifold is no metric at all in the natural manifold. The piece of evidence from which we start is that we know (or at least we think we know) that the universe, at scales of hundreds of Mpc or higher, is described by a Friedman-Lema\\^{\\i}tre-Robertson-Walker model, i.e. it has a Robertson-Walker (RW) symmetry (isotropic expansion of a homogeneous space). Since the presence of matter \\textit{per se} does not motivate the RW symmetry I shall attribute the symmetry fixing to the presence of a global defect, somewhere in the manifold. In practice the initial assumptions will be: \\begin{itemize} \\item an action integral of the empty space-time like the following:% \\begin{equation} \\int \\left[ R+\\frac{1}{2}\\lambda \\varepsilon ^{2}+\\mu \\varepsilon _{\\alpha \\beta }\\varepsilon ^{\\alpha \\beta }\\right] \\sqrt{-g}d^{4}x \\label{azione} \\end{equation}% where $\\lambda $ and $\\mu $ are the Lam\\'{e} coefficients of space-time (fixed parameters) and $\\varepsilon _{\\mu \\nu }=\\frac{1}{2}\\left( g_{\\mu \\nu }-E_{\\mu \\nu }\\right) $ is the strain tensor determined with respect to the symmetric tensor $E_{\\mu \\nu }$ that would correspond to the Euclidean metric on the reference manifold; it is also $\\varepsilon =\\varepsilon _{\\alpha }^{\\alpha }$ where all the raising and lowering of indices is performed by means of the unique metric tensor $g_{\\mu\\nu}$; no assumption is made about $\\varepsilon _{\\mu \\nu }$ being small or not with respect to $E_{\\mu \\nu }$; $R$ plays the role of dynamical term for the strain tensor components; \\item a defect inducing a global RW symmetry. \\end{itemize} Under these conditions we shall see that a Euclidean signature in a domain of the natural manifold matches a Lorentzian one in correspondence with an horizon in the manifold. We shall also see that the presence of matter/energy does not spoil the effect I have just mentioned, provided the additional ingredient, under the horizon, is in the form of a completely homogeneous field. Of course such a field should then undergo a phase transition giving rise to the ingredients of the universe we observe now. The SST whose essence has been outlined above is indeed a metric theory on a Riemannian manifold which admits everywhere, excluding singularities, a flat tangent space-time. In the Lorentzian domain the tangent space is Minkowskian; in practice this tells us that the theory preserves the principle of equivalence and recovers locally the special relativity. ", "conclusions": "As we have seen, it is possible to think of a four-dimensional manifold where both the Euclidean and the Lorentzian signature are present in different regions. The boundary between the two regions is smoothly crossed if the global manifold has a Robertson-Walker symmetry and the given symmetry can be there as a consequence of a texture defect. The simplest representation of the configuration of the manifold is obtained when the $\\tau $ variable is chosen as originating from the defect and increasing along world-lines perpendicular to the equal strain hypersurfaces of the manifold. The shape of the latter hypersurfaces depends on the geometry of the defect; if the space in our universe is flat, as apparently it is, so has to be the singular submanifold corresponding to the defect. The boundary between the Euclidean and the Lorentzian areas is the hypersurface corresponding to a null scale factor, $a=0$; this boundary is called a \"horizon\" from which the Lorentzian signature emerges, in analogy with the Schwarzschild horizon for a cylindrical stationary space-time (spherical in space). Describing the situation I have now and then used terms as \\emph{evolution}, \\emph{transition}, even \\emph{time}. Of course this way of depicting the manifold is, strictly speaking, inappropriate in the Euclidean domain. There, in fact, there is no time; all dimensions are space-like; nothing \\emph{propagates}. In the Euclidian domain the $\\tau $ variable is an affine parameter along incomplete geodesics bounded by the cosmic defect on one side. The geodesics, besides being incomplete, are chosen so that they are everywhere perpendicular to the equal strain hypersurfaces of the manifold. It is only on the Lorentzian side of the boundary between the two domains that $\\tau $ acquires the familiar role of \\emph{cosmic time}. By the way in a consistent fully geometric view any description in terms of \\emph{evolution} is somehow inappropriate: there is just one fourdimensional manifold described in terms of Gaussian coordinates, curved and locally warped according to a global symmetry and the local distribution of matter/energy. The manifold has two regions of different signature, matching each other in a reasonably smooth way on a three-dimensional boundary. The term \\emph{evolution} is what we use for the way we label $3+1$ foliations of the manifold in the Lorentzian domain. Even though we would not use \\emph{evolution} for it, a similar labeling for an analogous foliation is possible also on the Euclidean side: it would not be constrained by the light cones, but would be suggested by the symmetry. Of course this is so in an entirely classical approach, but so far the puzzle of the role of time and of some background in the attempts to quantize gravity remains unsolved. It is important to remark that, even though the action integral (\\ref{azione}% ) looks very much similar to the action for the so called \"massive gravity\", it is however different. In fact the original approach due to Firtz and Pauli \\cite{FP} viewed (\\ref{azione}) as a first order approximation of a perturbative treatment, whereas in our case (\\ref{azione}) is \"exact\". Further developments of the massive gravity theory, introduced in order to cure various inconveniences present in the original version, do consider also non-linear approaches where the perturbative treatment is extended, in principle, to all orders, however in practice they are bimetric theories, where to the dynamical metric tensor a background non-dynamical metric is added, and the latter is used for building many scalars of the theory \\cite% {hinter}. In the SSC instead, there exists just one metric and $E_{\\mu \\nu }$ is not a background and is not used as a metric tensor at all; raising and lowering of indices, then the construction of all scalars, are performed by means of the unique $g_{\\mu\\nu}$, just as in classical GR. Furthermore the problems whose presence is still a matter of debate in the massive gravity theories are absent from SSC, at least as far as the cosmic scale is concerned. An immediate example is the absence, in the cosmological solution, of the so called vDVZ (van Dam-Veltman-Zakharov \\cite{vDV}\\cite{zakh}) discontinuity: when letting $\\lambda$ and $\\mu$ go to zero ($B$ go to zero) in (\\ref{hub2}) and (\\ref{equaz}) the shear GR cosmological solutions are obtained. The features described so far pertain to the only manifold, without calling in matter fields of any sort: in this case on the Lorentzian side we have an empty space-time whose space expands for ever according to formula (\\ref{sola}). We have then considered the presence, in the Euclidean era, of matter/energy in the form of a fully homogeneous radiation-like field (in four dimensions). Provided the density of the field is small enough, the conclusion concerning the presence of a horizon does not substantially change, as we see in formula (\\ref{hk}). However if we want to recover the present universe we must allow for a phase transition occurring in the primordial field in correspondence of the horizon, so that from it not only the Lorentzian signature appears but also the abundance of fields and particles we experience today. It is evident that a strong analogy exists between our horizon and the Big Bang of the Standard Model. In our case, under the horizon (\"before\" the Big Bang) we find an Euclidean era. Afterwards the present approach based on the physical role of the strain of the manifold has been successfully tested on a number of standard cosmological tests \\cite% {test}." }, "1207/1207.2309_arXiv.txt": { "abstract": "Since the Voyager fly-bys of Uranus and Neptune, improved gravity field data have been derived from long-term observations of the planets' satellite motions, and modified shape and solid-body rotation periods were suggested. A faster rotation period ($-40$min) for Uranus and a slower rotation period ($+1$h20) of Neptune compared to the Voyager data were found to minimize the dynamical heights and wind speeds. We apply the improved gravity data, the modified shape and rotation data, and the physical LM-R equation of state to compute adiabatic three-layer structure models, where rocks are confined to the core, and homogeneous thermal evolution models of Uranus and Neptune. We present the full range of structure models for both the Voyager and the modified shape and rotation data. In contrast to previous studies based solely on the Voyager data or on empirical EOS, we find that Uranus and Neptune may differ to an observationally significant level in their atmospheric heavy element mass fraction $Z_1$ and nondimensional moment of inertia, \\nI. For Uranus, we find $Z_1\\leq 8$\\% and $\\nI=0.2224(1)$, while for Neptune $Z_1\\leq 65$\\% and $\\nI=0.2555(2)$ when applying the modified shape and rotation data, while for the unmodified data we compute $Z_1\\leq 17$\\% and $\\nI=0.230(1)$ for Uranus and $Z_1\\leq 54$\\% and $\\nI=0.2410(8)$ for Neptune. In each of these cases, solar metallicity models $(Z_1=0.015)$ are still possible. The cooling times obtained for each planet are similar to recent calculations with the Voyager rotation periods: Neptune's luminosity can be explained by assuming an adiabatic interior while Uranus cools far too slowly. More accurate determinations of these planets' gravity fields, shapes, rotation periods, atmospheric heavy element abundances, and intrinsic luminosities are essential for improving our understanding of the internal structure and evolution of icy planets. ", "introduction": "The outer planets Uranus and Neptune are mysterious in many ways. While their names \\emph{'ice giants'} suggest a composition of predominantly volatiles in ice phases such as water, methane, and ammonia ice, interior models instead predict a warm interior devoid of solid ices \\citep{Chau+11,Redmer+11}. Structure models of Uranus and Neptune are generally in agreement in predicting a small rock core, a deep interior of more than 70\\% heavy elements, and a significantly less enriched outer envelope with a transition at about 70\\% of the radius for both planets \\citep{Hubbard+95,FN10,Helled+11}. However, they also agree in failing to explain Uranus' measured intrinsic luminosity \\citep{Podolak+91}. Uranus' low luminosity is a riddle; even more so as the corresponding models for Neptune nowadays can reproduce its measured luminosity \\citep{Fortney+11} indicating that the ice giants are not so similar to each other as previously thought \\citep{Podolak+95}. Contrary to the efforts that have been made to measure the atmospheric composition, the Voyager~2 radio occultation data and ground-based observational microwave data have actually raised more questions than they were intended to solve. In both planetary atmospheres for instance, D/H appears enriched over the protosolar values \\citep{Bergh86,Bergh90} suggesting particle transport between the atmosphere and an ice-rich deep interior \\citep{Hubbard+95}, while the inferred helium abundance is consistent with the protosolar value and thus would suggest the opposite (i.e., inefficient particle transport) if hydrogen and helium occur in the deep interior where they might undergo phase separation \\citep{Hubbard+95,Podolak+95}. Moreover, for the interior of Uranus we cannot rule out a primary composition of silicates dissolved into hydrogen--helium envelopes as that might sufficiently accelerate Uranus' cooling time \\citep{HMacF80}. We just do not know how similar Uranus and Neptune really are in terms of composition and structure \\citep{Podolak+00}. In addition, Uranus and Neptune have complex multi-polar magnetic fields, and appear to have stronger atmospheric winds than Jupiter and Saturn, if the Voyager rotation periods are applied. Besides the natural desire to understand the giant planets in our solar system, learning about the 'ice giants' is important for the classification of extrasolar planets with similar masses and sizes, as many ones are observed \\citep{Borucki+11}, although not yet on similar orbital distances \\citep{Kane11}. Information on Uranus and Neptune interiors are typically derived from theoretical models which are designed to fit the observed physical data of the planets, such as their gravitational fields, masses, internal rotation, and radii. The physical data available for Uranus and Neptune are rather limited. In particular, the low-order gravitational harmonics $J_2$ and $J_4$ of Neptune have significant error bars, Neptune's equatorial 1-bar radius is actually not measured, and the shapes (flattening) of both Uranus and Neptune are not well known. Voyager 2 provided only one occultation radius at the 1 bar pressure level for each planet \\citep[see][for details]{Helled+10}. Although stellar occultations do provide information on the planetary oblateness, the shape is inferred for upper atmosphere (microbar pressure levels) and it is unclear whether this shape is consistent with the shape at the 1 bar pressure level. \\citet{Helled+10} have shown that minimization of wind velocities or dynamic heights of the 1 bar isosurfaces \\citep{AS07,Helled+09} constrained by Voyager 2 occultation radii and gravitational coefficients of the planets, leads to \\emph{modified solid-body rotation periods} of 16h 34m for Uranus and 17h 27m for Neptune, which is 40~min shorter than the Voyager rotation period for Uranus, and 1h20 longer for Neptune. \\citet{Helled+10} also state that both planets may have different rotation periods than the ones derived by the minimization method, and in addition, could be rotating differentially on cylinders. Non-solid body rotation can lead to a change in the calculated gravitational moments \\citep{Hubbard+91}. Our results are valid as long as differential rotation in Uranus and Neptune is \"shallow\", i.e., the region of differential rotation consists of a negligible mass, which is a preferred solution for Uranus and Neptune (Y. Kaspi, priv.~comm.). However, the suggested solid-body rotation periods better match the measured radii of the planets and result in more moderate atmosphere dynamics with wind velocities of $\\sim$ 150 m s$^{-1}$ for both planets. Based on these suggested rotation rates and occultation radii, \\emph{modified shapes} of the planets were derived. Since the shapes of Uranus and Neptune are not well constrained, and in addition, their rotation profiles must not be that of a solid-body, the modified solid-body rotation periods and shapes can be used to demonstrate the sensitivity of the interior models to the uncertainties in these physical properties. To obtain pressure-density relations for the interior of Uranus and Neptune, three different methods have been invoked. \\citet{Podolak+95} apply physical equations of state of hydrogen, helium, the ices H$_2$O, H$_2$S, CH$_4$, and NH$_3$, and rocks, assuming sharp transitions between a gaseous outer envelope, an inner icy envelope, and a rock core. As this assumption by itself restricts the possible internal density distributions, \\citet{Marley+95} and \\citet{Podolak+00} created random density distributions which, when reproducing the observed gravity field, could then be evaluated according to pressure-density relations of likely materials. Another possibility is to use an analytic function with sufficiently many free coefficients to adjust to the given constraints \\citep{Helled+10}. The resulting pressure-density relation can be considered an \\emph{empirical} EOS. The advantage of the two latter methods is that they can allow for, and one can constrain the locations of, continuous (rather than sharp) density gradients. On the other hand, the resulting random or fitted density distributions can be unrealistic in a sense of not representing any composition of real materials. We here use the physical EOS LM-REOS for H, He, and H$_2$O \\citep{N+08} for the envelope material and the rock EOS by \\citet{HM89} for the core. This combination was also applied to the Uranus and Neptune models in \\citet{FN10} (hereafter FN10). In this paper, we use the suggested modified periods and equatorial radii and the physical LM-R EOS to compute new three-layer interior and evolution models of Uranus and Neptune. In \\sect\\ref{sec:m} we describe our method of model construction and the observational data used. In \\sect\\ref{sec:r} we compare the resulting models obtained with the modified data to those with the Voyager data. The results are discussed and summarized in \\sect\\ref{sec:dc}. ", "conclusions": "\\label{sec:dc} \\subsection{Bulk composition} We had to make several simplifications to be able to calculate Uranus and Neptune models with state-of-the art methods. One such simplification is the representation of heavy elements in the envelopes by water and the confinement of rocks to the core. In real Uranus and Neptune, silicates may also occur in the envelopes. It is clear that our assumption of having no rocks in the envelopes leads to an overestimation of the envelope metallicities and the resulting ice to rock ratio (I:R). In addition, the smaller the core mass, the larger the I:R ratio. Models with I:R $\\gg 2.7$, the solar system value, may potentially invalidate this simplification and point to the presence of silicates in the envelopes. Example models are the Neptune models with pure water envelopes (N2b) and all our Uranus models. Interestingly, some Neptune models, e.g.~N1 and N2a, have a rather large core ($\\sim 3\\ME$) with a reasonable overall I:R ratio of 1.5 times the solar value. On the other hand, our representation of ices by a pure water EOS likely overstimates the density of the true mixture of ices, and thus somewhat underestimates the mass fraction of ices that would be composed of a mixture of H$_2$O, CH$_4$, and NH$_3$. Unfortunately, high-quality EOS of light ices at pressures higher than 2~Mbar for planetary modeling are not yet available. Under the assumptions and simplifications of this work, the bulk mass of heavy elements is $12.5\\:\\ME$ for Uranus and 14--$14.5\\:\\ME$ for the selected Neptune models. This is in good agreement with the empirical EOS based models by \\citet{Helled+11}, who can explain the polynomial density distributions of their Uranus (Neptune) models by $\\sim 11$--$13\\:\\ME$ ($\\sim 13$--$15\\:\\ME$) of heavy elements. \\subsection{Implications from a stable deep interior} If some part of the interior, for Uranus possibly 0.45--$0.5\\:\\MU$ (\\sect\\ref{ssec:evol}), is stable to convection so that heat cannot escape efficiently from the region below, then the super-adiabaticity of the temperature gradient there can be non-negligible \\citep{LC12}. A warmer deep interior will require a lower particle number density in order to conserve the pressure gradient. Otherwise, the induced higher warm-dense-matter pressure would cause a larger planet radius. For a given composition in the inner envelope (e.g. the $Z_2$ value of the adiabatic case), one might think of a lower particle number density to imply a lower envelope mean density (compared to the adiabatic case), resulting into a larger rock core mass to ensure mass conservation. However, the mean density in the inner envelope is roughly constrained by the measured $J_2$ value, see \\fig\\ref{fg:profilesR}. Therefore, the metallicity in the deep interior cannot have the same $Z_2$ value as in the adiabatic case. Indeed, if the stable region is caused by a compositional gradient \\citep{Hubbard+95}, the deep interior would have an average metallicity $Z_3>Z_2$, which may lead to a smaller rock core. Therefore, predictions on a change of the rock core mass are not possible without more sophisticated models that take into account both a compositional gradient and super-adiabaticity. At the current level of our models, the net effect of a stable interior would be the introduction of a third, deep envelope with $Z_3>Z_2$. In case of a significant super-adiabaticity in Uranus, this $Z_3$ might raise to 100\\%. Moreover, a H/He-free deep interior might then be possible even with a material density larger than that of water, i.e.~with an ice-rock mixture. Thus, a strongly superadiabatic deep interior (below about $0.5\\:\\MU$) may qualitatively allow for I:R ratios closer to the solar and at the same time for a H/He-free deep interior, resembling a traditional protoplanetary core of a few $\\ME$ that can accrete its gaseous envelope within a few million years \\citep{HoriIkoma10}. A stable deep (i.e., far below the outer/inner envelope boundary) interior offers an explanation for the measured atmospheric helium and deuterium abundances. Given that D/H can be enriched in the cold ices of the protosolar nebula, the observed enhanced atmospheric D:H ratio could result from upward-transport of deuterium from an ice-rich interior \\citep{Gautier+95} that is located between the inner/outer envelope boundary and the stable deep interior, where the magnetic field is believed to be generated, see \\sect\\ref{ssec:dc_lumi}. Assuming that hydrogen and helium \\emph{exist} in the deep interior at Mbar pressures where hydrogen is metallic, \\citet{Hubbard+95} suggest that helium could phase separate and rain down leading to a helium-depleted atmosphere. However, if the particle exchange with the upper regions is suppressed in the deep interior, the rained out helium will not be replenished from the reservoir above so that the atmospheric He:H ratio remains protosolar as observed. However, we caution against taking the observed D and He abundances for clear indications of a stable deep interior, since helium phase separation from hydrogen must not necessarily occur in a metallic environment. \\citet{Lorenzen+09,Lorenzen+11} calculated the H-He immiscibility regions in dependence on temperature, pressure, and helium concentration. In metallic hydrogen, demixing requires a sufficiently high helium concentration to be energetically preferred. For instance, at 2 Mbar our Uranus and Neptune models have a temperature of 3800--4500 K. The phase diagram of \\citet{Lorenzen+11} predicts demixing under these conditions only if the He concentration is above $\\sim 1\\%$. Using \\begin{equation} \\frac{N_{\\rm He}}{N} = \\left(1+\\frac{m_{\\rm He}}{m_{\\rm H}}\\,\\frac{X}{Y} + 3\\frac{m_{\\rm He}}{m_{\\rm H_2O}}\\frac{Z}{Y}\\right)^{-1}\\:, \\end{equation} where the $m_{i}$ denote the molar masses of H, He, and water, we calculate helium particle concentrations $N_{\\rm He}/N$ of 1.7--3\\% for $Z=0.9$--0.95. This is only slightly above the demixing condition in pure H-He mixtures. As the demixing behavior of helium in more complex mixtures such as H-He-H$_2$O is unknown, we cannot rule out miscibility of He the deep interior of Uranus and Neptune, be it stable (Uranus case) or unstable (Neptune case) as an explanation for the observed abundances. We encourage future work on the miscibility behavior of planetary mixtures, such as started by, e.g., \\citet{Chau+11,WM12a,WM12b}. \\subsection{Luminosity}\\label{ssec:dc_lumi} The modified shape and rotation data do not affect the cooling behavior. With Uranus, this is to be expected, since former calculations with various different structure models as allowed by the large error bars in $J_2$ and $J_4$ produced an uncertainty of about 1 Gyr in the evolution only, smaller than the uncertainty induced by the observational error bar of $\\Teff$ \\citep{Fortney+11}. With Neptune, this is of some surprise, since the new models are $\\sim 1000\\:$K colder in the deep interior than former models used for cooling curve calculations (FN10). However, the cooling curve depends on changes of temperature with time. These are similar right after the rapid cooling at young ages. As the heat stored in the interior after formation depends on the specific heat of the bulk material, where $c_v$ is smaller for rocks than for ices, future work should include the admixture of rocks into the envelopes and investigate the maximum possible shortening of Uranus' cooling time within the approach of homogeneous, adiabatic evolution. Magnetic field models of Uranus and Neptune require that 60--70\\% of the region interior to the ionic water layer, corresponding to 0.42--0.56 $R_p$ \\citep{Redmer+11}, is stable to convection \\citep{StanBlox06}. Using our interior models, this corresponds to 0.25--0.5 of the planet's mass. Indeed we can reproduce Uranus' measured luminosity if we assume that the heat flux from within 0.45--0.5$\\:M_p$ is a negligible fraction of the heat flux in case of a convective, adiabatic interior. This consistency was already noted by \\citet{Podolak+91}. However, this mass or radius level is not supported by any of our Uranus structure models. Even if a density gradient occurs continuously rather than as a sharp transition, our structure models for Uranus predict a location farther out at $0.9\\MU$. Also for Neptune, a stable interior up to 0.25--0.5\\:$M_p$ is not supported by our structure models, where the deepest strong density gradient can occur at 0.57$\\MN$. Neptune's luminosity is best explained by the absence of a stable region. Future work should aim at finding consistent solutions for the structure, thermal evolution, and and magnetic field generation of Uranus and Neptune. At the current stage, the envelope separation (density gradient) of our structure models does not seem to be directly related to the luminosity, nor to the magnetic field generation. \\subsection{Why a dichotomy?} Beside the low luminosity, Uranus differs from Neptune in having a high obliquity, dense narrow rings and its five largest satellites on regular orbits, while Neptune has a more Saturn-like obliquity and extended dusk disk with diffuse rings, and two major satellites (Triton and Nereide) on irregular orbits. These observed properties point to different formation histories and as a consequence, also to different internal structures. As discussed by \\citet{Stevenson86}, the stochastic process of impacts of various sizes and obliquities may have caused the differences we see today. A composition gradient in a heavy-element-rich giant planet could be a remnant of the formation process \\citep{Podolak+91} that has been disturbed in Neptune by a last big direct impact but not so in Uranus (an oblique impact). We are not the first authors to suggest a dichotomy between the internal structures of Uranus and Neptune. But our structure models are the first ones to confirm this pre-Voyager hypothesis of their different structures also on the basis of computed moment of inertia values, with an absolute difference of $d\\nI\\sim 0.03\\:(\\sim 0.01)$ for the models with the (un-)modified shape and rotation data. \\subsection{A Vote for improved observational data} New measurements of Uranus and Neptune's physical parameters could improve our knowledge of their bulk compositions and internal structures considerably. For instance, the outer envelope metallicity is rather sensitive to the assumed solid-body rotation period. While the modified periods used in this paper are perhaps not the correct ones for Uranus and Neptune, since differential rotation is possible and the winds may not be in a state of minimum energy, the presented models clearly indicate the need for more data. Recently, \\citet{Karkoschka11} analyzed a collection of atmospheric circulation data compiled from Voyager~2 and HST observations of Neptune's atmosphere over a time-span of 20 years and found a rotation period close to the Voyager value, based on the stability of the motion of high-altitude clouds. Better determination of the rotation periods, shapes, Uranus' luminosity, Neptune's gravity field, and of the envelope metallicities below the water cloud level are indispensable for a better understanding of the interior, evolution, and formation of the icy planets in the outer solar system. We encourage the application of various discovery methods in the future, and in particular, to fly space missions to Uranus and Neptune. \\subsection{A dichotomy in the atmosphere compositions?} One of the indications for a dichotomy in the interior structure is the possible difference in the outer envelope metallicities $Z_1$, which are at most 0.08 ($\\leq0.18$) for Uranus but up to 0.65 ($\\leq0.54$) for Neptune, using the (un-)modified shape and rotation data. In case of vertical mixing up into the troposphere, one would expect to see a signature from different outer envelope metallicities in the lower atmosphere. Over the past 60 years, several attempts of deriving the atmosphere abundances from measured spectra have been undertaken, see, e.g., the reviews of \\citet{Fegley+91,Gautier+95,GuiGau07}. However, that turned out to be quite challenging, mostly because of the numerous absorption lines of CH$_4$ and H$_2$, uncertainties in the line shapes and absorption coefficients, non-equilibrium chemistry, and the dependences of the derived abundances among each other. Today, carbon is the heavy element for which the interpretation of the available data gives the least ambiguous picture. A value of C:H=30--60$\\times$ solar for both planets has finally emerged \\citep{Gautier+95}. Assuming C:H$=30\\times$ solar and similar enrichments also of N and O, \\citet{Hubbard+95} estimate an outer envelope ice mass fraction of 0.08. While this is in good agreement with the possible $Z_1$ values of our Uranus and Neptune models, it does not support our high-$Z_1$ Neptune models as required for qualitatively different interiors. On the other hand, both CO and HCN have been detected in the atmosphere of Neptune but not of Uranus \\citep{Gautier+95}, indicating different processes at work in their atmospheres. To replace assumptions by real data, we encourage deep entry probe missions to Uranus and Neptune. \\subsection{Past and present structure models} As shown in figures \\ref{fg:obsJ2J4}--\\ref{fg:unZZ2}, more definite conclusions about the dissimilarity of the internal structures are prevented by the current observational uncertainty in the gravitational moments. This is at odds with the struggles of earlier modelers \\citep{Podolak+95,Marley+95} to find Uranus and Neptune models at all that matched the less accurately known gravity data in the past. Moreover, despite the tighter present constraints, our models seem to encompass many of the previously found models. In particular, \\citet{HM89} developed Uranus models with physical EOS, smooth transitions between an outer H-He rich envelope, an inner ice-rich envelope, and a rock core, and showed that models with 5--15\\% H-He in the inner envelope could give best agreement with the imposed constraints. Such an H-He abundance is also seen in our Uranus models. When the Voyager gravity and rotation rate data for Neptune became available, \\citet{Podolak+95} applied the three-layer approach and physical EOS to compute Uranus and Neptune models. Their Uranus models required some H-He to be present in the deep interior, too, whereas for Neptune this was found to be optional (non-conventional models). They could also find a different Neptune model with an inner envelope of pure water and transition deeper inside at $\\sim 0.7\\RN$, and a highly enriched outer envelope, indicating a possible difference in the interior structures of Uranus and Neptune. Our Neptune models in the upper right corners of Figs.~\\ref{fg:unZZ1},\\ref{fg:unZZ2} are similar to that latter one, while their non-conventional models are within the bulk of our Uranus and Neptune models. \\citet{Marley+95} again allowed for smooth density gradients and generated models with random density distributions. Nevertheless, they found that a rather sharp transition from a low-density outer envelope to a high-density inner envelope is necessary to fit the \\emph{Voyager} values of $J_2$ and $J_4$, where the pressure-density relation in the inner envelope could be that of an ice layer. Because of the admittance of smooth density gradients, they found for the first time that the transition could occur as deep as between $\\sim$ 60--65\\% of the planet's radius. Our model N2b shows the same properties, although within the three-layer approach. Finally, \\citet{Helled+11} were the first to use the accurate post-Voyager gravity data of \\citet{Jacobson07,Jacobson09}. Their polynomial density-profiles could well represent a metallicity gradient that rises from about solar metallicity at the surface to about 85\\% in the center, with slightly higher values preferred for Neptune but no indication of different internal structures. The found metallicities are within those of our models. According to the present work, the application of the modified shape and rotation data gives a more pronounced indication of different structure than seen so far. \\subsection{Conclusions} We present the full sets of three-layer interior models (with H/He/water envelopes and rocks confined to the core) of Uranus and Neptune for different solid-body rotation periods and flattenings, using the improved gravity data by \\citet{Jacobson07,Jacobson09}, and physical equations of state. We find that the resulting bulk composition is insensitive to the current level of uncertainty in the input data (observational constraints and equations of state) as our results are in good agreement with previous calculations (e.g.~\\citealt{Podolak+95, Marley+95, Helled+11}). However, our models with the modified rotation periods and shapes suggest that Uranus and Neptune could be quite different. Uranus would have an outer envelope with a few times the solar metallicity which transitions to a heavily enriched ($\\sim 90\\%$ by mass heavy elements) inner envelope at $0.9\\:\\MU$, giving a rather low moment of inertia of $\\sim 0.222$. In Neptune, this transition can occur deeper inside at $0.6\\:\\MN$ and be accompanied by a more moderate increase in metallicity, leading to a less centrally condensed planet with $\\nI\\sim 0.255$. While the observed magnetic fields of Uranus and Neptune are similar and can be reproduced by a rather narrow range of dynamo models, dissimilar interiors are required to explain the measured luminosities. We have presented a new indication for different internal structures based on the application of modified shape and rotation data. However, the density gradient in our models appears to be generally farther out than required by evolution and magnetic field models. The authors thank the two referees for providing fruitful comments that led us to consider the results in a wider context. NN acknowledges support from the DFG RE 881/11-1." }, "1207/1207.7006_arXiv.txt": { "abstract": "\\vspace{0.3cm} We study the primordial density perturbations and non-Gaussianities generated from the combined effects of an inhomogeneous end of inflation and curvaton decay in hybrid inflation. This dual role is played by a single isocurvature field which is massless during inflation but acquire a mass at the end of inflation via the waterfall phase transition. We calculate the resulting primordial non-Gaussianity characterized by the non-linearity parameter, $\\fNL$, recovering the usual end-of-inflation result when the field decays promptly and the usual curvaton result if the field decays sufficiently late. \\vspace{0.3cm} ", "introduction": "There are many mechanisms by which quantum fluctuations of light fields present during inflation could affect the primordial density perturbation \\cite{Bassett:2005xm, Lyth-Liddle}. Variations in the local value of an inflaton field, slow-rolling during inflation, lead to variations in the local duration of inflation and hence the resulting local density after inflation has ended. But local variations in other fields, not necessarily evolving during inflation, could also affect the local expansion either during or after inflation. Variations orthogonal to the background field evolution during inflation have been characterized as entropy field perturbations \\cite{Gordon:2000hv} which can alter gauge-invariant density perturbations on very large (super-Hubble) scales even after inflation has ended due to variations in the local equation of state and non-adiabatic pressure perturbations \\cite{Wands:2000dp}. While adiabatic field fluctuations in a single canonical inflaton field lead to Gaussian, adiabatic density perturbations after inflation, entropy field fluctuations during inflation may lead to a much richer phenomenology of non-Gaussian and non-adiabatic primordial density perturbations \\cite{Salopek:1988qh,Linde:1996gt,Langlois:1999dw,Lyth:2002my,Bartolo:2003jx,Bassett:2005xm}. A simple example of how isocurvature field perturbations can affect the primordial density perturbation is by altering the time at which inflation ends in hybrid inflation models where the end of inflation is triggered by an tachyonic instability in a ``waterfall'' field, leading to a rapid phase transition from false to true vacuum state \\cite{Linde:1991km,Linde:1993cn,Copeland:1994vg}. If the local variations of an isocurvature field, coupled to the waterfall field, alter the point at which the false vacuum becomes unstable then they will lead to a primordial density perturbation~\\cite{Bernardeau:2004zz,Lyth:2005qk,Dvali:2003em,Sasaki:2008uc, Naruko:2008sq, Huang:2009vk, Yokoyama:2008xw, Emami:2011yi, Lyth:2012br}. On the other hand, in the curvaton scenario~\\cite{Lyth:2001nq,Moroi:2001ct} the isocurvature field is only weakly-coupled to the inflaton and its decay products, so that it survives after inflation has ended and may source the primordial density perturbation when it decays into radiation some time after inflation has ended. In this paper we will study the effect of entropy fluctuations in a massless field during hybrid inflation, in a model originally considered by Lyth~\\cite{Lyth:2005qk}. We will show how primordial density perturbations can be generated from inflaton field fluctuations and entropy fluctuations in the isocurvature field which is coupled to the waterfall field. Entropy fluctuations therefore affect the surface of end of inflation, {\\em and} the field spontaneously acquires a mass at the waterfall transition. Therefore this massless field during inflation can play the role of a curvaton field which can have a significant energy density when the field oscillations decay some time after inflation. We calculate the primordial density perturbations from all these effects using the $\\delta N$ formalism \\cite{Starobinsky:1986fxa,Sasaki:1995aw, Wands:2000dp, Lyth:2005fi}. We identify the non-linear dimensionless density perturbation, $\\zeta$ \\cite{Bardeen:1983qw,Bardeen:1988hy,Malik:2008im}, with the perturbation in the local integrated expansion, $\\delta N$, from an initial uniform-curvature hypersurface up to a uniform-density hypersurface \\cite{Lyth:2004gb} in the long-wavelength limit~\\cite{Salopek:1990jq} where the local expansion is a function of the local field fluctuations, $\\delta\\phi_i$, on the initial hypersurface: \\ba \\label{p-q} \\zeta = \\sum_i N_{,i} \\delta\\phi_i + \\frac12 \\sum_{i,j} N_{,ij} \\delta\\phi_i \\delta\\phi_j + \\ldots \\,. \\ea where $N_{,i}\\equiv \\partial N/\\partial\\phi_i$. We consider the case of canonical light scalar fields, $\\phi_i$, where all these field fluctuations have a Gaussian distribution with power spectrum ${\\cal{P}}_{\\delta \\phi_i} \\simeq (\\frac{H_k}{2\\pi})^2$ when the wave-mode of interest, $k$, leaves the Hubble-horizon, $k=aH$, with the Hubble expansion rate $H_k$. The primordial power spectrum is then given at leading order by \\ba \\label{power0} {\\cal{P}}_{\\zeta} = \\left(\\frac{H_k}{2\\pi} \\right)^2 \\sum_i N_{,i}^2 \\ea and the lowest-order non-Gaussianity parameter, $f_{NL}$, is given by \\cite{Lyth:2005fi} \\ba \\label{fNL-def} \\frac{6}{5} f_{NL} = \\frac{\\sum_{,ij} N_{,i}N_{,j} N_{,ij}}{\\left(\\sum_j N_{,j}^2 \\right)^2} \\, . \\ea We introduce the hybrid inflation model in Section~II and discuss the inflaton field dynamics and perturbations. In Section~III we discuss the end of inflation and the density perturbations produced due to the inhomogeneous end of inflation. In Section IV we discuss how the same isocurvature field can act as a curvaton field, producing additional density perturbations when it finally decays some time after inflation. We present our results for observable quantities, such as the tilt of the primordial power spectrum and the non-linearity parameter, $\\fNL$, in Section~V. We conclude in Section~VI. ", "conclusions": "In this paper we have studied a model of hybrid inflation within which the primordial density perturbation can be produced either through an inhomogeneous end of inflation or through a curvaton scenario when the field decays, or from a combination of the two. The perturbations originate from vacuum fluctuations during inflation in a light isocurvature field, $\\sigma$. This field is coupled to the waterfall field, $\\chi$, whose tachyonic instability triggers the sudden end of hybrid inflation. Hence fluctuations, $\\delta\\sigma$, can affect the point at which the instability is triggered. Because the field $\\sigma$ acquires a mass at the phase transition, it will begin oscillating about its minimum when the mass is larger that the Hubble scale after inflation. If the oscillations are sufficiently long-lived then the energy density in the $\\sigma$ field may become significant and inhomogeneities in the energy density of the $\\sigma$ field are transferred to the primordial radiation density when the field decays, as in the curvaton scenario. \\footnote{Our model is very different from the recently proposed ``hybrid curvaton'' model of Dimopoulos et al \\cite{Dimopoulos:2012nj} where the curvaton field has a waterfall-type potential, so that oscillations of the curvaton field are triggered by a tachyonic instability.} We are able to recover the standard results for primordial perturbations from an inhomogeneous end of inflation when the $\\sigma$ field decays instantaneously at the end of inflation. We also recover familiar curvaton results in the limit where the field is sufficiently long-lived and curvature perturbations produced at the end of inflation are negligible. In both cases we express the resulting non-linearity parameter, $\\fNL$, in terms of the fractional density of the curvaton field, $\\Omega_\\sigma$, either at the end of inflation or when the curvaton decays. We find that for large non-Gaussianity we have \\be \\fNL \\approx \\left\\{ \\begin{array}{ll} w_\\sigma^2 ({5\\alpha}/{9\\Omega_{\\sigma,e}}) & {\\rm for}\\ \\Omega_{\\sigma,e} \\gg \\alpha\\Omega_{\\sigma,d} \\\\ w_\\sigma^2 ({5}/{3\\Omega_{\\sigma,d}}) & {\\rm for}\\ \\Omega_{\\sigma,d} \\gg \\Omega_{\\sigma,e}/\\alpha \\end{array} \\right. \\,. \\ee Any contribution to the overall primordial power spectrum from inflaton fluctuations, $w_\\sigma<1$, tends to suppress the primordial non-Gaussianity. Primordial perturbations from an inhomogeneous end of inflation are enhanced by the slow-roll parameter $\\alpha$. Nonetheless curvaton-type perturbations tend to dominate over the end-of-inflation effect when the decay is slow, $\\Gamma\\ll \\alpha^2 H_e$, unless the $\\sigma$-density is already significant immediately after the end of inflation, $\\Omega_{\\sigma,e}>\\alpha$. More generally we find that perturbations due to the inhomogeneous end of inflation place an upper bound on the non-Gaussianity \\be \\fNL \\leq \\frac{5\\alpha}{9\\Omega_{\\sigma,e}} \\,. \\ee We have used sudden-end-of-inflation and sudden-decay approximations to match the curvature and density perturbations across the end-of-inflation hypersurface and the curvaton-decay hypersurface. We expect this instantaneous matching to be a good approximation on length-scales much larger than the corresponding Hubble length. This long-wavelength approximation is usually an excellent approximation for scales relevant for the large-scale structure of our Universe \\cite{Deruelle:1995kd,Malik:2006pm,Sasaki:2006kq}. On the other hand we know that the tachyonic instability in the waterfall field at the end of inflation leads to a very inhomogeneous fragmentation of the inflaton and waterfall field due to tachyonic preheating \\cite{Felder:2000hj, Felder:2001kt}. If a curvaton field, $\\sigma$, is weakly coupled to the waterfall field so that its mass remains less than the Hubble scale (a light curvaton immediately after inflation) then it remains overdamped and we would not expect the coherent curvaton field to be significantly disrupted by preheating in the inflaton and waterfall fields. On the other hand if the curvaton is sufficiently strongly coupled so that its mass becomes larger than the Hubble scale immediately after inflation, then its dynamics may be considerably more complicated, and possibly highly non-linear. Previous studies \\cite{Enqvist:2008be,Chambers:2009ki} have considered resonant decay of the curvaton, but here we are considering possible parametric production of the curvaton itself. This has no effect on large scale density perturbations if all the fields rapidly thermalize, but short wavelength curvaton modes can alter the predictions in the curvaton scenario~\\cite{Sasaki:2006kq}. A full study of the effect on the curvaton production and decay after tachyonic preheating in this intermediate regime requires a careful numerical treatment which goes well beyond the present work." }, "1207/1207.2765_arXiv.txt": { "abstract": "Stellar models generally use simple parametrizations to treat convection. The most widely used parametrization is the so-called ``Mixing Length Theory'' where the convective eddy sizes are described using a single number, $\\alpha$, the mixing-length parameter. This is a free parameter, and the general practice is to calibrate $\\alpha$ using the known properties of the Sun and apply that to all stars. Using data from NASA's {\\it Kepler} mission we show that using the solar-calibrated $\\alpha$ is not always appropriate, and that in many cases it would lead to estimates of initial helium abundances that are lower than the primordial helium abundance. {\\it Kepler} data allow us to calibrate $\\alpha$ for many other stars and we show that for the sample of stars we have studied, the mixing-length parameter is generally lower than the solar value. We studied the correlation between $\\alpha$ and stellar properties, and we find that $\\alpha$ increases with metallicity. We therefore conclude that results obtained by fitting stellar models or by using population-synthesis models constructed with solar values of $\\alpha$ are likely to have large systematic errors. Our results also confirm theoretical expectations that the mixing-length parameter should vary with stellar properties. ", "introduction": "\\label{intro} Accurately treating convective heat transport in stellar models is difficult. The structure and evolution of most stars is related to convective transport processes in their outer layers. The transition from efficient convective transport in the deep envelope to the radiative atmospheric layers takes place in a region of inefficient convection where the temperature gradient is highly superadiabatic. The poorly known structure of this region remains one of the major uncertainties in stellar models. In one-dimensional calculations, this region is usually modeled using the ``mixing length theory,'' or MLT \\citep{boh58}. This prescription assumes that one can approximate the full range of turbulent eddy-sizes by a typical size, and that an eddy on average travels a distance determined by the eddy size before losing its identity. This distance, is known as the ``mixing length'' and is usually defined as $\\alpha H_p$, where $\\alpha$ is known as the ``mixing-length parameter,'' and $H_p$ is the local pressure scale height. Given the atmospheric structure, the approximation, when given the mixing length, fix the specific entropy in the deep convection zone, which in turn determines the radius of the model. { The model radius thus depends sensitively on the choice of the mixing-length parameter $\\alpha$, which is a free parameter that cannot be determined from the mixing length theory.} The common practice when modeling stars to determine their structure and evolution is to use the solar value of $\\alpha$. We know the mass, radius, luminosity and age of the Sun precisely. Solar models are constructed by searching for $\\alpha$ and the initial helium abundance $Y_0$ that yield a model with the correct radius and luminosity at the Sun's age. Since masses, radii and ages are usually unknown for other stars, the solar approach is unfeasible and, hence, the solar value of $\\alpha$ is used. The assumption of a fixed $\\alpha$ has no {\\it a priori} justification, and indeed, there is some evidence that the solar value of $\\alpha$ does not always work for other stars. \\citet{la84} first noted that the radius of $\\alpha$~Cen~A cannot be reproduced using the solar value of $\\alpha$. A combined astrometric and seismic study of the $\\alpha$~Cen system by \\citet{de86} confirmed this result as have more recent studies (e.g., Fernandes \\& Neuforge 1995; Miglio \\& Montalban 2005). Some studies of other binary systems have also suggested a mass-dependence of $\\alpha$ (e.g., Ludwig \\& Salaris 1999; Morel et al. 2000; Lebreton et al. 2001; Lastennet et al. 2003, etc.). In particular, Y{\\i}ld{\\i}z et al.~(2006) studied binaries in the Hyades cluster and suggested that $\\alpha$ increases with stellar mass. Numerical simulations of stellar convection also suggest that convective properties vary with stellar parameters (e.g., Ludwig et al. 1999; Trampedach 2007; Trampedach \\& Stein 2011). While studies of binaries and clusters have suggested that the solar $\\alpha$ is not always applicable, the situation for single field stars is not clear because of the lack of observational constraints. Asteroseismic data from space missions like CoRoT (Michel et al. 2008) and {\\it Kepler} (Borucki et al. 2010) allow us to place independent constraints on the mass and radius of single stars. These constraints, along with the classical constraints of $T_{\\rm eff}$ and metallicity, allow us to constrain $\\alpha$. We already know that the oscillation spectra of some CoRoT and {\\it Kepler} stars cannot be reproduced using the solar value of $\\alpha$ (e.g., Metcalfe et al. 2010; Deheuvels \\& Michel 2011, Mathur et al. 2012). Detailed asteroseismic modeling % of the solar analogs 16 Cyg A \\& B has also required the adoption of non-solar values of $\\alpha$ % (Metcalfe et al. 2012). % In this study we use asteroseismic data obtained by NASA's {\\it Kepler} mission to calibrate $\\alpha$ for a sample of dwarfs and subgiants. Observational data and our analysis technique are described in \\S~\\ref{sec:methods}. We present our results and discuss their implications in \\S~\\ref{sec:results}. It should be noted that $\\alpha$ is just a proxy for describing stellar convection, and in particular $\\alpha$ describes the entropy change between the surface and the deeper, isentropic, layers of efficient convection. As a result, the value of $\\alpha$ cannot be derived uniquely. It depends on the exact formulation of the mixing-length theory (see, e.g., Appendix A of Ludwig et al. 1999), as well as physics inputs that affect entropy e.g., atmospheric opacities, the temperature-optical depth relation ($T$-$\\tau$) in the atmosphere, the equation of state and processes such as the gravitational settling of helium and heavy elements. Thus, the value of $\\alpha$ in a given star needs to be examined in the context of the value of $\\alpha$ needed to construct a solar model with the same physics. { Note that conventional MLT assumes a fixed ratio between the distance traveled by an eddy and its size, and $\\alpha$ is the only free parameter. Some approximations (see e.g., Arnett et al. 2010) leave the ratio as an adjustable parameter. } We use the conventional B\\\"ohm-Virtense form of MLT and only adjust $\\alpha$. ", "conclusions": "\\label{sec:results} The initial helium abundance, $Y_0$, needed to construct models of the stars in our sample --- assuming that they all have the solar value of $\\alpha$ --- is shown in Fig.~\\ref{fig:yt}. Note that for $> 50$\\% of our sample the $Y_0$ estimate is less than the primordial value of $Y_{\\rm p}=0.2477\\pm 0.0029$ (Peimbert et al. 2007). { While the deficit is within $1\\sigma$ for some stars, there are stars with a $>3\\sigma$ deficit. The weighted average of $Y_0$ of the sample is $0.227\\pm 0.004$. The weighted median is $0.223$, both below $Y_{\\rm p}$.} It is, of course, highly unlikely that stars are born with less helium than was produced in the Big Bang. Given that the only parameter we could change in MLT is $\\alpha$, this implies that solar $\\alpha$ does not properly approximate convective heat transport in these stars. Note that a change of physics inputs to the models will change the solar value of $\\alpha$, and the exercise repeated with the new solar $\\alpha$ would give similar results. In Fig.~\\ref{fig:at}(a) we show the value of $\\alpha$ for our sample obtained assuming either the solar value of $Y_0$ (red points), or the simple chemical evolution model of $Y_0$ (black points). Note that $\\alpha$ for most of the stars is less than the solar value for both cases The average value of $\\alpha$ for this sample is $1.522$ for solar $Y_0$ and $1.597$ when the chemical-evolution model is used. In MLT, lower $\\alpha$ implies less efficient convection. Thus MLT predicts that in the superadiabatic layers, convective energy transport in our sample is generally less efficient than that in the Sun. Since the results with the two choices of $Y_0$ are similar, in the subsequent discussions we only use $\\alpha$ obtained with the chemical evolution model of $Y_0$. { Mathur et al. (2012) constructed detailed models to fit the mode frequencies of 22 {\\it Kepler} stars using the Asteroseismic Modeling Portal (AMP; Metcalfe et al. 2009) There are 16 stars in common with our sample. In Fig.~\\ref{fig:at}(b) we show the differences between the Mathur et al. $\\alpha$ values and the ones obtained in this work. The two $\\alpha$ estimates agree well, mostly within $1\\sigma$. Since the physics in the AMP models is different from ours, they obtain a solar $\\alpha$ of 2.12. Thus to compare their results with ours, we have scaled the AMP results to our value of the solar $\\alpha$. In Fig~\\ref{fig:at}(c) and (d) we show the variation of $\\alpha$ with $\\log g$ and [M/H]. } { In order to explore whether our $\\alpha$ estimates are correlated with stellar properties, we first determined the simple Spearman rank correlation between $\\alpha$ and different properties. The correlation coefficients are listed in Table~1. Also listed is the $p$-value, which is the probability that the correlation is a chance occurrence. A small $p$ therefore indicates a significant correlation. There thus appears to be significant correlation between metallicity and $\\alpha$. There also seems to be a mildly significant correlation between mass and $\\alpha$. The significance of the correlation of $\\alpha$ with $\\log g$ or $T_{\\rm eff}$ depends on whether or not we include the low-$\\log g$ stars in our analysis. Table~1 lists the coefficient obtained for the entire sample, as well as that obtained by removing the lowest $\\log g$ stars ($\\log g < 3.8$) in our sample. } { Since $\\alpha$ depends simultaneously on a number of parameters, to get a better estimate of the correlations we perform a trilinear fit to $\\alpha$ with the model \\begin{equation} \\alpha=a+b\\log g+c\\log T_{\\rm eff}+d{\\rm [M/H]}. \\label{eq:mod} \\end{equation} Table~1 lists the coefficients and $p$-values, and Fig.~\\ref{fig:avar} shows the residuals, and partial residuals, of the fit to Eq.~\\ref{eq:mod}. Note that the metallicity dependence is robust. The $\\log g$ and $T_{\\rm eff}$ correlations are small and less statistically significant when the entire sample is used; these increase in significance, but change signs, when the $\\log g$ cut is applied. The mass correlation seen in the Spearman correlation is most likely to be the result of the mass-$T_{\\rm eff}$ and mass-$\\log g$ correlation seen in Fig.~1 as indicated by the fact that the residuals of the fit to Eq.~\\ref{eq:mod} do not show any trend with mass (Fig.~\\ref{fig:avar}(a) and (e)). } { The metallicity dependence of $\\alpha$ is relatively easy to understand. It is most likely caused by the temperature sensitivity of the H$^-$ density and hence, of its dominant contribution to the optical continuum opacity. E.g., in the solar photosphere, $\\sim 50$\\% of the electrons that form H$^-$ are donated by metals, and the fraction increases steeply with height due to low-ionization potential elements like Na, Al, K, Ca and Cr. With a smaller amount of metals, the temperature sensitivity of the H$^-$ density will therefore increase. This in turn will increase the contrast between up- and down-flows, and especially increase the range of depths over which the down-flows will be cooled. This results in a larger entropy jump between the surface and the deeper layers, meaning a lower convective efficiency (a smaller $\\alpha$ in the context of MLT). } Since the average metallicity of our sample is sub-solar, we believe that this metallicity dependence accounts for the lower-than-solar average value of $\\alpha$ for our sample. The lack of a significant correlation between $\\alpha$ and $T_{\\rm eff}$ or $\\log g$ is surprising. This is most likely the result of the limited and skewed range of $\\log g$ of our sample, and is confirmed by the change of the sign of the correlation when the $\\log g$ cutoff is applied. A larger sample should resolve these issues, in particular, data on giants should help determine the $\\log g$ dependence properly. Piau et al. (2011), using a sample of red giants with radii known from interferometry, have shown that the red-giant models require sub-solar values of $\\alpha$ to explain the observations; however, they did not address the dependence of $\\alpha$ on stellar parameters. As noted earlier, $\\alpha$ is just a proxy for describing convection in stars. Although it is known that using such a proxy does not reproduce properties of the stellar near-surface layers correctly, MLT remains a practical tool in stellar modeling. An $\\alpha$ type parameter can also be derived from numerical simulations of stellar convection (e.g., Ludwig et al. 1999; Trampedach et al. 1999; Trampedach 2007). At present it is difficult to compare these results with ours since the simulations were for solar composition, and the metallicity dependence of our results is fairly strong. However, there do seem to be some differences between our findings and the simulations. At a given $\\log g$, { Ludwig et al. (1999) found $\\alpha$ to decrease with increasing $T_{\\rm eff}$ for dwarfs in their 2D simulations and they find $\\alpha$ to decrease with $\\log g$. A similar behavior was seen in the 3D simulations of Trampedach (2007). Our $\\alpha$-$T_{\\rm eff}$ correlation agrees with theirs, but the $\\log g$ one does not when we examine the entire sample; when we apply the $\\log g$ cutoff the reverse becomes true, the $\\log g$ correlation agrees, the $T_{\\rm eff}$ correlation does not. } As more detailed asteroseismic data become available from {\\it Kepler}, and they are modeled, we will be able to reduce the uncertainties in the mixing-length parameters needed to model the stars, and the dependence of the properties of near-surface convection, including the corresponding value of $\\alpha$, on stellar parameters will become clearer. As it is, our results have important implications for the different branches of astrophysics that depend on fitting stellar models. The usual way to determine stellar properties is through spectroscopic or photometric analyses of the star combined with fitting to grids of stellar models to obtain masses and radii (e.g., Takeda et al. 2007) or ages (e.g., J{\\o}rgensen \\& Lindegren 2005). At a given metallicity, a lower $\\alpha$ makes the evolutionary track of a given mass redder than its higher-$\\alpha$ counterpart, and thus the properties of a star obtained using models constructed with a solar $\\alpha$ will be quite different from those obtained with a sub-solar $\\alpha$. Although the $T_{\\rm eff}$ change with $\\alpha$ is really a result of a radius change, it appears as a change in the estimated mass of star being fitted (Basu et al. 2011). Consequently, we would be underestimating the mass of a sub-solar metallicity star if we use models constructed with the solar value of $\\alpha$. Stellar population and spectral synthesis models also use a single (usually the solar calibrated) value of $\\alpha$ (see e.g. Coelho et al. 2007), and our results now show that the uncertainties in $\\alpha$ need to be added to the error budget of results that use those models." }, "1207/1207.1599_arXiv.txt": { "abstract": "We intend to provide a comprehensive answer to the question on whether all Coronal Mass Ejections (CMEs) have flux rope structure. To achieve this, we present a synthesis of the LASCO CME observations over the last sixteen years, assisted by 3D MHD simulations of the breakout model, EUV and coronagraphic observations from \\textsl{STEREO} and \\textsl{SDO}, and statistics from a revised LASCO CME database. We argue that the bright loop often seen as the CME leading edge is the result of pileup at the boundary of the erupting flux rope irrespective of whether a cavity or, more generally, a 3-part CME can be identified. Based on our previous work on white light shock detection and supported by the MHD simulations, we identify a new type of morphology, the `two-front' morphology. It consists of a faint front followed by diffuse emission and the bright loop-like CME leading edge. We show that the faint front is caused by density compression at a wave (or possibly shock) front driven by the CME. We also present high-detailed multi-wavelength EUV observations that clarify the relative positioning of the prominence at the bottom of a coronal cavity with clear flux rope structure. % Finally, we visually check the full LASCO CME database for flux rope structures. In the process, we classify the events into two clear flux rope classes (`3-part', `Loop'), jets and outflows (no clear structure). We find that at least 40\\% of the observed CMEs have clear flux rope structures and that $\\sim29\\%$ of the database entries are either misidentifications or inadequately measured and should be discarded from statistical analyses. We propose a new definition for flux rope CMEs (FR-CMEs) as a coherent magnetic, twist-carrying coronal structure with angular width of at least 40$^\\circ$ and able to reach beyond 10 R$_{\\odot}$ which erupts on a time scale of a few minutes to several hours We conclude that flux ropes are a common occurrence in CMEs and pose a challenge for future studies to identify CMEs that are clearly \\textsl{not\\/} FR-CMEs. ", "introduction": "\\label{sec:Introduction} Since their detection in the early 1970s, Coronal Mass Ejections (CMEs) have been the subject of intense investigation with regard to their initiation mechanisms, their effects on the corona and their association with other coronal phenomena (eg., flares and prominences). This \\textsl{Topical Issue\\/} presents results from a Coordinated Data Analysis Workshop (CDAW) devoted to the question: `Do All Coronal Mass Ejections (CMEs) Have Flux Rope Structures?' Such a specific physics-based question shows that we have come a long way towards understanding the nature of these explosive events especially when we consider the original definition of a CME: `a relatively short scale white light feature propagating in a coronagraph's field of view' (paraphrasing \\opencite{1984JGR....89.2639H}). Traditionally, CMEs were observed with visible light coronagraphs and clues on their origin and nature were based on their morphology in those images \\cite{1979SoPh...61..201M,1985JGR....90.8173H,1993STIN...9326556B}. Despite the apparently large variation in the appearance of CMEs, two particular morphologies stand out: the `loop'-CME where a bright narrow loop-like structure comprises the CME front, and the `3-part'-CME \\cite{1985JGR....90..275I} where the bright front is followed by a darker cavity which frequently contains a bright core. It has become the archetypical morphology of a CME even though the `3-part' morphology could be identified in only about a third of the events \\cite{1979SoPh...61..201M}. It is still unclear whether the remaining variation is the result of projection effects due to the optically thin nature of the emission or not. It was recognized early that the cavity rather than the prominence in the core drove the CME \\cite{1987sowi.conf..181H}. An initial controversy on whether CMEs were planar (i.e., ejected loops) or three-dimensional (i.e., bubbles) structures was largely resolved by the end of the 1980's. \\inlinecite{1983SoPh...83..143C} demonstrated, using polarization analysis, that the loop front was indeed a bubble. The identification of halo CMEs by \\inlinecite{1982ApJ...263L.101H} with their quasi-circular appearance established their three-dimensional (3D) nature and led to the adaption of the 'ice-cream' model to describe and fit the kinematics of these events \\cite{1982ApJ...263L.101H,2002JGRA..107.1223Z,2004JGRA..10903109X,2005JGRA..11008103X}. A bubble or spherical structure is the intrinsic assumption behind this model which, by the way, is not a proper description as we will discuss later. As theories progressed towards a more physical basis for the CME initiation, they focused on 3D magnetic topologies that could account for the `3-part' morphology and the frequent association with prominences. This quickly led to scenarios of rising loop arcades, overlying a prominence, which underwent reconnection to form magnetic flux ropes (FR, hereafter; \\opencite{1982SoPh...79..129A}; \\opencite{1990JGR....9511919F}). Alternatively, the FR could pre-exist and rise under the driving of Lorenz forces \\cite{1974A&A....31..189K,1993GeoRL..20.2319C}. While the question on whether the FR is formed before or during the eruption remains open, the overwhelming majority of magnetohydrodynamics (MHD) models and simulations agree on one thing. Namely, the erupting structure is always a FR \\cite{2011LRSP....8....1C}. There is no physical mechanism that can produce a large-scale fast eruption from the corona without ejecting a fluxrope, to the best of our knowledge. At the same time, in-situ measurements of interplanetary CMEs (ICMEs) often encounter structures with smooth rotation in one, or more, components of the magnetic field which can be fitted with FR models (\\opencite{1982JGR....87..613K}; \\opencite{1990JGR....9511957L}; \\opencite{2011SoPh..273..205I}, to name a few). These so-called Magnetic Clouds (MCs) can be considered then as the interplanetary manifestations of the ejected FR predicted by theory and possibly detected as the cavity in the `3-part'-CMEs \\cite{1982GeoRL...9.1317B}. \\inlinecite{2003JGRA..108.1156C} found that 100\\% of ICMEs detected during solar minimum were MCs reducing to $<20\\%$ during solar maximum. So the CDAW question regarding the nature of CMEs, at least in the case of `3-part'-CMEs, seems to have been answered. A CME is simply the ejection of a magnetic FR structure from the lower corona which takes the form of a `3-part'-CME or a MC depending on the instrumenation used (images or in-situ, respectively) to detect it. But, if a FR is a necessary ingredient for an ejection, why not all CMEs show evidence for such structure? In other words, why all CMEs are not `3-part'-CMEs? Some have just a loop front while others appear as jets or structureless clouds or blobs. For example, \\inlinecite{1985JGR....90.8173H} categorized CMEs, between 6-10 R$_{\\odot}$, into ten morphological classes based on their appearance in \\textsl{Solwind\\/} observations. Why is there such a large variety of shapes? Could there be other types of magnetic structures, besides FRs, ejected from the Sun? If they do exist, they would suggest a major gap in our understanding of eruptive processes, given the prevalence of FR in our theories. Second, not all ICMEs exhibit MC signatures. Is this simply a result of `glancing' cuts between in-situ instruments and the ICME? Or do CME FRs lose their coherence as they travel in the interplanetary space, through reconnection with the ambient solar wind for example \\cite{2007SoPh..244..115D}? Third, many fast ICMEs are driving a shock followed by a sheath of post-shocked plasma. The resulting five-part ICME (shock, sheath, dense front, cavity, and dense plug) does not have a coronal counterpart. Where are the five-part CMEs or more precisely, where are the shock and sheath signatures in the coronagraph images? Shocks could deflect streamers and generally affect the ambient corona, ahead and at the flanks of a CME, thus creating complex brightness distributions in the images. Could such effects be responsible for misidentifications, and hence misinterpretations, of CME morphologies, kinematic profiles, and associations with structures in the low corona or the inner heliosphere? Fourth, and related point, the emission processes in both low (EUV) and middle (white light) corona are optically thin resulting in images that are projections on the plane of sky (POS). Do these projections affect our ability to properly interpret observations and how can we account for them? We will address this problem throughout this paper. The Large Angle and Spectrometric Coronagraph (LASCO; \\opencite{1995SoPh..162..357B}) project has accumulated the largest and longest database of coronagraphic observations of CMEs since 1996. Spanning more than % a complete solar cycle, it is reasonable to expect that events of every possible orientation, size, speed, mass, and morphologies have been captured. We should be in position to understand the role of projection effects on the images, identify the origin of the various features (CME or not) in a given LASCO image, and hence answer the question posed in this \\textsl{Topical Issue}. To accomplish this task comprehensively we have given this paper a relatively large scope. It represents a synthesis of the observational knowledge gained over the sixteen years of LASCO observations. In the following sections, we will provide: evidence for the FR structure within CME cavities (Section~\\ref{sec:3part}), evidence for the existence of white-light shock and tips on distinguishing the shock front from the CME front (Section~\\ref{sec:shock}), theoretical support for these interpretations using synthetic images from 3D MHD simulations (Section~\\ref{sec:models}), observations that clarify the connection between prominence and erupting cavity (Section~\\ref{sec:prom}), and finally statistics on the occurence of `3-part' or more precisely FR-CMEs, along with a discussion on the constrains of event lists (Section~\\ref{sec:stats}). We discuss and conclude in Section~\\ref{sec:discussion}. We will support several of our predictions and conclusions by using two-viewpoint imaging afforded by the Sun-Earth Connection Coronal and Heliospheric Investigation (SECCHI; \\opencite{2008SSRv..136...67H}) on-board the {\\it Solar TErrestial RElations Observatory (STEREO)} \\cite{2008SSRv..136....5K}. We will use the SECCHI observations as necessary but we want to focus on the single viewpoint from LASCO for two reasons. First, this article is part of a workshop devoted on the analysis of events observed with LASCO. Second, and more important, the \\textsl{STEREO\\/} mission has a finite lifetime. Budgetary and other concerns suggest that future observations (whether research or operationally oriented) will be obtained from a single vantage point. It is therefore crucial that future observers can interpret such single viewpoint observations accurately. ", "conclusions": "\\label{sec:discussion} Our aim is to provide convincing evidence of the CME as an erupting FR. To that end, we have used a variety of EUV and white light observations, MHD simulations, statistics, and have considered projection effects and theoretical predictions. Leaving the question of CME initiation aside, we found that the following picture can lead to a self-consistent interpretation of the observations across many wavelength ranges and is in agreement with the majority (if not all) of our current theoretical understanding of explosive energy release from the Sun. Basically, a CME is the eruption of a magnetic flux rope with its emission measure dominated by coronal temperature plasma, carrying a prominence along its bottom dips, piling up the overlying streamer plasma, and driving a wave ahead (if the acceleration is sufficently high). This interpretation has long been adopted for the '3-part'-CMEs, as we discussed earlier. The novelty in this work is the interpretation of the bright loop front as the pileup of material \\textsl{at the boundary of the flux rope} irrespective of the `3-part' appearance. The interpretation is supported strongly by the MHD simulations and straightforward physical reasoning (Section~\\ref{sec:models}). A FR structure propagating through plasma presents an extended obstacle against which the material is piled up and transported outwards. The narrow width and brightness of that front further suggests that the pileup occurs over a sharp boundary. Such a boundary is expected between the closed FR fields and the ambient magnetic field. The sharpness of the boundary may depend on the rate of magnetic field influx in the FR during its formation or the initial acceleration and starting height. Such effects have important connections to theories of CME initiation and can be investigated now. The other novelty is the introduction of the \\textsl{`two-front'\\/} morphology by pointing out the existence of faint, relatively sharp, fronts ahead of the bright loop front. The interpretation of the faint front as density compression by a wave (or shock, depending on speed) is again supported by MHD simulations, observations and physical expectations. The stark observational differences between the bright sharp front and diffuse front clearly point to a different origin. The diffuse fronts are: well-defined, faint, followed by diffuse emission, can be very extended, and envelope the sharp fronts. The sharp fronts, in turn, are: sharp, bright, followed by emission depletions, have well-defined extents, and are behind the diffuse fronts. The faint fronts appear only during fast eruptions and their characteristics, especially the weakness of their emission and lack of post-front depletion are strong indications that these fronts are results of local density compression and not of transported piled-up plasma. The, albeit few, 3D reconstructions of the density profile across the front \\cite{2009ApJ...693..267O} can readily explaining the profile as a result of LOS integration and recover compression ratios in agreement with theoretical expectations (less than 4). Besides its importance for understanding coronal shocks, the identification of the `two-front' morphology allows an understanding of the geometry of halo CMEs as it can help us distinguish among shock, streamer deflections and FR signatures in the coronagraph images. In that way, we can now obtain accurate outlines of the FR (or the shock, depending on the problem at hand) which should lead to better inputs to CME propagation models. The identification of these two features leads to a much simpler classification of CME white light morphologies. We used four categories (ignoring the `Unknown' category) compared to nine in \\inlinecite{1985JGR....90.8173H}. Two of them (`FR' and `Loop') refer to the same FR instrinsic structure as we have argued. Jet-CMEs also contain helical structures as recent research has shown \\cite{2008ApJ...680L..73P,2009ApJ...691...61P,2010AnGeo..28..687N}. Thus, our classification is essentially reduced to events with and events without \\textsl{apparent} helical topologies. The helical topology may not be visible in the latter for several reasons. They may propagate at large angles from the POS \\cite{2007ApJ...655.1142S} or through areas disturbed by previous events. They may be too compact to discern their cavity morphology without favorable projections \\cite{2006ApJ...650.1172W}. Finally some of these events do not appear to be CMEs in the first place failing to reach large distances in the corona (called `failed' CMEs by \\opencite{Vourlidas10}). A certain number of the remaining events appear to be related to H$\\alpha$ and/or 304 \\AA\\ surges similar to the event studied by \\inlinecite{2003ApJ...598.1392V}. The low coronal signatures of these events do not exhibit any particular morphology or geometry and hence tend to appear as semi-amorphous clouds, with the occassional traces of cool material. Our final estimate of 41\\% for the rate of occurrence of FR-CMEs in the LASCO data may not look very differernt from the widely quoted number of 30\\%. However, one must first consider the size of the event samples in past morphological works. \\inlinecite{1979SoPh...61..201M} reported a 26\\% occurence of `Loop'-CMEs in a sample of 77 SMM CMEs while \\inlinecite{1984ARA&A..22..267W} found loop and bubble CMEs in 80\\% of 65 SMM CMEs. Obviously, selection bias is important with such small event samples. The largest morphological study to date categorized 998 \\textsl{Solwind\\/} CMEs of which 31.3\\% belonged to an FR-CME class (we summed the statistics for the following structural classes, curved front, loop, streamer blowout, fan) \\cite{1985JGR....90.8173H}. We base our statistics here on a sample of 2970 events, 3$\\times$ larger than the \\textsl{Solwind\\/} sample and is still expanding. We will classify the full LASCO database in the near future. Therefore, we feel that our numbers in Table~1 are quite robust and a large improvement over past work. The central question of this \\textsl{Topical Issue\\/} is whether all CMEs are flux ropes. To provide a conclusive answer (to the extent possible in science), we attacked the problem in several ways: multiple viewpoint coronagraphic observations of CMEs, multi-thermal EUV observations of the pre-erupting structures, 3D MHD simulations, and large sample statistics. We summarize our findings as follows: \\begin{itemize} \\item The detection of a bright filamentary front in CMEs is a clear indication of the existence of a FR even if the event does not exhibit the classical 3-part morphology. \\item At least $41\\%$ of CMEs exhibit clear FR signatures (`3-part or `loop') in the coronagraph images. \\item The `two-front' morphology (faint front followed by a bright loop) is a reliable indicator of a CME-driven wave (or shock, depending on speed). \\item The FR can be separated from the shock signatures in images of halo CMEs at least in locations where the bright loop appears. \\item MHD simulations are able to capture the main structural properties of white light CMEs. \\item The prominence is not the cavity and is not the FR but is the core. The cool prominence material rests on the dips of the field lines comprising the FR (in the case of pre-existing FR, at least). \\item The majority of the prominence material either drains to the surface or is heated to coronal temperatures during the early phases of the eruption. This may be the reason for the scarcity of in-situ detections of cool material. \\item A typical fast CME comprises five parts: shock front, diffuse sheath, bright front, cavity, and core. \\end{itemize} Our discussion suggests that it is time to rethink the original definition for a CME \\cite{1984JGR....89.2639H}, as expressed in \\opencite{2006LRSP....3....2S}: \\textsl{``We define a CME to be an observable change in coronal structure that 1) occurs on a time scale of a few minutes and several hours and 2) involves the appearance (and outward motion) of a new, discrete, bright, white light feature in the coronagraph field of view.''} This definition manages to be broad (no mention of the physical origin or nature of the `structure') and narrow (CME is defined as a white light feature observed by a coronagraph) at the same time. It may have been an appropriate definition during the times of exploratory CME research, sparse wavelength coverage, and simplified physical models. But times have changed. We are regularly studying CMEs with multiple instruments and wavelengths, have accumulated CME observations spanning a full solar cycle, and are asking highly detailed questions with their modeling. Thus, it may be useful to derive a more precise CME definition using physically-based terms, at least for the events exhibiting clear FR structures (FR-CMEs). Based on the work presented here and in \\inlinecite{Vourlidas10}, we propose the following definition: \\textsl{We define an FR-CME to be the eruption of a coherent magnetic, twist-carrying coronal structure with angular width of at least 40$^\\circ$ and able to reach beyond 10 % R$_{\\odot}$ which occurs on a time scale of a few minutes to several hours}. The next challenge now is whether we can apply this definition to all CMEs (hence replace `FR-CME' with `CME' above). In other words, we propose that the proper questions we should be asking is not \\textsl{`are all CME flux ropes?'\\/} but rather \\textsl{`Are there \\underline{any\\/} CMEs that are not FR-CMEs?'} \\begin{acks} The work of AV and RAH is supported by NASA contract S-136361-Y to the Naval Research Laboratory. BJL and YL acknowledge support from AFOSR YIP FA9550-11-1-0048, NASA NNX11AJ65G, and NNX08AJ04G. We thank G. Stenborg for his continuing efforts to provide better quality solar images. SOHO is an international collaboration between NASA and ESA. LASCO was constructed by a consortium of institutions: the Naval Research Laboratory (Washington, DC, USA), the Max-Planck-Institut fur Aeronomie (Katlenburg-Lindau, Germany), the Laboratoire d'Astronomie Spatiale (Marseille, France) and the University of Birmingham (Birmingham, UK). The LASCO CME catalog is generated and maintained at the CDAW Data Center by NASA and The Catholic University of America in cooperation with the Naval Research Laboratory. The SECCHI data are produced by an international consortium of the NRL, LMSAL and NASA GSFC (USA), RAL and Univ. Bham (UK), MPS (Germany), CSL (Belgium), IOTA and IAS (France). \\end{acks}" }, "1207/1207.6693_arXiv.txt": { "abstract": "We investigate molecular evolution from a molecular cloud core to a first hydrostatic core in three spatial dimensions. We perform a radiation hydrodynamic simulation in order to trace fluid parcels, in which molecular evolution is investigated, using a gas-phase and grain-surface chemical reaction network. We derive spatial distributions of molecular abundances and column densities in the core harboring the first core. We find that the total of gas and ice abundances of many species in a cold era (10 K) remain unaltered until the temperature reaches $\\sim$500 K. The gas abundances in the warm envelope and the outer layer of the first core ($T \\lesssim 500$ K) are mainly determined via the sublimation of ice-mantle species. Above 500 K, the abundant molecules, such as H$_2$CO, start to be destroyed, and simple molecules, such as CO, H$_2$O and N$_2$ are reformed. On the other hand, some molecules are effectively formed at high temperature; carbon-chains, such as C$_2$H$_2$ and cyanopolyynes, are formed at the temperature of $>$700 K. We also find that large organic molecules, such as CH$_3$OH and HCOOCH$_3$, are associated with the first core ($r \\lesssim 10$ AU). Although the abundances of these molecules in the first core stage are comparable or less than in the protostellar stage (hot corino), reflecting the lower luminosity of the central object, their column densities in our model are comparable to the observed values toward the prototypical hot corino, IRAS 16293-2422. We propose that these large organic molecules can be good tracers of the first cores. ", "introduction": "It is well established that a star is formed by gravitational collapse of a molecular cloud core. The collapsing core is initially optically thin to the thermal emission of dust grains, and undergo isothermal run-away collapse as long as the cooling rate overwhelms the compressional heating. The isothermal condition breaks down when the central density reaches $\\sim$10$^{10}$ cm$^{-3}$, and the temperature starts rising. Increasing gas pressure decelerates the contraction, and the core come to the hydrostatic equilibrium, which is called a first hydrostatic core \\citep[e.g.,][]{larson69,masunaga98}. Although the first core is a transient object, it plays essential roles in determining the physical condition of a protostar; it fragments and forms binaries \\citep{matsumoto03}, and drives bipolar molecular outflows \\citep{tomisaka02}. When the temperature reaches $\\sim$2000 K and hydrogen molecules start to dissociate, the central region of the first core collapses again to form the protostar. Star formation processes include both the structure evolution as mentioned above, and molecular evolution. Rates of chemical reactions depend on various physical quantities. Among them, the temperature is crucial; e.g., rates of ice sublimation and reactions with potential barriers are proportional to $\\exp(-\\Delta E/kT)$, where $\\Delta E$ represents desorption energy or a height of an energy barrier. In addition, the molecular evolution in star forming regions is a non-equilibrium process. Therefore, accurate temperature determination via radiation hydrodynamic (RHD) simulations is important to investigate molecular evolution in star forming cores. There are a few previous works which combined a chemical reaction network model with a RHD model of star-forming cores. Aikawa et al. (2008, here after AW08) investigated chemistry from a molecular cloud core to a protostellar core adopting a spherically symmetric RHD model \\citep{masunaga00}, and mainly discussed the molecular evolution that occurs at $r \\gtrsim 10$ AU, where $T \\lesssim 300$ K. The spherical symmetry, however, should brake down inside the centrifugal radius of $\\sim$100 AU, where circumstellar disks will appear. \\citet{weeren09} adopted an axis symmetric RHD model \\citep{yorke99}, and investigated molecular evolution mainly at $T < 100$ K from a collapsing molecular cloud cores to a forming circumstellar disk. The spatial resolution of their physical model is $\\sim$3 AU in the central region. So the central first core and/or protostar is not resolved. In this paper, we combine a detailed chemical reaction network model of gas-phase and grain-surface chemistry with a three dimensional RHD simulation of a star-forming core for the first time, and investigate molecular evolution at $T$ = 10--2000 K. So far, a few groups have succcessfully performed three-dimensional R(M)HD simulations of a low mass star-forming core up to the first core stage \\citep{white06,commercon10,tomida10a, tomida10b}. On the other hand, evolution beyond first cores in three dimension is more challenging and investigated only very recently by \\citet{bate10, bate11}, \\citet{tomida12a}, and \\citet{tomida12b}. So we focus on chemistry in the first core stage in the present work. Our motivations are twofold. Firstly, first cores are good targets for Atacama Large Millimeter/submillimeter Array (ALMA). Although over 40 years have passed since \\citet{larson69} predicted the presence of first cores, they have not yet been detected. Observation of first cores is challenging; they are buried in dense envelopes, and have very compact structure ($\\sim$10 AU) and short lifetime (a few 1000 yr). Recently, some candidates have been reported: IRS2E in L1448 \\citep{chen10}, Per-Bolo 58 \\citep{enoch10} and L1541-mm \\citep{pineda11}. Observational properties of these objects, low bolometric luminosity $L < 0.1$ $L_{\\bigodot}$ and cold spectral energy distribution, are consistent with the theoretical predictions of first cores \\citep[e.g.,][]{masunaga98,saigo11}. On the other hand, IRS2E has a high velocity outflow ($\\sim$25 km/s), and Per-Bolo 58 is considerably luminous at 24 $\\mu$m and associated with a well-collimated outflow \\citep{dunham11}, which contradict the theoretical prediction that the outflow from first cores is slow ($\\lesssim$5 km/s) and not well collimated \\citep{machida08}. Confirmation and further investigations of the candidates require observation of molecular emission lines. Therefore, chemical models of first cores are highly desired. Secondly, recent hydrodynamic simulations have shown that the outer region of first cores might directly evolve to circumstellar disks, while the central region collapses to form protostars \\citep{saigo08,machida10,machida11,bate10,bate11}. In that case, the chemical composition of the first cores would be the initial composition of the circumstellar disks. The Molecular evolution during the formation and evolution of the circumstellar disks is still an open question. Although most chemical models of the circumstellar disks have used abundances of the molecular cloud cores as initial abundances, its validity is questioned \\cite[][]{weeren09,visser11}. In other words, there remains a missing link between chemistry in the molecular cloud cores and the circumstellar disks. Our detailed three-dimensional chemical-hydrodynamical model from the molecular cloud cores to the first cores is the first step to fill the gap. The rest of the paper is organized as follows. In Section 2, we briefly describe our RHD simulations and chemical reaction network model. In Section 3, we discuss trajectories of fluid parcels in our RHD simulations. We show molecular evolutions in the assorted fluid parcels, and the spatial distributions of molecular abundances in the first core stage. In Section 4, we discuss molecular column densities and uncertainties in our reaction network model. We compare our model with models of hot corinos in terms of physical and chemical properties. We also discuss the effect of accretion shocks on chemistry, and chemistry after the first core stage. We summarize our results in Section 5. ", "conclusions": "We have investigated molecular evolution in the first hydrostatic core stage. We performed the three-dimensional radiation hydrodynamic simulation from the molecular cloud core to the first core. We traced the fluid parcels, which follow the local flow of the gas, and solved molecular evolution along the trajectories, coupling gas-phase and grain-surface chemistry. We derived the spatial distributions of molecular abundances and column densities in the molecular cloud core harboring the first core. Our conclusions are as follows. 1. We found that the total molecular abundances of gas and ice of many species in the cold era (10 K) remain unaltered until the temperature reaches $\\sim$500 K. The gas abundances in the warm envelope and outer layers of the first core ($T \\lesssim 500$ K) are mainly determined via the sublimation of ice-mantle species. Above 500 K, the abundant molecules, such as H$_2$CO, start to be destroyed mainly via collisional dissociation or reactions with H atom, and converted to the simple molecules, such as CO, H$_2$O and N$_2$. On the other hand, some molecules are effectively formed; carbon-chains, such as C$_2$H$_2$ and cyanopolyynes, are formed at the temperature of $>$700 K. 2. Gaseous large organic molecules are associated with the first core ($r < 10$ AU). Although the abundances of large organic molecules in the first core stage are comparable or smaller than in the protostellar hot corino, their column densities in our model are comparable to the observed values toward the prototypical hot corino, IRAS 16293-2422. We propose that large organic molecules can be good tracers of the first cores. 3. The shock wave occurs when the envelope material accretes onto the first core. We found that the shock heating has little effect on chemistry, if the maximum temperature in the shock heating layer is below 1000 K. Even if the gas temperature reaches 2000 K, the heating has a limited effect on chemistry. While it significantly decreases the H$_2$CO abundance and increases the C$_2$H$_2$ abundance in the outer parts of the first core ($r \\lesssim 3$ AU), it has little effect on other species. 4. Recent hydrodynamic simulations have shown that outer region of the first core could directly evolve to the circumstellar disk. In that case, the chemical composition of the first core corresponds to the initial composition of the circumstellar disk; the total abundances of gas and ice at $r \\gtrsim 3$ AU ($T \\lesssim 500$ K) are mostly determined in the cold molecular cloud cores, while inner parts ($T \\gtrsim 500$ K) are dominated by simple molecules, such as CO, H$_2$O and N$_2$." }, "1207/1207.2529_arXiv.txt": { "abstract": "We investigate the properties of a dark matter sector where supersymmetry is a good symmetry. In this context we find that the stability of the dark matter candidate is possible even when R-parity is broken in the visible sector. In order to illustrate the idea we investigate a simple scenario where the dark matter candidate is the lightest scalar field in the dark sector which annihilates mainly into two sfermions when these channels are available. We study the relic density constraints and the predictions for the dark matter detection experiments. ", "introduction": "The possibility to describe the properties of the cold dark matter in the universe using a candidate in various particle physics scenarios has been studied for a long time. For a review of different candidates see Ref.~\\cite{Drees:2012sz}. One of the most popular dark matter candidates is the lightest supersymmetric particle in SUSY theories. In this type of scenario typically one considers the lightest neutralino ~\\cite{Goldberg:1983nd,Ellis:1983ew,Jungman:1995df} or the gravitino~\\cite{Buchmuller:2009fm} as dark matter candidates. Both candidates have been investigated in great detail by many experts in the field. Unfortunately, in these models one has a large number of free parameters and it is difficult to make unique predictions which can be tested in current or future dark matter experiments. It is well-known that in order to guarantee the stability of the lightest neutralino in supersymmetric models the so-called R-parity symmetry is assumed. The case of the gravitino is different because its lifetime can be large enough even if R-parity is broken~\\cite{Buchmuller:2009fm}. The possibility to understand the origin of R-parity conservation has been investigated by many groups. However, the simplest way to study this issue is to consider the B-L extensions of the Minimal Supersymmetric Standard Model (MSSM) where after symmetry breaking one can obtain R-parity as a symmetry at the low-scale. See Refs.~\\cite{Krauss:1988zc,Martin:1992mq,Aulakh:1999cd, Aulakh:2000sn,Babu:2008ep,Feldman:2011ms,Perez:2011zx} for the study of this problem in some supersymmetric scenarios and Refs.~\\cite{Basso:2012gz,O'Leary:2011yq,FileviezPerez:2011kd} for recent phenomenological studies of these models. Unfortunately, even if in these scenarios we can understand dynamically the origin of R-parity conservation it is difficult to make interesting predictions for dark matter experiments since we can have several dark matter candidates, the neutralinos or right-handed sneutrinos, and as in the MSSM there are many free parameters. In this Letter we investigate the properties of a dark matter sector where supersymmetry is a good symmetry before the breaking of the gauge symmetry. In this context we do not need to impose a discrete symmetry to guarantee the stability of the dark matter candidate and even if R-parity is broken in the visible sector he dark matter candidate is stable. To study this idea of having a supersymmetric sector we consider a simple scenario where in the visible sector we have the minimal B-L extension of the MSSM~\\cite{Barger:2008wn} and in the dark sector we have two chiral superfields with B-L quantum numbers. Here the link between the visible and dark sector is defined by the B-L gauge force which is broken in the visible sector by the vacuum expectation value (VEV) of the right-handed sneutrinos. We find that after the B-L breaking a mass splitting is induced in the dark sector and the lightest field is the only possible candidate for the cold dark matter in the universe. In this model the dark matter candidate annihilates mainly into two sfermions when these channels are available. We investigate the different scenarios where we can achieve the observed dark matter relic density and the possible predictions for dark matter experiments. We find that the current bounds from the Xenon100 experiment set strong constraints on this type of models where the elastic dark matter nucleon cross section is through a neutral gauge boson. This article is organized as follows: In Section II we define a simple scenario with a supersymmetric dark matter sector. In Section III we show the possible scenarios where one can achieve the relic density observed by the experiments. The constraints coming from the direct detection experiments are investigated in Section IV, while we summarize the main results in Section V. ", "conclusions": "In this Letter we have investigated a simple scenario for the cold dark matter in the universe where the sector responsible for dark matter has ``exact\" supersymmetry before symmetry breaking. In order to achieve this type of scenario we assume that supersymmetry breaking is mediated as in ``gauge mediation\", where the messenger fields only have quantum numbers of the visible sector, and the soft terms induced by gravity are very small. The SM singlet fields in the dark sector do not get large soft terms from gravity mediation and we can say that supersymmetry is a good symmetry in the dark sector. In order to illustrate our idea we consider the case where in the visible sector we have the simplest B-L extension of the minimal supersymmetric standard model while the dark sector is composed of two chiral superfields with B-L quantum numbers. We have found that in this case the dark matter candidate is the lightest scalar field in the dark sector and the B-L D-term induces a mass splitting after the symmetry is broken. We noticed that the dark matter candidate is stable even if R-parity is spontaneously broken in the visible sector. Since the link between the visible and dark sectors is through the B-L gauge force, the dark matter annihilates mainly into two sfermions when these channels are available. We have shown the allowed parameter space by the relic density and direct detection experiments in simplified scenarios where the annihilation is mainly into two sleptons. In the case when the dark matter candidate is below 100 GeV, the DM annihilation is mainly into two fermions at the one-loop level where inside the loops you have the sfermions and gauginos. The details of the scenario for light dark matter and the annihilation into photons will be investigated in a future publication~\\cite{Long-paper}. \\vspace{1.0cm} {\\textit" }, "1207/1207.4555_arXiv.txt": { "abstract": "We use numerical simulations to investigate how the statistical properties of dark matter (DM) haloes are affected by the baryonic processes associated with galaxy formation. We focus on how these processes influence the spin and shape of a large number of DM haloes covering a wide range of mass scales, from galaxies to clusters at redshifts zero and one, extending to dwarf galaxies at redshift two. The haloes are extracted from the OverWhelmingly Large Simulations (OWLS), a suite of state-of-the-art high-resolution cosmological simulations run with a range of feedback prescriptions. We find that the median spin parameter in DM-only simulations is independent of mass, redshift and cosmology. At $z = 0$ baryons increase the spin of the DM in the central region ($\\le 0.25 \\, r_{200}$) by up to 30 per cent when feedback is weak or absent. This increase can be attributed to the transfer of angular momentum from baryons to the DM, but is no longer present at $z = 2$. We also present fits to the mass dependence of the DM halo shape at both low and high redshift. At $z=0$ the sphericity (triaxiality) is negatively (positively) correlated with halo mass and both results are independent of cosmology. Interestingly, these mass-dependent trends are markedly weaker at $z=2$. While the cooling of baryons acts to make the overall DM halo more spherical, stronger feedback prescriptions (e.g. from active galactic nuclei) tend to reduce the impact of baryons by reducing the central halo mass concentration. More generally, we demonstrate a strongly positive (negative) correlation between halo sphericity (triaxiality) and galaxy formation efficiency, with the latter measured using the central halo baryon fraction. In conclusion, our results suggest that the effects of baryons on the DM halo spin and shape are minor when the effects of cooling are mitigated, as required by realistic models of galaxy formation, although they remain significant for the inner halo. ", "introduction": "\\label{shapesintro} A natural consequence of the standard hierarchical structure formation paradigm is that the shapes of dark matter haloes are triaxial, a property that is inherited from their progenitor density perturbations \\citep{bib:Bardeen86}. This additionally leads to aspherical growth as the halo accretes matter from preferential directions, associated with the surrounding sheets and filaments. The anisotropic accretion history of a halo also affects its angular momentum distribution, through the presence of non-zero torques. It is therefore clear that both the shape and spin of a dark matter halo are important diagnostics for an accurate determination of their structure and formation history. While the spin and shape of a dark matter halo are not directly observable, they have important consequences for the structure and dynamics of galaxies. For example, halo spin is an important parameter in galaxy formation models as it affects the size of the embedded galactic disc (e.g. \\citealt{bib:Mo98}, but see \\citealt{bib:Sales12}). Deviations from axisymmetry in elliptical galaxies are likely to influence the gas kinematics of the system \\citep{bib:deZeeuw89}, and may be responsible for exciting or sustaining warps and stabilising or deforming polar rings \\citep{bib:Steiman92}. Axisymmetry may also influence the fuelling efficiency of the central black hole \\citep{bib:Franx91}. Misalignment of the angular momentum of the halo and the galaxy may be responsible for the anisotropic distribution of subhaloes and satellite galaxies (\\citealt{bib:Holmberg74, bib:Knebe04, bib:Kang05, bib:Libeskind05, bib:Zentner05, bib:Libeskind07, bib:Knebe10}) and could cause galactic warps (\\citealt{bib:Ostriker89, bib:Debattista99, bib:Bailin04}). On cluster scales, asphericity in the dark matter halo will naturally correspond to asphericity in the gas density and will impact the shape of X-ray isophotes and the Sunyaev-Zel'dovich signal. Understanding the intrinsic shape of dark matter haloes is also important for weak lensing analysis (see, for example, the discussions in \\citealt{bib:Becker11, bib:Bett12}) and it is well known that intrinsic ellipticity can contribute significantly to a lensing halo's ability to form arcs \\citep{bib:Oguri03}. Several methods are used to constrain galaxy and halo shapes observationally (see, for example, \\citealt{bib:Sackett99, bib:Merrifield04}). Unfortunately studies performed to date do not yet reveal a consistent picture (see the discussion in \\citealt{bib:Obrien10}). The observations cover a large range of systems and vary in the extent to which the halo is probed, making a direct comparison somewhat difficult. Whether the discrepancies result from halo-to-halo scatter or from systematic errors in the observed estimates is unclear. However, given the rapidly accumulating number of data sets, ever increasing sophistication of the data analysis tools and the development of more realistic mock observations from simulations, one can soon expect the situation to change substantially. Theoretically, the predictions for the distribution of angular momentum and halo shapes have been studied extensively, primarily using cosmological $N$-body simulations (\\citealt{bib:Frenk88}; \\citealt{bib:Dubinski91}; \\citealt{bib:Warren92}; \\citealt{bib:Cole96}; \\citealt{bib:Bullock02}; \\citealt{bib:Jing02}; \\citealt{bib:Bailin05}; \\citealt{bib:Allgood06}; \\citealt{bib:Bett07}; \\citealt{bib:Maccio08}; \\citealt{bib:JeesonDaniel11}; \\citealt{bib:Vera-Ciro11}; \\citealt{bib:Bett12}; \\citealt{bib:Zemp12}). There is a general consensus that cold dark matter haloes have approximately log-normal spin distributions and are triaxial, with sphericities, $(c/a) \\simeq 0.5-0.8$ \\footnote{Where $a > b > c$ are the eigenvalues of the halo's inertia tensor.} and elongations, $(b/a)\\simeq 0.4-1$. Furthermore, haloes are generally found to be highly flattened and show a tendency toward prolate shapes $(c/b > b/a)$, especially in the inner regions. There is also general agreement that the sphericity decreases with increasing halo mass and that the spin is independent of mass (\\citealt{bib:Bullock02, bib:Jing02, bib:Springel04, bib:Hopkins05, bib:Bett07, bib:Maccio08, bib:JeesonDaniel11}). In addition, \\cite{bib:JeesonDaniel11} have shown that sphericity is strongly correlated with concentration, while both triaxiality and spin are anti-correlated with concentration. While these results are interesting, a significant uncertainty is how the dark matter halo is affected by the additional, non-gravitational processes acting on the baryons. Recent work has clearly established that the condensation of baryons to the centre of dark matter haloes tends to increase the central angular momentum of the halo (see for example, \\citealt{bib:Sharma05, bib:Tonini06a, bib:Kaufmann07, bib:Abadi10, bib:Bett10}) and to make the halo more spherical or axisymmetric (see, for example, \\citealt{bib:Katz91, bib:Dubinski94, bib:Evrard94, bib:Barnes96, bib:Tissera98,bib:Springel04,bib:Kazantzidis04, bib:Debattista08, bib:Pedrosa10, bib:Tissera10, bib:Bryan11,bib:Zemp12}). This result has been used to explain the discrepancy between the strongly-prolate triaxial shape found in $N$-body simulations and the more spherical central regions of observed systems. Incorporating baryonic physics in cosmological simulations is a non-trivial task and the computational cost of this process has placed limits on both the parameter space and the size of the sample of haloes explored to date. The detailed nature of the baryonic processes involved in galaxy formation and the precise influence of these processes on galaxies therefore remains largely uncertain. In this paper, we attempt to make progress on both fronts, by studying the spin and shape distributions for a large ($>1000$) sample of dark matter haloes, spanning a range of mass (from dwarf galaxies to clusters) and redshift ($z=0,1$ and $2$). We do this using the OverWhelmingly Large Simulations (OWLS; \\citealt{bib:Schaye10}) -- a suite of cosmological hydrodynamical simulations run with many different physical prescriptions for the baryons. By providing identical simulations run with different implementations of the subgrid physics, OWLS offers the opportunity to explore the effects of baryons under a range of physical conditions, for the same population of haloes. In particular, we use the OWLS data to quantify, in a statistically meaningful way, the influence of feedback processes (from no feedback, to feedback from stars and black holes) on the spin and shape distributions of dark matter haloes. The paper is organised as follows. The simulations used in this analysis, and the methods used to identify haloes and to estimate their spins and shapes are outlined in section \\ref{simshapes}. Our results are presented in section \\ref{results}, including fitting formulae for the predicted correlations between halo shape and mass. The robustness of our results is demonstrated via a resolution study, given in the appendix. Finally, we summarise our main results in section \\ref{summaryshapes}. ", "conclusions": "\\label{summaryshapes} In this paper, we have exploited a subset of runs from the OverWhelmingly Large Simulations (OWLS; \\citealt{bib:Schaye10}) to investigate the impact of baryons (through gas cooling, star formation and feedback from stars and black holes) on the spin and shape of dark matter haloes. Our results allow statistically meaningful conclusions to be drawn regarding the impact of baryons on these properties, due to the large number of haloes spanning a wide dynamic range in mass. We have also checked whether our results depend on cosmology and redshift. Our main results are summarised below. \\begin{enumerate} \\item The spin distribution of dark matter haloes in simulations without baryons is characterised by a log-normal curve, with best-fit values of $\\lambda_0= 0.036$ $(0.038)$ and $\\sigma = 0.62$ $(0.60)$ at $z = 0$ (2), in agreement with previous work \\citep{bib:Bullock01,bib:Bailin05,bib:Bett07,bib:Maccio08}. The distribution is very similar for the {\\it WMAP}\\,1, {\\it WMAP}\\,3 and {\\it WMAP}\\,5 cosmologies, suggesting that there is no strong dependence on $\\sigma_8$ (the parameter that varies the most between the three models). No significant dependence of spin with mass is seen, both at $z=0$ and $z=2$. At $z = 0$ the spin parameter remains essentially unchanged if computed using only mass within the central region ($0.25r_{200}$), as found by \\cite{bib:Bailin05}. However, at $z = 2$ the inner region of haloes has a higher mean spin than that computed over the whole halo. Restricting our sample to relaxed haloes causes a small (10-15 per cent) decrease in the mean value (in agreement with \\citealt{bib:Maccio06} and \\citealt{bib:JeesonDaniel11}). \\item At $z = 0$ the spin distribution of dark matter haloes extracted from the baryon runs is not significantly different to that of dark matter only haloes when computed using all dark matter particles within $r_{200}$. However, in the central regions ($0.25r_{200}$), where baryons are expected to play an important role, haloes in runs with absent or weak stellar feedback tend to have higher median spin values than those from stronger feedback runs (which are very similar to the dark matter only case). We showed that this is, at least in part, due to the transfer of angular momentum from the baryons to the dark matter in the former runs. At $z = 2$ the baryon runs exhibit slightly lower median spin values than the dark matter only case, an effect that is likely due to the increased circular velocity in weak feedback runs and decreased specific angular momentum within the central regions in the strong feedback runs. \\item Dark matter only haloes extracted from OWLS typically have sphericities of $\\sim 0.5 $ to $0.6$ and triaxialities of between 0.6 and 0.8 (indicating triaxial to prolate shapes) over the mass and redshift ranges we have explored. More massive haloes have less spherical and more prolate shapes. Again, we find that halo shape is insensitive to the choice of cosmological model. Galaxies and groups at $z = 2$ show the same trends with mass as the groups and clusters at $z = 0$, but weaker. \\item When baryons are included, we find that the mass dependent trends remain, and that the intercepts of the relation between sphericity (triaxiality) and mass slowly increase (decrease) with increasing galaxy formation efficiency. At $M_{200}=10^{12}h^{-1}{\\rm M}_{\\odot}$, baryons increase the dark matter shape parameters by around 10 per cent in the most extreme case (no feedback). A similar result is seen at higher redshift. Larger differences are again seen when we consider only the central regions of the halo. \\end{enumerate} In conclusion, we find that the baryons have a very minor effect on the spin and overall shape of the entire dark matter halo when the feedback is strong enough to match observed stellar mass fractions. In particular, the model with AGN feedback can reproduce several observational properties on galactic and groups scale at $z = 0$ by removing gas, suppressing the baryonic impact on the dark matter halo shape. It should therefore be safe to assume results from dark matter only simulations when considering the overall halo properties, at least on the scales resolved by our simulations. However, even when feedback is strong we find that baryons have a significant effect on the shape of the {\\it inner} halo." }, "1207/1207.6285_arXiv.txt": { "abstract": "{Several young supernova remnants (SNRs) have recently been detected in the high-energy (HE; 0.1 $<$ E $<$ 100 GeV) and very-high-energy (VHE; E $>$ 100 GeV) gamma-ray domains. As exemplified by \\rxj, the nature of this emission has been hotly debated, and direct evidence for the efficient acceleration of cosmic-ray protons at the SNR shocks still remains elusive.} {We study the broadband gamma-ray emission from one of these young SNRs, namely \\rcw, for which several observational lines of evidence indirectly point towards the presence of accelerated hadrons. We then attempt to detect any putative hadronic signal from this SNR in the available gamma-ray data, in order to assess the level of acceleration efficiency.} {We analyzed more than 40 months of data acquired by the Large Area Telescope (LAT) on-board the {\\emph{Fermi Gamma-Ray Space Telescope}} in the HE domain, and gathered all of the relevant multi-wavelength (from radio to VHE gamma-rays) information about the broadband nonthermal emission from \\rcw. For this purpose, we re-analyzed the archival \\xray~data from the \\asca/Gas Imaging Spectrometer (GIS), the \\xmm/EPIC-MOS, and the \\rxte/Proportional Counter Array (PCA).} {Beyond the expected Galactic diffuse background, no significant gamma-ray emission in the direction of \\rcw~is detected in any of the 0.1--1, 1--10 and 10--100 GeV \\fermi-LAT maps. The derived HE upper limits, together with the \\hess~measurements in the VHE domain, are incompatible with a standard E$_{p}^{-2}$ hadronic emission arising from proton-proton interactions, and can only be accommodated by a spectral index $\\Gamma \\leq$ 1.8, i.e. a value in-between the standard (test-particle) index and the asymptotic limit of theoretical particle spectra in the case of strongly modified shocks. In such a hadronic scenario, the total energy in accelerated particles is at the level of $\\eta_{CR}$ = E$_{{\\rm CR}}$/E$_{{\\rm SN}}$ $\\sim$ 0.07 d$^{2}_{{\\rm 2.5kpc}}$/\\neff$_{{\\rm cm-3}}$ (with the distance d$_{\\rm 2.5kpc}$ $\\equiv$ d/2.5 kpc and the effective density \\neff$_{{\\rm cm-3}}$ $\\equiv$ \\neff/1 cm$^{-3}$), and the average magnetic field must be stronger than 50 $\\mu$G in order to significantly suppress any leptonic contribution. On the other hand, the interpretation of the gamma-ray emission by inverse Compton scattering of high energy electrons reproduces the multi-wavelength data using a reasonable value for the average magnetic field of 15--25 $\\mu$G. In this leptonic scenario, we derive a conservative upper limit to $\\eta_{CR}$ of 0.04 d$^{2}_{{\\rm 2.5kpc}}$/\\neff$_{{\\rm cm-3}}$. We discuss these results in the light of existing estimates of the magnetic field strength, the effective density and the acceleration efficiency in \\rcw.} {} ", "introduction": "\\label{intro} Supernova remnants (SNRs) are thought to be the primary sources of the bulk of Galactic cosmic-ray (CR) protons observed at Earth, up to the knee energy at $\\sim$ 3~PeV. This paradigm mainly relies on the need to have sufficiently energetic sources that could provide the necessary power to maintain the Galactic CR energy density \\citep[\\eg][]{fields01}, and on the dominance of SNRs among sources of non-thermal radio emission. Our understanding of CR acceleration in SNRs mainly relies on the so-called Diffusive Shock Acceleration theory \\cite[in its non-linear version NLDSA, see][]{malkov01}, which is commonly invoked to explain several observational (though, indirect) lines of evidence for efficient particle acceleration at the SNR forward shocks up to very high energies. Among the requisites that must be fulfilled by this theory \\cite[see][for a recent review]{blasi10}, the observed broadband nonthermal emission of individual SNRs is of particular interest. This emission, arising from accelerated particles that emit photons through several channels (synchrotron [SC], inverse-Compton [IC], nonthermal bremsstrahlung, proton-proton interactions and subsequent $\\pi^0$~decay), offers new insights into the particle acceleration mechanisms at work in these sources, given the large number of multi-wavelength observations available nowadays towards many Galactic SNRs. In particular, recent observations of young SNRs in the high-energy (HE; 0.1 $<$ E $<$ 100 GeV) and very-high-energy (VHE; E $>$ 100 GeV) gamma-ray domains have raised several questions and triggered numerous theoretical investigations \\citep[\\eg][]{za10,berezhko10,fang11,caprioli11}. The critical issue regarding the nature of the observed gamma-ray emission has been intensively discussed in the literature in recent years \\citep[][and references therein]{ellison10}. Nevertheless, these joint HE/VHE observations can help us to discriminate spectrally between the leptonic (through IC emission) and hadronic (through $\\pi^0 \\rightarrow $2$\\gamma$) scenarios, as shown by the \\fermi-LAT \\citep{abdo11a} and \\hess~\\citep{hess07} observations towards the TeV-bright SNR \\rxj, which tend to support a leptonic model for the observed gamma-ray emission \\citep{ellison12,li11}. \\rcw~\\citep{rodgers60}, also known as G315.4$-$2.3 or \\msh~\\citep{mills61}, is a 42\\m~diameter Galactic SNR in the southern sky. There has been much controversy about the nature of the supernova (SN) progenitor, and the SNR age and distance. This has resulted from the difficulties in reconciling the young age of \\rcw, based on its connection with the first historical SN ever recorded in AD~185, with its large size, given the relatively large distance estimates, which place \\rcw~at 2.3--2.8 kpc \\citep{rosado96,sollerman03} near an OB stellar association. \\citet{williams11} recently reviewed all the arguments about the nature of the SN progenitor and critically examined the available observations of \\rcw. They suggest that the SNR is likely the remnant of a type Ia SN, and, from hydrodynamic simulations, that the off-center explosion occurred in a low-density cavity carved by the progenitor system \\citep[see also][]{vink97}. This scenario allows one to explain at once the young age and the large distance of \\rcw, which we scale in terms of d$_{2.5}$ $\\equiv$ d/(2.5 kpc) throughout the paper. The general outline of its nearly circular shell is similar in the radio \\citep{whiteoak96,dickel01}, infrared \\citep[IR,][]{williams11}, optical \\citep[\\eg][]{smith97}, and \\xray~\\citep{vink97,bamba00,bocchino00,borkowski01b,rho02,vink06} domains, although significant fine-scale differences have been reported \\cite[\\eg][ and Fig. \\ref{fig2}]{rho02}. \\xray~observations towards \\rcw~have revealed the presence of both thermal and nonthermal emission, with very distinct morphologies: while soft \\xrays~correlate with optical emission from non-radiative shocks and IR emission from collisionally heated dust, the continuum hard \\xray~emission is seen at the edges of radio emission. The high-temperature plasma, which mostly contains the strong Fe K$_{\\alpha}$ line emission, also shows a particular morphology extending towards the SNR interior, as revealed with \\suz, and is thought to originate from Fe-rich ejecta heated by the reverse shock \\citep{ueno07,yamaguchi08}. The global distribution of the Fe-rich plasma in the entire SNR has recently been mapped with \\suz~\\citep{yamaguchi11}, and the total Fe mass of $\\sim$ 1 \\msun~deduced from these observations, together with a relatively low ejecta mass of 1--2 \\msun, strengthens the scenario of a type Ia SN at the origin of \\rcw. Physical conditions (shock speed, ambient density, and magnetic field) vary greatly along the shell-like structure. In particular, slow shocks \\citep[$\\sim$ 600--800 km s$^{-1}$, see][]{long90,ghavamian01} and relatively high post-shock densities \\citep[$\\sim$ 2 cm$^{-3}$, see][]{williams11} have been measured in the southwest (SW) and northwest (NW) regions, while the northeast (NE) region exhibits much faster shocks \\citep[$\\gtrsim$ 2700 km s$^{-1}$ and 6000 $\\pm$ 2800 km s$^{-1}$, see][]{vink06,helder09} and lower densities \\citep[$\\sim$ 0.1--0.3 cm$^{-3}$, \\eg][]{yamaguchi08}. In this region, \\citet{helder09} argued that at least 50\\% of the total pressure is induced by CRs, based on the discrepancy between the measured shock velocity and the spectroscopically determined post-shock proton temperature \\citep[see also][]{vink10}. Moreover, \\citet{vink06} found evidence for a concave electron spectrum, as predicted by the NLDSA theory, in order to explain the radio and \\xray~SC emission observed from the same region. These two measurements, together with the recent detection of \\rcw~with the \\hess~experiment in the VHE domain \\citep{hess09}, seem to point towards an efficient CR source. However, complementary HE observations are needed to probe the nature of the gamma-ray emission, as discussed above. We here report on \\fermi-LAT observations and data analysis towards \\rcw~(section \\ref{lat}). In an attempt to constrain the broadband nonthermal emission from the SNR, we present a re-analysis of the archival \\asca/GIS, \\xmm/EPIC-MOS, and \\rxte/PCA \\xray~data (section \\ref{xray}), and collect the available information in the radio and VHE gamma-ray domains. We then discuss the constraints derived on the parent particle spectrum and on the subsequent acceleration efficiency, in the light of existing estimates (section \\ref{discu}). ", "conclusions": "\\label{conclu} The \\fermi-LAT upper limits in the HE domain derived in this work for the young SNR \\rcw~provide strong constraints on the injection spectrum of the primary population responsible for the extended VHE emission. A hadronic scenario can only reproduce the multi-wavelength data using a hard proton spectrum (spectral index $\\Gamma \\leq$ 1.8) and a total energy injected into hadrons residing in \\rcw~of $\\sim$ 7 $\\times$ 10$^{49}$(\\neff/1 cm$^{-3}$)$^{-1}$ d$_{2.5}^2$ erg. Given that the one-zone hadronic model suffers from several limitations (incompatible radio spectral index, very low $K_{ep}$ and extremely high B-field), only a two-zone model can fulfill all the observational constraints, though still with a low $K_{ep}$. In the one- and two-zone leptonic scenarios, the multi-wavelength data can be closely reproduced using electron spectral indices of $\\sim$ 2.0--2.3, a total energy injected in electrons of $\\sim$ 2 $\\times$ 10$^{49}$ erg, and a reasonable average magnetic field of 15--25 $\\mu$G. In such a case, the most conservative upper limit to the total energy injected into hadrons inside \\rcw~amounts to $\\sim$ 4 $\\times$ 10$^{49}$(\\neff/1 cm$^{-3}$)$^{-1}$ d$_{2.5}^2$ erg. These estimates of the total CR energy content in \\rcw~can be translated into CR pressure $P_{\\rm tot,CR}$ = E$_{\\rm tot,CR}$/3\\Veff, where \\Veff~is the effective CR volume within the SNR. Assuming that this volume is given by the best-fit parameters of the shell morphology observed with \\hess~\\citep[though not statistically significant over a uniform sphere, see][]{hess09}, we obtain \\Veff~= (6 $\\pm$ 2) $\\times$ 10$^{59}$ d$_{2.5}^{3}$ cm$^3$, and $P_{\\rm tot,CR}$ must be conservatively lower than 3.8$^{+2.0}_{-1.0}$ $\\times$ 10$^{-10}$ (E$_{\\rm tot,p}$/7 $\\times$ 10$^{49}$ erg) (\\neff/0.1 cm$^{-3}$)$^{-1}$ d$_{2.5}^{-1}$ erg cm$^{-3}$. This CR pressure, derived from the modeling of the broadband nonthermal emission of the {\\it entire} SNR (see section \\ref{discu}), is very close to the CR pressure found by \\citet{helder09} {\\it in the NE region}, and later confirmed by \\citet{vink10}, P$_{\\rm NE,CR}$ $\\gtrsim$ 3.7 $\\times$ 10$^{-10}$ (\\neff/0.1 cm$^{-3}$) (kT/2.3 keV) erg cm$^{-3}$. If we were to apply the same comparison for the other regions along the \\rcw~shell, which all exhibit higher plasma densities \\citep[0.5--2 cm$^{-3}$, see][]{yamaguchi11,williams11}, our upper limit to $P_{\\rm tot,CR}$ would most likely be inconsistent with the \\citet{helder09} estimate regarding the fractional CR pressure. However, the large uncertainties in the effective CR volume in the above calculations, among others, prevent us from drawing firm conclusions about the acceleration efficiency in \\rcw." }, "1207/1207.3375_arXiv.txt": { "abstract": "We present the accretion of a phantom scalar field into a black hole for various scalar field potentials in the full non-linear regime. Our results are based on the use of numerical methods and show that for all the cases studied the black hole's apparent horizon mass decreases. We explore a particular subset of the parameter space and from our results we conclude that this is a very efficient black hole shrinking process because the time scales of the area reduction of the horizon are short. We show that the radial equation of state of the scalar field depends strongly on the space and time, with the condition $\\omega = p/\\rho>-1$, as opposed to a phantom fluid at cosmic scales that allows $\\omega < -1$. ", "introduction": "In cosmological models at present time the universe is assumed to have various ingredients of exotic nature like the dark energy candidates, among which we find the phantom scalar field \\cite{Corasaniti}. The reason is that supernovae Ia data allow an equation of state of the dark energy component with $\\omega = p/\\rho <-1$, where $p$ and $\\rho$ are the pressure and the energy density of the fluid, which is a condition that a phantom scalar field satisfies at cosmic scales. This is a reason to start up an exploration of the consequences such type of scalar field might have at local scales, for instance on black holes. On the other hand, the behavior of the area of the horizon when the black holes is accreting matter is a very important property of black hole physics, because the accretion of such exotic material may impose important restrictions on the mass of black nowadays black hole candidates. In order to study this process in the full non-linear regime we start up with a model coupling the phantom scalar field and gravity. The model assumes the Lagrangian density of the phantom scalar field is given by \\begin{equation} {\\cal L} = R + \\frac{1}{2}g_{\\mu\\nu}\\partial^{\\mu}\\phi \\partial^{\\nu} \\phi - V(\\phi), \\label{eq:lagrangian} \\end{equation} \\noindent where $R$ is the Ricci scalar of the space-time, $g_{\\mu\\nu}$ is the space-time metric, $\\phi$ is a scalar field and $V(\\phi)$ is the potential of the field. The property defining a phantom scalar field is that the relative sign between the Ricci scalar and the kinetic term are the same. When the action constructed with such a Lagrangian density is varied with respect to the metric, the arising Einstein's equations are related to a stress energy-tensor that violates the null energy condition, that is $T_{\\mu\\nu}k^{\\mu}k^{\\nu} < 0$, where $k^{\\mu}$ is a null vector. The immediate implication of this violation is the violation of the weak energy condition too, which in turn implies that observers following time-like trajectories might measure negative energy densities. Although this property is at odds with the nowadays physics observed in laboratories, there are cosmological observations indicating the presence of such kind of matter\\cite{Corasaniti}. In this paper we explore the implications of this at astrophysical scales. The fact that the scalar field violates the null energy condition motivates the study of possible unusual implications in astrophysical scenarios, because the area increasing theorem does not apply in this case. For instance, using exact solutions corresponding to stationary accretion of a test phantom fluid, it was found that the mass of black holes decreases in a phantom energy dominated universe approaching the big rip \\cite{Babichev}. It has also been studied recently the behavior of a black hole apparent horizon (AH) in a FRW background, and the conditions under which a naked singularity can be formed due to the coincidence of the AH of the black hole and the cosmic horizon \\cite{Faraoni}; in such case the authors consider the effects of the back reaction of the scalar field on the space-time metric and indicate that cosmic censorship does not only forbid the existence of naked singularities but also the existence of a phantom field. Assuming that the null energy condition is satisfied the area of the horizon only increase. However, if this condition is violated, the area of the event horizon decreases and the black hole shrinks as shown recently in the full non-linear regime using numerical relativity in \\cite{Guzman}. The present paper is a detailed follow up of previous one, in which we also explore the effects of the scalar field potential on the accretion rates and final state of the black hole. Two important items are presented in this paper: i) the relation $p/\\rho < -1$ is not fulfilled at local scales by the scalar field (at least near to a black hole) although the null energy condition is not satisfied and ii) the accretion of such scalar field reduces the area of a black hole at similar rates for different types of potentials driving the scalar field. This paper is organized as follows. In section \\ref{sec:cauchy} we describe the 3+1 decomposition of the space-time, the evolution system of equations driving the evolution of the geometry and the construction of the initial data. In section \\ref{sec:results} we present the results obtained. Finally in section \\ref{sec:conclusions} we draw some conclusions and comments. ", "conclusions": "\\label{sec:conclusions} We present the full non-linear spherical accretion of a phantom scalar field with various potentials into a black hole. This accretion process is very efficient at reducing the black hole horizon area, which is a potentially important process that may impose restrictions on the existing black hole masses, or for instance primordial black holes. In our simulations we were able to reduce the mass of the initial black hole up to 50$\\%$. Attempts to reduce the area even further presents difficulties with our approach due to the steep gradients in the equations near the origin, and other techniques might be implemented in order to sort this problem out. We show that the final mass of the black hole's horizon is nearly independent of the potential used. This results is physically interesting because the process of area reduction will apply independently of the cosmologically motivated potential used. We show that even though the Lagrangian (\\ref{eq:lagrangian}) of the scalar field corresponds to a phantom field, at local scales the radial equation of state is always $p/\\rho \\ge -1$, and it depends on space and time, with values ranging from a cosmological constant equation of state $\\omega=-1$ up to a stiff mater $\\omega =1$. We want to stress the importance of the equation of state of the scalar field given by the Lagrangian (\\ref{eq:lagrangian}), and point out that perhaps the big rip scenario would strongly depend on the homogeneity and isotropy of the scalar field. We have shown how much the equation of state can depend on space and time for the model of a scalar field, at least in the strong gravitational field regime. The astrophysical possibility of this accretion process strongly depends on the chance of having a phantom scalar field located near a black hole with a rather sharp profile." }, "1207/1207.3719_arXiv.txt": { "abstract": "{We present the results of the \\xmm\\ observations of five hard X-ray emitters: \\jA,\\ \\jD,\\ \\jE,\\ \\jC,\\ and \\jB.\\ The first source is a confirmed supergiant high mass X-ray binary, the following two are candidates supergiant fast X-ray transients, \\jC\\ is a confirmed supergiant fast X-ray transient and \\jB\\ is one of the still unidentified objects discovered with \\inte.\\ The \\xmm\\ observations permitted the first detailed soft X-ray spectral and timing study of \\jA\\ and provided further support in favor of the association of \\jD\\ and \\jE\\ with the supergiant fast X-ray transients. \\jC\\ was not detected by \\xmm,\\ thus supporting the idea that this source reaches its lowest X-ray luminosity ($\\simeq$10$^{32}$~erg/s) around apastron. For \\jB\\ we identified for the first time the soft X-ray counterpart and proposed the association with a close-by radio object, suggestive of an extragalactic origin.} ", "introduction": "\\label{sec:intro} High-mass X-ray binaries (HMXBs) comprise a compact object orbiting a massive OB star. The compact object, usually a neutron star (NS), emits a conspicuous amount of X-ray radiations (up to luminosities of $\\sim$10$^{37}$~erg/s) due to the accretion of matter from the OB companion. Depending on the nature of the companion star, HMXBs can be divided into Be X-ray binaries (BeXBs) and supergiant X-ray binaries (SGXBs). In the former, the NS is in a wide eccentric orbit around a Be star. The X-ray luminosity is generally low ($\\sim$10$^{32}$-10$^{33}$~erg/s) when the NS is far away from periastron and accretes matter from low density regions. Remarkable increases in the X-ray luminosity ($\\Delta_{L_{\\rm X}}$$\\sim$100-1000) are displayed: (i) during the so-called ``Type-II'' X-ray outbursts, which last several orbital cycles and present few (if any) X-ray flux variations associated to the orbital phase; (ii) at the periastron ( ``Type~I'' outbursts), where the compact object is closer to the companion and accretion takes place through the low-velocity and high-density wind of the Be star \\citep{stella86,reig11}. In the SGXBs, the compact object moves around an early-type supergiant in a nearly circular orbit, and in some cases it is well embedded in its dense highly supersonic wind (the so-called ``highly obscured'' SGXBs). The X-ray luminosity is powered by the accretion of the strong stellar wind onto the compact star and displays, on-average, less pronounced changes along the orbit with respect to the BeXBs. Hydrodynamic instabilities in the wind of the supergiant star can give rise to significant density gradients in the environment around the NS, leading to aperiodic variations of the X-ray luminosity ($\\Delta_{L_{\\rm X}}$$\\sim$10-50) on time-scales ranging from few to thousands of seconds \\citep{negueruela10}. The accretion of particularly dense ``clumps'' in the wind can also induce bright short flares that can last for a few hours \\citep[see, e.g.,][and references therein]{kreykenbohm11}. A few peculiar SGXBs, sharing with the latter a remarkable number of similarities \\citep[orbital period, energy spectra, properties of the companion stars; see e.g. discussion in][]{bozzo10}, were discovered in the late 90s \\citep[see e.g.,][]{yamauchi95, smith98}, and collectively termed later ``supergiant fast X-ray transients'' \\citep[SFXTs;][]{sguera06,negueruela06}. At odds with the ``classical'' SGXBs, the SFXTs spend a large fraction of their time \\citep{romano11} in a quiescent state with a typical luminosity of 10$^{32}$-10$^{33}$~erg/s, and only sporadically undergo bright outbursts ($\\Delta_{L_{\\rm X}}$$\\sim$10$^4$-10$^5$) lasting a few hours and reaching peak luminosities comparable with those of the persistent SGXBs \\citep[$L_{\\rm X}$$\\sim$10$^{37}$~erg/s; see e.g.,][]{walter07}. The outbursts of SFXTs are associated to the accretion of particularly dense clumps as in other SGXBs, but the origin of the lower persistent luminosity and much more pronounced variability of these sources is still a matter of debate \\citep{zand05,walter07,grebenev07,bozzo08,bozzo11}. In this paper, we report on the \\xmm\\ observations of four HMXBs discovered with \\inte\\ and \\beppo.\\ Among these sources, IGR\\,J08262$-$3736 is classified as a classical SGXB, IGR\\,J17354$-$3255 and IGR\\,J16328$-$4726 are candidate SFXTs, and SAX\\,J1818.6$-$1703 is a confirmed SFXT. We also report on the \\xmm\\ observation of the still unclassified \\inte\\ source IGR\\,J17348$-$2045. A summary of the previous results for all the above sources is presented in Sect.~\\ref{sec:sources}, while our analysis and results of the \\xmm\\ observations are presented in Sect.~\\ref{sec:data}. Discussions and conclusions on the properties of the soft X-ray emission from the five objects are reported in Sect.~\\ref{sec:discussion}. ", "conclusions": "\\label{sec:discussion} We reported here on the observations of five hard X-ray emitters obtained with \\xmm.\\ Among these, four (\\jA,\\ \\jD,\\ \\jE,\\ \\jC) are HMXBs, whereas \\jB\\ is still unidentified. \\subsection{ \\jA\\ } \\label{sec:jAdiscussion} The \\xmm\\ observation of this source revealed a timing and spectral behaviour that are typical of classical SGXBs. The EPIC-pn lightcurve shows a relatively moderate variability on time scales of hundreds seconds with variations in the X-ray flux of a factor $\\sim$3-4. These variations are usually ascribed to changes in the mass accretion rate due to instabilities and density fluctuations in the supergiant wind (see Sect.~\\ref{sec:intro}). The EPIC spectra could not be fit with a simple absorbed power-law model, and clearly showed the presence of an excess in the residuals below 2--3~keV. The presence of this ``soft excess'', requiring an additional spectral component to obtain an acceptable fit to the data, is believed to be a ubiquitous feature of binary systems hosting accreting NSs \\citep{hickox04}. Its detectability is mostly related to the absorption column density in the direction of the source and its flux. In wind accreting binaries with X-ray luminosities comparable to that of \\jA\\ ($\\simeq$3$\\times$10$^{34}$~erg/s, assuming a distance of 6.1~kpc), the soft component is most likely originated by thermal X-ray photons close to the NS surface or by photoionized or collisionally heated diffuse gas in the binary system \\citep{hickox04}. The statistics of the EPIC spectra below 2--3~keV shown in Sect.~\\ref{sec:jA-xmm} were unfortunately too low to clearly distinguish between these possibilities. An acceptable fit could be obtained by using either a hot ($kT$$\\simeq$1.2~keV) compact ($\\simeq$0.1~km) BB component, or a colder ($kT$$\\simeq$0.1~keV) and more extended ($\\sim$130~km) thermal emission. The latter model would probably be more plausible in the case of \\jA\\ as low luminosity wind accreting systems are not expected to display hot emitting spots on the NS surface \\citep[see discussion in][and references therein]{bozzo10}. As an alternative interpretation, we also showed in Table~\\ref{tab:jAspec} that a partial-covering model would provide an acceptable description of the X-ray spectrum. In this case, the soft excess would be produced by the effect of partial obscuration of the emission from the NS by the surrounding high density material \\citep[see e.g.,][]{tomsick09b}. The combined fit of the EPIC spectra with that obtained from the long-term monitoring of the source with ISGRI, revealed that the models described above can also account for the higher energy emission recorded from the source (up to $\\sim$60~keV). The normalization constants between all the instruments were compatible with unity, thus suggesting that \\jA\\ is characterized by a virtually constant persistent X-ray flux. This is expected for a classical SGXB. \\subsection{ \\jD\\ } The \\xmm\\ observation detected only one of the two possible counterparts to the \\inte\\ source identified before with \\swift\\ and \\chan\\ (see Sect.~\\ref{sec:jD-xmm}). For the most likely counterpart (S1), the non-detection by \\xmm\\ sets a 3$\\sigma$ upper limit of 7$\\times$10$^{-14}$~\\ferg\\ on its observed 0.5--10~keV X-ray flux. This is a factor of $\\sim$6 lower than the previously determined upper limit (see Sect.~\\ref{sec:jD-xmm}). This observation thus increases significantly the dynamic range in the X-ray luminosity of \\jD,\\ supporting the classification of this source as a SFXT. We also note that the X-ray flux of the source S2 measured by \\xmm\\ is roughly compatible with that reported previously by \\citet{vercellone09}, thus supporting the idea that this is a persistent object unrelated to the \\inte\\ source. \\subsection{ \\jE\\ } The \\xmm\\ observation of this source revealed a variability on time scales of hundreds of seconds (see Fig.~\\ref{fig:jElcurve}) that is very reminiscent of that observed from the SFXT prototypes IGR\\,J17391$-$3021 and IGR\\,J08408$-$4503 \\citep{bozzo10}. This provides further support to the association of \\jE\\ with this class of objects. At odds with some other SFXTs observed in quiescence, we could not detect any soft excess in the X-ray spectrum of the source. According to the discussion in Sect.~\\ref{sec:jAdiscussion}, this spectral component might have gone undetected in \\jE\\ due to its relatively high absorption column density ($N_{\\rm H}$$\\sim$1.8$\\times$10$^{23}$~cm$^{-2}$) and the lower exposure time with respect to the \\xmm\\ observations of other quiescent SFXT sources with a similar X-ray luminosity \\citep{bozzo10}. As a final remark we note that the lowest X-ray flux measured by \\xmm\\ is compatible with the upper limit reported previously with \\swift\\,/XRT by \\citet{fiocchi10}, thus suggesting that this is the real quiescent emission level of the source. \\subsection{ \\jC\\ } This source is one of the confirmed SFXTs, and has been observed in outburst several times with \\inte\\ and \\swift.\\ The properties of its soft X-ray emission ($\\lesssim$10~keV) during outbursts were investigated in detail thanks to the observations performed with the XRT on-board \\swift,\\ but so far no observations with a focusing X-ray telescope endowed with a high sensitivity was able to detect the source in quiescence. The \\xmm\\ observation reported here and the one performed previously in 2006 \\citep{bozzo08b} were carried out close to the system apastron and provided a similar 3$\\sigma$ upper limit on the source X-ray luminosity of $\\simeq$2$\\times$10$^{32}$~erg/s (assuming a distance of 2.5~kpc). The nature of the prolonged quiescent state of \\jC\\ around the apastron is presently unknown, due to the lack of proper spectral information. As the orbital period of the system is $\\sim$30~days, the presence of X-ray eclipses during apastron would imply relatively strong constraints on the system inclination. A high eccentricity ($e$) might help to reduce the accretion rate at the apastron, but the value of $e$ inferred for \\jC\\ (0.3-0.4) can hardly explain the large dynamic range in the X-ray luminosity displayed by the source \\citep[][]{bird09,zurita09}. A further reduction of the mass accretion rate would thus require some additional mechanism to be at work at this orbital phase. Proposed models involve the inhibition of the accretion by the NS centrifugal and/or the magnetic barrier \\citep{grebenev07,bozzo08b}, or the presence of a highly structured wind from the companion star which is extremely rarefied around apastron and permits at this phase only accretion at a very low-level \\citep[see, e.g.][and references therein]{zurita09}. Further observations of \\jC\\ around orbital phase 0.4-0.6 with high sensitivity X-ray telescopes, like \\xmm\\ and \\chan,\\ are required to clarify the origin of the lowest X-ray luminosity state reached by this source. \\subsection{ \\jB\\ } This source was never observed before in the soft X-ray domain. The \\xmm\\ data reported here allowed us to identify the counterpart to the \\inte\\ source and revealed that the source X-ray emission is intrinsically highly absorbed ($\\simeq$10$^{23}$~cm$^{-2}$, see Sect.~\\ref{sec:jB-xmm}) and well described by a power-law model with $\\Gamma$$\\sim$1.5. A simultaneous fit with the long-term ISGRI spectrum also supports the idea that the source is a persistent hard X-ray emitter. The uncertainty in the EPIC-MOS position does not allow to firmly identify the counterparts at other wavelengths, but we remark here the possible association with the radio object NVSS\\,J173459$-$204533. If this association will be confirmed by future observations performed with higher spatial-resolution instruments, we propose that \\jB\\ could be a highly absorbed active Galactic nuclei (AGN). A number of these sources were recently discovered with \\inte\\ \\citep[see, e.g.][and references therein]{ricci11}. Further multi-wavelength observations of this source are required in order to establish its real nature." }, "1207/1207.6620_arXiv.txt": { "abstract": "To image faint substellar companions obscured by the stellar halo and speckles, scattered light from the bright primary star must be removed in hardware or software. We apply the ``locally-optimized combination of images'' (LOCI) algorithm to 1-minute Keck Observatory snapshots of GKM dwarfs in the Hyades using source diversity to determine the most likely PSF. We obtain a mean contrast of $10^{-2}$ at 0\\farcs01, $10^{-4}$ at $<$1'', and $10^{-5}$ at 5''. New brown dwarf and low-mass stellar companions to Hyades primaries are found in a third of the 84 targeted systems. This campaign shows the efficacy of LOCI on snapshot imaging as well as on bright wide binaries with off-axis LOCI, reaching contrasts sufficient for imaging 625-Myr late-L/early-T dwarfs purely in post-processing. ", "introduction": "Brown dwarfs have degenerate cores, convective interiors, and clear to cloudy atmospheres---intermediate properties between stars and planets. Their distribution is also intermediate: despite being easier to detect, brown dwarf counts\\cite{dwarfarchives.org} are already being overtaken by exoplanets\\cite{exoplanet.eu}. In order to better understand these ``failed stars,'' we seek to discover brown dwarfs in a known environment that can act as a controlled laboratory. Toward that end, we undertake here a survey for brown dwarfs as companions to sun-like stars in the Hyades star cluster. The Hyades is a nearby ($\\sim$46 pc) open star cluster with about 400 members\\cite{roeser2011}. Because its age and metallicity are known, stellar models can accurately determine the mass of a star given its luminosity. Because it is nearby and has a high proper motion ($\\sim$0\\farcs11/yr)\\cite{perryman1998}, direct-imaging surveys are fruitful in that companionship can be determined based on co-moving proper motion. Thus, we select 84 GKM Hyads for a direct-imaging search for low-mass companions, with sensitivity down to the L-T transition. In this survey, we employ Keck Observatory adaptive optics (AO) with the NIRC2 camera in natural guide star (NGS) mode. Images are acquired using ``snapshots'' of shallower (0.018~s $\\times$ 500 co-adds) and deeper (5~s $\\times$ 6 co-adds) exposures taken over the course of 3 nights in the first epoch (2005/08/23, 2006/01/15, and 2006/12/10), spending only 5--10 minutes total observing time (including overhead) per target. A second epoch of observations are undertaken in 2009--2012, with Lick/IRCAL and MMT/ARIES observations as well as Keck/NIRC2. The target stars are PSF-subtracted to reveal faint companions, using the locally-optimized combination of images\\cite{loci} (LOCI) algorithm. The LOCI algorithm determines the best estimate of the point-spread function (PSF) for a particular image, using a least-squares criterion. In this work we employ LOCI in two unique ways: 1) Snapshot LOCI: Short-exposure snapshots are taken with a fixed position angle, and the other target stars serve as the reference library for the image of interest. 2) Off-axis LOCI: We also find off-axis PSFs using LOCI on bright wide binaries, enabling us to search the secondary star for close faint companions. In this paper we describe implementation of the LOCI algorithm with fixed-position-angle AO snapshots, and present our preliminary results of the search for low-mass companions in the Hyades. ", "conclusions": "In this survey we implement snapshot LOCI to search GKM Hyads for low-mass companions, down to brown dwarfs at the L/T transition. We find that the stellar multiplicity in the Hyades is at least $\\sim$45\\% including 1 quad, 5 triples, and 23 binaries, complete to brown dwarfs beyond 20 AU projected separation. We have at least one new brown dwarf awaiting spectroscopic characterization. LOCI improves the sensitivity by 1--4 magnitudes at varying separations. With a rapid-fire observing strategy and no complicating coronagraphs nor ADI, the LOCI technique as employed in this survey results in contrasts of $10^{-2}$ at 0\\farcs01, $10^{-4}$ at $<$1'', and $10^{-5}$ at 5''. Using snapshot LOCI with a reference library consisting of a number of different stars imaged over several different nights, it is not necessary to normalize the flux; saturated and unsaturated PSFs can be used together in one reference library. Depending on the imaging system in use, if most aberrations are in the pupil plane then off-axis LOCI could work as well as on-axis LOCI. We find in the Keck case that off-axis LOCI works quite well to search for companions to bright wide binary stars, observed in position-angle mode: while off-axis LOCI resulted in lower contrasts, it resulted in similar or improved mass detection limits because the secondary stars are fainter. Therefore, snapshot LOCI can find the best PSF for both primary and secondary stars---the latter being impossible with typical coronagraphs or ADI. The snapshot LOCI technique used here can be implemented to search for faint companions among any set of well-correlated quasistatic-speckle-dominated PSFs." }, "1207/1207.1222_arXiv.txt": { "abstract": "KL Dra is a helium accreting AM CVn binary system with an orbital period close to 25 mins. Approximately every 60 days there is a 4 mag optical outburst lasting $\\sim$10 days. We present the most sensitive X-ray observations made of an AM CVn system during an outburst cycle. A series of eight observations were made using {\\sl XMM-Newton} which started shortly after the onset of an optical outburst. We find that X-rays are suppressed during the optical outburst. There is some evidence for a spectral evolution of the X-ray spectrum during the course of the outburst. A periodic modulation is seen in the UV data at three epochs -- this is a signature of the binary orbital or the super-hump period. The temperature of the X-ray emitting plasma is cooler compared to dwarf novae, which may suggest a wind is the origin of a significant fraction of the X-ray flux. ", "introduction": "Amongst the thousands of known cataclysmic variables (CVs) are a group of more than two dozen systems, known as `ultra-compact' or `AM CVn' binaries, which have orbital periods between $\\sim$5--70 minutes and spectra devoid of hydrogen (see Solheim 2010 for a recent review). These systems are composed of white dwarfs accreting from the hydrogen depleted cores of their companions, an idea supported by evidence for CNO processing (Marsh et al 1991). Like many CVs, a number of AM CVn binaries show outbursts in the optical and near UV when they brighten typically 2--5 magnitudes (Ramsay et al 2012). These outbursts are assumed to be similar to those seen in the hydrogen accreting dwarf novae, which show regular outbursts lasting days to weeks with recurrence times of weeks to months. In 2009 we started a survey of sixteen AM CVn systems to determine how often they go into outburst and whether the outbursting systems had orbital periods in the range 20--40 mins, as predicted by Tsugawa \\& Osaki (1997) (see also Cannizzo 1984). Observations of each target were made approximately once per week using the Liverpool Telescope. Our study found that roughly 1/3 of AM CVn systems showed at least one outburst and that the outbursting systems have an orbital period in the range 24--44 mins (Ramsay et al 2012) which was remarkably consistent with predictions. KL Dra is an AM CVn binary which has an orbital period close to 25 mins (Wood et al 2002). Our survey found that KL Dra undergoes an outburst every two months (Ramsay et al 2010, 2012) making it very similar to hydrogen accreting dwarf novae. We obtained near-UV and X-ray observations of KL Dra using the {\\sl Swift} satellite and found that although it was a strong UV source during optical outburst, we found no strong evidence for a variation in the X-ray flux over the outburst cycle (Ramsay et al 2010). With the greater sensitivity of the X-ray telescopes on board the {\\xmm} satellite compared to {\\swift}, we have obtained a series of eight `Target of Opportunity' (ToO) observations of KL Dra during an optical outburst to determine if the X-ray flux was suppressed during the outburst as has been observed in a number of dwarf novae outbursts (eg Wheatley et al 1996). ", "conclusions": "\\label{discussion} For dwarf novae in quiesence, material gets accumulated in the accretion disk as a result of Roche Lobe overflow from the late type main sequence star. Only a small fraction of the material in the disc gets accreted onto the white dwarf via a {\\sl boundary layer} between the accretion disk and the white dwarf. Although this layer is hot enough to generate X-rays and is optically thin to X-rays in quiescence, some fraction of X-rays detected in quiescence maybe due to a wind from the white dwarf (Perna et al 2003). Outbursts are thought to be due to a sudden increase in the mass transfer rate through the accretion disc due to a thermal instability originating in either the outer or inner parts of the accretion disc (Smak 1983). As a result, there is a sudden increase in the amount of material being accreted and the boundary layer becomes optically thick to X-rays. X-rays are therefore suppressed and replaced by emission at extreme UV wavebands (eg Wheatley, Mauche \\& Mattei 2003, Collins \\& Wheatley 2010). Eventually a cooling wave sweeps through the disc and the outburst ends (see Lasota 2001 for a review). How do the AM CVn binaries differ from the `classical' hydrogen dwarf novae? The most obvious contrast is their orbital period -- KL Dra has an orbital period close to 25 mins while dwarf novae which show regular and super-outbursts, tend to have an orbital period shorter than 2hrs. The binary separation for KL Dra is approximately three times smaller than a CV with an orbital period of 2 hrs, implying a smaller accretion disc. The timescale for outburst cycles will therefore be shorter compared to dwarf novae. The other factor is chemical composition -- AM CVn's are devoid of hydrogen. Unlike dwarf novae, whose outbursts are largely driven by the ionisation of hydrogen, the outbursts of AM CVn's are driven (at least partly) by the ionisation of helium (see Cannizzo (1984) and Kotko et al (2012)). Since the shock temperature, $T_{s}$, is proportional to the mean molecular mass, $\\mu$, we may expect the $T_{s}$ in AM CVn's to be twice as hot as CVs ($\\mu$=0.615 for Solar composition, and $\\mu$=4/3 for an ionised helium flow). However, the temperature would be cooler if the velocity of the accretion flow did not reach the free-fall velocity as is thought to be the case in intermediate polars (eg Saito et al 2010). On the other hand, given that the shock temperature of intermediate polars has been observed to be several ten's of keV, this does not seem to play an important role. On the other hand this apparent discrepancy between the expectation of a high temperature plasma and the observed rather cooler plasma could be explained if a significant fraction of the X-rays did not originate from an accretion shock. One possibility could be that a fraction of the X-rays originate in a wind from the white dwarf or the accretion disc (Kusterer, Nagel \\& Werner 2009) which would have a temperature less than that expected for X-rays emitted in a shock. KL Dra appears to be similar to most dwarf novae in that X-rays are suppressed during an optical outburst (U Gem appears to be the one exception, Mattei, Mauche \\& Wheatley 2000). There is weak evidence that in KL Dra the X-ray temperature is cooler during optical outburst compared to quiescence. In addition, while the softness ratio of the EPIC pn data shows some variation over the outburst, the X-ray spectrum is harder once the system is in quiescence compared to the first observations which were made in optical outburst. This is broadly similar to that seen in the dwarf novae SU UMa, whose X-ray spectrum during the optical outburst is significantly softer during the outburst (Collins \\& Wheatley 2010). In SS Cyg, the hardness changes in a complex manner, but during the outburst, the spectrum is (again) significantly softer (Wheatley et al 2003). With the detection of more examples of regular outbursting systems in wide field surveys (eg Levitan et al 2012) it maybe possible to expand the current sample of AM CVn systems which have X-ray coverage over the course of an outburst cycle." }, "1207/1207.6634_arXiv.txt": { "abstract": "For the first time, we study the evolution of the stellar mass-size relation for star-forming galaxies from $z\\sim4$ to $z\\sim7$ from Hubble-WFC3/IR camera observations of the HUDF and Early Release Science (ERS) field. The sizes are measured by determining the best fit model to galaxy images in the rest-frame 2100 \\AA \\ with the stellar masses estimated from SED fitting to rest-frame optical (from Spitzer/IRAC) and UV fluxes. We show that the stellar mass-size relation of Lyman-break galaxies (LBGs) persists, at least to $z\\sim5$, and the median size of LBGs at a given stellar mass increases towards lower redshifts. For galaxies with stellar masses of $9.5<\\log(\\mstar/\\msun)<10.4 $ sizes evolve as $(1+z)^{-1.20\\pm0.11}$. This evolution is very similar for galaxies with lower stellar masses of $8.6<\\log(\\mstar/\\msun)<9.5$ which is $r_{e} \\propto (1+z)^{-1.18\\pm0.10}$, in agreement with simple theoretical galaxy formation models at high $z$. Our results are consistent with previous measurements of the LBGs mass-size relation at lower redshifts ($z\\sim1-3$).\\\\ ", "introduction": "The size of a galaxy is a fundamental and important parameter to measure. Over the past decade, observations have revealed that sizes of galaxies at a given stellar mass were smaller at higher redshifts and change significantly with redshift. It has been shown that the sizes of galaxies correlate with their stellar masses and that this correlation exists at least up to $z\\sim3$ \\citep[e.g.,][]{franx2008, williams2009, mosleh2011, law2011}. There are many proposed scenarios to explain the physical processes of galaxy assembly that plausibly reproduce the observable stellar mass and size of galaxies at different redshifts (e.g., galaxy minor or major mergers \\citep[]{khochfar2006,khochfar2009, bell2006, naab2009} or gas accretion in outer regions and star formation \\citep[]{dekel2009,Elmegreen2008}. Accurate measurements of both \\textit{stellar masses} and \\textit{sizes} of galaxies over a wide redshift range are fundamentally important to constrain these galaxy formation models. To extend the mass-size relation of galaxies to the highest redshifts, we exploit the Lyman-break galaxies (LBGs) which are star forming galaxies with strong rest-frame UV emission and could be selected by photometric dropout techniques \\citep[e.g.][]{steidel2003, adelberger2004} . These galaxies can be identified out to very high redshifts (e.g., $z\\sim8$)\\citep[]{oesch2012, yan2012} and thus provide insight into the early evolution of the mass-size relation. Morphological studies of LBGs ($z\\sim2-6$) in rest-frame UV have shown that these galaxies are mostly compact sources however, multiple core LBGs have also been found \\citep[e.g.][]{ravindranath2006, law2007,conselice2009}. Analysing their size and structure could help to interpret the dominant mechanism for galaxy growth. The new Wide Field Camera 3 (WFC3) on board the Hubble Space Telescope (HST) can provide sizes of high-z galaxies in longer rest-frame wavelengths than Advanced Camera for Survey (ACS). Size studies of galaxies at redshifts $z>4$ using profile fitting techniques are rare \\citep{oesch2010b}. Here for the first time, we investigate the mass-size relation of dropout galaxies up to the very early stages of galaxy formation, using the advantages of wide-area HST surveys and high spatial resolution of WFC3. We measure the sizes of LBGs at approximately the same rest-frame wavelength at different redshifts, minimizing the effects of morphological K-correction. We utilize observations of both Hubble Ultra Deep Field (HUDF) and the deep wide-area Early Release Science (ERS) field to study the mass-size relation of the largest sample of LBGs at $z\\sim4-7$ so far. The cosmological parameters adopted throughout this paper are $\\Omega_{m}$ = 0.3, $\\Omega_{\\Lambda}$ = 0.7 and $H_{0} = 70$ $km$ $s^{-1}$ $Mpc^{-1}$.\\\\ \\begin{figure} \\includegraphics[width=3.4in]{fig1.eps} \\caption{From left to right, H-band postage stamps ($6''$ x $6''$), best-fit models from GALFIT, residual images and mask maps are shown for a z-dropout candidate (top row) and an i-dropout candidate (bottom row). The quality of the fits can be seen from the residual images.} \\label{fig1} \\end{figure} ", "conclusions": "For the first time, we have studied the stellar mass-size relation of LBGs out to $z\\sim7$ using ultra-deep WFC3/IR observations in the HUDF and ERS fields. We have shown that the mass-size relation of star-forming galaxies persists to very high redshifts, and that at fixed stellar mass, the sizes of galaxies increase significantly towards later cosmic time. The observed size growth of LBGs studied here - $r_{e}\\propto (1+z)^{-1.20\\pm0.11}$ for galaxies with stellar mass of $10^{9.5}$-$ 10^{10.4} \\msun$ - is in agreement with previous stellar mass-size studies at $z\\lesssim 3$ \\citep[e.g.,][]{dahlen2007,mosleh2011, nagy2011, law2011}. It is also consistent with the size evolution estimated by other studies based on luminosity-selected samples \\citep[e.g.,][]{ferguson2004, bouwens2004, bouwens2006,hathi2008a, oesch2010b}. \\\\ The redshift dependence of the size evolutions is very similar for both high and low stellar masses. Therefore, the galaxy size evolution might be written as a separable function of mass and redshift and this would be in agreement with simple models of galaxy formation developed for high redshift systems \\citep[see][]{wyithe2011}. However, our sample is not stellar mass complete; hence using deeper samples in future is needed for further investigations. In Figure \\ref{fig5} we compare our estimated size-redshift relation for LBGs with those in \\cite{law2011} (blue triangles) and \\cite{dahlen2007} (open squares). The sizes are normalized to a stellar mass of $10^{10}\\msun$. The size estimates from different studies are consistent with the best fit found in this study ($m=1.20\\pm0.11$, solid line). This suggests that LBGs may evolve into UV-bright galaxies at $z\\sim1$. At $z\\sim 6$ these galaxies are extremely compact: $r_{e}\\sim0.8$ kpc, at stellar mass of $10^{10}\\msun$. However, they grow by a factor of about 6 to $z\\sim0.85$. The $z\\sim6$ galaxies have a mass size relation close to those of the compact quiescent galaxies at $z\\sim2$; For example the normalized size ($r_{e}/(M/10^{10})^{0.3}$ ) of a sample of quiescent galaxies at $z\\sim2$ studied by \\cite{szomoru2012} is $\\sim0.6$ kpc. The scarcity of observed LBGs between $z\\sim0$ and $z\\lesssim1$, complicates the interpretation of the evolution of their stellar mass-size relation to the present time. The median size measured for local late-type (i.e., $n < 2.5$) SDSS galaxies by \\cite{shen2003} at the same stellar mass (black solid diamond in Figures \\ref{fig4} and \\ref{fig5}, corrected to the rest-frame UV using the analysis by \\cite{Gildepaz2007} and \\cite{Azzollini2009} and D. Szomoru (private communication)) is smaller than the size of $z\\sim 1$ UV-bright galaxies. We note that the SDSS late-type sample is most likely a different galaxy population than the UV-bright sources we study at $z\\gtrsim1$, and is therefore not relevant for direct comparison. \\cite{overzier2010} studied 30 local Lyman Break Analogs ($z<0.3$), however this sample was selected to have a surface brightness limit and is therefore not characteristic of star-forming galaxies in the nearby universe. In addition, there is some evidence that the evolution of the stellar mass-size relation for star-forming galaxies is slower between $z=1$ and $z=0$ \\citep[e.g.,][]{barden2005}. This needs further investigation using homogeneously-selected sample in future.\\\\ This work was funded in part by the Marie Curie Initial Training Network ELIXIR of the European Commission under contract PITN-GA-2008-214227." }, "1207/1207.6402_arXiv.txt": { "abstract": "Although the colour distribution of globular clusters in massive galaxies is well known to be bimodal, the spectroscopic metallicity distribution has been measured in only a few galaxies. After redefining the calcium triplet index--metallicity relation, we use our relation to derive the metallicity of 903 globular clusters in 11 early-type galaxies. This is the largest sample of spectroscopic globular cluster metallicities yet assembled. We compare these metallicities with those derived from Lick indices finding good agreement. In 6 of the 8 galaxies with sufficient numbers of high quality spectra we find bimodality in the spectroscopic metallicity distribution. Our results imply that most massive early-type galaxies have bimodal metallicity, as well as colour, distributions. This bimodality suggests that most massive galaxies early-type galaxies experienced two periods of star formation. ", "introduction": "Optical photometric studies of globular cluster (GC) systems have shown that almost all massive galaxies have bimodal GC colour distributions \\citep[e.g.][]{2001AJ....121.2950K, 2001AJ....121.2974L}. Since extensive spectroscopy has shown the majority of GCs are old ($\\geq 10$ Gyr) \\citep[e.g.][]{2001ApJ...563L.143F, 2005A&A...439..997P, 2005AJ....130.1315S}, this colour bimodality has usually been interpreted as a metallicity bimodality. Metallicity bimodality suggests each galaxy has experienced two periods of intense star formation \\citep{2006ARA&A..44..193B} which must be accommodated into any scenario of galaxy formation. However, both \\citet{2006BASI...34...83R} and \\citet{2006Sci...311.1129Y} showed that a strongly non-linear colour--metallicity relation can produce a bimodal colour distribution from a unimodal metallicity distribution. If the slope of the colour--metallicity relation flattens at intermediate colours, small differences in metallicity correspond to large differences in colour creating a gap in the colour distribution where none exists in the metallicity distribution. A unimodal GC metallicity distribution therefore only requires one period of continuous GC formation. Thus the true form of the GC metallicity distribution therefore has important ramifications both for globular cluster formation and for galaxy formation. Although near-infrared (NIR) photometric studies \\citep[e.g.][]{2007ApJ...660L.109K, 2008MNRAS.389.1150S, 2012A&A...539A..54C} generally support the interpretation of colour bimodality as metallicity bimodality, large samples of spectroscopic metallicities are required to settle the question of the form of GC metallicity distributions. Unfortunately for only a few galaxies have large numbers of GCs been studied spectroscopically. Although the Milky Way has been long known to host a bimodal metallicity distribution \\citep{1979ApJ...231L..19H, 1985ApJ...293..424Z}, the GC metallicity distribution in M31 is less clear. \\citet{2000AJ....119..727B} studied the colour and metallicity distribution of M31 GCs. They found colour bimodality in $(V - K)$ but not in $(V - I)$. Using 125 GCs with spectroscopic metallicities they found that the metallicity distribution is bimodal. They showed that given the large reddening and photometric uncertainties, a unimodal colour distribution would be observed from their spectroscopic metallicity distribution about 75\\% of the time. Using Lick indices \\citet{2009A&A...508.1285G} derived the metallicities of 245 M31 GCs; they found that the metallicity distribution is either bimodal or trimodal. \\citet{2011AJ....141...61C} measured the metallicities of 282 GCs in M31 using the strength of iron spectral indices. Combining their spectroscopic metallicities with 22 GC metallicities obtained using resolved colour magnitude diagrams, they found no statistically significant metallicity bimodality (or trimodality). \\citet{2008MNRAS.386.1443B} measured the metallicity of 207 GCs in NGC 5128, the closest, easily observable giant early-type galaxy, using Lick indices and found a bimodal metallicity distribution. However the study of \\citet{2010ApJ...708.1335W}, which measured the metallicity of 72 GCs in NGC 5128 using Lick indices, did not find statistically significant metallicity bimodality despite finding bimodality in the metallicity sensitive [MgFe]$'$ index and colour. \\citet{2011MNRAS.417.1823A} found evidence of metallicity bimodality in a sample of 112 GCs in the early-type spiral M104 (the Sombrero Galaxy) which they studied spectroscopically. \\citet{1998ApJ...496..808C} determined the metallicities of 150 GCs in M87 using Lick indices and claimed evidence of bimodality. From this data set, \\citet{2006Sci...311.1129Y} noted that the metallicity sensitive index Mg $b$ had a skewed single peak distribution rather than a bimodal distribution. Using 47 Lick index measurements by \\citet{2003ApJ...592..866C} of the massive elliptical M49's GCs, \\citet{2007AJ....133.2015S} found metallicity bimodality. Although these previous spectroscopic studies suggest GC metallicity bimodality is common in massive galaxies, a larger galaxy sample is required to confirm it. Due to the large amounts of telescope time required to acquire significant samples of GCs, a more efficient observational technique is required. The calcium triplet (CaT), at 8498 \\AA{}, 8542 \\AA{} and 8662 \\AA{} in the rest frame, is one of the strongest spectral features in the optical and NIR. In individual stars the CaT becomes stronger with both increasing metallicity and lower surface gravity \\citep{2002MNRAS.329..863C}. Since giant stars dominate the NIR light of an old population and the surface gravity of metal rich giants is lower than metal poor giants, the strength of the CaT in integrated light increases with metallicity. The CaT was first used by \\citet{1988AJ.....96...92A} to measure the metallicities of Milky Way GCs using integrated light and has been successful at determining the metallicities of individual stars \\citep{1997PASP..109..907R, 2008MNRAS.383..183B}. Single stellar population models \\citep[][hereafter V03]{2003MNRAS.340.1317V} show that the CaT index of integrated starlight is sensitive to metallicity while being insensitive to ages greater than 3 Gyr. Although the CaT index is only weakly sensitive to the IMF for the \\citet{1955ApJ...121..161S} IMF and bottom light IMFs such as the \\citet{2001MNRAS.322..231K} IMF \\citepalias{2003MNRAS.340.1317V}, it is affected by more bottom heavy IMFs such as those found in the most massive elliptical galaxies by \\citet{2010Natur.468..940V} and by \\citet{2012Natur.484..485C}. Combining a wide field, red sensitive spectrograph (e.g. \\textsc{deimos}, \\citealt{2003SPIE.4841.1657F}) and the CaT potentially allows large samples of spectroscopic GC metallicities to be assembled. Unfortunately the CaT has not yet been proven as a reliable metallicity indicator. While \\citet{1988AJ.....96...92A} found a linear relationship between the strength of the CaT and metallicity over the range of metallicity of Milky Way GCs ($-2 < $[Fe/H]$ < 0$), in the models of \\citetalias{2003MNRAS.340.1317V} the CaT saturates at metallicities greater than [Fe/H] $\\sim -0.5$. \\citet[hereafter F10]{2010AJ....139.1566F} was the first to use the CaT to determine the metallicity of a large number of extragalactic GCs. They used \\textsc{deimos} to measure the metallicities of 144 GCs in NGC 1407. Although they found a bimodal metallicity distribution, the fraction of GCs in the metal poor peak was different than the fraction of GCs in the blue colour peak. Their measurements of the CaT appeared to be non-linear with colour. They suggested this conflict is possibly due to saturation of the CaT, a non-linear colour--metallicity relationship or the effects of Paschen lines from hot horizontal branch stars on the CaT measurements. In addition, they found the brightest blue and red GCs to have similar CaT values. More recently \\citet[hereafter F11]{2011MNRAS.415.3393F} used the same technique to study the metallicites of 57 GCs around NGC 4494. Unlike in NGC 1407, NGC 4494 showed a linear relation between colour and the CaT index. Although showing colour bimodality, the CaT strength appeared single peaked in NGC 4494. A larger sample of CaT measurements in multiple galaxies is thus required to establish whether the CaT can be reliably used to measure metallicity of extragalactic GCs. ", "conclusions": "We have used the strength of the calcium triplet (CaT) to measure the metallicities of 903 globular clusters (GCs) in 11 galaxies ranging from brightest group ellipticals to isolated lenticulars. This is the largest sample of spectroscopic GC metallicities ever assembled and the first to provide large numbers ($> 50$) of metallicities for multiple galaxies. We showed that the CaT is a viable method of measuring the metallicity of extragalactic GCs with our CaT derived metallicities agreeing with literature Lick index-based metallicities. Using literature metallicities we defined a new colour--metallicity relation (Equation~\\ref{eq:colourmetal}). Although agreement is seen between the metallicities predicted by this relation and CaT-based metallicities in half of our galaxies, the remaining galaxies show some disagreement either at low metallicity or at high metallicity. It is unclear whether these differences are caused by our heterogeneous photometry or real differences in the colour--metallicity relation between galaxies. Differences in the colour--metallicity relation could be driven by differences in the age or the IMF of GCs between galaxies. We found spectroscopic metallicity bimodality in six of the eight galaxies with more than forty GC measurements, while evidence of trimodality was seen in one of the remaining galaxies (NGC 4494) and we can not rule out bimodality in the other (NGC 2768). We also found that the brightest GCs in NGC 1407 posses a different, unimodal metallicity distribution than the fainter, bimodal GCs. We caution that studies that rely on the properties of the brightest GCs to infer the properties of the GC system may give misleading results. These bimodality results are similar to what has been seen in previous spectroscopic studies of the GC metallicity distributions of early-type galaxies where three galaxies (NGC 5128, \\citealt{2008MNRAS.386.1443B}, NGC 4594, \\citealt{2011MNRAS.417.1823A} and M49, \\citealt{2007AJ....133.2015S}) show evidence of bimodality while the situation in a fourth galaxy \\citep[M87][]{1998ApJ...496..808C} is inconclusive. When the previous studies are combined with our own results, nine of the twelve galaxies show evidence for bimodality while in the remaining three galaxies bimodality can not be ruled out. The spatial distributions and kinematics of GCs system also favour multiple, distinct subpopulations. Red GCs are more centrally concentrated than blue GCs \\citep[e.g.][]{1996AJ....111.1529G, 2006MNRAS.367..156B, 2009ApJ...703..939H, 2011MNRAS.416..155F, Duncan2012}. In addition the colour subpopulations often show different kinematics \\citep{2003ApJ...591..850C, 2010A&A...513A..52S, 2010ApJ...709.1083L, 2011ApJS..197...33S, Vince2012} with sharp changes in rotation occurring at the blue/red dividing line. NIR photometric studies \\citep[e.g.][]{2007ApJ...660L.109K, 2008MNRAS.389.1150S, 2012A&A...539A..54C} also see evidence for bimodality in several galaxies although there are unresolved issues with the connection between metallicity and NIR photometry \\citep{2012ApJ...746...88B}. Our observed GC metallicity distributions, together with other observations of GC systems, favour a picture where most massive galaxies have bimodal GC systems. Since GC metallicity bimodality appears to be common most early-type galaxies should have experienced two periods of intense star formation. \\begin{table*} \\caption{Calcium Triplet Measurements and Metallicities} \\label{tab:data} \\begin{tabular}{l c c c c c c} Name & RA. & Dec. & $(g - i)$ & $i$ & CaT & [Z/H] \\\\ & [deg] & [deg] & [mag] & [mag] & [\\AA] & [dex] \\\\ (1) & (2) & (3) & (4) & (5) & (6) & (7) \\\\ \\hline NGC4494\\_GC100 & 187.811375 & 25.779614 & $0.93 \\pm 0.03$ & $21.04 \\pm 0.02$ & $7.44_{-0.64}^{+0.54}$ & $-0.32_{-0.29}^{+0.25}$ \\\\ NGC4494\\_GC101 & 187.855092 & 25.783281 & $0.98 \\pm 0.08$ & $20.81 \\pm 0.06$ & $6.52_{-0.54}^{+1.47}$ & $-0.75_{-0.25}^{+0.68}$ \\\\ NGC4494\\_GC102 & 187.873400 & 25.783786 & $1.04 \\pm 0.03$ & $21.84 \\pm 0.02$ & $9.46_{-0.94}^{+0.64}$ & $0.60_{-0.43}^{+0.30}$ \\\\ \\ldots & \\ldots & \\ldots & \\ldots & \\ldots & \\ldots & \\ldots \\\\ \\hline \\end{tabular} \\medskip The full version of this table is provided in a machine readable form in the online Supporting Infromation. \\emph{Notes} Column (1): Globular cluster (GC) IDs. The IDs are the same as in \\citet{Vince2012} except for the GCs in NGC 4494 where 'NGC4494\\_' has been appended to the IDs from \\citetalias{2011MNRAS.415.3393F}. Column (2) and (3): Right ascension and declination in the J2000.0 epoch, respectively. Column (4): Adopted $(g - i)$ colour. Column (5): Adopted $i$ magnitude. Column (6): CaT index measurement corrected for S/N bias. Column (7): Total metallicity. \\end{table*} \\begin{table*}\\caption{Bimodal Metallicity Distribution Results using GMM}\\label{tab:gmm-metal} \\begin{center} \\begin{tabular}{l c c c c c c c c c c c} Galaxy & Sample & $N$ & $\\mu_{poor}$ & $\\sigma_{poor}$ & $\\mu_{rich}$ & $\\sigma_{rich}$ & $f_{poor}$ & $p$ & $D$ & $k$ & Bi \\\\ (1) & (2) & (3) & (4) & (5) & (6) & (7) & (8) & (9) & (10) & (11) & (12) \\\\ \\hline NGC 1407 & $(g-i)$ & 202 & $-1.01 \\pm 0.08$ &$0.19 \\pm 0.04$ & $-0.27 \\pm 0.10$ &$0.38 \\pm 0.05$ & $0.34 \\pm 0.11$ & 0.999 & $2.50 \\pm 0.28$ & -0.96 & Y \\\\ & CaT & 202 & $-1.52 \\pm 0.26$ &$0.23 \\pm 0.12$ & $-0.53 \\pm 0.36$ &$0.52 \\pm 0.11$ & $0.16 \\pm 0.22$ & 0.946 & $2.46 \\pm 0.66$ & -0.40 & Y \\\\ \\hline NGC 2768 & $(g-i)$ & 49 & $-0.64 \\pm 0.28$ &$0.34 \\pm 0.10$ & $0.48 \\pm 0.48$ &$0.11 \\pm 0.14$ & $0.96 \\pm 0.36$ & 0.645 & $4.40 \\pm 1.51$ & 0.37 & N \\\\ & CaT & 49 & $-1.12 \\pm 0.44$ &$0.64 \\pm 0.20$ & $0.07 \\pm 0.43$ &$0.40 \\pm 0.17$ & $0.78 \\pm 0.29$ & 0.141 & $2.24 \\pm 0.68$ & -0.78 & N \\\\ \\hline NGC 3115 & $(g-i)$ & 122 & $-1.33 \\pm 0.04$ &$0.32 \\pm 0.04$ & $-0.23 \\pm 0.03$ &$0.23 \\pm 0.03$ & $0.49 \\pm 0.05$ & 0.999 & $3.94 \\pm 0.37$ & -1.06 & Y \\\\ & CaT & 122 & $-1.19 \\pm 0.10$ &$0.50 \\pm 0.09$ & $-0.07 \\pm 0.06$ &$0.32 \\pm 0.05$ & $0.57 \\pm 0.08$ & 0.998 & $2.67 \\pm 0.50$ & -0.75 & Y \\\\ \\hline NGC 3377 & $(g-i)$ & 84 & $-1.24 \\pm 0.41$ &$0.49 \\pm 0.12$ & $-0.45 \\pm 0.24$ &$0.16 \\pm 0.12$ & $0.80 \\pm 0.27$ & 0.903 & $2.13 \\pm 0.72$ & -0.22 & Y \\\\ & CaT & 84 & $-1.29 \\pm 0.29$ &$0.54 \\pm 0.15$ & $-0.16 \\pm 0.15$ &$0.35 \\pm 0.12$ & $0.56 \\pm 0.17$ & 0.910 & $2.49 \\pm 0.62$ & -0.66 & Y \\\\ \\hline NGC 4278 & $(g-i)$ & 150 & $-1.08 \\pm 0.07$ &$0.34 \\pm 0.05$ & $-0.22 \\pm 0.08$ &$0.25 \\pm 0.05$ & $0.66 \\pm 0.09$ & 0.995 & $2.86 \\pm 0.37$ & -0.82 & Y \\\\ & CaT & 150 & $-1.48 \\pm 0.15$ &$0.42 \\pm 0.09$ & $-0.48 \\pm 0.05$ &$0.32 \\pm 0.06$ & $0.41 \\pm 0.11$ & 0.999 & $2.65 \\pm 0.53$ & -0.47 & Y \\\\ \\hline NGC 4365 & $(g-i)$ & 131 & $-1.26 \\pm 0.07$ &$0.21 \\pm 0.04$ & $-0.36 \\pm 0.06$ &$0.33 \\pm 0.04$ & $0.30 \\pm 0.06$ & 0.999 & $3.26 \\pm 0.31$ & -0.97 & Y \\\\ & CaT & 131 & $-1.47 \\pm 0.23$ &$0.25 \\pm 0.10$ & $-0.51 \\pm 0.12$ &$0.40 \\pm 0.07$ & $0.18 \\pm 0.16$ & 0.951 & $2.88 \\pm 0.53$ & -0.58 & Y \\\\ \\hline NGC 4494 & $(g-i)$ & 53 & $-0.99 \\pm 0.18$ &$0.35 \\pm 0.08$ & $-0.31 \\pm 0.12$ &$0.13 \\pm 0.05$ & $0.71 \\pm 0.16$ & 0.954 & $2.55 \\pm 0.67$ & -0.58 & Y \\\\ & CaT & 53 & $-0.90 \\pm 0.38$ &$0.62 \\pm 0.19$ & $0.26 \\pm 0.32$ &$0.19 \\pm 0.17$ & $0.80 \\pm 0.27$ & 0.819 & $2.51 \\pm 0.53$ & -1.08 & N \\\\ \\hline NGC 5846 & $(g-i)$ & 54 & $-2.79 \\pm 0.91$ &$0.14 \\pm 0.32$ & $-0.63 \\pm 0.33$ &$0.55 \\pm 0.16$ & $0.02 \\pm 0.30$ & 0.841 & $5.37 \\pm 2.20$ & 1.19 & N \\\\ & CaT & 54 & $-1.42 \\pm 0.31$ &$0.35 \\pm 0.14$ & $-0.41 \\pm 0.13$ &$0.34 \\pm 0.08$ & $0.24 \\pm 0.18$ & 0.925 & $2.93 \\pm 0.99$ & -0.43 & Y \\\\ \\hline \\end{tabular} \\medskip \\emph{Notes} Column (1): Galaxy name. Column (2): Colour or CaT-based metallicity. Column (3): Number of GCs with CaT measured. Column (4): Mean metallicity of the metal poor subpopulation. Column (5): Dispersion of the metal poor subpopulation. Column (6): Mean metallicity of the metal rich subpopulation. Column (7): Dispersion of the metal rich subpopulation. Column (8): Fraction of the clusters in the metal poor population. Column (9): p-value that a bimodal fit is preferred over a unimodal fit. Column (10): Separation of the GMM peaks normalised by their width. Column (11): The kurtosis of the sample. Column (12): Bimodality of the sample. A p-value greater than 0.9, a separation of $D > 2$ and a negative kurtosis are all required for a sample to be considered bimodal. We did not run GMM on NGC 821 (17 GCs), NGC 1400 (34 GCs) and NGC 7457 (7 GCs) due to the low number of GCs with CaT measurements in these galaxies. \\end{center} \\end{table*} \\begin{table*}\\caption{Bimodal Colour Distribution Results using GMM}\\label{tab:gmm-colour} \\begin{center} \\begin{tabular}{l c c c c c c c c c c c} Galaxy & Sample & $N$ & $\\mu_{blue}$ & $\\sigma_{blue}$ & $\\mu_{red}$ & $\\sigma_{red}$ & $f_{blue}$ & $p$ & $D$ & $k$ & Bi \\\\ (1) & (2) & (3) & (4) & (5) & (6) & (7) & (8) & (9) & (10) & (11) & (12) \\\\ \\hline NGC 821 & All & 110 & $0.81 \\pm 0.05$ &$0.04 \\pm 0.03$ & $1.03 \\pm 0.10$ &$0.16 \\pm 0.04$ & $0.18 \\pm 0.26$ & 0.969 & $1.87 \\pm 0.70$ & -0.48 & N \\\\ \\hline NGC 1400 & All & 1010 & $0.80 \\pm 0.01$ &$0.06 \\pm 0.01$ & $1.09 \\pm 0.02$ &$0.10 \\pm 0.01$ & $0.53 \\pm 0.05$ & 0.990 & $3.53 \\pm 0.37$ & -1.16 & Y \\\\ \\hline NGC 1407 & All & 3312 & $0.82 \\pm 0.00$ &$0.06 \\pm 0.00$ & $1.12 \\pm 0.01$ &$0.11 \\pm 0.01$ & $0.48 \\pm 0.02$ & 0.990 & $3.36 \\pm 0.12$ & -1.13 & Y \\\\ & CaT & 202 & $0.86 \\pm 0.02$ &$0.05 \\pm 0.01$ & $1.08 \\pm 0.03$ &$0.11 \\pm 0.01$ & $0.34 \\pm 0.11$ & 0.999 & $2.50 \\pm 0.27$ & -0.97 & Y \\\\ \\hline NGC 2768 & All & 195 & $0.70 \\pm 0.04$ &$0.05 \\pm 0.02$ & $1.00 \\pm 0.03$ &$0.16 \\pm 0.01$ & $0.17 \\pm 0.09$ & 0.999 & $2.51 \\pm 0.23$ & -0.78 & Y \\\\ & CaT & 49 & $0.97 \\pm 0.08$ &$0.10 \\pm 0.03$ & $1.29 \\pm 0.14$ &$0.03 \\pm 0.04$ & $0.96 \\pm 0.36$ & 0.645 & $4.40 \\pm 1.51$ & 0.37 & N \\\\ \\hline NGC 3115 & All & 180 & $0.77 \\pm 0.01$ &$0.09 \\pm 0.01$ & $1.07 \\pm 0.01$ &$0.07 \\pm 0.01$ & $0.54 \\pm 0.04$ & 0.999 & $3.88 \\pm 0.26$ & -1.24 & Y \\\\ & CaT & 122 & $0.79 \\pm 0.01$ &$0.06 \\pm 0.01$ & $1.09 \\pm 0.01$ &$0.07 \\pm 0.01$ & $0.47 \\pm 0.05$ & 0.999 & $4.74 \\pm 0.39$ & -1.38 & Y \\\\ \\hline NGC 3377 & All & 71 & $0.75 \\pm 0.02$ &$0.05 \\pm 0.02$ & $0.94 \\pm 0.03$ &$0.07 \\pm 0.02$ & $0.52 \\pm 0.13$ & 0.997 & $3.07 \\pm 0.54$ & -1.02 & Y \\\\ & CaT & 84 & $0.79 \\pm 0.02$ &$0.08 \\pm 0.01$ & $1.01 \\pm 0.03$ &$0.06 \\pm 0.01$ & $0.67 \\pm 0.10$ & 0.994 & $3.32 \\pm 0.46$ & -1.01 & Y \\\\ \\hline NGC 4278 & All & 700 & $0.82 \\pm 0.01$ &$0.08 \\pm 0.01$ & $1.10 \\pm 0.01$ &$0.10 \\pm 0.01$ & $0.63 \\pm 0.03$ & 0.999 & $3.28 \\pm 0.20$ & -0.82 & Y \\\\ & CaT & 150 & $0.80 \\pm 0.03$ &$0.05 \\pm 0.02$ & $1.02 \\pm 0.04$ &$0.11 \\pm 0.02$ & $0.40 \\pm 0.13$ & 0.999 & $2.56 \\pm 0.49$ & -1.02 & Y \\\\ \\hline NGC 4365 & All & 2159 & $0.79 \\pm 0.00$ &$0.06 \\pm 0.00$ & $1.05 \\pm 0.01$ &$0.12 \\pm 0.00$ & $0.36 \\pm 0.02$ & 0.990 & $2.82 \\pm 0.10$ & -1.01 & Y \\\\ & CaT & 131 & $0.80 \\pm 0.01$ &$0.04 \\pm 0.01$ & $1.05 \\pm 0.01$ &$0.10 \\pm 0.01$ & $0.28 \\pm 0.06$ & 0.999 & $3.34 \\pm 0.31$ & -1.08 & Y \\\\ \\hline NGC 4494 & All & 127 & $0.86 \\pm 0.02$ &$0.09 \\pm 0.01$ & $1.07 \\pm 0.03$ &$0.03 \\pm 0.01$ & $0.80 \\pm 0.09$ & 0.999 & $3.07 \\pm 0.32$ & -0.97 & Y \\\\ & CaT & 53 & $0.86 \\pm 0.02$ &$0.08 \\pm 0.01$ & $1.06 \\pm 0.02$ &$0.04 \\pm 0.01$ & $0.65 \\pm 0.11$ & 0.986 & $3.20 \\pm 0.53$ & -1.14 & Y \\\\ \\hline NGC 5846 & All & 894 & $0.76 \\pm 0.02$ &$0.07 \\pm 0.01$ & $1.01 \\pm 0.02$ &$0.14 \\pm 0.01$ & $0.34 \\pm 0.08$ & 0.999 & $2.29 \\pm 0.16$ & -0.89 & Y \\\\ & CaT & 54 & $0.89 \\pm 0.09$ &$0.13 \\pm 0.05$ & $1.05 \\pm 0.13$ &$0.14 \\pm 0.05$ & $0.50 \\pm 0.30$ & 0.022 & $1.18 \\pm 1.14$ & -0.24 & N \\\\ \\hline NGC 7457 & All & 46 & $0.85 \\pm 0.04$ &$0.17 \\pm 0.05$ & $1.43 \\pm 0.18$ &$0.04 \\pm 0.08$ & $0.98 \\pm 0.22$ & 0.551 & $4.70 \\pm 1.49$ & 0.44 & N \\\\ \\hline \\end{tabular} \\medskip \\emph{Notes} Column (1): Galaxy name. Column (2): Whether all GC candidates or just GCs with CaT measurements. Column (3): Number of GCs. Column (4): Mean colour of the blue subpopulation. Column (5): Dispersion of the blue subpopulation. Column (6): Mean colour of the red subpopulation. Column (7): Dispersion of the red subpopulation. Column (8): Fraction of the clusters in the blue population. Column (9): p-value that a bimodal fit is preferred over a unimodal fit. Column (10): Separation of the GMM peaks normalised by their width. Column (11): The kurtosis of the sample. Column (12): Bimodality of the sample. A p-value greater than 0.9, a separation of $D > 2$ and a negative kurtosis are all required for a sample to be considered bimodal. \\end{center} \\end{table*}" }, "1207/1207.4541_arXiv.txt": { "abstract": "{By adopting the differential age method, we utilize selected 17832 luminous red galaxies (LRGs) from Sloan Digital Sky Survey Data Release Seven (SDSS DR7) covering redshift $00$ the model has no ghosts. It is shown in \\cite{Dimopoulos:2009am, Dimopoulos:2009vu} that the curvature perturbations generated in this model attain the scale-invariant statistically-isotropic power spectrum, if the varying functions have a specific time dependence, i.e. $f \\propto a^{-4}$ and $m \\propto a$, where $a$ is the scale factor, and if the vector mass is initially light and is heavy by the end of inflation (the physical vector mass $m/\\sqrt{f}$ grows considerably during inflation). In a field theory, an external time dependence should be understood as the vev of a field. Given that the functions $f$ and $m$ needs to present a nontrivial evolution during inflation, this field cannot be trivially integrated out, and, in particular, its perturbations cannot be disregarded. The field acts as ``clock''; therefore, the most minimal choice is to identify this field with the inflaton field. Ref. \\cite{Dimopoulos:2010xq} already provided some specific examples where $f$ and $m$ are functions of the inflaton, and showed that, for a generic inflaton potential that satisfies the slow-roll conditions, the attractor solutions of the background evolution lead to small anisotropy in the expansion and to the required time dependence for the background values of $f$ and $m$. However, ref. \\cite{Dimopoulos:2010xq} does not compute cosmological perturbations in this set-up, but refers to the results of \\cite{Dimopoulos:2009am, Dimopoulos:2009vu}, in which the functions $f$ and $m$ are only external classical quantities. The non-minimal coupling through the kinetic and mass terms modulated by inflaton inevitably induces an interaction between the scalar and the vector perturbations. It is natural to think that this interaction modifies the evolution of perturbations in a non-trivial way. As we will show later in this paper, this interaction, which is present already in the linearized level, can actually be directionally biased due to the background anisotropy. Moreover, since the system of perturbations is now a coupled one, consistent quantization has to be done in the matrix form, the formulation first developed in \\cite{Nilles:2001fg} and summarized in Section \\ref{sec:perturbations}. In this paper, we treat the full coupled system consistently, and show that the scalar-vector interaction produces direction-dependent effects in the power spectrum of curvature perturbations, in the framework of the vector curvaton scenario. We find that the near scale invariance of the spectrum is a generic feature of the model. The (near) statistical isotropy can also be achieved; however, it requires a specific choice of the coupling constant entering in the functions $f$ and $m$, and appropriate initial conditions. This is in contrast with the approximated computation of \\cite{Dimopoulos:2009am, Dimopoulos:2009vu}, in which the statistical isotropy appeared to be a generic result, independent of the functional form of $f$ and $m$. We stress that, despite introducing an inflaton, we still want to work under the assumption of \\cite{Dimopoulos:2009am,Dimopoulos:2009vu} that this is a model of vector curvaton, so that the vector field should be responsible for the cosmological perturbations. This can for instance be achieved by assuming that, after inflation, the inflaton quickly decays into relativistic fields, and that the energy density of the decay products become completely negligible with respect to that of the curvaton (or its decay products). We do no follow the details of reheating here, but we simply compute the vector field perturbations until they freeze out, and we assume that they are the source of cosmological perturbations, according to the curvaton mechanism hypothesis. As also done in \\cite{Dimopoulos:2009am,Dimopoulos:2009vu}, we still disregard (for technical reasons, as the computation is very involved) the metric perturbations in the analysis. These are the same working assumptions of the realizations of this mechanism suggested in \\cite{Dimopoulos:2010xq}. In addition to \\cite{Dimopoulos:2010xq}, we however consistently include the interaction with the inflaton perturbations induced by the two functions $f$ and $m$ that characterize this model. We show that this drastically changes the curvaton power spectrum in this model. The paper is organized as follows. In Section \\ref{sec:background}, we analyze the background dynamics including the vev of the scalar and vector fields and the anisotropic expansion of the universe. The background attractor of this model is derived. In Section \\ref{sec:perturbations}, we discuss the evolution of the perturbations. We focus on the 2D scalar modes, which are the ones that contribute to the energy density and thus to the curvature perturbations. We quantize the system consistently for a coupled system and derive the equations of motion to evolve the system of perturbations in the matrix form. In Section \\ref{sec:observables}, we numerically compute the power spectrum of the vector density fluctuations and relate it to that of curvature perturbations. Results are shown for some different values of parameters. In Section \\ref{sec:conclusions}, we discuss the results and conclude. Throughout the paper, natural units are used, $\\hbar = c = 1$, and the Einstein notation is assumed for repeated indices. ", "conclusions": "the statistical isotropy of the curvature power spectrum is not in general attainable without finely-tuned parameters and initial conditions. Although this model is certainly not the only possibility to realize the vector curvaton scenario with varying kinetic and mass terms of the desired time dependence, we have studied a concrete, realistic example that consistently takes into account the vector-scalar interaction in the presence of the vev of vector field. Curvaton mechanism was first introduced in order to separate the generation of curvature perturbations from the detail of inflation dynamics and to relax the requirements for inflaton \\cite{Lyth:2001nq}. This mechanism is favorable in this sense, if the fields are minimally coupled. However, when multiple fields couple to each other in a non-minimal way, as required in \\cite{Dimopoulos:2009am, Dimopoulos:2009vu}, the interaction non-trivially modifies the situation. In such models, calculation in the decoupled limit is not sufficient, and the consistent treatment of the whole system is crucial." }, "1207/1207.2151_arXiv.txt": { "abstract": "We present the hydrodynamic BL~Herculis-type models which display a long-term modulation of pulsation amplitudes and phases. The modulation is either strictly periodic or it is quasi-periodic, with the modulation period and modulation pattern varying from one cycle to the other. Such behaviour has not been observed in any BL~Her variable so far, however, it is a common property of their lower luminosity siblings -- RR~Lyrae variables showing the Blazhko effect. These models provide a support for the recent mechanism proposed by Buchler \\& Koll\\'ath to explain this still mysterious phenomenon. In their model, a half-integer resonance that causes the period doubling effect, discovered recently in the Blazhko RR~Lyrae stars, is responsible for the modulation of the pulsation as well. Although our models are more luminous than is appropriate for RR~Lyrae stars, they clearly demonstrate, through direct hydrodynamic computation, that the mechanism can indeed be operational. Of great importance are models which show quasi-periodic modulation -- a phenomenon observed in Blazhko RR~Lyrae stars. Our models coupled with the analysis of the amplitude equations show that such behaviour may be caused by the dynamical evolution occurring in the close proximity of the unstable single periodic saddle point. ", "introduction": "Nonlinear modelling is one of the key methods to study the large amplitude variability of classical pulsators: RR~Lyrae stars and Cepheids. One of the notable successes of nonlinear pulsation theory is the explanation of resonant effects in classical pulsators. One of the effects is the bump progression in the light/radial velocity curves of classical Cepheids pulsating in the fundamental mode (so-called Hertzsprung progression). The effect is caused by the 2:1 resonance between the fundamental mode and the second overtone, the latter mode being resonantly excited to high amplitude \\citep[e.g.][]{ss76,kovb89,bmk90,bms00}. Due to resonant nonlinear phase synchronisation the second overtone is not visible separately, but manifests itself as a distortion in the light/radial velocity curve. A similar effect is present in classical Cepheids pulsating in the first overtone \\citep[see e.g.][]{kienzle,fbk00}. Another interesting resonant effect, which we understand thanks to nonlinear pulsation theory, is a period-doubling behaviour -- alternating deep and shallow minima in the light/radial velocity curves. It is a characteristic feature of RV~Tau variables, a group of pulsators often considered to be a subgroup of type-II Cepheids \\citep[e.g.][]{wallerstein,szabados}. The period doubling effect is caused by the half-integer resonances, as analysed by \\cite{mb90}. In case of RV~Tau variables, the 5:2 resonance between the fundamental mode and the second overtone is crucial. Period doubling was also discovered in Blazhko RR~Lyrae stars observed with the satellite mission {\\it Kepler} \\citep{kol10, szabo10} and explained by \\cite{kms11} as a result of a 9:2 resonance between the fundamental mode and the 9th overtone. Recently \\cite{sosz_cep_blg} and \\cite{ssm12} reported the discovery of the period doubling effect in a BL~Her star, demonstrating the predictive power of the nonlinear pulsation theory. Existence of period-doubled BL~Her stars was in fact predicted twenty years earlier by \\cite{bm92} \\citep[see also][]{bb94}, who found the effect in their radiative hydrodynamic models. \\cite{ssm12} confirmed that the 3:2 resonance between the fundamental mode and the first overtone is the cause of the alternations, as analysed earlier by \\cite{bm92}. In \\cite{ssm12} we have studied the period doubling effect in BL~Her models with our state-of-the-art convective hydrocodes \\citep{sm08a}. During the test computations with the decreased eddy viscosity we identified a class of models showing the modulation of pulsation amplitude and phase, which is either strictly periodic or quasi-periodic. Such behaviour has not been detected in any BL~Her star so far. Also, because of the strongly reduced eddy-viscosity, the pulsations are violent and spurious spikes appear in the light curves. Still, the models are of great importance for the lower luminosity cousins of BL~Her stars -- RR~Lyrae pulsators, in which the more or less periodic modulation of pulsation amplitude and phase -- the Blazhko effect -- is a common property \\citep[e.g.][]{kk08}. The amplitude modulation in RR~Lyrae stars is one of the most disturbing problems of stellar astrophysics. Although it was discovered more than century ago \\citep{blazhko}, its origin is still mysterious. During the recent years extensive ground-based observation campaigns \\citep[e.g.][]{kol06,jj09} and a top quality and nearly continuous observations of space missions {\\it CoRoT} and {\\it Kepler} allowed to rule out the two models proposed to explain the Blazhko effect, the magnetic oblique rotator/pulsator model and the non-radial resonant rotator/pulsator model \\citep[see][for a review]{GezaSF}. One of the main drawbacks of these models is the predicted clock-work modulation, while it became clear, with the advent of satellite data, that the Blazhko effect can be a very irregular phenomenon \\citep[see e.g.][]{eg11,eg12,kol11}. The recent idea of \\cite{st06}, which assumes that modulation of turbulent convection by transient magnetic fields causes the modulation of pulsation was also questioned \\citep{smk11,mks12}. The discovery of period doubling effect in some Blazhko variables led to the new model behind the Blazhko modulation -- the radial mode resonance model proposed by \\cite{bk11}. Using the amplitude equations' formalism \\citep[AE, see e.g.][]{bg84} \\cite{bk11} showed that the same resonance that causes the period doubling in Blazhko variables, i.e. the 9:2 resonance between the fundamental mode and the ninth overtone, can also cause either periodic or chaotic modulation of the pulsation. The AE formalism is a powerful tool to study the nonlinear pulsation, specifically the possible limit cycles and their stability. The nature of the computed limit cycles, however, depends on the arbitrarily assumed values of the saturation and resonant coupling coefficients, for which realistic values are too difficult to compute \\citep[see e.g.][]{klapp,nowakowski}. Consequently, it is not clear whether the specific solution of the AEs in which pulsation is modulated, can be represented in the real stars. This must be confirmed with realistic nonlinear hydrodynamic models. Attempts to find modulation of pulsation in hydrodynamic models of RR~Lyrae stars did not lead to success so far, likely because of the surface nature of the involved 9th order overtone, which is difficult to model \\citep{bk11}. In this paper we present the hydrodynamic models of BL~Her variables, more luminous, but otherwise (masses, chemical composition, evolution history) siblings of RR~Lyrae stars. In these models a low order 3:2 resonance, between the fundamental mode and the first overtone causes the modulation of pulsation. Our calculations demonstrate for the first time that the mechanism proposed by \\cite{bk11} can be indeed operational in hydrodynamical models. In addition, they illustrate how the quasi-periodic modulation may arise. Thus, they provide a strong support for the radial mode resonance model of the Blazhko effect. We first present our hydrodynamic models which display modulation of pulsation (Section~\\ref{sec.hydro}) and next we demonstrate that the computed behaviour can be captured and understood with the AE formalism (Section~\\ref{sec.ae}). In Section~\\ref{sec.concl} we comment on the reliability of the models and mechanisms responsible for the wealth of detected behaviours. Finally, in Section~\\ref{sec.impl} we discuss the implications for understanding the Blazhko effect in RR Lyrae stars. ", "conclusions": "\\label{sec.concl} \\begin{itemize} \\item {\\it Origin of the period doubling and of modulation of pulsation.} In Fig.~\\ref{fig.hyd.hr} loci of several half-integer resonances are plotted. In principle each of them may cause the period doubling effect, provided that the resonant mode is not damped too strongly, and that respective model falls close to the resonant centre \\citep{mb90}. These conditions rule out all the resonances displayed in Fig.~\\ref{fig.hyd.hr} except the 3:2 resonance between the fundamental mode and the first overtone. In Fig.~\\ref{fig.hyd.orig} we plot the amplitude of the period doubling (amplitude of the highest peak at around $f_0/2$) vs. the mismatch parameter for the three half-integer resonances. Only models of $136\\LS$ are plotted as for this luminosity we conducted a model survey through the full period doubling domain. It is clear that only for the 3:2 resonance all the models are close to the resonance centre, with $|\\Delta_{3:2}|<0.06$ (where $\\Delta_{3:2}=\\omega_1/\\omega_0-1.5$). Considering e.g. the 7:2 resonance with the fourth overtone (middle panel of Fig.~\\ref{fig.hyd.orig}), although some models are located at direct proximity of the resonance centre, other models within the period doubling domain may be located as far as $|\\Delta_{7:2}|>0.5$, which rules out this resonance as a possible cause of period doubling. The same holds for the 9:2 resonance with the seventh overtone. The domain with the modulation of pulsation lays within the period doubling domain. Its location with respect to the mismatch parameters is also plotted in Fig.~\\ref{fig.hyd.orig}, in which we used the amplitude of the highest modulation peak in the vicinity of $f_0$ ($f_0+f_{\\rm m}$ or $f_0-f_{\\rm m}$) to characterise the strength of modulation. For clarity, only the domain for $L=136\\LS$ is plotted, the domains for other luminosities fall roughly at the same location. Already from Fig.~\\ref{fig.hyd.hr} it is clear that the domain with modulation of pulsation follows (at lower luminosities) the loci of the 7:2 resonance with the fourth overtone. Indeed, in Fig.~\\ref{fig.hyd.orig} the resonance centre falls in the middle of the modulation domain and all of the models lay in close proximity to the resonance centre -- as one may read from Fig.~\\ref{fig.hyd.orig} ($|\\Delta_{7:2}|<0.05$ for all the models). This is however true also for the 3:2 resonance that causes the period doubling. Although the models are at offset with respect to the resonance centre, they all fall within $0.02<\\Delta_{3.2}<0.04$. As the modulation domain arises within the period doubling domain, which is caused by the 3:2 resonance, it is therefore natural to assume that the same resonance causes the modulation. It is confirmed with the analysis of the amplitude equations presented in Section~\\ref{sec.ae} \\citep[and also in][]{bk11}. No additional resonance is needed to cause the modulation. It appears within the period doubling domain, not at its centre but at positive $\\Delta$ (Fig.~\\ref{fig.ae_mp}; note a different definition of $\\Delta$ at this plot). Note that the positive mismatch as defined in eq.~(\\ref{eq.delta}) in Section~\\ref{sec.ae} corresponds to $\\Delta_{3:2}>0$, just as we have for our hydromodels (Fig.~\\ref{fig.hyd.orig}). Taking into account that all the behaviours we have found in our hydrodynamic models (period doubling, periodic and quasi-periodic modulation) can be reproduced with AEs for the 3:2 resonance only, we conclude that this resonance is the only cause of both period doubling and amplitude modulation in our hydrodynamic models. \\begin{figure} \\centering \\resizebox{\\hsize}{!}{\\includegraphics{pap.hyd.orig2.eps}} \\caption{Amplitude of the sub-harmonic frequency, $A_{1/2f}$, and amplitude of the highest modulation peak at $f_0$, $A_{\\rm m}$, plotted versus the mismatch parameters for (top panel) 3:2 resonance with the first overtone, $\\Delta_{3:2}=\\omega_1/\\omega_0-1.5$, (middle panel) 7:2 resonance with the fourth overtone, $\\Delta_{7:2}=\\omega_4/\\omega_0-3.5$ and (bottom panel) 9:2 resonance with the seventh overtone, $\\Delta_{9:2}=\\omega_7/\\omega_0-4.5$.} \\label{fig.hyd.orig} \\end{figure} \\item {\\it Reliability of the computed models.} Our models have strongly decreased eddy viscosity as compared to the standard models of BL~Her stars, i.e. $\\alpha_{\\rm m}=0.05$ to be compared with $\\alpha_{\\rm m}=0.2$ we used in \\cite{ssm12}. Consequently the pulsation is violent and spurious spikes appear in the luminosity curve. Still the models may represent the phenomena that may occur in real BL~Her stars. Certainly it is so for the period doubling discovered in these stars only recently, 20 years after the effect was predicted based on hydrodynamic models (see Introduction). No modulation of pulsation was observed in BL~Her stars so far. The turbulent convection model that we use \\citep{ku86} is a very simple 1D formula with several free parameters. The treatment of viscous effects seem to be the weakest point of this and similar models, which is not surprising. The viscous dissipation of turbulent energy takes place over many length scales and its one parameter description is, out of necessity, provisional. We note that the current convective models cannot reproduce a double-mode pulsation unless unphysical neglect of negative buoyancy below the convective zone leads to artificial viscous damping there \\citep{sm08b,sm10}. With the purely radiative models double-mode pulsation couldn't be reproduced either, unless the artificial dissipation (replaced in convective codes by viscous dissipation) was decreased \\citep{kovb93}. Similarly, the period doubling in RR~Lyrae models of \\cite{kms11} was produced assuming lower values of eddy viscosity. Thus, discovery of modulation of pulsation in BL~Her stars in the future cannot be excluded. \\item {\\it Models with chaotic modulation of pulsation.} In case of $L=136\\LS$ we computed additional models across the instability strip. This allowed us to detect another domain in which pulsation is modulated, this time however, in a very chaotic manner. In these models we observe a wealth of dynamical behaviours. In particular we identify several stability windows within the chaotic regime, with stable period $k$ cycles (e.g. with $k=3$, $k=7$ or $k=9$). They arise due to the tangent bifurcation and, as the control parameter ($T_{\\rm eff}$) is increased, undergo a series of pitchfork bifurcations (sub-harmonic cascade). These models will be subject of a separate paper (Smolec \\& Moskalik, in prep.). \\end{itemize}" }, "1207/1207.2367_arXiv.txt": { "abstract": "We present a new implementation of the numerical integration of the classical, gravitational, $N$-body problem based on a high order Hermite\u2019s integration scheme with block time steps, with a direct evaluation of the particle-particle forces. The main innovation of this code (called HiGPUs) is its full parallelization, exploiting both OpenMP and MPI in the use of the multicore Central Processing Units as well as either Compute Unified Device Architecture (CUDA) or OpenCL for the hosted Graphic Processing Units. We tested both performance and accuracy of the code using up to 256 GPUs in the supercomputer IBM iDataPlex DX360M3 Linux Infiniband Cluster provided by the italian supercomputing consortium CINECA, for values of $N \\leq 8$ millions. We were able to follow the evolution of a system of 8 million bodies for few crossing times, task previously unreached by direct summation codes. ", "introduction": "The study of the motion of $N$ point-like masses having initial positions and velocities $\\mathbf{r}_{i0},\\mathbf{v}_{i0}, i=1,2,\u2026,N$ interacting through a pair-wise force that depends only on their positions is known as the \\textit{N-body problem}. Its applications can be found on both small and large scales starting from nuclear physics up to astrophysical problems. In this latter case, the interaction force is gravity, and, when adopting newtonian gravitational interaction, the problem is referred as the \\textit{classical, gravitational, newtonian N-body problem}. Building an appropriate mathematical model for this problem is a crucial task for astrophysicists who want to give an exhaustive representation of objects from planetary systems to galaxy clusters. Actually, an explicit mathematical solution by series, as provided by \\citep{wa}, exists (upon some conditions) for an arbitrary $N$ but it requires such an enormous amount of terms to approach convergence that deprives it of any pratical usability. As a matter of fact, the gravitational $N$-body problem is explicitly solvable only for $N< 3$ while, for $N\\geq 3$, the procedure to get a solution is exclusively numerical. \\par The first numerical simulation of a self gravitating $N$-body system was carried out by \\citet{holm}. Accounting for the similar inverse square law scaling with distance of newtonian gravity and the light intensity coming from a light source, he evaluated the gravitational interaction between two model galaxies represented as two groups of 37 lamps. This required entire weeks to accomplish this 74-body simulation. The first numerical simulations of the movement of an $N$-body system were performed by \\citet{vho}. Faster numerical integrations were not possible until the years 1960s, when the first digital computers were introduced and Aarseth carried out simulations up to 250 stars \\citep{aars}. From those times, Aarseth has been doing a deep work on high precision simulations of $N$-body systems, producing a series of codes (from NBODY0 up to the last NBODY7 release \\citep{nb7}); a good general reference to direct summation codes is Aarseth's book \\citep{grav}. The well known (at least in the field of astrophysics) GrAvityPipE \\textit{GRAPE project} by Sugimoto, Hut and Makino around the end of years 80's (see \\url{http://www.astrogrape.org/}) constituted a significant improvement for the $N$-body simulations. The GRAPE boards constitute actual \\lq gravity accelerators\\rq~ attached to a host workstation and they are still in use (GRAPE-6) allowing high performances for a relatively low cost. Nevertheless, over the last 10 years, the \\textit{Graphics Processing Units} (GPUs) are replacing Central Processing Units (CPUs) and GRAPE boards. Actually, the purpose of the GPUs, until a few years ago, was to manage graphics with high levels of data parallelism. Actually, on a screen, many pixels must be updated at the same time and each information is completely independent from each other. A newer GPU concept was born mostly thanks to \\textit{CUDA} (Compute Unified Device Architecture), which was developed by Nvidia (2006) and can be considered an extension of some standard high level programming languages like C, C++ and Fortran. It was the begin of a new movement called \\textit{General Purpose computing on GPUs} (GPGPUs, see also \\url{http://gpgpu.org/}) which is in continuous development and growth because graphic cards are, in general, easy to program and very cheap, exhibiting a huge computational power both in single and double precision precision floating point operations keeping power consumption low. A general review on the hardware and software developments since the 1960s that led to the successful application of Graphics Processing Units (GPUs) for astronomical simulations is given by \\citep{bed12}. This paper is organized as follows: in Sect. 2 we present the main aspects of the classic, gravitational $N$-body problem; Sect. 3 contains the description of our new $N$-body Hermite's integrator (HiGPUs); in Sect. 4 we illustrate the hardware resources we used for benchmarking the code, whose results are extensively presented and discussed in Sect. 5. Conclusions are drawn in Sect. 6. In the following we will use the words {\\it body}, {\\it star} and {\\it particle} indifferently because we will consider objects belonging to an $N$-body system always as point-like masses. ", "conclusions": "Composite architectures based on computational nodes hosting multicore Cnetral Processing Units connected to one or more Graphic Processing Units have been shown by various authors to be a clever and efficient solution to supercomputing needings in different scientific frameworks. Actually, these architectures are characterized by a high ratio between performance and both installation cost and power consumption. A practical proof of this is that some of the most powerful systems in the TOP 500 list of world's supercomputer are based on that scheme. They are indeed a valid alternative to massively parallel multicore systems, where the final computational power comes by the use of a very large number of CPUs, although each of them has a relatively low clock frequency. It is quite obvious that a full exploit of the best performance of the CPU+GPU platforms requires codes that clearly enucleate a heavy computational kernel, to be assigned in parallel to the GPUs acting as \\lq number crunchers\\rq~ which release, periodically, their results to the hosts. In physics, the study of the evolution of systems of objects interacting via a potential depending on their mutual distance falls into this category. In this paper we presented and discussed a new, high precision, code apt to simulating the time evolution of systems of $N$ point masses interacting with the classical, pair, newtonian force. The high precision comes from both the evaluation by direct summation of the pairwise force among the system bodies and by a proper treatment of the multiple space and time scales of the system, which means resorting to an individudal time-stepping procedure and resynchronizations, as well as using a high order (6th) time integrator. In this paper we discussed the implementation of our fully parallel version of a direct summation algorithm whose O($N^2$) computational complexity is dealt with by GPUs acting as computational accelerators in the hosting nodes where multicore CPUs are governed and linked via MPI directives. The code, called HermiteIntegratorGPUs (HiGPUs, available to the scientific community on \\url{astrowww.phys.uniroma1.it/dolcetta/HPCcodes/HiGPUs.html} or in the frame of the AMUSE project on \\url{amusecode.org}) shows a very good performance both in term of scaling and efficiency in a good compromise between precision (as measured by energy and angular momentum conservations) and speed. We performed an extensive set of test simulation as benchmarks of our code using the PLX composite cluster of the CINECA italian supercomputing inter-university consortium. We find that the integration of an $N = 8,000,000$ bodies is done with an 80\\% efficiency, that is a deviation of just 20\\% from the linear speedup when using 256 nVIDIA Tesla M2070. This corresponds to less than 10 hours of wall clock time to follow the evolution of the 8M body system up to one internal crossing time, performance, at our knowledge, never reached for such kind of simulations. This means that with HiGPUs it is possible to follow the evolution of a realistic model of Globular Cluster (a spherical stellar system orbiting our and other galaxies and composed by about 1,000,000 stars packed in a sphere of about 10 pc radius) with a 1:1 correspondence between number of real stars in the system and simulating particles over a length of about 10 orbital periods around the galactic center in few days of simulation. These kind of simulations will allow, for instance, a thorough investigation of open astrophysical questions that may involve, in their answer, the role of globular clusters and globular cluster systems in galaxies. We cite the open problem of the origin of Nuclear Clusters as observed in various galaxies, like our Milky Way. Some authors (e.g., \\citep{Mil04} and \\citep{Bek07}) suggested a dissipational, gaseous origin while others (\\citep{CD93}, \\citep{CDV97}) indicate as more realistic a dissipationless origin by orbital decay and merging of globular clusters, as already numerically tested in \\citep{CDM08} and \\citep{Aetal12}). One limit in the use of our code is the GPU memory: with a 6GB RAM, as in the case of nVIDIA Tesla M2070, the upper limit in $N$ is $\\sim 8,400,000$, which is, anyway, a number sufficiently large to guarantee excellent resolution in the simulation of most of the astrophysically interesting cases involving stellar systems. A the light of the results obtained in this paper, we are convinced that it is worth testing some other commercially available high-end GPUs apt to work in a scientific environment, like, for instance, the AMD of the HD6970 and 7970 series, which we showed (\\url{astrowww.phys.uniroma1.it/dolcetta/HPCcodes/HiGPUs.html}) to be absolutely competitive in terms of performance/cost ratio. This suggests as clever the idea of adopting for department sized high end computing platforms the solution of few tens of nodes composed by exacore CPUs connected to 2 up to 4 HD7970 GPUs acting as computing accelerators; all this at a very reasonable cost, both on the purchase cost and power consumption." }, "1207/1207.2325.txt": { "abstract": "{We reconsider Lorentz Violation (LV) at the fundamental level. We show that Lorentz Violation is intimately connected with gravity and that LV couplings in QFT must always be fields in a gravitational sector. Diffeomorphism invariance must be intact and the LV couplings transform as tensors under coordinate/frame changes. Therefore searching for LV is one of the most sensitive ways of looking for new physics, either new interactions or modifications of known ones. Energy dissipation/Cerenkov radiation is shown to be a generic feature of LV in QFT. A general computation is done in strongly coupled theories with gravity duals. It is shown that in scale invariant regimes, the energy dissipation rate depends non-triviallly on two characteristic exponents, the Lifshitz exponent and the hyperscaling violation exponent. } ", "introduction": "Lorentz invariance is one of the most important pillars of modern day physics, experimentally unchallenged for more than a hundred years. It has been tested to unprecedent accuracy reaching { $10^{-21}$} in several contexts, \\cite{tests} and down to { $10^{-29}$} recently in the neutron sector, \\cite{romalis}. Its reign coincided with the demise of ``aether\". Michelson and Morley first in a series of experiments and Einstein's special relativity in 1905 eventually banished aether as a physical theory. Few physicists however realized that a new kind of ``aether\" made it back into physics after 1915, with the general theory of relativity: the gravitational field. Theorists rarely question Lorentz Invariance of physics and its universal grip. Such investigations emerge from time to time. An interesting idea was put forward by H. Nielsen and his collaborators, \\cite{nielsen}, suggesting that Lorentz Invariance is a low-energy phenomenon. Lorentz symmetry behaves as a conventional symmetry, and suggests that vanishing Lorentz iolating couplings are always fixed points of the renormalization group (RG). The nontrivial question is whether they are attractive fixed points. Nielsen and Ninomiya showed this in a special case, \\cite{nn}. However, the situation at large remains unclear. Related ideas were applied to other low-energy symmetries, under the name ``anti-grand-unification\", \\cite{int}, but were not explored further than the original papers. A different turn was taken in the late 80's. Kostelecky and Samuel argued that in open string theory, unusual tachyon vevs can trigger vevs of tensors and therefore Lorentz Violation, (LV) \\cite{ks}. Since then a lot has been done in understanding the dynamics of tachyons in open strings , \\cite{sen} and this understanding indicates that the original expectations in \\cite{ks} do not bear out. However, the conjectured phenomenon is dynamical Lorentz violation and so far there is no hint of it in quantum field theory, (QFT). Kostelecky has followed-on in the analysis of the implications of LV, and several detailed parametrizations or LV in the Standard Model were introduced and studied, \\cite{kc}. In the ninenties, several phenomenological Lorentz-violating dispersion relations have been contrasted to particle physics and astrophysical data starting with the works of Coleman and Glashow, \\cite{Coleman} and later \\cite{aemns}, putting new stringent bounds on LV. Such dispersion relations were expected to be generated among others from ``quantum gravity\" or ``spacetime foam\" models, but it is fair to say that such expectations are still beyond the reach of understood techniques and frameworks. In an unexpected turn, in the late nineties, and after the proposal of the AdS/CFT correspondence, \\cite{malda}, the physics of probe D$_3$ branes was studied in a black D$_3$ brane background with the goal of computing the sYM effective potential at finite temperature and strong coupling, \\cite{sym}. It was noted that the effective speed of light on the brane was variable, depended on the radial position of the brane and could become much smaller than the bulk speed of light (gravitational waves)\\footnote{It was shown in full generality later in \\cite{gh} that the open string metric light-cone is always inside the closed string metric light-cone.}, \\cite{sym}. It was suggested that such a setup could be used as an alternative to inflation, by utilising the fact that the bulk speed of light was larger than the speed of light on the brane, \\cite{sym,freese}. Brane models with a variable speed of light were subsequently constructed by stabilising branes in stable orbits around black branes, \\cite{kir1,stephon,quevedo}. This idea can be brought to its logical conclusion by advocating a planetary brane universe, \\cite{kir2}. It became also obvious that in orientifold constructions of the standard model (SM), not all SM ingredients emerge from the same D-brane. This affects certainly issues of unification, \\cite{unif}, but it also implies that in a generic bulk gravitational background such realizations would provide SM particles with different light speeds. This would be particularly prominent for right-handed neutrinos as they would naturally emerge from the ``bulk\" in such constructions\\footnote{Here ``bulk\" is used liberally to indicate branes that wrap the large spacetime dimensions. The fact that bulk neutrinos would be naturally light in the context of generic braneworlds was suggested earlier in \\cite{neu}.}, \\cite{unif,unif2}. Similar observations on LV were made on RS braneworlds in \\cite{grosjean}. The idea that branes in nontrivial gravitational backgrounds have a varying speed of light, was brought to its natural conclusion in \\cite{mirage}. There it has been argued that a probe brane in motion in a gravitational background (possibly generated by other branes) undergoes a cosmological evolution on its world volume that is not due to the brane energy density. It just reflects the presence of bulk gravitational fields. This phenomenon has been termed ``mirage cosmology\". A similar effect on RS braneworlds was described in \\cite{kraus}, explaining ``dark radiation\" in RS cosmology as a bulk effect. The bulk geometries that were utilized to generate varying speed of light theories involved essentially black-brane geometries with regular horizons. Such geometries have a non-trivial entropy associated to the regular horizon. Gubser, \\cite{gubser} has revisited the question of a varying speed of light and provided examples where the bulk geometry, has zero entropy. Subsequently, many such solutions were found as a byproduct of the holographic description of strongly coupled systems at finite density, \\cite{taylor,cgkkm,gk}. Lorentz violation in QFT was revisited in \\cite{anselmi}. The key point was that an UV scale invariant QFT with Lifshitz scaling symmetry has a different power counting structure, controlled by the dynamical exponent $z$. If $z$ is sufficiently larger than the Lorentz invariant value $z=1$, then the standard hierarchy problem in a scalar sector may be vacuous. Moreover in such cases, four fermion interactions may be renormalizable, and can be used as a substitute for electroweak symmetry breaking, \\cite{wadia}. Ho\\v rava proposed a similar idea in the gravitational sector, \\cite{horava}, using a UV theory of the gravitational field that breaks diffeomorphism invariance but has $z=3$ Lifshitz scaling symmetry. This can provide a power-counting renormalizable theory in four dimensions\\footnote{Of course this does not make the theory renormalizable (UV complete). No calculation of the relevant $\\beta$-functions has been done so far to substantiate that the important interactions are UV marginally relevant.}. It was pointed out, \\cite{kk,mukohyama}, that a $z=3$ Lifshitz invariant UV theory of Ho\\v rava gravity would contain the main ingredients to solve the horizon and flatness problems of cosmology and generate scale-invariant perturbations without inflation. The original Ho\\v rava theory had several phenomenological problems in order to be considered as a valid substitute for standard gravitation, \\cite{hlreview}. A modified version, \\cite{blas} was proposed that does not seem to have any obvious experimental conflicts. The OPERA experiment in September 2011, \\cite{opera}, has made the extraordinary claim, that neutrinos traveling from CERN to the GranSasso arrived earlier than expected suggesting that they moved with superluminal speeds. The news made the tour of the world prompting journalists to hail the demise of Einstein and physicists to ``remain sceptical to such an upheaval of special relativity\". Since then, several further tests have been made and the final experimental scrutiny indicated that the result was incorrect: A cable fault and a coarse timer were responsible for the unusual first result. The neutrino sector is more prone to the discovery of new physics as it is the most difficult to make measurements due to the weakness of interactions. Previous limits of Lorentz invariance include at a much lower energy (around 10 MeV), a stringent limit of ${v - c\\over c} < 2\\times 10^{-9}$ from the observation of (anti)neutrinos emitted by the SN1987A supernova, \\cite{longo}. With a baseline analogous to that of OPERA, but at lower neutrino energies (E peaking at 3 GeV with a tail extending above 100 GeV), the MINOS experiment reported a measurement of ${v-c\\over c} = (5.1\\pm 2.9)\\times 10^{-5}$, \\cite{minos}. In the past, a high-energy ($E > 30$ GeV) and short baseline experiment, \\cite{1}, was able to test deviations down to ${v -c\\over c} < 4\\times 10^{-5}$. The new (corrected) OPERA result is an order of magnitude better at ${v -c\\over c} < 10^{-6}$, \\cite{opera2}. In this paper we re-examine the issue of Lorentz violation in QFT, and we pair it with the fact that QFT is coupled to gravity. Our main conclusions are as follows: \\begin{itemize} \\item LV is intimately interlaced with gravity. LV couplings in QFT are fields in a gravitational sector. Diffeomorphism invariance is intact, and the LV couplings transform as tensors under coordinate/frame changes. \\item Searching for LV is one of the most sensitive ways of looking for new physics: either new interactions or modifications of known ones. \\item Energy dissipation/Cerenkov radiation is a generic feature of LV. \\item A general computation can be done in strongly coupled theories with gravity duals. IN a scaling regime, the energy dissipation rate depends non-triviallly on two characteristic exponents, the Lifshitz exponent and the hyperscaling violation exponent. \\end{itemize} ", "conclusions": "" }, "1207/1207.1634_arXiv.txt": { "abstract": "{The high frequency peaked BL Lac PKS 2155-304 with a redshift of \\textit{z}=0.116 was discovered in 1997 in the very high energy (VHE, E$>$100\\,GeV) $\\gamma$-ray range by the University of Durham Mark VI $\\gamma$-ray Cherenkov telescope in Australia with a flux corresponding to 20\\% of the Crab Nebula flux. It was later observed and detected with high significance by the Southern Cherenkov observatory H.E.S.S. establishing this source as the best studied Southern TeV blazar. Detection from the Northern hemisphere is difficult due to challenging observation conditions under large zenith angles. In July 2006, the H.E.S.S. collaboration reported an extraordinary outburst of VHE $\\gamma$-emission. During the outburst, the VHE $\\gamma$-ray emission was found to be variable on the time scales of minutes and with a mean flux of $\\sim$7 times the flux observed from the Crab Nebula. Follow-up observations with the MAGIC-I standalone Cherenkov telescope were triggered by this extraordinary outburst and \\mbox{PKS\\,2155-304} was observed between 28 July to 2 August 2006 for 15 hours at large zenith angles.} {We studied the behavior of the source after its extraordinary flare. Furthermore, we developed an analysis method in order to analyze these data taken under large zenith angles.} {Here we present an enhanced analysis method for data taken at high zenith angles. We developed improved methods for event selection that led to a better background suppression.} {The quality of the results presented here is superior to the results presented previously for this data set: detection of the source on a higher significance level and a lower analysis threshold. The averaged energy spectrum we derived has a spectral index of ($-3.5\\pm0.2$) above 400\\,GeV, which is in good agreement with the spectral shape measured by H.E.S.S. during the major flare on MJD 53944. Furthermore, we present the spectral energy distribution modeling of \\mbox{PKS\\,2155-304}. With our observations we increased the duty cycle of the source extending the light curve derived by H.E.S.S. after the outburst. Finally, we find night-by-night variability with a maximal amplitude of a factor three to four and an intranight variability in one of the nights (MJD 53945) with a similar amplitude. } {} ", "introduction": "The blazar \\mbox{PKS\\,2155-304} is the so-called lighthouse of the Southern hemisphere. The high frequency peaked BL Lac \\mbox{PKS\\,2155-304}, at a redshift of \\textit{z}=0.116, was discovered in the VHE $\\gamma$-ray range by the University of Durham Mark VI $\\gamma$-ray Cherenkov telescope (Australia) in 1997 with a flux corresponding to $\\sim$0.2 times the Crab Nebula flux \\citep{durham}. \\mbox{PKS\\,2155-304} was confirmed as a TeV $\\gamma$-ray source by the H.E.S.S. group after observations in 2002 and 2003 \\citep{hess1}. In July 2006, the H.E.S.S. collaboration reported an extraordinary outburst of VHE $\\gamma$-emission \\citep{hess2155}. During this outburst, the $\\gamma$-ray emission was found to be variable on time scales of minutes with a mean flux of $\\sim$7 times the flux observed from the Crab Nebula for E$>$200\\,GeV. Large amplitude flux variability at these time scales implies that the TeV emission originates from a small region due to the requirement that light travel times must be sufficiently short in the frame of the emitting region. Follow-up observations of the outburst by the MAGIC telescope were triggered in a Target of Opportunity program by an alert from the H.E.S.S. collaboration \\citep{atel}. The results of this campaign are presented in this paper. The CANGAROO group also observed the source immediately after the flare, obtaining a significance of 4.8\\,$\\sigma$ and an averaged integral flux above 660\\,GeV that corresponds to $\\sim$45\\% of the flux observed from the Crab Nebula \\citep{saka}. The H.E.S.S. collaboration continued observations and detected 44 hours later again a major VHE flare. The data were taken contemporaneously with the Chandra satellite and a strong correlation between the X-ray and the VHE $\\gamma$-ray bands was found \\citep{hess_2}. MAGIC observed on six consecutive nights following the trigger and here we present the final results of the data set. Due to observational constraints, MAGIC did not observe the source during the major flares, but in two cases data were taken immediately afterwards, part of the data being simultaneous with H.E.S.S. and Chandra data. Two years later, in 2008 another multi-wavelength campaign was performed providing simultaneous MeV-TeV data taken by the Fermi Gamma-ray Space Telescope and the H.E.S.S. experiment \\citep{hess_fermi}. With these data the low state of the source could be modeled, including high energy data for the first time. All these observations establish this source as the best studied Southern TeV blazar. Blazars are Active Galactic Nuclei whose relativistic plasma jets nearly point towards the observer. The overall (radio to $\\gamma$-ray) spectral energy distribution (SED) of these objects shows two broad non-thermal continuum peaks. For high energy peaked BL Lac objects (HBLs), the first peak of the SED covers the UV/X-ray bands whereas the second peak is in the multi GeV band. There are various models to explain this spectral shape. They are generally divided into two classes: leptonic and hadronic. Both models attribute the peak at keV energies to synchrotron radiation from relativistic electrons (and positrons) within the jet, but they differ on the origin of the TeV peak. The leptonic models advocate the inverse Compton scattering mechanism, utilizing synchrotron self Compton (SSC) interactions and/or inverse Compton interactions with an external photon field, to explain the VHE emission, (e.g. \\citet{maraschi,dermer,sikora}). On the other hand, hadronic models account for the VHE emission through initial p-p or p-gamma interactions or via proton synchrotron emission (e.g. \\citet{mannheim,aharonian,pohl}). Blazars often show violent flux variability, which may or may not be correlated between the different energy bands. Strictly simultaneous observations are crucial to investigate these correlations and understand the underlying physics of blazars. The structure of the paper is the following: In section~\\ref{magic} we introduce the MAGIC telescope. In section~\\ref{data} we present a new analysis method optimized for large zenith angle (ZA) observations. We test the method on Crab Nebula data taken under large zenith angles (60$^\\circ$ to 66$^\\circ$) in section~\\ref{crab_1}. In section~\\ref{2155_1} we apply the method to the \\mbox{PKS\\,2155-304} data set and in addition, we model the spectrum in section~\\ref{sed}. We summarize the results in section~\\ref{conc}. ", "conclusions": "\\label{conc} A study of the high-zenith angle performance of the MAGIC telescope (operating in single-telescope mode) was carried out with observations of the source \\mbox{PKS\\,2155-304} during a high state, conducted between ZA 60$^\\circ$ - 66$^\\circ$. A new analysis procedure was used in this work in order to enhance the sensitivity of the observations under these special conditions, for instance, included information about the signal arrival times to improve the image cleaning and background subtraction. We tested this new analysis method on a Crab data sample and obtained a sensitivity of 5.7\\% of the Crab Nebula flux for 50\\,hrs of observations at high ZA above 0.4\\,TeV. The differential energy spectrum of the Crab Nebula is in excellent agreement with the published data at lower zenith angles. This improved analysis is used to reanalyze data of \\mbox{PKS\\,2155-304} taken with MAGIC in 2006. The energy spectrum of the whole data set from 400\\,GeV up to 4\\,TeV has a spectral index of ($-3.5\\pm0.2$). As it agrees with the index derived by the H.E.S.S. collaboration during the flaring state of the source, we conclude that the spectrum does not show any change in its spectral slope with flux state above 400\\,GeV. Furthermore, we corrected the measured spectrum for the effect of the EBL absorption using the model of \\citet{franc} and made an SED modeling of simultaneous data in the VHE and X-ray energy range. The light curves derived with MAGIC show a significant variability on daily as well as on intra-night time scales. The MAGIC observations immediately after the extraordinary flare measured by H.E.S.S. on MJD 53944 indicate that the source remained essentially constant for the rest of that night. The measurements of the MAGIC and the H.E.S.S. experiments are generally in good agreement. Finally, we conclude that high zenith angle observations with the MAGIC telescope have proven to yield high quality spectra and light curves above 300\\,GeV. With these observations we could extend the duty cycle of \\mbox{PKS\\,2155-304} observations, which is clearly convenient for the study of any flaring source. High-zenith angle observations, although challenging, generally allow for more uninterrupted coverage of highly variable sources. Also, future high ZA observations of some objects allow unique spectral measurements at higher energies than is possible at lower ZA by virtue of the larger effective area." }, "1207/1207.1919_arXiv.txt": { "abstract": "We present O, Na, and Fe abundances, as well as radial velocities, for 113 red giant branch (RGB) and asymptotic giant branch (AGB) stars in the globular cluster M13. The abundances and velocities are based on spectra obtained with the WIYN--Hydra spectrograph, and the observations range in luminosity from the horizontal branch (HB) to RGB--tip. The results are examined in the context of recent globular cluster formation scenarios. We find that M13 exhibits many key characteristics that suggest its formation and chemical enrichment are well--described by current models. Some of these observations include: the central concentration of O--poor stars, the notable decrease in [O/Fe] (but small increase in [Na/Fe]) with increasing luminosity that affects primarily the ``extreme\" population, the small fraction of stars with halo--like composition, and the paucity of O--poor AGB stars. In agreement with recent work, we conclude that the most O--poor M13 giants are likely He--enriched and that most (all?) O--poor RGB stars evolve to become extreme HB and AGB--manqu{\\'e} stars. In contrast, the ``primordial\" and ``intermediate\" population stars appear to experience standard HB and AGB evolution. ", "introduction": "For many years globular clusters (GCs) were viewed as prototypical simple stellar populations containing stars of a single age and chemical composition. However, a detailed examination of GC chemistry revealed large star--to--star abundance variations of the light elements carbon through aluminum (e.g., see reviews by Kraft 1994; Gratton et al. 2004; 2012). While the anticorrelated behavior of carbon and nitrogen with increasing luminosity along the red giant branch (RGB), attributed to first dredge--up (e.g., Iben 1965) and ``canonical extra mixing\" (e.g., Denissenkov \\& VandenBerg 2003), was observed in both cluster and field stars, a peculiar pattern of enhanced N, Na, and Al abundances coupled with depleted O and Mg seemed to be only found in some cluster stars. The simultaneous anticorrelation of O and Mg with Na and Al pointed to high temperature proton--capture burning as the likely source. Unfortunately, it was not immediately clear if the processed material found in the photospheres of GC RGB stars was due to \\emph{in situ} mixing or pollution from a previous generation of more massive stars. The comprehensive GC abundance survey by Carretta et al. (2009a,b) verified that these light element ``anomalies\", in particular the O--Na anticorrelation, are likely present in \\emph{all} Galactic GCs. Additionally, several authors have now shown that the large star--to--star light element abundance variations found on the RGB are also present along the lower RGB, subgiant branch (SGB), and main--sequence turn off (e.g., Gratton et al. 2001; Cohen \\& Mel{\\'e}ndez 2005; Bragaglia et al. 2010). This observation indicates that the unique abundance patterns of GCs are the result of the cluster formation and subsequent evolution rather than \\emph{in situ} processing. Recent photometric observations have discovered that many (all?) GCs exhibit multiple evolutionary sequences in their color--magnitude diagrams (e.g., see reviews by Piotto 2009; Gratton et al. 2012). Since most clusters exhibit a $<$0.1 dex spread in [Fe/H]\\footnote{[A/B]$\\equiv$log(N$_{\\rm A}$/N$_{\\rm B}$)$_{\\rm star}$--log(N$_{\\rm A}$/N$_{\\rm B}$)$_{\\sun}$ for elements A and B.} (e.g., Carretta et al. 2009c), except for a few notable cases with significant [Fe/H] dispersion, the multiple photometric sequences are believed to be driven by He abundance differences. While the source of He, and subsequently the light element variations, is not known, plausible candidates include: $\\sim$5--9 M$_{\\rm \\odot}$ asymptotic giant branch (AGB) stars (e.g., Ventura \\& D'Antona 2009; 2011), rapidly rotating massive main--sequence stars (e.g., Decressin et al. 2010), and massive binary stars (e.g., de Mink et al. 2009). Although recent GC formation models incorporating winds from intermediate mass and massive stars mixed with pristine gas are able to reproduce many of the light element abundance trends unique to the cluster environment (e.g., Decressin et al. 2010; Valcarce \\& Catelan 2011; D'Ercole et al. 2012), the very low oxygen abundances ([O/Fe]$<$--0.4) found in some GC RGB stars seems to require additional processing. While the old paradigm that the O--Na anticorrelation is entirely driven by \\emph{in situ} deep mixing in cluster RGB stars is clearly incorrect, the discovery that many GC stars are also He--rich has an important consequence for resurrecting a modified deep mixing scenario. D'Antona \\& Ventura (2007) showed that it is possible for a metal--poor star that is both He--rich \\emph{and} initially moderately O--poor ([O/Fe]$\\sim$--0.2) and Na--rich to further deplete oxygen down to [O/Fe]$\\sim$--1, without a significant change in the [Na/Fe] ratio. In light of this, M13 is a particularly illuminating case. It has long been known that M13 hosts some of the most O--poor and Na/Al--rich RGB stars of any cluster, and that these stars appear to be found preferentially near the RGB--tip (e.g., Kraft et al. 1992; Pilachowski et al. 1996; Kraft et al. 1997; Cavallo \\& Nagar 2000; Sneden et al. 2004; Cohen \\& Mel{\\'e}ndez 2005; Johnson et al. 2005). However, the sample size of stars for which [O/Fe] has been determined is $\\sim$5 times less than that for which [Na/Fe] and [Al/Fe] have been measured. Since the [O/Fe] ratio may be the most sensitive indicator of deep mixing (D'Antona \\& Ventura 2007), in this work we have measured [O/Fe] (and [Na/Fe]) abundances for $>$100 RGB and AGB stars ranging in luminosity from the RGB bump to the RGB--tip. We now use this extended sample to examine how M13's extreme O--Na anticorrelation extension fits into the modern picture of GC formation and evolution. ", "conclusions": "As mentioned in $\\S$1, recent models predict that GCs likely form in (at least) two distinct episodes. In this scenario, the first star formation event produces stars with halo--like composition (the primordial population), and then the $\\ga$5 M$_{\\rm \\odot}$ progeny of the first generation pollute the cluster with material heavily processed by high temperature proton--capture burning, including newly synthesized He. This new material may be funneled to the cluster core (e.g., D'Ercole et al. 2008) where the second generation stars (intermediate and extreme populations) form; however, it appears that some dilution with pristine gas is required to reproduce the observed light element abundance trends (e.g., Prantzos et al. 2007; D'Ercole et al. 2011). Some implications of this scenario are that: (1) the primordial population is preferentially stripped relative to the second generation stars, (2) the second generation stars may be significantly He--enhanced and more centrally concentrated, and (3) the extra He, in addition to producing multiple evolutionary sequences in cluster color--magnitude diagrams, may cause some stars to experience \\emph{in situ} deep mixing above the RGB bump and/or cause the most He--rich stars to become RGB--manqu{\\'e}, AGB--manqu{\\'e}, or extreme blue horizontal branch (HB) stars. As we discuss below, M13 appears to exhibit many of these characteristics.\\footnote{Interestingly, multiple sequences in M13 color--magnitude diagrams have yet to be found (see Sandquist et al. 2010 for a recent update.)} \\subsection{Supporting Observations of Globular Cluster Formation Models} We noted in $\\S$3 that the primordial population constitutes a considerably smaller fraction of stars in M13 (15$\\%$) than the intermediate and extreme populations. This is consistent with the cluster formation scenario mentioned above where a significant percentage (up to $\\sim$90$\\%$) of first generation but not second generation stars are lost early in the formation process. Although our estimate is somewhat lower than the 34$\\%$ primordial fraction determined by Carretta et al. (2009a; their Table 5), we note that there is typically no clear separation between the primordial and intermediate populations. However, the dominance of the intermediate population in M13 strongly suggests that its formation and chemical enrichment followed the same path as other halo GCs. Similarly, we show in Figure \\ref{f2} that the extreme population appears to be marginally more centrally concentrated than the primordial and intermediate stars. This is supported by the results of two--sided KS tests, which indicate that the primordial and intermediate populations trace the same radial distribution (KS--prob=0.9018) but the extreme population is different than both (KS--prob$_{\\rm P,E}$=0.1603; KS--prob$_{\\rm I,E}$=0.0935). Although the statistical significance is marginal, we note that similar results have been found in a few other clusters where the central concentration of extreme stars is supported by independent observations of radial changes in the color--magnitude diagram (e.g., Carretta et al. 2010; Lardo et al. 2011). While the dynamical evolution of a GC is expected to smear out the radial profile and uniformly mix the various populations after a Hubble time (e.g., Decressin et al. 2010; but see also Bekki 2010), the fact that M13 and other clusters still show a semblance of the extreme stars being centrally concentrated is evidence in support of current cluster formation models. In this light, $\\omega$ Cen is a particularly illustrative example. Since the core relaxation time is similar to the cluster age, $\\omega$ Cen likely preserves early formation history clues. In fact, Johnson \\& Pilachowski (2010) and Gratton et al. (2011) find a clear composition dependence on radial location, with the extreme stars being the most centrally concentrated and the primordial stars the least. \\subsection{Connecting to the New Deep Mixing Model} Although we now know that the historical argument relating \\emph{in situ} deep mixing and the O--Na anticorrelation is incorrect, a modified version has recently been resurrected to explain cluster RGB stars with [O/Fe]$\\la$--0.4 (e.g., D'Antona \\& Ventura 2007). As mentioned in $\\S$1, predicted yields from both $>$5 M$_{\\rm \\odot}$ AGB and massive main--sequence stars generally fail to produce second generation stars with [O/Fe]$\\la$--0.4 and thus a secondary process is required. Interestingly, our M13 observations (and those of past authors) appear to verify the predictions of the D'Antona \\& Ventura (2007) model (see also Figure \\ref{f1}): (1) all of the known stars with [O/Fe]$\\la$--0.4 are located well above the RGB bump, (2) in general there is a monotonic decrease in [O/Fe] with increasing luminosity for the extreme population, and (3) at the highest luminosity there is a large difference in $\\langle$[O/Fe]$\\rangle$ between the extreme and intermediate populations but only $\\sim$0.1 dex increase in $\\langle$[Na/Fe]$\\rangle$ for the extreme stars. We believe the requirements to induce deep mixing (enhanced He and initially low [O/Fe]) are also met for the extreme M13 giants. Figure \\ref{f1} shows that (with one exception) the lowest [O/Fe] ratio found at log(L/L$_{\\rm \\odot}$)$<$2.8 is consistently at [O/Fe]$\\sim$--0.3. Note that this is consistent with the Cohen \\& Mel{\\'e}ndez (2005) observations that do not find stars below the RGB bump with [O/Fe]$<$--0.2. This supports the idea that the low [O/Fe] values found only in the brightest M13 RGB stars is an evolutionary effect and that significant O--depletion does not occur at low RGB luminosities. With regard to He--enhancement, we do not have direct He measurements for these stars but note that the most Na/Al--rich (and thus O--poor) stars in $\\omega$ Cen (Dupree et al. 2011) and NGC 2808 (Pasquini et al. 2012) have enhanced He. We also find ancillary evidence, similar to that found by Carretta et al. (2006) in NGC 2808, in support of He--enrichment from the increase in [Fe/H] from --1.58 in the intermediate population to --1.54 in the extreme population.\\footnote{We caution the reader on this point because the [Fe/H] difference is small and several of the extreme population stars are known to be variable.} However, we note that Sandquist et al. (2010) does not find significant evidence for He--enrichment in M13. On the other hand, if the O--poor stars are He--rich then the fact that deep mixing appears to be activated at a single luminosity (log(L/L$_{\\rm \\odot}$)$\\sim$2.8) may be evidence in support of the extreme stars having a small He spread. This is qualitatively in agreement with photometric studies that often find discrete populations in cluster color--magnitude diagrams rather than a spread (e.g., Piotto 2009). \\subsection{Composition and Post--RGB Evolution} In the scenario described above, the most He--rich stars likely undergo deep mixing that has the observational effect of decreasing [O/Fe]; however, it also increases the envelope He abundance to as much as Y=0.5 (e.g., D'Antona \\& Ventura 2007). Since He--enhancement may be strongly manifest in HB and AGB evolution, we can look to these stars for clues regarding He--enhancement and RGB evolution. One of the most notable features of Figure \\ref{f1}, which has been shown previously with Na abundances (e.g., Pilachowski et al. 1996), is the lack of extreme stars on the AGB.\\footnote{With the present data we are unable to differentiate AGB and RGB stars near the RGB--tip. However, we expect most, if not all, of the stars near the RGB--tip to be first ascent giants because of the short evolutionary timescale of AGB stars.} Since we find the extreme stars to constitute $\\sim$20$\\%$ of M13's total population, we should expect to find $\\sim$2--3 super O--poor AGB stars in our sample. Interestingly, we find that only the primordial and intermediate AGB stars are present in about the same proportion as on the RGB. Following similar results in other GCs (e.g., Norris et al. 1981; Campbell et al. 2010; Gratton et al. 2010), we conclude that in M13 only the primordial and intermediate populations undergo standard HB and AGB evolution. What about the fate of the extreme population? M13 is known to contain a bimodal and extreme blue HB (e.g., see Sandquist et al. 2010 and references therein). Circumstantial evidence supports the idea that the ``faint peak\" population of HB stars, which have very high T$_{\\rm eff}$, were also once the most O--poor giants on the RGB. In particular, the fraction of faint peak relative to total HB stars is about equal to the fraction of extreme to total RGB stars. The faint peak stars were also found by Sandquist et al. (2010) to be more centrally concentrated than the ``intermediate\" and ``bright peak\" populations. Furthermore, the Sandquist et al. (2010) data indicate that: (1) the fraction of AGB--manqu{\\'e} to total AGB stars is $\\sim$23$\\%$ and (2) the origin of the AGB--manqu{\\'e} stars is likely the bluest part of the HB. Additionally, Peterson et al. (1995) provide [O/Fe] abundances for cool HB stars in M13 and do not find any with [O/Fe]$<$0. All of these observations suggest that the extreme RGB stars evolve from the RGB to the bluest end of the HB and then become AGB--manqu{\\'e} stars. \\subsection{Final Thoughts} The results presented here have allowed us to re--examine M13 in light of recent advances in our understanding of GC formation. M13 may be well explained by the new ``standard\" picture in which first generation stars with halo--like composition are preferentially lost early in the cluster evolution, and a second, more enriched population forms in the cluster center from gas processed and ejected by $>$5 M$_{\\rm \\odot}$ first generation stars. For M13, this has the effect of instigating \\emph{in situ} deep mixing in the most He--rich giants and perhaps causing them to terminate their evolution before ascending the AGB. In fact, proper modeling of the warmest HB stars in clusters like M13 may require considering composition changes to the RGB envelope due to \\emph{in situ} mixing. However, two outstanding issues remain: (1) will precise photometry in the inner part of the cluster finally reveal multiple populations and (2) are some M13 stars actually He--rich?" }, "1207/1207.3407_arXiv.txt": { "abstract": "{We report the energy dependence of normal branch oscillations (NBOs) in Scorpius X-1, a low-mass X-ray binary Z-source. Three characteristic quantities (centroid frequency, quality factor, and fractional root-mean-squared (rms) amplitude) of a quasi-periodic oscillation signal as functions of photon energy are investigated. We found that, although it is not yet statistically well established, there is a signature indicating that the NBO centroid frequency decreases with increasing photon energy when it is below 6-8 keV, which turns out to be positively correlated with the photon energy at the higher energy side. In addition, the rms amplitude increases significantly with the photon energy below 13 keV and then decreases in the energy band of 13-20 keV. There is no clear dependence on photon energy for the quality factor. Based on these results, we suggest that the NBO originates mainly in the transition layer. ", "introduction": "The bright persistent neutron star low-mass X-ray binary (LMXB) Scorpius X-1 traces a \"Z\" track in the X-ray color-color diagram (CCD), which consists of three branches --- the horizontal branch (HB), the normal branch (NB), and the flaring branch (FB, Hasinger \\& van der Klis 1989). The typical Fourier power spectra in Scorpius X-1 are composed of distinct aperiodic variabilities in most of these branches, including noise components (broad structure) and quasi-periodic oscillations (QPOs, narrow feature). There are three distinct types of QPOs observed by the Rossi X-Ray Timing Explorer (RXTE), i.e., the normal branch oscillation (NBO) at $\\sim$ 6 Hz, the horizontal branch oscillation (HBO) at $\\sim$ 45 Hz with a harmonic at about 90 Hz, and the kiloherz (kHz) QPOs (van der Klis et al. 1996, 1997). The NBOs in Scorpius X-1, which occur on the mid- and lower NB, show some correlations with the variability of the $\\sim$ 45 Hz HBO and the twin kHz QPOs (van der Klis et al. 1996; Yu 2007), and therefore imply some kind of coupling between the three types of QPOs (van der Klis et al. 1996; Dieters \\& van der Klis 2000). The NBO frequency remains approximately constant at about 6 Hz, in general. However, the timing properties along the Z track show that the NBO frequency extends to at most $\\sim$ 21 Hz and its position moves smoothly into the lower part of the FB (Priedhorsky et al. 1986; Dieters \\& van der Klis 2000). This smooth transition indicates that the NBO and flaring branch oscillation (FBO) in Scorpius x-1 are physically related to each other (Kuulkers \\& van der Klis 1995; Dieters \\& van der Klis 2000; Casella et al. 2006). Properties of both the twin kHz QPO and 45 Hz HBO depend on the NBO flux. The upper kHz QPO frequency is anticorrelated with the NBO flux, and the lower kHz QPO becomes stronger when the NBO flux is low (Yu et al. 2001). Significant HBOs are detected during the NBO phase of high flux, while the HBO disappears at the low flux of NBO phase (Yu 2007). The coupling between the properties of kHz QPO and HBOs and the phase of the NBO makes the NBO a unique phenomenon. The photons from different regions of the accretion disk carry different energies and present distinct physical properties. So the QPOs detected in different energy band may have varying characteristics. Inspired by the idea that the photon energy dependence of QPOs may provide some additional information on the nature of QPOs, and limited by the available data that provide sufficient spectral information and at the same time include QPO signals, we performed a systematical investigation of the energy dependence of NBOs. We describe our data reduction and analysis in section 2. Results are presented in section 3. In section 4, we discuss the physical implications of our result. Section 5 contains a short summary. ", "conclusions": "The mechanism of NBOs was suggested to be related to the radial oscillation during radiation feedback (Fortner et al. 1989; Miller \\& Park 1995) or the different modes of disk oscillations (Alpar et al. 1992; Titarchuk et al. 2001; Bildsten \\& Cutler 1995; Bildsten et al. 1996; Hor\\'{a}k et al. 2004). The standard explanation attributes the $\\sim$ 6 Hz NBO to part of the accretion with a near Eddington accretion rate, in an approximately spherically symmetric radial inflow at a radius of $\\sim$ 100 km (van der Klis 1995; Hasinger et al. 1990). According to the radiation hydrodynamic model (Fortner et al. 1989), NBO originates from a radiation-force/opacity feedback loop within a spherical flow region at about 300 km from the neutron star surface. The frequency of the oscillations depends on the luminosity. Alternatives are the acoustic oscillation in a thick disk due to the density and optical-thickness perturbations caused by the rotating medium in a subsonic region of the accretion disk (the frequency is determined by the rotation rate and vorticity, Alpar et al. 1992), g-mode oscillations arising from the thermal buoyancy in the ocean on rotating neutron stars (Bildsten \\& Cutler 1995; Bildsten et al. 1996), the low-frequency modulation in a nonlinear resonance oscillation of the relativistic disk (Hor\\'{a}k et al. 2004), and the acoustic oscillations of a spherical shell surface within $\\sim$ 20 km of the neutron star surface because of the change of accretion geometry in the transition zone at high accretion rate (Titarchuk et al. 2001). All of these interpretations focus on the $\\sim$ 6 Hz frequency and its dependence on the accretion rate. None of them can explain the energy dependence of the centroid frequency (if it is further confirmed) and of the rms amplitude. In the NBOs of Scorpius X-1, the rms amplitude has a strong energy dependence and the centroid frequency is also likely to have a `V' shape energy dependence. The centroid frequency changes from anticorrelation to a positive correlation at about 6-8 keV. This nonmonotonic energy dependence of the centroid frequency implies that the emission zone of the NBOs experiences a radial variation during the accretion process. The rms amplitude increases monotonically and significantly with the photon energy below $\\sim$ 13 keV. When the photon energy is higher than 13 keV, the rms amplitude seems to stop increasing, which may indicate that the NBO signals reach the strongest strength at 13-20 keV. Higher energy photons are often emitted from the inner region of an accretion disk. Consequently, the most likely region responsible for such physics can be referred to the transition zone between the inner boundary of accretion disk and the magnetosphere, in which the transition of the properties for accretion flow occurs (Titarchuk et al. 1998; Titarchuk \\& Osherovich 1999; Titarchuk et al. 1999). Normal branch oscillation occurs at high accretion rates, i.e., near the Eddington rate. The high-mass flux deposits increasingly more matter in the transition zone, leading to the expansion and the radial scale change of this region. However, the pile of more deposits with greater viscosity can contribute to the radiation enhancement, as well as to the emission of photons with higher energies. As a result, the radial scale change of the transition zone may be responsible for the observed nonmonotonic energy dependence of the centroid frequency. In addition, because of the limited region of the transition zone, more deposits enhance the viscosity and thus suppress the further monotonically increasing strength of the NBO signal, which can explain the energy dependence of rms amplitude at high energy. We therefore suggest that the NBOs are a type of oscillations in the transition zone. This kind of oscillation may originate in the transition layer due to the viscosity of clumps (e.g. Titarchuk et al. 2001). On the other hand, this oscillation may also carry the information of frequency oscillation modes in the accretion disk and manifest itself as modulation of the accretion rate during the course when disk matter penetrates the transition layer (Paczynski 1987; Nowak \\& Wagoner 1993; Abramowicz et al. 2007)." }, "1207/1207.1772_arXiv.txt": { "abstract": "\\noindent Inflation is the leading paradigm for explaining the origin of primordial density perturbations. However many open questions remain, in particular whether one or more scalar fields were present during inflation and how they contributed to the primordial density perturbation. We propose a new observational test of whether multiple fields, or only one (not necessarily the inflaton) generated the perturbations. We show that our test, relating the bispectrum and trispectrum, is protected against loop corrections at all orders, unlike previous relations. ", "introduction": " ", "conclusions": "" }, "1207/1207.0068_arXiv.txt": { "abstract": "Galaxies are not distributed randomly in the cosmic web but are instead arranged in filaments and sheets surrounding cosmic voids. Observationally there is still no convincing evidence of a link between the properties of galaxies and their host structures. However, by the tidal torque theory (our understanding of the origin of galaxy angular momentum), such a link should exist. Using the presently largest spectroscopic galaxy redshift survey (SDSS) we study the connection between the spin axes of galaxies and the orientation of their host filaments. We use a three dimensional field of orientations to describe cosmic filaments. To restore the inclination angles of galaxies, we use a 3D photometric model of galaxies that gives these angles more accurately than traditional 2D models. We found evidence that the spin axes of bright spiral galaxies have a weak tendency to be aligned parallel to filaments. For elliptical/S0 galaxies, we have a statistically significant result that their spin axes are aligned preferentially perpendicular to the host filaments; we show that this signal practically does not depend on the accuracy of the estimated inclination angles for elliptical/S0 galaxies. ", "introduction": "Disc-dominated galaxies are the morphologically dominant class of galaxies in the present-day Universe \\citep{Bamford:09}. As these galaxies are rotationally supported, it is of vital importance to understand how disc galaxies acquire their angular momentum. The tidal-torque theory predicts alignment effects of disc galaxies, since the acquisition of the angular momentum is partially governed by environmental effects, such as tidal shearing produced by the neighbouring primordial matter distribution and the moment of inertia of the forming protogalaxy \\citep[for a review see][]{Schafer:09}. Hence, testing the alignment of the orientation of galaxies with that of their surrounding environment (e.g., filaments) provides vital information about how galaxies form in the cosmic web. It is known that angular momenta of neighbour galaxies are correlated \\citep{Slosar:09, Lee:11, Andrae:11}, indicating that the environment influences the acquisition of the angular momentum. However, \\citet{Andrae:11} argue that this correlation is plausible but not statistically significant. During the last years, correlations between dark matter haloes and their host filaments have been detected in $N$-body simulations \\citep{Hatton:01, Faltenbacher:02, Bailin:05, Altay:06, Brunino:07, Wang:11}. More specifically, \\citet{Aragon-Calvo:07}, \\citet{Trowland:12}, and \\citet{Codis:12} showed that the orientation of the halo spin vector is mass-dependent. Spin axes of low-mass haloes tend to be aligned parallel to the filaments, whereas high mass haloes have an orthogonal alignment. In this picture, low-mass haloes form through the winding of flows in cosmic sheets/walls; hence at the intersection of these walls (in filaments), the spin axes of the haloes are parallel to the filaments. On the other hand, massive haloes are typically the product of mergers along such a filamentary structure, and thus their spin axes are perpendicular to the spine of the filaments. Previous studies based on observed galaxy catalogues have also indicated that spin axes of galaxies tend to be correlated with the structures in which they are embedded \\citep{Kashikawa:92, Navarro:04a, Trujillo:06}. These studies concentrate on sheets and voids and show that the spin axes of galaxies are preferentially oriented perpendicular to the host structures. In a more recent study, \\citet{Jones:10} used a small sample of edge-on galaxies from the SDSS and showed that the spin axes of these galaxies are aligned perpendicular to the spine of the parent filament. These objects are also less bright, and so their results are in contradiction with the results of $N$-body simulations. The picture is even more complicated since the spin axes of galaxies and their host haloes are not strictly parallel; \\citet{van-den-Bosch:02} find a median misalignment of approximately $30\\degr$. Furthermore, minor mergers can significantly disturb the angular momenta of galaxies \\citep{Moster:10}. To add even more controversy to the picture, \\citet{Slosar:09a} claimed recently that in contrast to previous studies, there are no correlations between galaxies and the voids in which they are located. Furthermore, \\citet{Varela:12} used the SDSS data and the morphological classification from the Galaxy Zoo project and found that galaxy spins tend to be perpendicular to the void walls they are located in, which is in contradiction with the results of \\citet{Navarro:04a} and \\citet{Trujillo:06}. Thus, the results remain contradictory, mostly because of the difficulty to measure the inclination angles of galaxies and/or to properly define the large-scale filamentary structures. In the present study, we use the full SDSS spectroscopic sample to study the correlation between the spin axes of galaxies and the orientation of their host filaments. We build a field of orientations for the filaments. We use a three-dimensional galaxy model to estimate the inclination angles of galaxies. To have good estimates of the inclination angles, we limit our sample strictly to galaxies with the bulge-to-disc ratio less than three: e.g. spirals and elliptical/S0 galaxies. The structure of the present paper is as follows: in Sect.~\\ref{sect:data} we briefly describe the used data. In Sect.~\\ref{sect:model} we describe the three-dimensional galaxy model and explain how we selected the galaxy sample for the correlation study. The algorithm to build the filament orientation field is described in Sect.~\\ref{sect:fil} and the correlation estimator is defined in Sect.~\\ref{sect:cor}. Finally, in Sect.~\\ref{sect:results} we give the results and discuss them. Throughout this paper we assume the following cosmology: the Hubble constant $H_0 = 100\\,h\\,\\mathrm{km\\,s^{-1}Mpc^{-1}}$, the matter density $\\Omega_\\mathrm{m}=0.27$ and the dark energy density $\\Omega_\\Lambda=0.73$ \\citep{Komatsu:11}. ", "conclusions": "We have examined the orientation of spin axes of galaxies with respect to the cosmic filaments. We used 3D photometric modelling of galaxies to restore the inclination angles of galaxies with a sufficiently high accuracy, especially for disc-dominated galaxies. The alignment between galaxy spins and the axis of filaments was characterised by the shape of the probability distribution of $\\cos{\\theta}$ where $\\theta$ is the angle between the two vectors. We studied the correlation for spiral and elliptical/S0 galaxies separately. We found a strong and significant correlation between the spin axes of elliptical/S0 galaxies and filaments; these galaxies tend to be aligned perpendicular to filaments, whereas more luminous elliptical/S0 galaxies have a stronger orthogonal alignment. We showed that this finding is not influenced by the fact that the inclination angles for elliptical/S0 galaxies cannot be estimated accurately. Within the caveats of a limited statistical significance, we found that the spin axes of bright spiral galaxies tend to be aligned parallel to the host filaments. The significant alignment we have found is due to finding the inclination angles of galaxies by 3D modelling: using a simple inclination angle estimate by the visible axial ratio produces a weaker correlation. This emphasises the benefits of having a good estimate for inclination angles when studying the correlation of galaxies with the cosmic web. Our findings for spiral and elliptical/S0 galaxies are in agreement with the recent results of $N$-body simulations, suggesting that spirals form through peaceful accretion of matter and ellipticals are the results of mergers that mostly occur along the filaments. We hope that our results provide a better understanding of the formation of galaxies in the context of their host large-scale structures, which is one of key questions for galaxy formation theory. A precise knowledge of correlated galaxy orientations is of high importance also for future high-precision weak lensing studies of dark energy. So, it is of high importance to improve the filament finding algorithms and the 3D modelling of galaxies to have a better observational insight into this problem. Future work is planned in this direction." }, "1207/1207.7077.txt": { "abstract": "We present evidence of large-scale outflows from three low-mass (log(M$_{*}$/M$_{\\odot})\\sim9.75$) star-forming (SFR $>4$ M$_{\\odot}$ yr$^{-1}$) galaxies observed at $z=1.24$, $z=1.35$ and $z=1.75$ in the 3D-HST Survey. Each of these galaxies is located within a projected physical distance of 60 kpc around the sight line to the quasar SDSS J123622.93+621526.6, which exhibits well-separated strong (W$_{r}^{\\lambda2796}\\gtrsim0.8$\\AA) \\brittion{Mg}{II} absorption systems matching precisely to the redshifts of the three galaxies. We derive the star formation surface densities from the H$\\alpha$ emission in the WFC3 G141 grism observations for the galaxies and find that in each case the star formation surface density well-exceeds 0.1 M$_{\\odot}$ yr$^{-1}$ kpc$^{-2}$, the typical threshold for starburst galaxies in the local Universe. From a small but complete parallel census of the $0.650.8$\\AA~ \\brittion{Mg}{II} covering fraction of star-forming galaxies at $10.4$\\AA~\\brittion{Mg}{II} absorbing gas around star-forming galaxies may evolve from $z\\sim2$ to the present, consistent with recent observations of an increasing collimation of star formation-driven outflows with time from $z\\sim3$. ", "introduction": "Outflowing galactic winds have been widely observed in galaxies from z$\\sim$6 to the present and appear to play a fundamental role in galaxy evolution, regulating galactic star formation \\citep[e.g.,][]{Sanders88, DiMatteo05, Hopkins05} and enriching the intergalactic medium (IGM) at high redshift \\citep[e.g.,][]{Madau01, Scannapieco02}. However, the enclosed gas mass, physical extent, and physical conditions required to trigger galaxy-scale outflows remain to be better quantified before a complete understanding of the contribution of winds to observed co-evolution of galaxies and the IGM may be achieved. %Quasar absorption lines produced by intervening galaxies are particularly powerful probes of the galaxy evolution, as they enable the detection of galaxies by virtue of their gas cross-sections, irrespective of their stellar luminosities, from the earliest epochs of galaxy formation to the present. Bright, background quasars provide one of the most effective tools for probing the cold gas content of galaxies and their extended halos \\citep[e.g.,][]{BS69}, and metal absorption features in quasar spectra have been shown to probe enriched gas in a wide range of galaxy environments, extending to $\\sim150$ kpc around foreground galaxies. The low-ionization \\brittion{Mg}{II} doublet ($\\lambda\\lambda$2796,2803\\AA), a tracer of T $\\sim10^{4}$ K photo-ionized gas, is particularly prolific, and is observable in the range $0.350.4$), the individual galactic hosts tend to be too faint or too close to the quasar to be properly resolved with ground-based observations. Theoretically \\brittion{Mg}{II} absorption is capable of tracing the interstellar medium of the host galaxy \\citep[e.g.,][]{PW97}, star formation-driven outflows \\citep[e.g.,][]{Norman96,Nulsen98}, and cold mode accretion from gas in the extended halos of galaxies \\citep[e.g.,][]{Kacprzak10, Stewart11}. The currently favored paradigm implies that the absorber rest-frame \\brittion{Mg}{II} equivalent width, W$_{r}$ (2796\\AA), can be used to discriminate between these origins, with the highest equivalent widths being associated with disks and star formation-driven winds, and the lowest W$_{r}$ systems indicating more extended virialized gas. At intermediate redshifts, the observed anti-correlation between clustering strength (halo mass) and W$_{r}$ argues against the interpretation that the Mg II absorption systems are virialized \\citep{B06, L09, Gauthier09}; instead, the observed correlation between W$_{r}$ and star-formation rate (SFR) from stacking analyses implies that these systems can be explained by outflowing winds \\citep{Zibetti07, NSM10, Menard11}. This scenario is also consistent with observations of blue-shifted \\brittion{Mg}{II} absorption features in the spectra of star-forming galaxies at similar redshifts \\citep{Weiner09, Rubin10, Erb12, Martin12, Kornei12}. However, the aforementioned statistical studies of \\brittion{Mg}{II} quasar absorption lines are not based on the direct detection of \\brittion{Mg}{II} absorber host galaxies, and it has been suggested that both the stacking and clustering results could have other explanations \\citep{TC08, Chen10a}. Still, the detection of a strong azimuthal dependence of high-W$_{r}$ \\brittion{Mg}{II} absorption within 50 kpc of disk-dominated galaxies at $z<1$ has added to the support for an outflow origin \\citep{Bordoloi11, Kacprzak12, B12b}, given the similar geometry of bipolar outflows observed in local starbursts \\citep[e.g.,][]{Heckman90, LH96, Cecil01, SH09}. A number of studies have used deep imaging or spectroscopy to try to either directly or statistically identify \\brittion{Mg}{II} host galaxies at $z>0.5$ \\citep[e.g.,][]{LeBrun93, Steidel97, Nestor07, B07, Straka10, Chun10, Nestor11} resulting in a few hundred individual detections. While \\brittion{Mg}{II} has been found in association with a wide range of galaxy types, the ultra-strong (W$_{r}>$2\\AA~) absorbers seem to predominantly trace galaxies with high SFRs, at least at high redshift. However, many galaxy hosts of ultra-strong \\brittion{Mg}{II} remain undetected in even the deepest ground-based data. In the most sensitive survey to date, \\citet{B07} detected just 66\\% of the host galaxies of ultra-strong \\brittion{Mg}{II} absorbers at $z\\sim1$ using ultra-deep SINFONI IFU observations, and this number dropped to 20\\% at $z\\sim2$ \\citep{B12a}. These results could indicate that either the \\brittion{Mg}{II} host galaxies have SFRs below the detection limit, are beyond the SINFONI field of view (40 kpc at $z=1$), or lie too close to the quasar to be detected at ground-based resolution. In the low redshift universe ($z\\lesssim0.5$), studies of directly detected \\brittion{Mg}{II} absorber host galaxies have shown that the origins of \\brittion{Mg}{II} are even less clear in the case of the more common population of absorbers with W$_{r}\\la$2\\AA. Analyses of galaxies around detected \\brittion{Mg}{II} absorbers \\citep{Kacprzak11b} and of the absorption properties of galaxies close to quasar sight lines \\citep{Chen10a, Chen10b} each find an inverse relation between W$_{r}$ and impact parameter at $z\\lesssim0.4$ with a large ($\\sim1$ dex) scatter, which can be reduced by accounting for host galaxy inclination and luminosity. The SFRs of the host galaxies in each study provide little support for models in which winds driven by star formation in the host galaxy produce \\brittion{Mg}{II} absorption with W$_{r}\\la1$\\AA, suggesting that the bulk of \\brittion{Mg}{II} absorbers probes infalling gas from disk-halo processes in normal galaxies. While much of the evidence described above supports a picture in which high equivalent width metal absorbers trace star-forming galaxies at small angular separations from quasar sight lines, the high W$_{r}$ systems have typically been studied at high redshift where only the brightest most rigorously star-forming systems are likely to be identified; whereas the local universe studies, which probe galaxies to fainter limits, are statistically restricted to lower W$_{r}$ systems as well as galaxies with lower average specific SFRs. This potential bias combined with the uncertainties outlined above suggest that the interpretation of a correlation between W$_{r}$ and host SFR requires a deeper examination. In an effort to begin resolving this problem, we have harnessed the peerless sensitivity and resolving power of the WFC3 G141 grism to study the absorption properties of a complete sample of galaxies surrounding the only known bright $z>2$ quasar within the 3D-HST survey \\citep{PvD11,Brammer12}. These data enable us to study, for the first time, the properties (redshifts, morphologies, azimuthal angles, SFRs, and stellar masses) of a small but complete sample of Mg II-selected galaxies at $z>1$. Furthermore, we are also able to study the W$_{r}>0.2$\\AA~absorption properties for a volume-limited sample of galaxies in the foreground of the quasar at virtually unlimited impact parameters (5 $\\lesssim \\rho$ [kpc] $\\lesssim$ 450) and to extremely low SFRs (2 M$_{\\odot}$ yr$^{-1}$), thus eliminating the most prohibitive biases of previous high redshift studies of the cold gas content of normal galaxies. The observations contributing to this work are described in Section 2, and details of our analysis are given in Section 3. We present a discussion of our results in Section 4, along with a discussion of the implications of these results, with regard to the evolving distribution of \\brittion{Mg}{II} around galaxies from $z\\sim2$. Throughout this paper, we assume a flat $\\Lambda$--dominated CDM cosmology with $\\Omega_m=0.3$, $H_0=70$ km s$^{-1} $Mpc$^{-1}$, and $\\sigma_8=0.8$ unless otherwise stated. \\begin{table*}[t!] \\centering \\begin{minipage}{0.95\\textwidth} \\begin{center} \\caption{Ancillary Deep Broad-Band Photometry\\label{tbl-bbphot}} \\begin{tabular}{rrrrr} \\tableline\\tableline Filter & Telescope & Instrument & Observations & Reference \\\\ \\tableline\\tableline \\hline U & KPNO 4m & MOSAIC & 9-13 March 2002 & \\citet{Capak04} \\\\ G & Keck I & LRIS B & 3 April 2003 & \\citet{Steidel03} \\\\ Rs & Keck I & LRIS R & 3 April 2003 & \\citet{Steidel03} \\\\ B, V, I, Z & {\\it HST} & ACS & {\\it HST} Cycle 12 9583 & \\citet{Giavalisco04} \\\\ J, H, Ks & Subaru & MOIRCS & 2006-2008 & \\citet{Kajisawa11} \\\\ F140W & {\\it HST} & WFC3 & {\\it HST} GO-11600 & Weiner et al. (in prep.) \\\\ 3.6, 4.5, 5.8, 8.0$\\mu$m & {\\it Spitzer} & IRAC & 2004 & \\citet{Dickinson03} \\\\ \\tableline \\end{tabular} \\end{center} \\end{minipage} \\end{table*} ", "conclusions": "Our essentially unbiased search for the galaxy hosts of \\brittion{Mg}{II} absorption with W$_{r}^{\\lambda2796}>0.4$\\AA~in the 3D-HST Survey area produces a 100\\% detection rate in the available redshift range of the G141 grism. In each case, the \\brittion{Mg}{II} host galaxies are located within 60 kpc of the quasar sight line and have resolved emission lines and SEDs implying SFRs $>4$ M$_{\\odot}$ yr$^{-1}$. All of the galaxies appear to be isolated with no evidence of recent disruption. Thus there is also no ambiguity in the absorber-galaxy pairing due to multiple candidate galaxy hosts at similar redshifts. \\subsection{Properties of \\brittion{Mg}{II} Host Galaxies} For the three galaxies at $z>1$ exhibiting strong \\brittion{Mg}{II} absorption in the quasar spectrum, we measure effective radii along the semi-major axis for these galaxies ranging from 1$-$3 kpc and S{\\'e}rsic indices ranging from $0.5\\lesssim n\\lesssim 2.7$. We note that the true uncertainties in the structural parameter estimates are larger than those reported by {\\it GALFIT} \\citep[e.g.,][]{Haussler07}. If we assume that our small but randomly-selected sample is representative of the larger \\brittion{Mg}{II} host galaxy population at $12$ \\citep[0.1 M$_{\\odot}$ yr$^{-1}$ kpc$^{-2}$;][]{Heckman02}. Given that outflow velocities exhibit a stronger correlation with $\\Sigma$SFR than with SFR \\citep{Kornei12}, it is perhaps not surprising that we find evidence of winds extending to large scales around each of these galaxies. If we assume a wind velocity of 400 km s$^{-1}$, typical of the observations of star-formation-driven winds at $z\\sim1.4$ \\citep{Weiner09}, the impact parameters of the galaxies with respect to the quasar sight line imply that the winds were launched at least $\\sim50$ and $\\sim150$ Myr earlier. The fact that each of the galaxies is still forming stars at a high rate suggests that we may be observing a prolonged burst of star formation on these same timescales. SED modeling of the galaxy G$-$2, which is observed in absorption but at too high a redshift for H$\\alpha$ detection in the G141 spectrum, indicates a SFR approximately equal to that of G$-$1 and G$-$3. We therefore estimate the $\\Sigma$SFR to be $\\sim2$ M$_{\\odot}$ yr$^{-1}$ kpc$^{-2}$, again well above the same starburst threshold. The \\brittion{[O}{III]}/\\brittion{H}{$\\beta$} ratio for this galaxy from the 1D G141 spectrum is measured to be 2.83$\\pm^{0.98}_{0.68}$, consistent with a normal star-forming galaxy. As shown in Figure~\\ref{fig:galfig}, the resolved \\brittion{[O}{III]} emission is extended, tracing the rest-frame optical morphology and thus disfavoring a dominant flux contribution from an AGN. The three \\brittion{Mg}{II} host galaxies we detect at $10.8$\\AA~ \\brittion{Mg}{II} halo mass derived from clustering measurements at $z\\sim1$, 1.8$\\pm^{4.2}_{1.6}\\times$10$^{12}$M$_{\\odot}$ \\citep{L11}. This value is also consistent with multiple clustering measurements at $z\\sim0.6$ \\citep{B06, L09, Gauthier09}, indicating that \\brittion{Mg}{II} host galaxies occupy similarly massive haloes at all redshifts $z\\lesssim2$. %While our sample is unavoidably small, our detection sensitivity (both in terms of absorption strength and galaxy luminosity) makes this galaxy-blind analysis of quasar absorption lines truly the first of its kind outside of the local Universe. The WFC3 G141 has allowed us to probe impact parameters previously out of reach by the best ground-based instruments. The ability to resolve objects with impact parameters to quasar sight lines extending from 5 h$^{-1}$kpc to several hundreds of kpc ensures that our galaxy sample is complete. %\\begin{figure}[t!] %\\epsscale{1.2} \\begin{figure*}[t!] \\centering \\begin{center} \\epsscale{0.85} %\\vspace{5cm} %\\special{psfile=newchen.eps hoffset=-20 voffset=-70 hscale=45 vscale=45 angle=0} %\\special{psfile=newchen_sSFRfit.eps hoffset=220 voffset=-70 hscale=45 vscale=45 angle=0} \\plotone{newchen.eps} \\caption{The observed relation between the rest-frame equivalent width of \\brittion{Mg}{II} absorption, W$_{r}^{\\lambda2796}$, and galaxy impact parameter, $\\rho$ in this work and from the largest existing surveys. Measurements from the literature are overplotted for comparison at low and intermediate redshift (circles, \\citet{Kacprzak11a}; squares, \\citet{Chen10a}; triangles, \\citet{Kacprzak11b}). The shaded gray region indicates the limiting observable impact parameter for this work. Red horizontal lines provide the significance of absorption line detection in the DEIMOS quasar spectrum we examine. \\label{fig:newchen}} \\end{center} \\end{figure*} Despite our unavoidably small sample, a 100\\% detection rate of star-forming galaxy hosts of W$_{r}^{\\lambda2796}>0.8$\\AA~\\brittion{Mg}{II} absorption at $z>1$ may be interesting given the results from recent ground-based analyses, which have reported a large fractional absence of strong \\brittion{Mg}{II} host galaxy detections to sensitive limits in the same redshift range. A search for the galactic counterparts of 20 ultra-strong (W$_{r}^{\\lambda2796}>2$\\AA) \\brittion{Mg}{II} absorbers at $z=2$ using SINFONI undertaken by \\citet{B12a} had only 20\\% success in identifying absorber host galaxies, at a reported SFR sensitivity limit of 2.9 M$_{\\odot}$yr$^{-1}$. This detection rate represents a drop from 66\\% at $z\\sim1$ with a similar search \\citep{B07}. The SINFONI field of view limits a search for companions at $z=2$ to impact parameters less than $\\sim40$ kpc, but the substantial observational evidence indicating that ultra-strong \\brittion{Mg}{II} absorbers are restricted to small impact parameters implies that the missing galaxies are not simply beyond the field of view. \\citet{B12a} also rule out scenarios in which the galaxy hosts are obscured by the quasar at very small impact parameters. Thus, the authors hypothesize that the SFRs of W$_{r}>2$\\AA~\\brittion{Mg}{II} host galaxies may be substantially lower than expected at $z\\sim2$, causing them to lie below the detection threshold of the survey. The strongest \\brittion{Mg}{II} absorber in our sample (G$-$1, $z=1.357$) meets the definition for ultra-strong absorbers studied in the analyses of \\citet{B07} and \\citet{B12a}. With a projected separation of 22.8 kpc and a SFR $\\sim7$ M$_{\\odot}$yr$^{-1}$, the host galaxy we detect is also within the field-of-view and SFR limitations of SINFONI observations and would most certainly have been detected by the H$\\alpha$ surveys of Bouch{\\'e} et al. The second-strongest absorber in our sample (G$-$2, $z=1.749$) has a \\brittion{Mg}{II} equivalent width of W$_{r}=1.82\\pm0.09$\\AA, slightly weaker than the classic definition of ultra-strong \\brittion{Mg}{II}. Still, the projected separation between the host galaxy and the quasar sight line is 19.5 kpc, which is again well within the SINFONI field of view, and the estimated SFR is also above the detection threshold of the SINFONI observations. While our success rate at detecting these host galaxies at small impact parameters hints at a contradiction with the low detection rates from Bouch{\\'e} et al., a sample of two objects cannot provide optimal constraints on the typical SFRs of \\brittion{Mg}{II} host galaxies at $z>1$. A greater number of G141 observations centered on quasar sight lines with high redshift \\brittion{Mg}{II} absorption would be required to better examine whether lower-than-predicted SFRs can account for the large number of host galaxies undetected in other surveys. \\begin{figure*}[ht!] \\centering \\epsscale{0.8} \\plotone{azangle_hist.eps} \\caption{The distribution of the azimuthal angle ($\\Phi$) of galaxies, relative to the quasar sight lines in which \\brittion{Mg}{II} absorption is detected at W$_{r}>0.1$\\AA. We compare our three individual high redshift measurements of $\\Phi$ (stars) with an ensemble of observations from two surveys at lower redshift. The lowest redshift data from \\citet{Kacprzak11a} and \\citet{B12b} consist of 11 galaxy-absorber pairs at $z\\sim0.1$. The intermediate redshift data from \\citet{Kacprzak12} includes 88 galaxy-absorber pairs drawn from a compilation of studies \\citep{Chen10a, Kacprzak11a, Kacprzak11b, Churchill12}, shown in full (solid) and for late-type galaxies only (dashed). The data hint at an evolving azimuthal distribution of gas contained in outflows and traced by \\brittion{Mg}{II}, which is consistent with the observed evolution in the collimation of outflows from star-forming galaxies from $z\\sim3$ \\citep{Law12}. \\label{fig:azdist}} %\\caption{Measurements of W$_{r}^{\\lambda2796}$ vs. $\\rho$, normalized by galaxy inclination, $\\theta_{i}$, relative to the quasar from \\citet{Kacprzak11b}. Our new measurements are overplotted, which indicate a fine agreement with the best-fit relation at $z\\sim0.5$ (dashed line) to $z\\sim2$. \\label{fig:norminc}} \\centering \\end{figure*} %\\vspace{10 mm} %\\clearpage While this work is primarily focused on galaxies with z$>$1, we identify two galaxies at lower redshift with impact parameters of $\\sim100$ kpc (G$-$5, $z=0.831$; G$-$6, $z=0.856$) and corresponding 3$\\sigma$ detections of \\brittion{Mg}{II} with W$_{r}\\sim0.25$\\AA. We note that due to the small equivalent widths of the \\brittion{Mg}{II} 2796\\AA~ transition in each case, no additional absorption lines are detected above the 3$\\sigma$ threshold at the same redshift, making the absorption identifications less certain than the higher equivalent width \\brittion{Mg}{II} systems previously discussed. However, both of these systems exhibit an absorption line with 1$\\sigma$ significance at the expected location of the 2803\\AA~ transition of the \\brittion{Mg}{II} doublet, indicating that despite the weakness of the features the detections are real. As shown in Figure ~\\ref{fig:newchen}, the equivalent widths and impact parameters measured for these absorbing galaxies also agrees well with the previously determined relations at lower redshift. The galaxy identified as G$-$5 has a well-identified H$\\alpha$ emission line in its G141 spectrum, with a luminosity implying a SFR $=1.22$ M$_{\\odot}$yr$^{-1}$. Galaxy G$-$6 exhibits no observable emission lines, despite having a well-determined photometric redshift. The SPS modeling of its SED implies a SFR $\\sim0.1$ M$_{\\odot}$yr$^{-1}$, consistent with the lack of H$\\alpha$ emission. Both galaxies fall well below the $\\Sigma$SFR threshold believed to launch large-scale winds. Given the large projected distances at which the \\brittion{Mg}{II} is observed around these galaxies, it is possible that the enriched gas was launched during a burst of star formation as much as 300 Myr earlier. While the lower redshift absorption features are not the focal point of this work, it is interesting to note that their \\brittion{Mg}{II} host galaxies, which coincidentally are detected at larger impact parameters ($\\sim$100 kpc) relative to the quasar sight line, exhibit a seemingly wider range in SFRs but similar stellar masses. Numerous studies at $z\\lesssim1$ have reported the detection of \\brittion{Mg}{II} around galaxies with a diverse range of morphological and spectral types \\citep[e.g.,][]{Steidel94}, including quiescent galaxies, similar to G$-$6 \\citep[e.g.,][]{GC11}. Hence, the apparent host diversity we find is not unusual. \\subsection{The Spatial Distribution of \\brittion{Mg}{II} Around Galaxies} The inverse relationship between \\brittion{Mg}{II} absorption equivalent width and galaxy impact parameter assumed in the survey design of \\citet{B12a} is well-established in the low and intermediate redshift Universe \\citep[e.g.,][]{LB90, Chen10a, Chen10b} and has been argued to drive the strong observed correlation between \\brittion{Mg}{II} equivalent width and dust extinction in background quasars \\citep{Menard08}. A tight inverse relation of W$_{r}$ and $\\rho$ could also explain the strong correlation between \\brittion{Mg}{II} equivalent width and [\\brittion{O}{II}] emission observed in stacks of SDSS absorption spectra \\citep{Menard11}. Significant evolution in the observed scaling of the W$_{r}-\\rho$ relation could result in the non-detection of \\brittion{Mg}{II} host galaxies in surveys with a limited field of view \\citep[e.g.,][]{B07,B12a}. While our data set is small, its unbiased selection and completeness enables us to look for any obvious departures from these scaling relations at lower redshift. In Figure~\\ref{fig:newchen}, we compare our measurements at $z\\sim1.5$ to those in the largest existing low and intermediate redshift surveys \\citep{Chen10a, Chen10b, Kacprzak11a, Kacprzak11b}. While our observations appear to skirt the upper edge of the measurements from the absorber-blind survey of \\citet{Chen10a,Chen10b}, they agree quite well with the absorption-selected galaxy observations of \\citet{Kacprzak11a, Kacprzak11b}. Thus, we find no signs of dramatic evolution in this relation from $z\\sim1.5$. The large scatter in the W$_{r}-\\rho$ relation, which is equivalent to $\\sim1$ dex along each axis, is also well-established \\citep[e.g.,][]{Chen10a, Chen10b,Kacprzak11b}. Variance in one or more of the many properties of absorber host galaxies likely lies at the root of this scatter, and recent work by \\citet{Chen10b} indicates that the stellar mass of the host galaxy is the fundamental driver. This fact is surprising, given the substantial body of evidence linking the strongest \\brittion{Mg}{II} absorbers to starburst galaxies at $z>1$ \\citep[e.g.,][]{B07,Nestor11}, and photometric stacks of residual light around quasars in the SDSS indicate that stronger (W$_{r}>1$\\AA) \\brittion{Mg}{II} absorbers preferentially trace young, star-forming galaxy populations \\citep{Zibetti07}, compared to (W$_{r}<1$\\AA) \\brittion{Mg}{II} absorbers. The various recent studies linking strong \\brittion{Mg}{II} absorbers to star-forming galaxies at intermediate redshifts suggests that a substantial fraction of these absorbers originate in star formation-driven winds. As such outflows are expected to propagate parallel to the minor axis of galaxies \\citep[e.g.,][]{Strickland04}, a signature of a wind-driven origin is expected in the azimuthal distribution of absorbers around galaxies. The first compelling evidence of the preferential distribution of strong \\brittion{Mg}{II} along the minor axis of star-forming galaxies was reported in the stacked absorption line profiles of intermediate redshift galaxies in zCOSMOS \\citep{Bordoloi11}. Using deep {\\it HST} images to resolve the morphologies and position angles of individual \\brittion{Mg}{II} galaxies, \\citet{Kacprzak11b, Kacprzak12} confirmed a bimodal azimuthal distribution of \\brittion{Mg}{II} absorption in the halos of $z\\sim0.5$ host galaxies, finding that the bulk of these absorbers are aligned within 20$^{\\circ}$ of the major or minor axis of the host. Investigating a smaller sample of 11 galaxies, \\citet{B12b} report an even more striking bimodality at z$\\sim$0.1. To investigate the persistence of this azimuthal dependence at higher redshift, we examine the spatial distribution of the \\brittion{Mg}{II} absorption around the $z>1$ galaxies in this work. In Figure~\\ref{fig:azdist}, we compare our measurements to the azimuthal distributions of absorber-galaxy pairs with W$_{r}>0.3$\\AA~ from \\citet{B12b} and W$_{r}>0.1$\\AA~ from \\citet{Kacprzak12}. While the size of our sample limits what significant conclusions may be drawn from the distribution of $\\Phi$ in this analysis, our data are the first measurements of this kind above z$\\sim$1 and therefore merit discussion. At $160^{\\circ}$) within which the bulk of \\brittion{Mg}{II} has been shown to inhabit at $z<1$, particularly around star-forming galaxies \\citep{Bordoloi11, Kacprzak12} and at low redshift \\citep{B12b}. The collimation of outflows along the minor axis of galaxies observed in local starbursts \\citep[e.g.,][]{Heckman90, LH96, Cecil01, SH09} implies a correlation between the inclination of galaxies and their wind velocities, which can be measured using the blueshifts of low-ionization absorption features of galaxies hosting outflows. This correlation has been observed locally, where galaxy morphologies are readily resolved \\citep{Heckman00,YChen10}, and with deep multi-band imaging at $z=1$ \\citep{Kornei12}. However, observations at higher redshift ($z=1.4$) by \\citet{Weiner09} failed to reproduce the relation, possibly due to the difficulty of resolving galaxy morphologies at higher redshift, even with {\\it HST} {\\it I}-band imaging. A recent study by \\citet{Law12} incorporated deep {\\it HST} WFC3/IR imaging to better resolve the rest-frame optical morphologies of star forming galaxies at $z>2$. Their findings indicate that the correlation between outflow velocity and inclination is indeed absent in low-mass galaxies at $z\\sim2-3$, suggesting that the typically more irregular and ``puffy'' high redshift galaxies have poorly collimated outflows. This increasing collimation of star formation-driven winds with redshift is an expected result of galaxies evolving with time from low-mass dispersion-dominated systems to more massive, stable disks. If the typical geometry of star formation-driven outflows evolves with time to become highly collimated in the local universe, a signature of this evolution should be expected in the azimuthal distribution of metal-enriched gas in the circumgalactic medium surrounding star forming galaxies. Despite the small sample size, our data are consistent with an evolutionary trend in which the collimation of galaxy-scale outflows increases with time. Moreover, the fact that the typical halo masses of \\brittion{Mg}{II} host galaxies appears not to evolve from z$\\sim$2 suggests that the trend of increasing collimation with time holds for star-forming galaxies at constant mass (i.e., the outflows of star-forming galaxies with log(M$_{*}$/M$_{\\odot}$) $=9.75$ are randomly oriented at $z>2$ and more tightly collimated at lower z). We stress that a more extensive survey of \\brittion{Mg}{II} host galaxies with {\\it HST} would be required to determine the significance of this finding. \\begin{figure*}[ht!] \\centering \\epsscale{0.9} \\plotone{galseds.eps} \\caption{The broad-band SEDs of the eight galaxies identified within 150 kpc of the quasar sight line, provided with arbitrary units of $F_{\\lambda}$. In each case the best-fit galaxy template, fit to the redshift determined from the grism spectrum, is overplotted. Redshift probability distribution functions, $P(z)$, are inset within each panel. The green curve provides the broad-band photometric solution of the $P(z)$. The blue curve provides the $P(z)$ determined from the G141 data. \\label{fig:galseds}} \\centering \\end{figure*} \\subsection{Absorption Properties of All Proximate Galaxies} In a parallel search for absorption around all 3D-HST galaxies identified within 150 kpc of the quasar sight line, we detect 8 galaxies with well-determined redshifts in the range $0.650.9$ are coincidentally located within 60 kpc of the quasar sight line, while the galaxies with lower redshifts are detected at larger impact parameters. The separations in the redshifts of these objects ensures that there is no physical correlation between any of these galaxies or with the quasar. Thus, the apparent angular clustering of high redshift galaxies around the quasar sight line should be understood simply as a coincidence and not distract from the other underlying physical correlations of absorbers and host galaxies in this study. Three of the four galaxies with $\\rho<60$ kpc exhibit \\brittion{Mg}{II} absorption with W$_{r}>0.8$\\AA, as discussed in length in Section 4.1. The remaining galaxy within this radius is undetected in \\brittion{Mg}{II} to our observational limit of W$_{r}>0.4$\\AA. Of the galaxies detected in absorption, all have SFR $>4$ M$_{\\odot}$ yr$^{-1}$ and $\\Sigma$SFR $>0.3$ M$_{\\odot}$yr$^{-1}$kpc$^{-2}$, uncorrected for dust extinction. The remaining galaxy, G$-$4, which is undetected in absorption at a similarly close separation of 57 kpc, exhibits no measurable emission lines in the G141 data at the best-fit combined grism/photometric redshift of 0.959. Although we only measure an upper-limit for the \\brittion{H}{$\\alpha$} emission in the 1D grism spectrum (W$_{r}\\sim6$\\AA), SPS modeling of the SED results in a SFR $\\sim2$ M$_{\\odot}$yr$^{-1}$. The possibility remains that the best-fit photometric redshift for G$-$4 could be incorrect, which might simultaneously explain the lack of emission lines in the G141 spectra and the absence of \\brittion{Mg}{II} absorption in the quasar spectra where they are expected. However, the redshift probability distribution function derived from the SED of the galaxy indicates a $<$1\\% probability that the observable wavelength of H$\\alpha$ emission lies outside of the range of the G141 spectrum. Thus, the fact that we find no strong evidence for rigorous ongoing star formation in the only galaxy within 60 kpc that is undetected in \\brittion{Mg}{II} absorption suggests that the SFR of the host galaxy is indeed correlated with the large-scale distribution of cold gas in the circumgalactic medium. Taken together, our data indicate that the covering fraction ($f_{c}$) of cold gas with W$_{r}^{2796}>0.8$\\AA~ may be as large as unity to at least 60 kpc around star-forming galaxies at $1\\lesssim z\\lesssim2$. Given the small number of objects available for this study, we are unable to place tight constraints on these measurements. However, it is interesting to note that this covering fraction estimate is quite high relative to measurements at lower redshift. A recent examination of the \\brittion{Mg}{II} absorption properties of galaxies in the DEEP2 survey estimated $f_{c}=0.5$ around typical galaxies at $z=1$ \\citep{L11}. At $z=0.5$, the typical covering fraction varies from $f_{c}=0.25$ to $f_{c}=1$, depending on the chosen limits of \\brittion{Mg}{II} equivalent width and impact parameter \\citep{BB91, Bechtold92, Steidel94, TB05}. In the local Universe, $f_{c}$ appears to drop below 0.25 for all galaxies \\citep{BC09}, but remains high among those with high star-formation rates \\citep{Bowen95}. Thus, we seem to be observing a trend of a declining mean $f_{c}$ with time. This is, of course, what one might expect given the fact that the global SFR density is also dropping dramatically from $z\\sim2$. As the fraction of galaxies with haloes rich in cold gas is depleted with time, the average $f_{c}$ should also be expected to decline. If we include all galaxies detected within 150 kpc in this study, regardless of redshift or SFR, we estimate that the covering fraction for W$_{r}^{2796}>0.2$\\AA~ is $f_{c}$(60 kpc) $=0.75$, $f_{c}$(100 kpc) $=0.66$, and $f_{c}$(150 kpc) $=0.63$." }, "1207/1207.2292_arXiv.txt": { "abstract": "{We review the current status of liquid noble gas radiation detectors with energy threshold in the keV range, which are of interest for direct dark matter searches, measurement of coherent neutrino scattering and other low energy particle physics experiments. Emphasis is given to the operation principles and the most important instrumentation aspects of these detectors, principally of those operated in the double-phase mode. Recent technological advances and relevant developments in photon detection and charge readout are discussed in the context of their applicability to those experiments.} ", "introduction": "\\label{sec:Introduction} Liquefied noble gases have attracted the attention of experimental physicists since the middle of the 20th century \\cite{Hutchinson48,Marshall53,Northrop56,Northrop58}. The unique combination of their scintillation properties with the fact that electrons released in the ionization process can remain free to drift across long distances favorably distinguishes these liquids from other dense detector media. Another notable property of the `noble liquids' is the possibility of extracting electrons to the gas phase, where the ionization signal can be amplified through secondary scintillation or avalanche mechanisms. These properties have been extensively studied over the years and significant progress has been made in understanding the underlying physics as well as on development of the associated technologies, notably gas purification, material cleaning, cooling, photon and charge detection, and low noise electronics. A wide spectrum of applications has been considered involving both precise particle tracking (owing to the low electron diffusivity) and calorimetry/spectroscopy (good energy resolution): in high energy physics, $\\gamma$-ray astronomy, neutrinoless $\\beta$$\\beta$-decay, medical imaging and, more recently, dark matter (DM) searches and coherent neutrino scattering (CNS) detection. The latter two applications have much in common from the instrumentation point of view, as we shall discuss below: a low energy threshold and low intrinsic background are important requisites in both cases. In addition, both are low energy processes which benefit from coherence of the scattering across all nucleons in the nucleus, thus enhancing expected event rates very significantly. For that reason, we shall consider them in close connection, despite the fact that direct DM searches are more mature than CNS experiments: while large scale DM detectors using liquid xenon and argon are already running or in advanced stage of construction, the possibility of detecting CNS with those media is still under consideration, and several important questions remain open. Although this review gives a reasonably comprehensive account of efforts to measure low energy signals with liquefied noble gas detectors, its main aim is to show how the noble liquid technology works for this purpose and to discuss some recent advances and improvement attempts. The authors apologize in advance for inevitable incompleteness. More information can be found in other recently published monographs and review papers \\cite{BarabashBolozdynya93,LopesChepel05,AprileBook06,AprileDoke09,BolozdynyaBook10,Akimov11}. The article is organized as follows. After motivating the search for particle dark matter and coherent neutrino-nucleus scattering in Section~\\ref{sec:ScatteringDMnu}, we examine in the following section the main characteristics of the signals expected in liquid xenon and liquid argon targets and how those interactions might be detected. We devote Section~\\ref{sec:RelevantProperties} to a detailed description of the physics involved in the response mechanisms. In Section~\\ref{sec:StateOfTheArt} we review the devices used to detect scintillation light and ionization charge, both those which are well established as well as others under development. In Section~\\ref{sec:DirectDMExperiments} we examine how the different DM search programs have implemented the noble liquid technologies and indicate their status at the time of writing. A brief overview of ongoing efforts towards a first detection of CNS is presented in Section~\\ref{sec:NeutrinoDetection}. We conclude in Section~\\ref{sec:Conclusion}, highlighting the progress in sensitivity achieved by direct DM searches with noble liquids in the last two decades. ", "conclusions": "\\label{sec:Conclusion} Liquefied noble gas technology has been developed into well established radiation and particle detection techniques, having found a wide range of application. Technical challenges have been overcome over the last decade to demonstrate that these media can provide low energy threshold, low background, target mass scalability and stable operation over long periods. In particular, double-phase liquid/gas detectors are now at the forefront of low energy particle physics. They are top contenders for a first direct detection of WIMP dark matter as well as of coherent neutrino-nucleus elastic scattering. In dark matter searches, xenon detectors have made the most progress so far, but liquid argon experiments are also close to data-taking underground; liquid neon setups are being considered. We present in Figure~\\ref{fig:DMResultsSI} WIMP-nucleon exclusion limits reported by liquid noble gas experiments. These have covered over three orders of magnitude in scalar cross-section sensitivity in just over a decade, and there are signs that this rate is accelerating as self-shielding begins to pay off with the larger target masses now being deployed. XENON100 offers presently the tightest experimental constraint from any technology. \\begin{figure}[ht] \\vspace{-5mm} \\centerline{\\includegraphics[width=0.7\\textwidth]{plots/DMResultsSI.pdf}} \\vspace{-3mm} \\caption{Spin-independent WIMP-nucleon scattering cross-section limits (90\\% CL) published from liquid xenon (black) and liquid argon experiments (red); in order of publication: DAMA/LXe~\\cite{Bernabei98}, \\mbox{ZEPLIN-I}~\\cite{Alner05}, \\mbox{ZEPLIN-II}~\\cite{Alner07}, \\mbox{WARP~2.3-l}~\\cite{Benetti08}, XENON10~\\cite{Angle08,Aprile09} (low energy analysis~\\cite{Angle11}), \\mbox{ZEPLIN-III}~\\cite{Akimov12a} and XENON100~\\cite{Aprile12a}. The parameter space favored by Constrained MSSM after 1~fb$^{-1}$ of LHC data is shown in green~\\cite{Buchmueller11}. The DAMA/NaI annual modulation result~\\cite{Bernabei08} interpreted as a nuclear recoil WIMP signal in~\\cite{Gondolo09} is also shown (in blue) for reference.} \\label{fig:DMResultsSI} \\end{figure} Prospects for a first detection of CNS have also improved due to the excellent sensitivity afforded by the ionization response in double-phase detectors. This quest goes hand in hand with the search for light WIMPs, both involving sub-keV detection thresholds in scalable target technologies. A first measurement of CNS may be close, although some challenges remain in controlling neutron backgrounds at the neutrino sources. Practically all current dark matter experiments using the double-phase technique rely on measurement of the ionization signal via secondary scintillation in a uniform electric field detected by a large number of photomultiplier tubes. The quest for sensitivity stimulated development of new PMTs capable of operating at cryogenic temperatures and having high quantum efficiency for xenon VUV light. In the last two decades, the quantum efficiency at these wavelengths has doubled. Progress in reducing the radioactivity background has been even more remarkable --- a factor of $\\sim$1,000 has been achieved over this period. Alternative readout techniques, such as solid state photon detectors and micro-pattern detectors for charge amplification, have continued to be developed. The case for these approaches over the traditional PMT readout has not been made so far --- at least at low energies. However, at higher energies and especially for very large detectors new solutions may be advantageous or even necessary (e.g.~LAr TPCs for neutrino detection). The success of experiments that rely on the double-phase technique depends on a good understanding of the physics involved in the detection process. In our view most processes are now well understood; nevertheless, some knowledge gaps remain in the keV and sub-keV energy regions, in particular regarding the response to nuclear recoils. Exploration of this energy regime will be essential to enable a first measurement of CNS, and these experiments are now contributing to this effort. In the year 2012 we celebrated the silver jubilee of direct dark matter searches~\\cite{Ahlen87}. The noble liquids have now taken a clear lead in trying to answer this most important of scientific questions. The authors conclude by reaffirming their belief that the double-phase technique will maintain a strong position in low energy particle physics in the future." }, "1207/1207.2127.txt": { "abstract": "{}{}{}{}{} % 5 {} token are mandatory \\abstract % context heading (optional) % {} leave it empty if necessary {The spin rate of stars evolves substantially during their lifetime, due to the evolution of their internal structure and to external torques arising from the interaction of stars with their environments and stellar winds.} % aims heading (mandatory) {We investigate how the evolution of the stellar spin rate affects, and is affected by, planets in close orbits, via star-planet tidal interactions.} % We investigate for two different star masses if two extreme spin evolutions can lead to noticeably different orbital evolutions for planets.} % methods heading (mandatory) {We used a standard equilibrium tidal model to compute the orbital evolution of single planets orbiting both Sun-like stars and very low-mass stars (0.1 $\\Msun$). We tested two stellar spin evolution profiles, one with fast initial rotation (1.2 day rotation period) and one with slow initial rotation (8 day period). We tested the effect of varying the stellar and planetary dissipation and the planet's mass and initial orbital radius. } % results heading (mandatory) {For Sun-like stars the different tidal evolution between initially rapidly and slowly rotating stars is only evident for extremely close-in gas giants orbiting highly dissipative stars. However, for very low mass stars the effect of initial rotation of the star on the planet's evolution is apparent for less massive ($1 \\Mearth$) planets and for typical dissipation values. We also find that planetary evolution can have significant effects on the stellar spin history. In particular, when a planet falls on the star it makes the star spin up.} % conclusions heading (optional), leave it empty if necessary {Tidal evolution allows to differentiate the early behaviors of extremely close-in planets orbiting either a rapidly rotating star or a slowly rotating star. The early spin-up of the star allows the close-in planets around fast rotators to survive the early evolution. For planets around M-dwarfs, surviving the early evolution means surviving on Gyr timescales whereas for Sun-like stars the spin-down brings about late mergers of Jupiter planets. In light of this study, we can say that differentiating between one spin evolution from another given the present position of planets can be very tricky. Unless we can observe some markers of former evolution it is nearly impossible to distinguish the two very different spin profiles, let alone intermediate spin profiles. Though some conclusions can still be drawn from statistical distributions of planets around fully convective M-dwarfs . However, if the tidal evolution brings about a merger late in its history it can also entail a noticeable acceleration of the star in late ages, so that it is possible to have old stars that spin rapidly. This raises the question of better constraining the age of stars.} ", "introduction": "The spin rate is an important quantity for the evolution of a star and also for the evolution of any planets orbiting close-in. The parameter that governs the direction of tidal evolution for a planet orbiting a star (or a satellite orbiting a planet) is the initial semi-major axis with respect to the corotation radius, the orbital radius where the orbital period matches the central body's spin period. For a planet interior to the corotation radius the planet's mean motion is faster than the primary's rotation, so the tidal bulge raised by the planet on the primary lags behind the position of the planet. The planet feels a drag force that slows it down and causes its orbital radius to shrink, in some cases leading to an eventual merger with the primary. However, for a planet exterior to the corotation radius, the tidal bulge on the star is in advance with respect to the position of the planet and tidal forces push the planet outward. %Knowing the spin evolution of the parent star is therefore crucial to understanding the orbital evolution of planets \\textbf{ that are tidally interacting with their host stars}. Coupling the evolution of radius and spin of the central body to the orbital evolution of planets has been done for brown dwarfs by \\citet{Bolmont2011}. %In this work wind braking was considered negligible on account of the combination of very low ionization fraction and high densities, which results in very large resistivities and thus efficient magnetic field diffusion \\citep{Mohanty2002}. However, \\citet{ReinersBasri2008} claims that wind braking is still important for BDs. %Here the study is conducted for fully convective low mass stars and Sun-like stars. The rotational evolution of Sun-like stars can be described in 3 main stages: the pre-main sequence (PMS) stage, the zero-age main sequence (ZAMS) approach and the main sequence (MS) relaxation \\citep{Bouvier2008}. During the PMS stage the young stars are observed to have a range of spin periods, typically from a few to $\\sim 10$ days, and there is evidence that a highly efficient braking mechanism is at work \\citep{Herbst2007}. It is still not clear what mechanisms are responsible for the observed distribution and the angular momentum loss, but these may be due to the interaction between the star and surrounding accretion disk \\citep[e.g.,][]{GhoshLamb1978, Shu1994, MattPudritz2005, Matt2010}, powerful stellar winds \\citep{Hartmann1982, Hartmann1989, ToutPringle1992, PaatzCamenzind1996,MattPudritz2005a,MattPudritz2008,Matt2012}, or other processes. Toward the end of the PMS phase, the fastest stars in the observed distributions appear to spin up in a way consistent with angular momentum conservation, while the rotation rates of the slowest rotators does not appear to change significantly. Thus, near the ZAMS, the spin period distributions are the widest, typically ranging from a few hours to $\\sim 10$ days \\citep[e.g.,][]{Bouvier1997, Bouvier2008}. Once on the main sequence, the stellar structure evolves slowly enough that the torque from ordinary stellar winds becomes important. Thus, on gigayear timescales, the average spin rates decrease \\citep{Skumanich1972}, and the range of observed spin rates narrows. To bracket the range of observed stellar spin rates, we consider here two populations: initially fast rotators, whose evolution follows the upper envelope of the observed spin rate distributions, and slow rotators that follow the lower envelope. %\\bseanrm{There are 3 main stages in the rotational evolution of Sun-like stars: the pre-main sequence (PMS) stage, the zero-age main sequence (ZAMS) approach and the main sequence (MS) relaxation \\citep{Bouvier2008}. During the PMS stage the young star is magnetically connected to the accretion disk. This connection removes angular momentum from the star and prevents it from spinning up as it contracts \\citep{GhoshLamb1978,Shu1994,MattPudritz2005}. During this \"disk locking\" period the rotation rate of the star remains constant as has been observed \\citep{Rebull2004,Rebull2006}. When the disk dissipates, the star spins up as it contracts and can reach rotation periods as short as $3$~hours \\citep{Bouvier1997}. Depending on the initial rotation period and the disk lifetime, a star can have a wide range of rotation periods when it approaches the ZAMS. Observations of young stars clusters showed that the stellar rotation period at $1$~Myr can be as short as $\\sim 1$~day and as long as $\\sim 10$~days \\citep{Bouvier2008}. To bracket this range we consider here two populations: initially fast rotators whose rotation period is of $1.2$~days and initially slow rotators whose rotation period is of $8$~days. Finally, when the star begins its MS stage, stellar magnetized winds become strong enough to remove angular momentum from the star \\citep{Skumanich1972,Hartmann1982,Hartmann1989} \\citep{ToutPringle1992,PaatzCamenzind1996,Matt2010}.} %In this work, we also address the subject of M-dwarfs. During the PMS and approach to ZAMS, the observed spin period distributions of M-dwarfs is qualitatively similar to that of Sun-like stars. Observations of young clusters constrain the rotation period of low-mass stars younger than a few $\\times 100$~Myr \\citep{Stassun1999,Herbst2001,Irwin2008b} but that approach fails for old clusters due to the faintness of old M-dwarfs. Nonetheless, old slowly-rotating M-dwarfs have been detected \\citep{Benedict1998,Kiraga2007,Charbonneau2009}. Contrary to Sun-like stars that are mostly radiative except for a small (in terms of mass) convective region at the surface, very low mass stars ($M_\\ast<0.35\\Msun$) are entirely convective \\citep{ChabrierBaraffe1997}. For Sun-like stars with a radiative core, the interface between the core and convective envelope is thought to be important for the magnetic dynamo, whereas in fully convective low mass stars other mechanisms have to be invoked to explain their observed magnetic activity \\citep{ReinersBasri2007}. For example, \\citet{ChabrierKuker2006} showed that mean field modeling can produce a $\\alpha^2$ dynamo, which creates large-scale nonaxisymmetric fields and \\citet{Browning2008} showed that three-dimensional nonlinear magnetohydrodynamic simulations of the interiors of fully convective M-dwarfs can also produce a large-scale dynamo. %A star which is mostly radiative and a star which is mostly convective will not have the same magnetic field features \\citep{Morin2010} % % %The number of observations of young cluster has increased over the past decade allowing to constrain the rotation period of objects of less than a few $100$~Myr \\citep{Stassun1999,Herbst2001,Irwin2008b}, observations of old clusters are scarce due to the faintness of old M-dwarfs . However old slow rotating M-dwarfs have been detected \\citep{Benedict1998,Kiraga2007,Charbonneau2009}. The coupling between stellar spin history and tidal evolution has been studied by \\citet{Zahn1994} for close binaries and by \\citet{Dobbs-Dixon2004} for short period planets. Individual systems where tidal interactions are thought to have played a role have also been the subject of various studies. \\citet{Lin1996} proposed that the planet orbiting 51 Peg stopped its disk-induced inward migration because of its presence outside corotation before the disk dispersal. Individual systems of the OGLE survey have been studied by \\citet{Paetzold2004}. Some studies give constraints for stellar dissipation, such as \\citet{CaronePaetzold2007} for OGLE-TR-56b and by \\citet{Lanza2011} for the CoRoT-11 system. In this study, we try to have a more general and systematic approach of the effect of the stellar spin evolution on the tidal evolution of close-in planets. To this end, we couple stellar evolutionary models \\citep{Baraffe1998}, wind parametrization \\citep{Bouvier2008,Irwin2011} and tidal evolution. We consider two limiting cases for the stellar spin evolution that correspond to a star whose initial rotation is either very fast or very slow. These different evolutionary paths can be seen in Figure 1 of \\citet{Bouvier2008} for Sun-like stars, or in Figures 13 to 15 of \\citet{Irwin2011} for M-dwarfs . The slowly rotating stars begin with a rotation period of $8$~days and the fast rotating stars with a period of $1.2$~days. Both fast and slow rotators evolve as explained in \\citet{Bouvier1997}, where the loss of angular momentum due to the stellar wind is quantified given different star-dependent parameters among which is the rotation rate of the star. The higher the spin of a star, the more active the star, the stronger the winds and the stronger the braking. Here we use a standard equilibrium model to study the tidal evolution of planets orbiting stars. Our paper is structured as follows. The tidal and star evolutionary models are briefly discussed in Section \\ref{model1}. Some preliminary analysis based on order of magnitude study are made in Section \\ref{wtf} before giving the results of the tidal evolution of planets around the two types of stars considered here in Section \\ref{Results}. Finally, in Section \\ref{Discussion} we discuss the difficulty of linking the results of tidal evolution and observations and the important effect of late mergers on the rotation rate of the stars. %%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%% %%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%% ", "conclusions": "\\label{Discussion} %%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%% We have shown that the stellar spin history affects the tidal evolution of close-in planets, albeit in a confined region of parameter space. In cases where the stellar spin history does matter, the difference between two close-in planets -- one orbiting an initially fast-rotating star and the other orbiting an initially slow-rotating star -- comes from an early phase of outward tidal migration for the planet around the fast rotator (caused by the star's closer-in corotation distance). This phase of outward migration does not occur for the planet orbiting the slow rotator. This outward migration and the decrease of the radius of the star weaken the later tidal evolution and effectively delay or sometimes even prevent the later in-spiralling of planets onto their stars. At later times both slow and fast rotators spin down due to stellar winds \\citep{Skumanich1972,Bouvier1997,Bouvier2008,Irwin2011}, so the orbital history of close-in planets orbiting old stars depends on something that is not directly observable: the stellar spin evolution. For Sun-like stars the stellar spin history affects tidal evolution only in relatively extreme circumstances. In particular, the spin history has a strong effect if stars are very dissipative, with dissipation rates $\\sigma_\\ast$ of about $1000$ times the fiducial value \\citep{Hansen2010}. For strongly dissipative stars, there are differences in tidal evolution for planets orbiting initially slowly- vs. rapidly rotating stars for very close-in ($a \\lesssim 0.05$~AU), massive planets ($M_p \\gtrsim M_J$). In contrast, the stellar spin history plays a role in a much wider region of parameter space for $0.1 \\Msun$ stars, mainly because these stars are fully convective and so we think they are much more dissipative, with dissipation factors assumed to match those of brown dwarfs. For these stars the differences in tidal evolution for planets orbiting initially slowly- vs. fast-rotating stars are apparent for mean dissipation values, for planets out to $\\sim 0.02$~AU, and for planets with masses as small as $1 \\Mearth$. %For both $0.1 \\Msun$ and $1 \\Msun$ stars the difference between the two spin profiles -- if there is one -- can be seen for very close-in planets. In the case of a fast rotator, the planet being closer to the corotation radius will have more chance to eventually cross it as the star spins up due to contraction. However for a slow rotator, the corotation distance is located farther so close-in planets will fall on the star due to the stellar tide. Low mass stars and Sun-like stars have different dissipation factors and different radii, so the evolution timescales are different and evolve differently (see Figs. \\ref{ordermag_bis} and \\ref{ordermag3}). The early tidal interaction is stronger for planets around very low mass stars, and the difference between fast rotator profile and slow rotator profile is apparent for planets of one Earth mass with mean dissipation values. After some time, the system freezes in a given state because the radius of the star has shrunk too much for the tidal evolution to occur on less than $10$~Gyr timescale. For Sun-like stars, the tidal evolution for mean dissipation factor occurs on very long timescales, which is why much more massive planets and higher stellar dissipation rates are needed to produce stellar spin-driven differences in tidal evolution. %So in order to see some difference between the two profiles, one has to consider massive planets such a Jupiter-mass planet and very dissipative stars. %For M-dwarfs and Sun-like stars we can observe a significant effect of the planets on the spin of the star. The inward migration of a Jupiter orbiting inside the corotation radius of an initially slow rotating $0.1 \\Msun$ star can lead to a significant spinning-up of the star. By the transfer of angular momentum from the planet's orbit, an initially slow rotating star can become a fast rotating star. This effect is more dramatic if the planet actually falls on the star. \\citet{Irwin2011} used different wind parametrization for fast or slow rotators so they can infer from the present spin rate if the star was initially fast rotating or not. However we show here that it might not be that straight forward. If the star experienced a merger with a planet it can modify the rotation rate of the star and change the slow rotator into a fast rotator. Massive planets orbiting very low-mass stars with high dissipation rates ($\\sigma_\\ast \\times 1000$) can create systems in perfect synchronization where the spin of the star is equal to the spin of the planet (Fig. \\ref{JUP}). However, the equilibrium is not stable and the system departs from it as the star spins up due to contraction or spins down due to the stellar winds. This strongly alters the stellar spin profile because the star can be efficiently spun down by a planet initially located outside the corotation radius or spun up by a planet interior to corotation. Unfortunately, hot Jupiters around M-dwarfs are extremely rare due to the inefficiency of the planets formation processes around low mass stars \\citep{Laughlin2004,IdaLin2005,Kennedy2008}. Only three hot Jupiters are known to exist around stars with masses less than $0.7 \\Msun$ \\citep{Pepe2004,Hellier2011,Borucki2011,Johnson2012} and none around stars with masses as small as $0.1 \\Msun$. Hot Jupiters around very low mass stars remain to be detected. %For Sun-like stars, in order to see the difference between the two spin profiles one need massive planets such as Jupiter very close to the star and strong stellar dissipations. The cases for Earth and Uranus mass planets showed little difference between the two spin profiles, the planets do not experience this extreme behavior of both outwards and then inwards migration. The tides raised by those planets on the star are never going to be strong enough to have noticeable effects in less than $10$~Gyr timescales. For $0.8\\Msun$ stars, the results are qualitatively the same although the stellar tide is weaker for the same planet mass, the effect is thus weaker but exists nonetheless for planets of mass $M_p \\geq \\Mjup$. A statistical distribution of planets around fully convective M-dwarfs could constrain the tidal dissipation factor $\\sigma_\\ast$. Specifically, from the location of the inner edge of the planetary semi-major distribution one can infer a inferior limit for the dissipation factor. The more distant the inner edge, the more dissipative the star. However, in order to draw these conclusions, one needs a good estimate of the stellar ages. For Sun-like stars such conclusions cannot be made because, if the dissipation rate is high enough to affect the orbital evolution at early times, significant tidal evolution still takes place at late times, as well. Slow rotators and fast rotators have similar evolutions after a few $10^8$~yrs so the observation of a hot Jupiter orbiting a star of known age and known dissipation would not allow us to infer if the stellar spin history. Indeed, in both cases, different initial semi-major axis can lead to the same observed semi-major axis. One can imagine trying to infer a planet's orbital evolution from the composition of its atmosphere to know if the planet came from a \"cold\" region ($0.04$~AU) or a \"hot\" region ($0.03$~AU). Unfortunately, this exercise is fraught with uncertainties in both the expected atmospheric composition of planets that form at different orbital distances and the tidal parameters ($\\Omega_{\\ast,0}$, $\\sigma_\\ast$, the age of observed stars\u00c9). Nonetheless, we emphasize the planets crashing on the star at late ages can entail a significant spin-up of the star and create a population of old fast rotating stars. The spin-up of the star due to a merger has been pointed out in \\citet{Levrard2009}, where they found that planets falling on their host star due to tides never reach a tidal equilibrium. \\citet{Jackson2009} also addressed the problem of tidally induced mergers and the effect of these mergers on the parent star. They also found that a considerable spin-up is to be expected and also a change in stellar composition. Planet-star mergers thus may confuse stellar age determinations. In general, fast rotators are thought to be young, although we have shown that a merger can lead to old, fast rotating stars that would mimic many of the characteristics of young stars. An independent determination of the age of observed stars is therefore very important, especially for fast rotators. %It would be difficult to determine which kind of spin evolution profile a star underwent given the distribution of the planet's semi-major axis. Unless we can know the age of the star with enough precision, and even then if we observe a Jupiter around a star of a known age and a known dissipation, we could not differentiate if the star was a fast rotator or a slow rotator. Indeed, in both cases, different initial semi-major axis can lead to the same observed semi-major axis. Maybe, we could infer from the composition of the atmospheres the previous evolution, if the planet came from a \"cold\" region ($0.04$~AU) or a \"hot\" region ($0.03$~AU). Unfortunately, this would require to know the dissipation factor of the star quite precisely which is far from being the case. There are too many unknown quantities ($\\Omega_{\\ast,0}$, $\\sigma_\\ast$, the age of observed stars...) to be able to draw conclusions about this question. %\\textbf{The influence of tides-induced spin up has been studied by \\citet{Pont2009}, where he considers transiting systems for which the star shows excess rotation. The spin-up of the star is due to the tidal interaction between the planet and the star. The systems for which the star rotate too fast will evolve towards a merger, making their rotation decrease even more. Most hot Jupiters are located inside the corotation radius so they are scheduled for merger eventually. } %\\citet{Chambers1999} %%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%% %%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%% %" }, "1207/1207.0722_arXiv.txt": { "abstract": "Close-in exoplanets with highly eccentric orbits are subject to large variations in incoming stellar flux between periapse and apoapse. These variations may lead to large swings in atmospheric temperature, which in turn may cause changes in the chemistry of the atmosphere from higher CO abundances at periapse to higher CH$_{4}$ abundances at apoapse. Here we examine chemical timescales for CO$\\rightleftarrows$CH$_{4}$ interconversion compared to orbital timescales and vertical mixing timescales for the highly eccentric exoplanets HAT-P-2b and CoRoT-10b. As exoplanet atmospheres cool, the chemical timescales for CO$\\rightleftarrows$CH$_{4}$ tend to exceed orbital and/or vertical mixing timescales, leading to quenching. The relative roles of orbit-induced thermal quenching and vertical quenching depend upon mixing timescales relative to orbital timescales. For both HAT-P-2b and CoRoT-10b, vertical quenching will determine disequilibrium CO$\\rightleftarrows$CH$_{4}$ chemistry at faster vertical mixing rates ($K_{zz}>10^7$ cm$^2$ s$^{-1}$), whereas orbit-induced thermal quenching may play a significant role at slower mixing rates ($K_{zz}<10^7$ cm$^2$ s$^{-1}$). The general abundance and chemical timescale results -- calculated as a function of pressure, temperature, and metallicity -- can be applied for different atmospheric profiles in order to estimate the quench level and disequilibrium abundances of CO and CH$_{4}$ on hydrogen-dominated exoplanets. Observations of CO and CH$_{4}$ on highly eccentric exoplanets may yield important clues to the chemical and dynamical properties of their atmospheres. ", "introduction": "For planets with high eccentricity, the large variations in flux received from their host stars may yield substantial variations in atmospheric temperature and dynamical behavior during the course of an orbit \\citep[][]{langton2008,laughlin2009,iro2010,cowan2011,kane2011,rauscher2012,lewis2012}. In some cases, orbit-induced temperature variations may be large enough to produce a significant shifts in the chemical behavior of the planet. For example, the swing between high atmospheric temperatures at periapse and lower atmospheric temperatures at apoapse may shift equilibrium chemistry predictions from relatively higher CO abundances at periapse (toward a CO-dominated atmosphere) to relatively higher CH$_{4}$ abundances at apoapse (toward a CH$_{4}$-dominated atmosphere). Large changes in the abundances of CO and CH$_{4}$ are of particular interest because these compounds strongly influence the spectral properties of exoplanet atmospheres \\citep[e.g.,][]{seager2000,burrows2005apjl,charbonneau2007,charbonneau2008,fortney2007,barman2008,swain2008,swain2009apj,swain2009apjl,swain2010,desert2009,madhusudhan2009,madhusudhan2011,madhusudhan2011nature,Stevenson2010,tinetti2010,tinetti2010faraday,beaulieu2011,knutson2011,lee2012,shabram2011,waldmann2012}. Phase-dependent variations in the planetary spectrum \\citep[e.g.,][]{barman2005,fortney2006dyn,knutson2007nature,knutson2012,showman2008,showman2009,cowan2011,lewis2012} may therefore also reflect changes in chemical composition of highly eccentric transiting exoplanets (see Table \\ref{tab: exoplanets}), particularly at wavelengths sensitive to CO and CH$_{4}$ \\citep[e.g.,][]{lewis2012} However, the extent of temperature-dependent variations in carbon chemistry throughout the orbit -- and whether equilibrium chemistry at a given altitude prevails over orbital timescales -- also depends upon the rate of CO$\\rightleftarrows$CH$_{4}$ interconversion relative to the time elapsed between periapse and apoapse. Thermochemical equilibrium can be maintained throughout the orbit only if chemical timescales are less than orbital timescales. To study this effect, \\citet{iro2010} estimated the CO$\\rightleftarrows$CH$_{4}$ interconversion timescale for HD 80606b and HD 17156b using the kinetics of \\citet{bezard2002} and found that orbital timescales are generally much shorter than chemical timescales -- indicative of disequilibrium chemistry -- at pressure levels where orbit-induced temperature variations are expected to be significant. This behavior will occur on objects which are expected to have relatively low atmospheric temperatures (and therefore sluggish reaction kinetics) near apoapse, leading to orbit-induced thermal quenching wherein an equilibrium composition achieved near periapse survives to become a disequilibrium composition at apoapse. The observed properties of highly eccentric exoplanets are subject to numerous variables including thermochemical and photochemical reaction rates, convective transport, and horizontal dynamical and radiative timescales. Here, we focus specifically on the role of thermochemical quench chemistry in response to vertical transport and eccentricity-induced atmospheric temperature variations, as this behavior may strongly influence observable CO and CH$_{4}$ abundances (even if photolysis occurs at higher altitudes). For simplicity, the relevant timescale for the temperature variation is taken to be $0.5p$ (where $p$ is the orbital period), which describes the time elapsed between temperature swings at periapse and apoapse. We first calculate the abundances of CO and CH$_{4}$ and CO$\\rightleftarrows$CH$_{4}$ chemical timescales as a function of pressure, temperature, and metallicity in a solar-composition gas, using updated kinetics for CO$\\rightleftarrows$CH$_{4}$ interconversion \\citep{visscher2010icarus,moses2011,visscher2011}. The results are then compared to orbital and mixing timescales and pressure-temperature profiles of individual objects to estimate the quench levels and abundances of CO and the CH$_{4}$ as disequilibrium species. The primary objective of this study is to discuss a convenient method for exploring thermal quenching processes in the CO$\\rightleftarrows$CH$_{4}$ system in exoplanet atmospheres. Although we focus on CO-abundant HAT-P-2b and CH$_{4}$-abundant CoRoT-10b as specific examples, the abundance and timescale results presented here may in principle be applied to any H$_{2}$-dominated substellar object that is subject to orbit-induced temperature variations. \\begin{deluxetable}{lccr@{.}lccc} \\tablecaption{Highly-Eccentric ($e>0.3$) Transiting Exoplanets} \\tablewidth{0pt} \\tablehead{ \\colhead{Object} & $a$(AU) & $e$ & \\multicolumn{2}{c}{$p$(days)} & $p_{s}$(days) & $M{_\\textrm{J}}$ & $R_{\\textrm{J}}$} \\startdata HD 80606b & 0.447 & 0.934 & 111&44 & 1.7 & 3.9 & 1.03\\\\ HD 17156b & 0.163 & 0.682 & 21&22 & 3.6 & 3.3 & 1.02\\\\ CoRoT-10b & 0.105 & 0.530 & 13&24 & 4.3 & 2.8 & 0.97\\\\ HAT-P-2b & 0.068 & 0.517 & 5&63 & 1.9 & 8.9 & 1.16\\\\ HAT-P-34b & 0.068 & 0.440 & 5&45 & 2.4 & 3.4 & 1.20\\\\ HAT-P-17b & 0.088 & 0.346 & 10&34 & 5.9 & 0.5 & 1.01\\\\ WASP-8b & 0.080 & 0.310 & 8&16 & 5.1 & 2.1 & 1.04\\\\[-1.5mm] \\enddata \\tablecomments{Data from \\citet{wright2011}. The pseudo-synchronous rotation period $p_{s}$ was calculated using the expressions of \\citet{hut1981} as presented in \\citet{iro2010}.} \\label{tab: exoplanets} \\end{deluxetable} ", "conclusions": "Equilibrium abundances for CO and CH$_{4}$ and chemical timescales for CO$\\rightleftarrows$CH$_{4}$ interconversion were calculated as a function of pressure, temperature, and metallicity. A comparison of the abundance and timescale plots with atmospheric pressure-temperature profiles can be used to estimate the quench levels and abundances of CO and CH$_{4}$ for a variety of substellar objects. In principle, this approach can be applied for any substellar object with an H$_{2}$-dominated atmosphere that is subject to eccentricity-induced temperature variations. Overall, orbit-induced thermal quenching tends to favor CO over CH$_{4}$ because CO is the higher-temperature species and because chemical reactions proceed more rapidly at periapse than at apoapse. Moreover, equilibrium can be maintained to higher altitudes on warmer CO-dominated objects than on cooler CH$_{4}$-dominated objects, with respect to orbital and vertical mixing timescales. Whether vertical quenching or orbit-induced thermal quenching governs disequilibrium CO$\\rightleftarrows$CH$_{4}$ chemistry depends upon the orbital timescale relative to the mixing timescale (as a function of $K_{zz}$) along the atmospheric profile. In some cases, vertical quenching along the periapse profile may govern the disequilibrium abundances of CH$_{4}$ and CO throughout the upper atmosphere over the entire orbit. For both HAT-P-2b and CoRoT-10b, the effect of orbit-induced thermal quenching is roughly equivalent to transport-induced quenching assuming $K_{zz}\\sim10^{7}$ cm$^{2}$ s$^{-1}$. For lower $K_{zz}$ values ($<10^{7}$ cm$^{2}$ s$^{-1}$), thermal quenching may have a significant effect, whereas quenching via vertical transport will determine disequilibrium abundances of CO and CH$_{4}$ at higher $K_{zz}$ values ($>10^{7}$ cm$^{2}$ s$^{-1}$). For CO$\\rightleftarrows$CH$_{4}$ chemistry, differences in the quenching mechanism may result in large differences in the abundances of disequilibrium carbon-bearing species. Refinements to atmospheric structure models and improved observational estimates of CO and CH$_{4}$ on highly eccentric exoplanets may yield important clues to the chemical and dynamical behavior of their atmospheres. Further development of the chemical models, including the consideration of photochemical production and loss rates throughout the orbit, may likewise provide improved abundance estimates of disequilibrium species throughout the upper atmospheres of highly eccentric transiting exoplanets." }, "1207/1207.2727_arXiv.txt": { "abstract": "We report the results of a study exploring the stellar populations of 13 luminous (L $>$ 1.2L$^*$), spectroscopically confirmed, galaxies in the redshift interval $5.5$ 300 Myr) are inferred for these objects, indicating a formation redshift at $z>8$. This property was first reported for two $i-$drop galaxies in \\cite{Eyles2005}, with their spectral energy distribution (SED) fitting indicating ages in the range 250-650 Myr, and stellar masses of order 20\\% of current-day $L^*$ galaxies. A later analysis of a larger sample showed that 40\\% of their sample displayed evidence for substantial 4000\\AA/Balmer spectral breaks \\citep{Eyles2007}. More recently, \\cite{RichardJohan2011} reported the discovery of a lensed galaxy at $z=6.027$ which also appears to have a strong Balmer-break, consistent with a mature stellar population. If confirmed, a substantial population of galaxies at $z>6$ with mature stellar populations has important consequences for star formation at $z>8$ and its contribution to reionization. SED-fitting techniques, however, often employ templates that only account for stellar emission (eg. \\citealt{Bruzual2003,Maraston2005}) and recently several groups have investigated how the inclusion of nebular emission in the fitting can affect the derived parameters (e.g. \\citealt{Robertson2010a,Schaerer2010a,Labbe2010,Ono2010a}). For example, \\cite{Labbe2010} reported difficulties in reconciling the observed Balmer break with a very blue ultraviolet (UV) continuum slope ($\\beta\\sim-3$; $f_{\\lambda}\\propto \\lambda^{\\beta}$) observed in a stack of the photometry for the faintest galaxies in a sample of $z\\sim7$ LBGs ($H_{160, AB}>27.5$) taken from the Hubble Ultra-deep Field (HUDF) and Early Release Science (ERS) field. They suggest that episodic star formation may be required to match the colours of these objects (being simultaneously blue in the rest-frame UV and red in colours spanning the 4000\\AA/Balmer break), as well as the contribution of nebular emission lines that are not included in the standard population synthesis codes. An illustration of the potentially dramatic effect that nebular emission can have on derived stellar masses and ages is provided by \\cite{Ono2010a}, in their fitting of stacked multi-wavelength observations of Lyman alpha emitters (LAEs) at $z\\sim5.7$ and $z\\sim6.6$ within the Subaru/ XMM-Newton Deep Survey (SXDS) field. The derived masses from models with maximum nebular emission (escape fraction of ionising photons, $f_{esc}=0$), are more than an order-of-magnitude lower than those based on only stellar emission (e.g. $M_*\\sim3\\times10^7\\,\\mathrm{M}_{\\odot}$ compared to $M_*\\sim5\\times10^8\\,\\mathrm{M}_{\\odot}$ for the $z\\sim5.7$ stack) and the derived ages are also much younger ($\\sim3$ Myr compared to $\\sim300$ Myr). Recent results from cosmological hydrodynamic simulations of reionisation predict that Lyman-break selected galaxies have young ages, low intrinsic reddening and sub-solar metallicities at high redshifts (e.g. $\\sim50-150$ Myr at $6$$2\\sigma$) are detected for about 55\\% of the sources. These shifts are even better aligned with the jet direction, deviating from the latter by less than $30\\degr$ in over 90\\% of the cases. There is an indication that the core shift decreases with increasing redshift. Magnetic fields in the jet at a distance of 1 parsec from the central black hole, calculated from the obtained core shifts, are found to be systematically stronger in quasars (median $B_1\\approx0.9$~G) than those in BL~Lacs (median $B_1\\approx0.4$~G). We also constrained the absolute distance of the core from the apex of the jet at 15~GHz as well as the magnetic field strength in the 15~GHz core region.} {} ", "introduction": "\\label{intro} Bipolar relativistic outflows (jets) in active galactic nuclei (AGN) are formed in the immediate vicinity of the supermassive central black hole and become detectable at distances of $\\gtrsim$$100$ gravitational radii ($R_g=GM_{bh}/c^2$) at millimeter wavelengths \\citep{Junor_M87,Lobanov_07,Hada_M87}. The jets take away a substantial fraction of the energy and angular momentum stored in the accretion flow \\citep{Hujeirat_03} and spinning central black hole \\citep{Koide_02,Komissarov_05}. As discussed by \\cite{Vlahakis_04}, a poloidal-dominated magnetic field embedded in the accretion disk or in the black hole ergosphere is wound-up into toroidal loops that may provide effective jet collimation via hoop stress and accelerate the flow by magnetic pressure gradient up to a distance of $\\sim$$10^3-10^5$~$R_g$. Very Long Baseline Interferometry (VLBI) observations provide us with the perfect zoom-in tool to explore AGN jets with a milliarcsecond angular resolution corresponding to parsec-scale linear resolution. Typically, the parsec-scale radio morphology of a bright AGN manifests a one-sided jet structure due to Doppler boosting \\citep[e.g.,][]{BlandfordKonigl79,Kellermann_07,MOJAVE} that enhances the emission of the approaching jet. The apparent base of the jet is commonly called the ``core'', and it is often the brightest and most compact feature in VLBI images of AGN. The VLBI core is thought to represent the jet region, located at the distance $r_\\mathrm{core}$ to the central engine, at which its optical depth reaches $\\tau_\\nu\\approx1$ at a given frequency. At short mm-wavelengths the core may also be the first recollimation shock downstream of the $\\tau = 1$ surface instead of the surface itself. This does not affect our analysis, which uses longer wavelengths. Thus, the absolute position of the radio core is frequency-dependent and varies as $r_\\mathrm{core}\\propto\\nu^{-1/k_r}$ \\citep{BlandfordKonigl79,Koenigl81}, i.e., it shifts upstream at higher frequencies and downstream at lower frequencies (the so-called ``core shift'' effect). The first core shift measurement from VLBI observations was performed by \\cite{Marcaide84}. Recent multi-frequency studies of the core shift effect \\citep{Sullivan_09cs,Fromm10,Sokolovsky_11cs,Hada_M87} showed that $k_r\\approx1$ in most sources and epochs. This is consistent with the \\cite{BlandfordKonigl79} model of a synchrotron self-absorbed conical jet in equipartition between energy densities of the magnetic field and the radiating particle population. Nonetheless, departures in $k_r$ from unity are also possible and can be caused by pressure and density gradients in the jet or by external absorption from the surrounding medium \\citep{L98,Kadler_04}. The frequency-dependent offsets of the core positions can be used for astrophysical studies of ultra-compact AGN jets to calculate the magnetic fields, synchrotron luminosities, total (kinetic and magnetic field) power, maximum brightness temperature and geometrical properties of the jet \\citep{L98}. The core shift effect also has immediate astrometric applications. A typical shift between the radio (4~cm) and optical (6000~\\AA) domains for distant quasars is estimated to be at the level of 0.1~mas \\citep{Kovalev_cs_2008}, which is comparable with the expected positional accuracy of the {\\it GAIA} astrometric mission \\citep{Lindegren96}. Thus, the core shifts are likely to influence not only the positional accuracy of the radio reference frame but also an alignment of optical and radio astrometry catalogs. Moreover, it is natural to expect that opacity properties are variable on a time scale from months to years due to the continuous emergence of new jet components, and especially during strong nuclear flares. Therefore, as discussed by \\cite{Kovalev_cs_2008}, a special coordinated program is required to perform multi-frequency and multi-epoch VLBI observation of a pre-selected source sample to investigate the problem of core shift variability. A major difficulty in measuring the core shift is the accurate registration of the VLBI images taken at different frequencies. The problem stems from the loss of absolute position information in the standard VLBI data reduction path, which involves self-calibration of the station phases. Several approaches have been presented to overcome this difficulty and measure core shifts. One of them is based on relative VLBI astrometry, i.e., phase-referencing to a calibrator source \\citep[e.g.][]{Marcaide84,Lara94,Guirado95,Ros01cs, Bietenholz04,Hada_M87}. This particular technique is resource-consuming and has been used for a limited number of sources only. Another approach is the self-referencing method \\citep{L98,Kovalev_cs_2008,Sokolovsky_11cs}, in which the core shift is derived by referencing the core position to bright optically thin jet features whose positions are expected to be achromatic. Although this method has provided the majority of known core shift measurements, it has a certain limitation. It cannot be applied for faint or smooth jets that lack compact bright feature(s) well separated from the core at different frequencies. A proper alignment of the optically thin parts of the jet can also be accomplished by two-dimensional cross-correlation of the images, initially suggested and performed by \\cite{Walker_2D} for multi-frequency VLBA observations of 3C~84. The algorithm was also discussed by \\cite{Croke_2D} and Fromm et al. (in prep.), and applied by \\cite{Sullivan_09cs} to obtain core shifts in four BL~Lac objects. This approach, in conjunction with source model fitting, presents a more widely applicable method for deriving core shifts (see Sect.~\\ref{s:method} for detailed discussion), which we use in this paper. Another alternative indirect method recently proposed by \\cite{Kudryavtseva_11_cs} is based on an analysis of time lags of flares monitored with single-dish observations. Although it has obvious limitations on the epoch at which the core shift can be measured, the method is promising for highly compact sources, which pose problems for other opacity study methods due to the lack of optically thin jet structure. It is noteworthy that all of the aforementioned techniques provide a comparable accuracy level. To date, only two core shift studies \\citep{Kovalev_cs_2008,Sokolovsky_11cs} have been carried out on large samples. They have shown that the effect is significant for many sources. In this paper, we measure frequency-dependent shifts in the absolute core positions and study the statistical properties of the detected core shift vectors by using a large sample of sources from the MOJAVE (Monitoring Of Jets in Active galactic nuclei with VLBA Experiments) program \\citep{MOJAVE}. We also analyze systematics and discuss the uncertainties of the two-dimensional cross-correlation technique, investigating its properties for different jet morphologies. We constrain the basic physical properties of the jets, such as the magnetic field strength in the core region and at the true base of the flow, the distance from the jet apex to the radio core, as well as the estimate of the central black hole mass from the derived core shifts. Throughout the paper, we assume the power index $k_r=1$, i.e., $r_\\mathrm{core}\\propto\\nu^{-1}$ (see model assumptions for this case above). We use the $\\Lambda$CDM cosmological model with $H_0=71$~km~s$^{-1}$~Mpc$^{-1}$, $\\Omega_m=0.27$, and $\\Omega_\\Lambda=0.73$ \\citep{Komatsu09}. All position angles are given in degrees from north through east. ", "conclusions": "\\label{s:summary} We have implemented a method for measuring the frequency-dependent shift in absolute position of the parsec-scale core and applied it to multi-frequency (8.1, 8.4, 12.1, and 15.4~GHz) VLBA observations of 191 sources performed during 2006 within the MOJAVE program. The method is based on results from (i) image registration achieved by a two-dimensional cross-correlation technique and (ii) structure model fitting. It has proved to be very effective and provided the core shifts in 163 sources (85\\%), with a median of 128~$\\mu$as between 15 and 8~GHz, and 88~$\\mu$as between 15 and 12~GHz. Despite the moderate separation of the observing frequencies, the derived core shifts are significant ($>$$2\\sigma$) in about 55\\% of cases, given an estimated typical uncertainty of 50~$\\mu$as between 15 and 8~GHz, and 35~$\\mu$as between 15 and 12~GHz. The errors are dominated by uncertainties from the two-dimentional cross-correlation procedure, because the relative positional uncertainties of the compact bright cores are at a level of a few microarcseconds. The significant core shift vectors are found to be preferentially aligned with the median jet direction, departing from it by less than $30\\degr$ in more than 90\\% of cases. We used the measured core shifts for constraining magnetic field strengths and core sizes for 89 sources. The magnetic field at a distance of 1~pc from the jet injection point is found to be $\\sim$0.9~G for quasars and $\\sim$0.4~G for BL~Lacs. Extrapolating all the way back by assuming a $B\\propto R^{-1}$ dependence, the magnetic field in the close vicinity of the black hole is about $2\\times10^3$~G. The core sizes, i.e., the distances from the true jet base to its apparent origin at 15.4~GHz are statistically larger in quasars than in BL~Lacs, with a median of 13.2 and 4.0~pc, respectively. At these distances, the magnetic field has a median of 0.07~G for quasars and 0.1~G for BL~Lacertae objects. Future multi-epoch and multi-frequency VLBI observations (including phase-referencing) of a pre-selected sample of sources with prominent core shifts are needed to address the question of the core shift variability, which can be used not only for astrophysical studies but also for astrometric applications. Ideally, these observations should be performed during and after strong nuclear flares that can be detected in advance in higher energy domains, e.g., optical or gamma-ray, and cover a wide range of observing frequencies, extending down to 5 or 2~GHz, where synchrotron self-absorption is essential." }, "1207/1207.1332_arXiv.txt": { "abstract": "In this paper, we continue to examine the fundamental basis for the Friedmann-Robertson-Walker (FRW) metric and its application to cosmology, specifically addressing the question: What is the {\\it proper} size of the visible universe? There are several ways of answering the question of size, though often with an incomplete understanding of how far light has actually traveled in reaching us today from the most remote sources. The difficulty usually arises from an inconsistent use of the coordinates, or an over-interpretation of the physical meaning of quantities such as the so-called proper distance $R(t)=a(t)r$, written in terms of the (unchanging) co-moving radius $r$ and the universal expansion factor $a(t)$. In this paper, we prove for the five non-trivial FRW metrics with constant spacetime curvature that, when the expansion began from an initial singularity, the visible universe today has a proper size equal to $R_{\\rm h}(t_0/2)$, i.e., the gravitational horizon at half its current age. The exceptions are de Sitter and Lanczos, whose contents had pre-existing positions away from the origin. In so doing, we confirm earlier results showing the same phenomenon in a broad range of cosmologies, including $\\Lambda$CDM, based on the numerical integration of null geodesic equations through an FRW metric. ", "introduction": "Recent efforts aimed at providing a better understanding of the fundamental basis for the Friedmann-Robertson-Walker (FRW) metric and its application to cosmology have uncovered several previously unrecognized properties relevant to the interpretation of cosmological data. The standard model of cosmology ($\\Lambda$CDM) is only marginally consistent with these developing theoretical considerations, reflected in the growing tension between its predictions and what is actually observed, both in the cosmic microwave background (CMB) and the unexpected early appearance of quasars and galaxies at high redshift, and in the matter distribution, gamma-ray burst rate and Type Ia supernovae in the nearby Universe. For example, the use of Birkhoff's theorem and its corollary \\cite{birkhoff23} has shown that the Universe possesses a gravitational horizon (with radius $R_{\\rm h}$) coincident with the better known Hubble sphere emerging empirically from the observed universal expansion \\cite{melia07}. This new insight has allowed us to consider the impact of strictly adhering to the requirements of both the Cosmological principle and Weyl's postulate \\cite{weyl23}, which together force $R_{\\rm h}$ to always be equal to $ct$, the distance light could have traveled during a time $t$ since the big bang \\cite{melia12}. $\\Lambda$CDM agrees with this constraint only partially, oddly very early in the universal expansion close to the Planck time, and more recently, where the various observations are telling us that $R_{\\rm h}(t_0)\\approx ct_0$ today---but not in between. For a summary of how the current cosmological data compare with the condition $R_{\\rm h}=ct$ and the predictions of the standard model, see references \\cite{spergel03, copi09,melia12a,melia12b,melia13,meliamaier13}. This fundamental approach to the study of the cosmological spacetime has also allowed us to examine the nature of cosmological redshift $z$ in FRW metrics with constant spacetime curvature. We recently showed that the interpretation of $z$ as due to the `stretching' of space is coordinate dependent \\cite{melia12c}. An equally important outcome of this study has been a greatly improved understanding of how null geodesics behave in FRW, allowing us to better appreciate which sources are actually observable today. We recently confirmed the importance of $R_{\\rm h}$ in delimiting the size of the observable universe \\cite{bikwa12,melia12d} by proving that in all cosmologies with an equation-of-state parameter $w\\ge -1$, where the pressure $p$ and density $\\rho$ are related by the expression $p=w\\rho$, no light rays reaching us today could have ever attained a proper distance $R(t)$ greater than $R_{\\rm h}(t_0)$. A principal motivation for the present paper is actually another interesting result that emerged from the numerical integration of the null geodesics in reference \\cite{bikwa12}. There, we showed that for a broad range of cosmologies, including $\\Lambda$CDM, no null geodesics reaching us today (at time $t_0$) could have ever started from, or reached, a proper distance greater than $\\sim ct_0/2$ away from us. Our purpose here is to examine the fundamental basis for this constraint, and we will prove that in FRW metrics with a constant spacetime curvature, the most distant sources we see today---particularly the CMB---emitted their light at time $(1/2)t_0$ from a proper distance $R_{\\rm h}(t_0/2)$ away, which therefore defines the size of the visible universe today. Applied to the CMB, this result may seem paradoxical because the time $t_e$ at recombination was presumably much earlier than $(1/2)t_0$. Needless to say, this issue has itself caused confusion over the years, with some workers believing that light must have therefore traveled a proper distance $c(t_0-t_e)$ in reaching us. For example, a recent recalibration (by $\\Delta t\\sim 2$ Gyr) of the age of extragalactic eclipsing binaries was used to stretch the cosmic distance ladder by $\\sim c\\,\\Delta t$ \\cite{bonanos06}. Similarly, conclusions concerning the Universe's topology are often based on how far light has traveled since the big bang \\cite{cornish03,vaudrevange12}. And an older (often cited) publication on distance measures makes several incorrect assocations between how far light could have traveled and the inferred distance to horizons \\cite{davis04}. But it is easy to demonstrate that the proper distance to a source is not equal to the light-travel distance, and that the difference is merely due to the time dilation between frames moving at relative speeds close to $c$. In other words, we shall see that whereas $t_e$ may be close to $0$ (for, say, the CMB), the corresponding time on clocks at rest with respect to us was dilated significantly to a value $\\sim(1/2)t_0$ ($\\gg t_e$). ", "conclusions": "To fully understand and appreciate the results we have presented in this paper, one must acknowledge the critical role played by the choice of coordinates in describing the expansion of the Universe. We had already seen an example of this, based on how the choice of frames impacts our interpretation of the cosmological redshift $z$ \\cite{melia12c}. We proved earlier that, although $z$ is conventionally calculated directly from the expansion factor $a(t)$, its origin cannot be attributed to an expansion of space when viewed in terms of the FRW metric written in stationary form. We found that $z$ is actually the cosmological version of a lapse function encountered more typically in the context of the Schwarzschild and Kerr metrics. That is, $z$ is simply due to the combined effects of the kinematic expansion and the gravitational acceleration---but only in terms of the proper velocity, calculated using the proper distance and proper time for an individual observer. In this paper, we have expanded our study of the fundamental aspects of the cosmic spacetime by using these alternative sets of coordinates to address another issue that sometimes gives rise to confusion and ambiguity: what is the true size of the visible universe? The question itself is fraught with ambiguity because it goes without saying that to measure a size, one must have a precise definition of distance. General relativity is founded on the basic principle that $c$ is invariant and is measured to have the same value for all observers. But what is often overlooked or forgotten is that in order to make the measurements consistent with this tenet, distances and times must be determined with devices at rest with respect to the observer. Only then can he claim that $c$ is an upper limit to all speeds and that light travels at speed $c$ under all circumstances and at all times. These notions are particularly important to the question we have addressed in this paper, especially for cosmologies that begin their expansion from a singularity at time $t=0$. The reason for this is rather straightforward. In these cosmologies, all the worldlines of sources we see today started from the same location---very near the same co-moving point we ourselves are now occupying. Clearly, to suggest that the light they emitted has traveled a distance $c(t_0-t_e) \\rightarrow ct_0$ since the big bang is quite non-sensical. The correct statement is that the most distant sources we see today are precisely those moving at close to proper lightspeed, which reached a proper distance $R_{\\rm h}(t_0/2)$ before emitting the light that is just now reaching us at time $t_0$. It is remarkable---though obvious in retrospect---how elegantly and beautifully this simple result emerges from the properties of the FRW metric itself written in stationary form, when we take the limit $t_e\\rightarrow 0$ for the time at which the light from the most distant sources was emitted. One of the principal results of our analysis has been the demonstration that even though $t_e\\approx 0$ for these sources, the time measured on our clocks at rest with respect to us was actually $T_e=(1/2)t_0$. And now we understand that this effect is entirely due to the time dilation between us and sources receding at proper speed $c$ when they emitted this light. These conclusions do come with a caveat, however, because most of these results are based on the use of FRW metrics with a constant spacetime curvature, allowing us to find an alternative set of coordinates to write them in stationary form. Without this option, we would not yet know how to evaluate distances and times in such a way as to demonstrate without any doubt how far sources could have traveled before emitting the light we see today. One ought to expect the proper size of the visible universe to be measurable against the gravitational horizon even in cases where the spacetime curvature is not constant, but only future work can establish this result conclusively." }, "1207/1207.1935_arXiv.txt": { "abstract": "We have investigated Saturn's core formation at a radial pressure maximum in a protoplanetary disk, which is created by gap opening by Jupiter. A core formed via planetesimal accretion induces the fragmentation of surrounding planetesimals, which generally inhibits further growth of the core by removal of the resulting fragments due to radial drift caused by gas drag. However, the emergence of the pressure maximum halts the drift of the fragments, while their orbital eccentricities and inclinations are efficiently damped by gas drag. As a result, the core of Saturn rapidly grows via accretion of the fragments near the pressure maximum. We have found that in the minimum-mass solar nebula, kilometer sized planetesimals can produce a core exceeding 10 Earth masses within two million years. Since Jupiter may not have undergone significant type II inward migration, it is likely that Jupiter's formation was completed when the local disk mass has already decayed to a value comparable to or less than Jovian mass. The expected rapid growth of Saturn's core on a timescale comparable to or shorter than observationally inferred disk lifetime enables Saturn to acquire the current amount of envelope gas before the disk gas is completely depleted. The high heat energy release rate onto the core surface due to the rapid accretion of the fragments delays onset of runaway gas accretion until the core mass becomes somewhat larger than that of Jupiter, which is consistent with the estimate based on interior modeling. Therefore, the rapid formation of Saturn induced by gap opening of Jupiter can account for the formation of multiple gas giants (Jupiter and Saturn) without significant inward migration and larger core mass of Saturn than that of Jupiter. ", "introduction": "\\label{sc:intro} Since Jupiter resides at 5.2\\,AU, it is likely that Jupiter did not undergo significant type II migration \\citep[e.g.,][]{ida_lin08}. The mass of the disk surrounding Jupiter decayed to a level comparable to or less than Jupiter mass before Jupiter completed its formation and opened up a gap. Since the disk could not push Jupiter, its type II migration is negligible. The formation of Saturn becomes problematic along this line. With an assumption that cores grow through the collisional accretion of surrounding planetesimals, the accretion timescale for Saturn's core may be 5--10 times longer than that for Jupiter's core. Therefore, since Jupiter forms on a timescale comparable to a disk depletion timescale, a few Myrs, Saturn's core formed on a timescale of 10--30 Myrs. However, when the core formed, the disk mass should have been so severely depleted that the core cannot accrete disk gas as massive as the present envelope mass. To reconcile the inconsistency, Saturn's core should have formed on a timescale shorter than several Myrs after Jupiter's formation. \\rev{ \\revs{ Note that a scenario by \\citet{walsh} where Jupiter first migrates inward, and then outward \\citep[e.g.,][]{morbidelli_crida}, also requires rapid formation of Saturn. } } \\revs{A single planetary embryo is formed from collisional coagulation of planetesimals in an annulus of the disk along the embryo's orbit and further grows through collisions with surrounding remnant planetesimals.} An embryo reaches the critical core mass, $\\sim 10 M_\\oplus$, to start gas accretion for Saturn formation \\citep{mizuno80,bodenheimer86,ikoma00}. However, since massive embryos enlarge the random motion of planetesimals, collisions between planetesimals are destructive. Embryos also grow by the accretion of fragments that result from such collisions. The eccentricities and inclinations of fragments are well damped by gas drag, and fragment accretion accelerates embryo growth. On the other hand, since such small fragments quickly drift inwards, the final mass of an embryo is much smaller than the critical core mass to start gas accretion \\citep{kobayashi+10,kobayashi+11,ormel_kobayashi}. However, Jupiter may have opened up a gap in a disk to truncate gas inflow to Jupiter at the present Jupiter's mass, which forms a radial pressure maximum near the edge of the gap. Radial drift of fragments is stalled at around the pressure maximum \\citep{adachi76}. In addition, type I migration of a core is also stalled there \\citep{tanaka,masset}. Since a core grows via fragment accretion without the loss of fragments and the core itself, rapid formation of Saturn's core may be realized under these conditions. In order to investigate the rapid formation of Saturn core near the edge of the gap, we perform simulations, which can derive accurate solutions in both limits dominated by collisional fragmentation \\citep{kobayashi10} or coagulation \\citep{kobayashi+10}. In Section \\ref{sc:pm}, we investigate the radial drift of bodies near the pressure maximum. We model a disk for a simulation and briefly explain the method of the simulation in Section \\ref{sc:model}. The core growth derived by the simulations is shown in Section \\ref{sc:result}. We discuss our findings in Section \\ref{sc:disc} and present our conclusion in Section \\ref{sc:conc}. ", "conclusions": "\\label{sc:conc} We have investigated the core formation of Saturn at the pressure maximum caused by Jupiter's gap opening in the solar nebula. Small particles feel strong gas drag and drift radially. The loss of bodies due to the drift stalls embryo growth before embryos reach the critical core mass \\citep{kobayashi+10,kobayashi+11}. Although bodies stirred by a large embryo can drift from the pressure maximum, the drift halts near the pressure maximum due to positive radial slope of gas pressure; hence an embryo at the pressure maximum effectively grows without loss of surrounding bodies (see Section \\ref{sc:pm}). Starting from monodisperse planetesimals, planetary embryos are generated via collisional evolution and stir remnant planetesimals, resulting in fragmentation of planetesimals. Fragments collide with each other and they accumulate at a radius $r_{\\rm f}$, given by Equation~(\\ref{eq:rf}). The random velocity of bodies with $r_{\\rm f}$ is well damped by gas drag; thereby embryos rapidly accrete such bodies. Fragments produced in the outer disk move to the pressure maximum, which contribute to further embryo growth. Due to the rapid accretion, a core as massive as 10 Earth masses forms in several million years, for kilometer-sized initial planetesimals in an MMSN disk, while the core formation needs a disk with 3 times larger solid surface density for 10\\,km planetesimals and with 10 times solid surface density for 100\\,km planetesimals. The growth is almost independent of gas density. Since the rapid formation in a disk with moderate mass is consistent with insignificant type II migration of Jupiter and larger core of Saturn than that of Jupiter, Saturn's core may have formed in the pressure maximum after Jupiter opened up a gap in the solar nebula. \\vspace{1cm} We thank S. Inutsuka for his comments and encouragement and H. Tanaka and S. Okuzumi for useful discussion. We also acknowledge D. Lin and E. Kokubo for inspiring us to consider the calculations of induced formation model of Saturn. \\rev{ For C.W.O. support for this work was provided by NASA through Hubble Fellowship grant \\#HST-HF-51294.01-A awarded by the Space Telescope Science Institute, which is operated by the Association of Universities for Research in Astronomy, Inc., for NASA, under contract NAS 5-26555. }" }, "1207/1207.4818_arXiv.txt": { "abstract": "Using a simple two-dimensional, zero-$\\beta$ model, we explore the manner by which reconnection at a current sheet releases and dissipates free magnetic energy. We find that only a small fraction (3\\%--11\\% depending on current sheet size) of the energy is stored close enough to the current sheet to be dissipated abruptly by the reconnection process. The remaining energy, stored in the larger-scale field, is converted to kinetic energy in a fast magnetosonic disturbance propagating away from the reconnection site, carrying the initial current and generating reconnection-associated flows (inflow and outflow). Some of this reflects from the lower boundary (the photosphere) and refracts back to the X-point reconnection site. Most of this inward wave energy is reflected back again, and continues to bounce between X-point and photosphere until it is gradually dissipated, over many transits. This phase of the energy dissipation process is thus global and lasts far longer than the initial purely local phase. In the process a significant fraction of the energy (25\\%--60\\%) remains as undissipated fast magnetosonic waves propagating away from the reconnection site, primarily upward. This flare-generated wave is initiated by unbalanced Lorentz forces in the reconnection-disrupted current sheet, rather than by dissipation-generated pressure, as some previous models have assumed. Depending on the orientation of the initial current sheet the wave front is either a rarefaction, with backward directed flow, or a compression, with forward directed flow. ", "introduction": "Magnetic reconnection has been frequently proposed as the mechanism whereby magnetic energy, stored in the solar corona, is rapidly released in a solar flare. In most current models, fast magnetic reconnection occurs at a current sheet where magnetic field lines of differing connectivity are brought into close enough proximity for a small-scale process, such as Ohmic diffusion or various kinetic effects, to forge new connections between them. The current sheet, by carrying a net current, also stores the magnetic energy which the reconnection liberates. This energy is not, however, co-located with the current sheet itself so the mechanism responsible for the reconnection electric field will not be the same one responsible for releasing or dissipating the magnetic energy. Many investigations have been focussed on the former, local mechanism (flux transfer) while far fewer have addressed the latter, global mechanism (energy release). A concrete illustration of the above issue is provided by the simple, two-dimensional quadupolar coronal field shown in \\fig\\ \\ref{fig:qpl}. A pair of sources, $P2$ and $N2$, have emerged underneath an older bipole, $P1$--$N1$. If not all of the new flux is able to reconnect with the overlying flux, the state of minimum magnetic energy, $\\bvec(\\xvec)$, will contain a current sheet separating new from old flux as shown in \\fig\\ \\ref{fig:qpl}a \\citep[see][]{Heyvaerts1977,Priest2000}. The state of lowest possible energy is a potential field, $\\bvec_0(\\xvec)$, shown in \\fig\\ \\ref{fig:qpl}b, in which additional flux $\\Delta\\psi$ (grey regions) interconnects the new and old polarities. A current sheet, carrying current $I_{\\rm cs}$, exists to maintain a connectivity different from the potential field, as in \\fig\\ \\ref{fig:qpl}a. For this reason the net current $I_{\\rm cs}\\sim\\Delta\\psi$ \\citep{Longcope2001b,Longcope2004}. The current sheet will be reduced, or eliminated entirely, as flux is transferred across the sheet into domain connecting $P1$--$N2$ and $P2$--$N1$ (the shaded regions). It is in this way the local topological changes at the current sheet may have energetic consequences. \\begin{figure}[htb] \\epsscale{0.8} \\centerline{(a)\\plotone{f1a.ps}} \\centerline{(b)\\plotone{f1b.ps}} \\centerline{(c)\\plotone{f1c.ps}} \\caption{The field from four photospheric, magnetic sources. (a) Equilibrium $\\bvec(\\xvec)$ state after emergence and incomplete reconnection of the bipole $P2$--$N2$. The current sheet is shown as a bold arc connected to thinner lines: the separatrices. (b) The potential field, $\\bvec_0(\\xvec)$, accessed by reconnection transfering flux $\\Delta\\psi$ into connection between $P2$--$N1$ and$P1$--$N2$; the new field lines fill the shaded regions. (c) The non-potential field, $\\bvec-\\bvec_0$, due to the current sheet.} \\label{fig:qpl} \\end{figure} The free magnetic energy of the topologically constrained magnetic field is computed by subtracting the energy of the unconstrained (i.e.\\ potential) field, \\begin{eqnarray} \\Delta E_M &=& {1\\over8\\pi}\\int|\\bvec|^2\\, d^3x ~-~{1\\over8\\pi}\\int|\\bvec_0|^2\\, d^3x \\nonumber \\\\ &=& {1\\over8\\pi}\\int|\\bvec-\\bvec_0|^2\\, d^3x ~~, \\end{eqnarray} where the integrals are over the entire coronal volume ($z>0$).\\footnote{A cross term involving the integral of $(\\bvec-\\bvec_0)\\cdot\\bvec_0$ can be seen to vanish after expressing $\\bvec_0=-\\nabla\\chi$ and integrating by parts.} The final expression shows that the free energy is equivalent to the energy of the {\\em non-potential} component, $\\bvec(\\xvec)-\\bvec_0(\\xvec)$, due to the current sheet with homogeneous conditions at the lower boundary ($B_z-B_{0z}=0$). This energy will naturally decrease as the current $I_{\\rm cs}$ is decreased by reconnection. The non-potential field, shown in \\fig\\ \\ref{fig:qpl}c, extends far beyond the sheet itself. At large distances it decreases inversely with distance similar to the field from a simple wire. It is evident from the form of the non-potential field (\\fig\\ \\ref{fig:qpl}c) that changes in the current sheet must propagate great distances in order to release the stored magnetic energy. Much of the volume over which the free magnetic energy is stored is not in magnetic contact with the current sheet, so neither Alfv\\'en waves nor slow magnetosonic waves (nor shocks) can change the field directly. Yet if the current in the sheet is to change, that change must be reflected in the distant field, in accordance with Amp\\`ere's law. How this occurs is unlikely to depend too critically on the mechanism responsible for the flux transfer which must, according to all current theories, occur on small scales. Thus it should be possible to study the mechanism of energy release using a arbitrary form of rapid flux transfer at the current sheet. Such an analysis was performed by \\citet{Longcope2007e} using a current sheet situated at a two-dimensional X-point, in a zero-$\\beta$ plasma, whose reconnection was effected by a sudden increase in Ohmic resistivity. They found that the disruption of the current sheet launched a cylindrical current shell at the leading edge of a fast magnetosonic wave (FMW). The shell contained almost all of the current formerly carried by the current sheet. It therefore left in its wake (i.e.\\ inside the shell) a nearly potential magnetic field. The non-potential magnetic energy of the initial field was converted to kinetic energy of the wave's flow. Remarkably, the Ohmic dissipation responsible for initiating the wave, and thus releasing the stored magnetic energy, directly dissipated a rather small fraction of the stored energy. This puzzling result, that large resistivity does not result in large Ohmic losses, can be anticipated from the arguments above: even after it diffuses, the current sheet occupies a very small volume and therefore has access to very little magnetic energy. The flow field in the FMW has a quadrupolar structure familiar in steady reconnection models: inflows along one axis (the vertical axis in \\fig\\ \\ref{fig:qpl}) and outflows along the other (horizontal). Previous studies of transient magnetic reconnection models have found fast magnetosonic rarefaction waves to be the drivers of inflow \\citep{Lin1994,Heyn1996,Nitta2001}. These analyses, set on infinitely long current sheets, focussed on the inner reconnection region rather than the front of the wave. Since \\citet{Longcope2007e} considered a sheet carrying finite current they were able to analyze its full global propagation and energy release. In particular, their model revealed that the kinetic energy of the inflow must be supplied by decreasing the free magnetic energy stored in the initial equilibrium. In their unbounded domain this kinetic energy is not subsequently converted to heat. The finite current sheet studied by \\citet{Longcope2007e} was situated in an unbounded domain, which therefore contained an infinite amount of free magnetic energy. The FMW would propagate indefinitely, converting this stored energy to kinetic energy at a uniform rate. The wave energy could therefore become arbitrarily large in comparison to the energy directly dissipated at the current sheet. In order to make contact with previous work by \\citet{Craig1991} and \\citet{Hassam1992}, \\citet{Longcope2007e} briefly considered the effect of a boundary: a concentric cylindrical conductor. This reflected (perfectly) the outward-propagating FMW back inward. The inward wave collapsed on the X-point, as described by \\citet{Craig1991} and \\citet{Hassam1992} permitting still more Ohmic dissipation to occur there. The majority of the wave's energy was, however, {\\rm reflected} once more by the X-point. The wave was therefore trapped between a perfectly reflecting outer conductor and an imperfectly reflecting X-point. The only losses were at the X-point, so eventually all energy was in fact dissipated at the X-point through Ohmic dissipation. The characteristic time for dissipation depended on the round-trip transit time between reflectors. Owing to the exponential increase of the Afv\\'en speed with distance, this transit time scales logarithmically with the dissipation scale, and thus logarithmically with the resistivity. The ability of an X-point to Ohmically dissipate magnetic energy in logarithmic time was the most significant result of the studies \\citet{Craig1991} and \\citet{Hassam1992}, and appears to suggest that Ohmic dissipation alone is capable of all magnetic energy dissipation in a flare. The upshot of this reasoning is that the initial disruption of the current sheet by Ohmic diffusion will dissipate a small fraction of its magnetic energy, but that repeated reflections can result in the complete dissipation of all stored energy at the X-point. To achieve the latter end, however, previous authors have assumed a conducting boundary completely surrounding the X-point. In a more realistic geometry, such as that of \\fig\\ \\ref{fig:qpl}, there is a conducting boundary at the photosphere ($z=0$) but it does not completely surround the X-point. \\citet{McLaughlin2006} studied the linearized, $\\beta=0$ dynamics of a quadrupolar magnetic field with a planar lower boundary like \\fig\\ \\ref{fig:qpl}. Rather than initiate a wave by reconnection, they launched a FMW by {\\em fiat} from the lower boundary. They found that a fraction of the wave (they report 40\\%) was refracted into the null, while the remainder continued to propagate away. This suggests that the photospheric lower boundary may indeed be less effective at mediating energy dissipation than are the concentric cylindrical conductors, although it is not clear that the same fraction would apply to a wave launched outward from the null rather than from the lower boundary. \\citet{McLaughlin2006} omitted diffusion from their model but assumed its effect would be to dissipate all the wave energy refracted into the null. \\citet{Craig1991} and \\citet{Hassam1992} found to the contrary that the main effect of resistivity is to {\\em reflect } the wave back outward from the null. It was only through repeated reflection that the wave energy could be ultimately dissipated. If only 40\\% were reflected back at each step (as suggested by \\citet{McLaughlin2006}), then the final dissipation would be far less than 40\\% initially directed toward the null. The lack of 100\\% reflection back to the null point appears to pose a difficulty for the system reaching a potential state. It was found by \\citet{Longcope2007e} that flux transfer continues at the X-point long after the majority of the current sheet had disappeared. As a result $\\Delta\\psi$ continues to decrease, {\\em below zero}, bringing the system {\\em away} from the potential field state ($\\Delta\\psi=0$). When the wave reflected from the cylindrical boundary, however, it reversed this flux transfer, bringing $\\Delta\\psi$ back upward. While it continued to overshoot the potential value, numerous reflections led to its gradual convergence to zero. If the initial wave is only partially reflected back to the null point, as the results of \\citet{McLaughlin2006} suggest, then it is unclear how $\\Delta\\psi$ would ever converge to zero, and the system achieve a potential field. To more fully understand energy release and dissipation we seek to determine what fraction of the FMW launched from the null point will be reflected back to the X-point in this more realistic configuration. We must also know how much of the partially reflected wave will ultimately dissipate at the X-point and what fraction will reflect once more. This analysis, performed below, shows that the lowest frequency components of the wave are reflected almost completely by the photospheric boundary. This results in the current at the X-point being entirely eliminated, and $\\Delta\\psi\\to0$, after numerous reflections. It also results in the direct dissipation of a significant fraction of the free energy (between 40\\% and 75\\% in one example we consider). This does not occur immediately through the local reconnection mechanism, Ohmic diffusion in our case, but rather it requires repeated reflections from the photosphere and thus takes many transit times to achieve. The remainder of the free energy (25\\% -- 60\\%) is emitted by the reconnection as FMW at higher frequencies. These waves form a number of pulses propagating primarily vertically upward, away from the photosphere. The flow in these wave forms is either a rarefaction or a compression depending on the orientation of the initial current sheet (horizontal or vertical, respectively). Such fast magnetosonic disturbances have long been known to accompany flares, but until now no quantitative reconnection model has predicted what fraction of the free magnetic energy they accounted for. We present the model calculation for a two-dimensional, quadrupolar field, like the one in \\fig\\ \\ref{fig:qpl}. In the next section we specify the geometry of the model field, and quantify the free magnetic energy stored in advance of reconnection. In the following section we analyze the dynamical behavior in the vicinity of the X-point, and describe how this is matched to the external field, including the photospheric boundary. Section \\ref{sec:num} presents numerical solutions to the external response, including the emission of a FMW upward. The numerical solutions cannot be continued for the many FMW-transit times of the full solution. Instead we use them to characterize the reflection from the photosphere and then use the reflection coefficient to synthesize a solution for long times in \\S \\ref{sec:long_time}. This long time solution is used to quantify the ultimate fate of the free energy. ", "conclusions": "We have used a simplified model to study the response of the large-scale coronal magnetic field to reconnection at a current sheet. The reconnection reduces (almost to zero) the current in the sheet by transferring magnetic flux across it. As reported by \\citet{Longcope2007e}, the current removed from the sheet is carried away at the front of a fast magnetosonic wave. A portion of this wave reflects from the photospheric boundary, and is refracted back toward the X-point. Subsequent reflections of the wave between the X-point and the photosphere lead eventually to the elimination of all current, and with it all free magnetic energy. The number of reflections required to reach the potential field is roughly the logarithm of the global Lundquist number, corroborating the findings of \\citet{Craig1991} and \\citet{Hassam1992}. The total reconnection time is therefore $\\sim\\ln^2 S$ times a single Alfv\\'en transit (i.e.\\ $1/\\oma$). This final approach resembles the case of perfectly reflecting concentric cylindrical boundary because the photospheric boundary has a much higher reflection coefficient than the resistive X-point at very low frequencies. During early reflections, however, the FMW contains frequencies above $\\oma/\\ln S$, of which a sizable fraction are directed vertically upward from the photosphere. We thus find, at least for this simplified model, that an appreciable fraction of the initial free energy is carried from the current by fast magnetosonic waves; in the case we explored in detail 25\\% -- 60\\% of the energy was converted to waves, depending on the size of the initial current sheet. We expect some of the basic elements of this result to hold in models more sophisticated that the one we used for our detailed study. A common element in all models of fast magnetic reconnection is that the reconnection electric field, i.e.\\ the flux transfer, occurs on small scales, while free magnetic energy is stored over vastly larger scales. Information about the flux transfer, and associated current reduction, must be transmitted to the larger corona. This includes significant volumes which are not magnetically linked to the reconnection site, for which transmission must occur through fast magnetosonic modes. Previous studies of unsteady fast reconnection have shown fast magnetosonic waves propagating into the unreconnected flux to be responsible for creating the reconnection inflow, the same inflow assumed as a boundary condition or driver in steady models \\citep{Lin1994,Heyn1996,Nitta2001}. Our model contains similar fast mode rarefaction waves, and reveals the fraction of energy they contain. One of the most dubious simplifications is our use of uniform classical resistivity, $\\eta$, as the means of generating a reconnection electric field. The simple mathematical form of this effect permits a more thorough analysis than would more complex physical effects. Since our goal was to study the response of the large-scale field, where the electric field is irrelevant, we chose the simplest possible form for the small-scale effect. We expect that a more accurate treatment of the small scales would reveal complex flows, whose overall effect might be characterized as an effective diffusivity. It is this kind of turbulent $\\eta$ we use in applying our results to the reconnection of global currents, as in \\fig\\ \\ref{fig:Xpoint_circ}. Were the turbulent diffusivity computed with any degree of self-consistency, it is unlikely to remain constant through the many repeated reflections of the FMW. Remarkably, we find that its actual value has no effect on the magnitude of the X-point reflection (see \\eq\\ [\\ref{eq:Ri}]), critical to the ultimate energy dissipation. Thus we expect the basic scenario revealed in our simple model would be found in more sophisticated ones. The equations governing the large-scale response, namely \\eqs\\ (\\ref{eq:momentum}) and (\\ref{eq:induction}), are linear, purely two-dimensional (with no magnetic field in the ignorable direction; no ``guide field'') and assume zero pressure. The assumption of linearity is the most easily justified, since the initial current sheet contributes a small correction to the field far from itself. \\citet{Longcope2007e} analyze in detail the requirements for linearity in the internal solution, and find that it is a reasonable approximation for many times $1/\\oma$, by which time the reflected wave comes into play. Significant pressure can alter the nature of the equilibrium from which the energy must be released. A recent numerical solution found that only a tiny fraction ($\\sim3\\times10^{-9}$) of the magnetic energy in a pressure-dominated equilibrium was released by suddenly enhancing the resistivity \\citep{Fuentes-Fernandez2012}. They found that resistive diffusion does redistribute the current, as in the $\\beta=0$ case, but that the new distribution of Lorentz forces are compensated by a new distribution of pressure. This redistribution process is localized to the current sheet so that instead of an axisymmetric FMW reducing the free energy, a small wave with $m=4$ is launched, with negligible energetic consequence. A still more recent experiment at lower values of $\\beta$ continued to observe reconnection leading to pressure re-distribution rather than a decrease in the sheet's current \\citep{Fuentes-Fernandez2012b} At the other end of the process, \\citet{McLaughlin2006b}, found that pressure in the vicinity of the null kept the fast magnetosonic speed above zero, thus mitigating the focusing of wave energy there. Pressure thus reduces the amount of energy initially released and the fraction of that dissipated. By focusing on the zero-pressure limit we are finding the maximum possible energetic effect in solar flares, which pressure would almost certainly reduce. We believe our assumption of two-dimensionality places the most severe limitations on the applicability of our model. In the absence of a guide field ($B_x$) all shear Alfv\\'en modes are polarized in the ignorable direction ($\\xhat$) and are completely uncoupled from the reconnection dynamics. It is for this reason that the response propagates isotropically from the X-point as a FMW. Adding any amount of an ignorable component to the system will result in a global response including both FMWs, propagating in all directions, and shear Alfv\\'en waves confined to the separatrix field lines. While the FMW will reflect from the photosphere specularly, as they do in our model, the reflected Alfv\\'en waves will travel back along the same field lines to reach the X-point again. This far more complex, and more interesting, system will need to be studied in the future. The present study can be taken as a limiting case, and can provide an outline, for such a future investigation. We expect the FMW component of that system to behave in a manner similar to our model, but to account for a smaller portion of the energy, since Alfv\\'en waves will contain some as well. Many past observations have provided evidence for both FMW and shear Alfv\\'en modes initiated by solar flares. Evidence of fast mode shocks are found in coronal type II radio bursts \\citep{Payne-Scott1947} and chromospheric Moreton waves \\citep{Moreton1960,Moreton1960b}, while shear Alfv\\'en waves are observed in post-flare loop oscillations \\citep{Aschwanden1999,Nakariakov1999}. In both cases their timing strongly suggests these disturbances are initiated by flares or by CMEs, although the details of the initiation remain unclear. The wave front in our model includes a ``skirt'' along the lower boundary, evident in \\fig\\ \\ref{fig:WKB}b, which might manifest as a Moreton wave \\citep{Uchida1968}. The upper portion of the wave, if correctly oriented, could shock to produce metric type II radio signatures. A common model of these observed waves has been that direct energy dissipation creates a pressure pulse driving a FMW outward in all directions \\citep[see][and references therein]{Vrsnak2008}. Our model provides a detailed picture of the initiation mechanism which differs in several important respects form pressure-pulse models. Our primary conclusion is that, contrary to the common assumption, the broad initial distribution of free magnetic energy makes it impossible to create a local pressure pulse through its rapid dissipation. We find instead that rapid diffusion results in a current redistribution whose newly unbalanced Lorentz forces produce the FMW. We admittedly discard the pressure which would make possible the alternative wave-generation mechanism. Were it included, however, the thermal energy could not exceed the magnetic energy directly dissipated. We find this to be 3\\%--11\\% of the total, and therefore smaller than the wave energy produced through Lorentz forces. The plasma flow direction is one potentially observable difference between our model and the pressure-pulse waves proposed previously. The pressure pulse or piston driver leads to outward (compressive) flows in all directions, while our wave has both outward and inward (rarefaction) flows along different portions of the cylindrical front. The outward portion is an extension of the inner reconnection outflow jet and thus related to the fast magnetosonic termination shock predicted in various models \\citep{Forbes1986c}. The inward portion is the reconnection inflow, but its speed is a function of the external field rather that the inner reconnection rate. The relative locations of these two components depend on the orientation of the initial current sheet, and thus on the sign of $I_{\\rm cs}$. Moreover, since all its energy is introduced at the outset, the pressure-driven wave has an amplitude that decreases with distance from the source. In contrast, the FMW continues to draw energy from the large scale magnetic field and thus decays much more slowly; in the vicinity of the X-point its amplitude remains constant even as its net energy increases exponentially \\citep{Longcope2007e}. It remains to compare the amplitude profiles of observed FMWs to the predictions of these different models. Finally, our model shows how the dissipation of magnetic energy at the smallest scales, i.e.\\ the current sheet itself, requires repeated reflection from distant boundaries (i.e.\\ the photosphere), and is therefore not entirely local. The initial phase of flux transfer dissipates only a very small fraction of the energy, as found by \\citet{Longcope2007e}, and expected from arguments based on the finite energy density and very small volume. Reflected waves will focus back onto the X-point due to its vanishing Alfv\\'en speed. As first predicted by \\citet{Craig1991} and \\citet{Hassam1992}, this focusing leads in the end to the dissipation of significant energy within the small region after many reflections. The energy dissipation thus persists far longer than the time taken for the initial flux transfer that we call reconnection. It is tempting to see in this persistent dissipation a possible explanation for flare durations longer than free cooling following a single impulsive energy release \\citep{Warren2006}. Before doing so we must confront the effect, alluded to above, of three-dimensional geometry. With a guide field component the fast magnetosonic speed no longer vanishes at the X-point, and we would expect far less of the wave energy to focus back there. Alfv\\'en waves, on the other hand, will be confined to the field lines and will thus reflect repeatedly back to the reconnection site. This kind of wave reflection and dissipation, sometimes invoked to explain the persistent energy release in flares, is indirectly related to the fast mode version we have studied. Genuine fast mode focusing would, however, occur at a coronal null point in a three-dimensional magnetic field, as it does in our two-dimensional version. \\bigskip This work was supported by a grant from the NSF/DOE Plasma Sciences partnership. Aaron Schye aided the effort with a preliminary WKB computation. DWL thanks M.D.\\ Ding and the Nanjing University, School of Astronomy and Space Science for hosting a visit during which some of this work was done. \\appendix" }, "1207/1207.5266.txt": { "abstract": "We have investigated the post-merger signatures of red-sequence galaxies in rich Abell clusters at $z \\lesssim$ 0.1: A119, A2670, A3330 and A389. Deep images in $u'$, $g'$, $r'$ and medium-resolution galaxy spectra were taken using MOSAIC 2 CCD and Hydra MOS mounted on a Blanco 4-m telescope at CTIO. Post-merger features are identified by visual inspection based on asymmetric disturbed features, faint structures, discontinuous halo structures, rings and dust lanes. We found that $\\sim$ 25\\% of bright ($M_r < -20$) cluster red-sequence galaxies show post-merger signatures in four clusters consistently. Most ($\\sim$ 71\\%) of the featured galaxies were found to be bulge-dominated, and for the subsample of bulge-dominated red-sequence galaxies, the post-merger fraction rises to $\\sim$ 38\\%. We also found that roughly 4\\% of bulge-dominated red-sequence galaxies interact (on-going merger). A total of 42\\% (38\\% post-merger, 4\\% on-going merger) of galaxies show merger-related features. Compared to a field galaxy study with a similar limiting magnitude \\citep{van05}, our cluster study presents a similar post-merger fraction but a markedly lower on-going merger fraction. The merger fraction derived is surprisingly high for the high density of our clusters, where the fast internal motions of galaxies are thought to play a negative role in galaxy mergers. The fraction of post-merger and on-going merger galaxies can be explained as follows. Most of the post-merger galaxies may have carried over their merger features from their previous halo environment, whereas interacting galaxies interact in the current cluster in situ. According to our semi-analytic calculation, massive cluster haloes may very well have experienced tens of halo mergers over the last 4--5 Gyr; post-merger features last that long, allowing these features to be detected in our clusters today. The apparent lack of dependence of the merger fraction on the clustocentric distance is naturally explained this way. In this scenario, the galaxy morphology and properties can be properly interpreted only when the halo evolution characteristics are understood first. ", "introduction": "The formation of massive early-type galaxies in the universe is still in question. The hierarchical galaxy formation scenario is widely accepted at present, which predicts that massive galaxies form through hierarchical galaxy mergers. If each galaxy merger induces star formation and consequent stellar mass growth in the united galaxy, it would lead to a large scatter in the age and metallicity of stellar populations in massive early-type galaxies. However, the observational characteristics of early-type galaxies, such as their red colors tightly correlated in optical color-magnitude relations (CMRs) and their high $\\alpha-$elements ratios, imply that most of the stellar contents in massive early-type galaxies formed in a short timescale at an early epoch ($z >$ 1). Furthermore, the red-sequence in the optical CMR appears to be established by $z \\sim$ 1 \\citep{tan05}, and recent observations with 8--10m-class telescopes reveal the appearance of massive red-sequence galaxies ($M \\gtrsim 10^{11} $M$_{\\sun}$) at $z \\sim$ 2 -- 3 \\citep{kod07,dro08,kan09}. These observational results suggest that most massive early-type galaxies had almost completed their star formation and mass aggregation by $z \\sim$ 1 and then virtually passively evolved. There are other observational clues which indicate a significant increase in the stellar mass density in massive red-sequence galaxies since $z \\sim$ 1, which is not allowed according to the passive evolution of blue galaxies alone. \\citet{bel04} showed that the $B$-band luminosity density of massive red-sequence galaxies does not evolve much in the redshift range of $0 < z \\leq 1.1$. Because the $B$-band light should dim as stars get old, they argued that some young populations should be provided regularly from blue galaxies to the red-sequence during this period. However, the simple fading of massive blue galaxies cannot be applied, as the brightest red galaxy is always brighter than the brightest blue galaxy throughout the redshift range. After all, they suggested that galaxy mergers are an important process in the formation of luminous red galaxies since $z \\sim$ 1. The result was supported again by \\citet{fab07}. Residual star formation (RSF) in early-type galaxies provides another clue about galaxy formation. RSF has been extensively studied, recently using the UV data from space telescopes such as the Galaxy Evolution Explore ($GALEX$) and the Hubble Space Telescope ($HST$). The latest work on the UV upturn phenomenon of early-type galaxies showed that RSF fraction among cluster elliptical galaxies is as high as $\\sim$ 30\\%; it is even higher in the field environment \\citep{yi11}. Several studies have claimed that the RSF in early-type galaxies is most likely related to galaxy mergers or interactions \\citep{kav07,kav10b}. Because the UV bright phase of a young stellar population lasts only $\\sim$ 1 Gyr, we can estimate that the RSF detected at less than $z =$ 0.1 is stimulated at a relatively low redshift, $z \\sim$ 0.2 -- 0.3. Again, this is indirect evidence of the substantial frequency of galaxy mergers given a low redshift. \\citet{van05} showed tidally disturbed features around field elliptical galaxies at $z \\sim$ 0.1 with optically deep images from NDWFS (NOAO Deep Wide Field Survey) and MUSYC (MUlti-wavelength Survey of Yale and Chile). That paper suggested that $\\sim$ 70\\% of field elliptical galaxies were assembled through dry mergers (i.e., mergers of gas-poor, bulge-dominated systems) in the recent epoch. \\citet{kav10a} also found disturbed features as well as dust lanes from 25\\% of luminous early-type galaxies ($M_r < -20.5$, $z <$ 0.05) in the `Strip 82' fields of Sloan Digital Sky Survey (SDSS) data, which have much longer integration times than general SDSS fields. These are very meaningful results because mass growth via continuous galaxy mergers were up to that point mostly supported observationally by statistical analyses of photometric colors instead of by visual evidence of galaxy mergers. Moreover, the merger fraction is often estimated from galaxy pair fractions, considering them as pre-mergers \\citep[e.g.,][]{bun09,der09}. Although \\citet{van05} and \\citet{kav10a} did not deal with complete volume-limited samples, they searched a substantial number of galaxies with unprecedentedly deep optical images and presented direct evidence of galaxy mergers related to the formation of red, early-type galaxies. Our interest moved to galaxy clusters, which are dominated by massive early-type galaxies. It has been expected that galaxy clusters are not a likely environment for frequent merger events to take place due to the high peculiar motions of the galaxies within the cluster. Moreover, it is commonly held that massive galaxies in clusters formed faster at an earlier epoch than did galaxies in a field environment (Gunn \\& Gott 1972; Dressler 1980; Tanaka et al. 2005 and references in it). Therefore, it was presumed that it would be difficult to find post-merger signatures from massive early-type galaxies in galaxy clusters. Naturally, there has been less effort to find post-merger signatures among cluster galaxies. To find post-merger signatures of early-type galaxies in galaxy clusters, we carried out optical deep imaging and multi-object spectroscopic observations of \\st{four} rich Abell clusters at 0.04 $< z <$ 0.11 with the Blanco 4-m telescope at the Cerro Tololo Inter-American Observatory (CTIO). We picked four rich Abell clusters (A119, A2670, A3330 and A389) as our targets. They have been observed by GALEX at least in the medium-depth imaging mode (MIS: Medium Imaging Survey) and are optimally visible from the CTIO. Their GALEX NUV exposure times range from $\\sim$ 1 hours for A119 to $\\sim$ 17 hours for A3330. UV light is very sensitive to the existence of young stellar populations and thus plays an important role in the study of the recent star formation history that may have been caused by mergers. The UV properties will be published in an upcoming paper following this one, in which we initially present the optical colors and morphologies of cluster galaxies. We performed a visual inspection of all the red-sequence galaxies brighter than $M_{r'} = -20$ in each galaxy cluster using deep optical images. This result will enable us to compare the recent galaxy merger histories between cluster and field environments. Morphological indexes such as the bulge-to-total ratio and asymmetry of the galaxies were also measured to examine their morphological characteristics. In this paper, the observations and data reductions are presented in Section~\\ref{sec:obs} and Section~\\ref{sec:data}. Section~\\ref{sec:sample} describes the galaxy sample selection process and Section~\\ref{sec:vi} shows the scheme and the result of the visual inspection using the deep optical images. Morphological examination will be presented in Section~\\ref{sec:mi}. We will conclude this paper with a summary and discussion in Section~\\ref{sec:disc}. \\begin{figure} \\center{ \\includegraphics[scale=0.4]{fig1.jpg} \\caption{The field-of-views of the MOSAIC2 and Hydra data are superimposed over the DSS image for each cluster. The square in each panel indicates the field of the stacked deep image while the gray-dashed circles represent the observed Hydra FOVs. Their consistent FOVs enabled us to perform an almost complete spectroscopic survey of the cluster galaxies.\\label{field}} } \\end{figure} ", "conclusions": "\\label{sec:disc} We investigated the post-merger signatures of red-sequence galaxies in rich Abell clusters at $z \\lesssim$ 0.1. Deep images in $u'$, $g'$, $r'$ and medium-resolution galaxy spectra were taken for A119, A2670, A3330 and A389 with a MOSAIC 2 CCD and a Hydra MOS mounted on the Blanco 4-m telescope at CTIO. Post-merger signatures were identified by visual inspection of their disturbed features, e.g., asymmetric structures, faint features, discontinuous halo structures, rings and dust lanes. Most ($\\sim$ 71\\%) of the featured galaxies were found to be bulge-dominated systems (E/S0 galaxies) from the radial surface brightness profile fitting. On average, the asymmetry of the post-merger galaxies is always slightly larger than the mean of normal galaxies within the same bulge-to-total ratio range. We noted that cluster red-sequence galaxies went through galaxy mergers until the recent epoch. In this work, $\\sim$ 25\\% of the bright ($M_r < -20$) cluster red-sequence (RS$_{sp}$) galaxies exhibited post-merger signatures in four target clusters consistently. The fraction increases to $\\sim$ 38\\% with the bulge-dominated RS$_{sp}$ galaxies (B/T $>$ 0.4). In addition, brighter galaxy samples show higher fractions of post-merger signatures in the clusters. \\begin{deluxetable}{cccc} \\tablecolumns{4} \\tabletypesize{\\footnotesize} \\tablewidth{0pt} \\tablecaption{Comparisons between cluster and field \\label{compfield}} \\tablehead{ \\colhead{} & \\colhead{Class} & \\colhead{Cluster} & \\colhead{Field\\tablenotemark{a}} } \\startdata & PM & 25 $\\pm$ 3\\% & 35\\% \\\\ Red\\tablenotemark{b} & I & 5 $\\pm$ 1\\% & 18\\% \\\\ & Total & 30 $\\pm$ 4\\% & 53\\% \\\\[3pt] \\hline \\\\[-3pt] & PM & 38 $\\pm$ 5\\% & 49\\% \\\\ Bulge-dominated\\tablenotemark{c}& I & 4 $\\pm$ 1\\% & 21\\% \\\\ & Total & 42 $\\pm$ 6\\% & 70\\% \\enddata \\tablenotetext{a}{The fractions for the field environment were adopted from \\citet{van05}.} \\tablenotetext{b}{Fractions for the cluster are derived with the RS$_{sp}$ galaxies in this paper, while the field red galaxies are denoted with $B - R$ colors.} \\tablenotetext{c}{Fractions with only bulge-dominated galaxies among the red galaxies. For the cluster, galaxies with B/T $>$ 0.4 are included while visually classified E/S0 galaxies are considered for the field.} \\end{deluxetable} In Table~\\ref{compfield}, we compared the average of fractions from the four Abell clusters with the result of field environment from \\citet{van05} . The clusters and field environments were compared using two subsets of galaxy samples, red galaxies and bulge-dominated red galaxies. Our RS$_{sp}$ galaxies were compared with the luminous red galaxies selected using $B-R$ colors in \\citet{van05}, and the bulge-dominated galaxies (B/T $>$ 0.4) were compared with the visually classified field E/S0 galaxies in the same paper. We presented the fractions of post-merger galaxies and interacting galaxies separately as well as the combined fractions which indicate the fraction of disturbed galaxies related to past or ongoing galaxy mergers. The van Dokkum sample does come with spectroscopic redshift; spectroscopic redshifts were available only on 5 galaxies out of 126 and so no k-correction was applied. Based on a small subsample he argues that his galaxies are roughly at redshift of 0.1, which is comparable to ours. The colors of his bulge-dominated galaxies are also comparable to ours after proper color conversions using \\citet{lup05}. The table demonstrates that the overall fractions in clusters are lower than those in the field. We find it difficult to understand the ``interacting\" galaxies in these clusters. They appear to be beginning their interaction/merger, and so we classified them as ``interacting\" or pre-merger. Considering that interacting systems with small companions are likely missed in our search, it is likely that our fraction for interacting galaxies is a hard lower limit. For example, some of our ``faint companion galaxies\" (which were excluded in our analysis) could be ``interacting\". Direct satellite-satellite mergers are supposed to be extremely rare in large halo environments, and so it is difficult to understand them. The fractions of post-merger galaxies are not very different between the cluster and the field. This is not compatible with a simple theoretical expectation based on merger timescale as a function of relative speeds of galaxies. Therefore, we suggest that many massive red-sequence galaxies formed in a less dense environment through galaxy mergers and entered into the cluster through dark matter halo mergers. An important caveat exists on this comparison between van Dokkum's work and ours. Merger timescales and thus probability heavily depends on mass ratio of colliding systems, but it is very difficult to extract the mass ratio information from observed images. We thought it would be the best attempt to use similar instruments for similar depth (exposure) on galaxies at similar distances. But even in this case, we are not free from a possible environmental (cluster vs field) dependence on mass ratios of galaxy mergers. \\begin{figure} \\center{ \\includegraphics[scale=0.55]{fig14.pdf} \\caption{Fractions of PM galaxies along the projected distances from the BCGs. The size of distance bin is 0.2$R_{200}$ in each galaxy cluster. Contrary to the general assumption, which would predict more frequent mergers at the outskirts of a cluster, the fraction of PM galaxies does not change much along the clustocentric distances. \\label{distance}} } \\end{figure} The radial distribution of the post-merger galaxies in the clusters provides another hint supporting this scenario. The fractions of post-merger galaxies are derived along the projected distance from the BCG in the scale of $R_{200}$ of each galaxy cluster, as shown in Figure~\\ref{distance}. We found that the fractions do not change much along the clustocentric distances. In situ mergers can scarcely explain this result because it is difficult to expect frequent mergers in the central region of clusters compared to the outskirts. This implies again that a large fraction of massive red-sequence galaxies with post-merger signatures likely merged in a less dense environment and then fell into the central region of the cluster. Is it possible for a galaxy to maintain the post-merger features produced in the outskirts until the galaxy arrives at the cluster center? It has been suggested for a long time that tidally ejected materials during galaxy mergers may return over many Gyr \\citep{her92, hib95}. \\citet{jen08} also proposed that the features of interacting galaxies can be observable for up to $\\sim$ 2 Gyr by applying stellar population models to the simulated equal-mass gas-rich disc mergers. We note that the duration of merger feature is highly sensitive to galaxy type, merger geometry, mass ratio, and last but not least the observing surface brightness limit\\footnote[1]{~The suggestion of 2 Gyr visibility of merger features by \\citet{jen08} was based on the surface brightness, $\\mu =$ 25 assumption. The visibility time is likely longer in our case because our imaging was deeper.}. The dynamical friction timescale of a satellite galaxy to the central cluster halo would be another parameter to compare with the observable timescale of post-merger features. In this work the dynamical friction timescale is defined as the timescale between the epochs of when the halo of satellite galaxy is started to merge into the cluster halo at the cluster virial radius and when the satellite galaxy have finally merged into the BCG in the model. If a dynamical friction timescale for a galaxy is much longer than the timescale of the post-merger features, our results should be interpreted in a different light. By taking advantage of a semi-analytic approach, we obtained the relationship between the merger mass ratio and the merger timescale of galaxies to the BCG. Our semi-analytic model (SAM) is based on N-body backbone dark matter merger trees and physical ingredients that govern the baryonic properties of galaxies. The merger timescales in the model are calculated using the formulae of \\citet{jia08}, which provides a modified form of Chandrasekhar's formula. Given that Chandrasekhar's formula generally underestimates merger timescales, \\citet{jia08} fits the formula into merger timescales derived from hydrodynamic simulations \\citep[see also][]{kim11}. \\begin{figure} \\center{ \\includegraphics[scale=0.45]{fig15.pdf} \\caption{Merging timescale of galaxies in a cluster of 5$\\times10^{14}$M$_{\\sun}$ of our SAM along the mass ratio. If we assume that a central galaxy is the BCG, bright satellite galaxies (M$_{sat} >$ 0.25M$_{cen}$) merge into the BCG within $\\sim$ 4 Gyr at $z <$ 0.5. \\label{tmerge}} } \\end{figure} We selected a cluster in the model which has a cluster mass, 5$\\times10^{14}$M$_{\\sun}$, comparable to our cluster samples. Figure~\\ref{tmerge} shows the merging timescales of all satellite galaxy mergers along the mass ratio between the BCG and the satellite, consequently forming a cluster of 5$\\times10^{14}$M$_{\\sun}$ at present. The redshift of each merger was expressed using different symbols, showing that the merger timescale was shorter at a higher redshift for the same mass ratio in the model. This arises because the size of the cluster was smaller at an early epoch. We can assume that the oldest observable post-merger features in our target clusters are produced at $z \\sim$ 0.5, where the difference in the lookback time from $z \\sim$ 0.1 is about 4 Gyr. The figure shows that the dynamical friction timescale of a satellite galaxy as massive as M$_{sat}$/M$_{cen} >$ 0.25 will be $\\lesssim$ 4 Gyr below the redshift, $z =$ 0.5. If we assume that the magnitude difference between the BCG and the second BCG is 1 magnitude, then the mass ratio would be $\\sim$ 0.4 (M$_{sat}$:M$_{cen} =$ 1:2.5). The figure indicates that then the second BCG will merge into the BCG within 3 Gyr at $z <$ 0.5. Therefore, it appears that massive post-merger galaxies can fall into the cluster center and maintain their disturbed features, even in low redshift ($z <$ 0.5). The consistent fractions of the featured galaxies along the clustocentric distance most likely arise from on-going halo mergers which started during various epochs. The average count of post-merger galaxies in our cluster samples was 17 $\\pm$ 6 galaxies (13, 25, 13 and 17 galaxies for A119, A2670, A3330 and A389, respectively). We also counted the number of halo mergers related to the current cluster halo in the SAM, to check whether a cluster can experience sufficient halo mergers to absorb massive galaxies since $z=$ 0.5. In the model of a 5$\\times10^{14} $M$_{\\sun}$ cluster, we found 21 satellite halos merging into the cluster at $z <$ 0.5, with galaxies brighter than $M_{r'} = -20$. It theoretically supports that a massive cluster may have enough halo mergers to take recently merged galaxies in at recent epoch. In this preliminary analysis, we assumed that merger features from the previous halo environment last roughly the same time as the subhalo's dynamical friction timescale. This is certainly an over-simplification. A more realistic calculation would require accurate merging halo mass ratios, merging galaxy mass ratios, and merger feature timescales throughout cluster evolution history. In conclusion, we found that $\\sim$ 42\\% of massive, bulge-dominated red-sequence galaxies in galaxy clusters have continued their mass assembly through galaxy mergers until the recent epoch, as $\\sim$ 70\\% of the field galaxies have done. Although the fractions of post-merger galaxies in clusters are lower than that in the field, it is still too high compared to the expectation from apparent fractions of galaxy interaction in galaxy clusters. Therefore, we propose that most of those post-merger galaxies were assembled in a low-density region and fell into the current cluster via halo mergers. We have supported this scenario with theoretical predictions using a semi-analytical model. We do not speculate the progenitors of the post-merger galaxies in this paper, regardless of whether they are remnants of dry mergers or wet mergers. This will be investigated in depth in an upcoming paper which will discuss on the ultraviolet properties of the galaxies most likely related to the residual star formation induced by galaxy mergers. \\ %% In a manner similar to \\objectname authors can provide links to dataset %% hosted at participating data centers via the \\dataset{} command. The %% second curly bracket argument is printed in the text while the first %% parentheses argument serves as the valid data set identifier. Large %% lists of data set are best provided in a table (see Table 3 for an example). %% Valid data set identifiers should be obtained from the data center that %% is currently hosting the data. %% %% Note that AASTeX interprets everything between the curly braces in the %% macro as regular text, so any special characters, e.g. \"#\" or \"_,\" must be %% preceded by a backslash. Otherwise, you will get a LaTeX error when you %% compile your manuscript. Special characters do not %% need to be escaped in the optional, square-bracket argument. %% In this section, we use the \\subsection command to set off %% a subsection. \\footnote is used to insert a footnote to the text. %% Observe the use of the LaTeX \\label %% command after the \\subsection to give a symbolic KEY to the %% subsection for cross-referencing in a \\ref command. %% You can use LaTeX's \\ref and \\label commands to keep track of %% cross-references to sections, equations, tables, and figures. %% That way, if you change the order of any elements, LaTeX will %% automatically renumber them. %% This section also includes several of the displayed math environments %% mentioned in the Author Guide. %% The \\notetoeditor{TEXT} command allows the author to communicate %% information to the copy editor. This information will appear as a %% footnote on the printed copy for the manuscript style file. Nothing will %% appear on the printed copy if the preprint or %% preprint2 style files are used. %% The eqnarray environment produces multi-line display math. The end of %% each line is marked with a \\\\. Lines will be numbered unless the \\\\ %% is preceded by a \\nonumber command. %% Alignment points are marked by ampersands (&). There should be two %% ampersands (&) per line. %% Putting eqnarrays or equations inside the mathletters environment groups %% the enclosed equations by letter. For instance, the eqnarray below, instead %% of being numbered, say, (4) and (5), would be numbered (4a) and (4b). %% LaTeX the paper and look at the output to see the results. %% This section contains more display math examples, including unnumbered %% equations (displaymath environment). The last paragraph includes some %% examples of in-line math featuring a couple of the AASTeX symbol macros. %% The displaymath environment will produce the same sort of equation as %% the equation environment, except that the equation will not be numbered %% by LaTeX. %% If you wish to include an acknowledgments section in your paper, %% separate it off from the body of the text using the" }, "1207/1207.3579_arXiv.txt": { "abstract": "Combined data from gamma-ray telescopes and cosmic-ray detectors have produced some new surprising insights regarding intergalactic and galactic magnetic fields, as well as extragalactic background light. We review some recent advances, including a theory explaining the hard spectra of distant blazars and the measurements of intergalactic magnetic fields based on the spectra of distant sources. Furthermore, we discuss the possible contribution of transient galactic sources, such as past gamma-ray bursts and hypernova explosions in the Milky Way, to the observed flux of ultrahigh-energy cosmic-rays nuclei. The need for a holistic treatment of gamma rays, cosmic rays, and magnetic fields serves as a unifying theme for these seemingly unrelated phenomena. ", "introduction": "Gamma rays from Active Galactic Nuclei (AGN) are studied extensively using ground-based atmospheric Cherenkov telescopes (ACT), as well as Fermi Space Telescope and other instruments. Their signals reveal important information about the sources, as well as about extragalactic background light (EBL) and intergalactic magnetic fields (IGMF) along the line of sight. The same sources are expected to accelerate cosmic rays, although it is more difficult to associate cosmic rays with their sources because the local, galactic magnetic fields alter the arrival directions of cosmic rays. \\subsection{Secondary gamma rays from the line-of-sight interactions of cosmic rays} It was recently proposed that the hardness (and uniform redshift-dependent shape) of gamma-ray spectra of distant blazars can be naturally explained by the line-of-sight interactions of cosmic rays accelerated in the blazar jets~\\cite{Essey:2009zg,Essey:2009ju,Essey:2010er,Essey:2010nd,Essey:2011wv,Murase:2011cy,Razzaque:2011jc,Prosekin:2012ne}. The cosmic rays with energies below $10^{17}-10^{18}$~eV can cross large distances with little loss of energy and can generate high-energy gamma rays in their interactions with cosmic background photons relatively close to the observer. Such {\\em secondary} gamma rays can reach the observer even if their energies are well above TeV. In the absence of cosmic-ray contribution, some unusually hard intrinsic spectra\\cite{Lefa:2011xh} or hypothetical new particles~\\cite{De_Angelis:2007dy,Hooper:2007bq} have been invoked to explain the data. As long as the IGMFs are smaller than $\\sim$10 femtogauss, secondary gamma rays come to dominate the signal from a sufficiently distant source. One can see this from the way the flux scales with distance for primary and secondary gamma rays~\\cite{Essey:2010er}: \\begin{eqnarray} F_{\\rm primary,\\gamma}(d)& \\propto & \\frac{1}{d^2} e^{-d/\\lambda_\\gamma} \\label{exponential} \\end{eqnarray} \\begin{eqnarray} F_{\\rm secondary,\\gamma}(d)& \\propto & \\frac{\\lambda_\\gamma}{d^2}\\Big(1-e^{-d/\\lambda_\\gamma}\\Big) \\\\ & \\sim & \\left \\{ \\begin{array}{ll} 1/d, & {\\rm for} \\ d \\ll \\lambda_\\gamma, \\\\ 1/d^2, & {\\rm for} \\ d\\gg \\lambda_\\gamma . \\end{array} \\right. \\end{eqnarray} Obviously, for a sufficiently distant source, secondary gamma rays must dominate because they don't suffer from the exponential suppression as in Eq.~(\\ref{exponential}). The predicted spectrum turns out to be similar for all the distant AGN, depending only on their redshift. These predictions are in excellent agreement with the data~\\cite{Essey:2009ju,Essey:2009zg,Essey:2010er}. \\begin{wrapfigure}{r}{0.6\\textwidth} \\begin{center} \\includegraphics[width=0.55 \\textwidth]{deltagamma_rev3.eps} \\caption{ Spectral index change $\\delta \\Gamma = \\Gamma_{\\rm GeV} - \\Gamma_{\\rm TeV}$ as a function of redshift. While the low-redshift blazars agree with the Stecker--Scully relation\\cite{Stecker:2006wq}, the data indicate the existence of an additional, distinct population with a weak redshift dependence at redshifts 0.15 and beyond\\cite{Essey:2011wv}. In particular, the recently measured redshift\\cite{Landt:2012it} of PKS 0447-439 is in agreement with the trend. \\label{fig:delta}} \\end{center} \\end{wrapfigure} One can see the transition from primary to secondary gamma rays in Fig.~\\ref{fig:delta}, which shows the spectral index difference for blazar spectra as a function of their redshifts. At small redshifts, the data confirm the Stecker -- Scully relation\\cite{Stecker:2006wq}, but, at redshifts 0.15 and beyond, there is clearly a new population of blazars, whose observed spectral index shows only a weak dependence on the redshift. The nearby population is obviously the blazars from which primary gamma rays are observed. The distant blazars are observed in secondary gamma rays, which are produced in line-of-sight cosmic ray interactions. These secondary gamma rays are produced relatively close to the observer, regardless of the distance to the source. Hence, their redshift dependence is much weaker\\cite{Essey:2011wv}. Finally, there is an intermediate population around redshift 1.2 which is composed of some blazars seen in primary gamma rays and some seen in secondary gamma rays. A recent redshift measurement of PKS 0447-439 redshift\\cite{Landt:2012it} further strengthens our interpretation. Gamma rays with energies above 1~TeV have been observed from this blazar by HESS\\cite{Zech:2011ym}. The spectral properties agree with the trend (Fig.~\\ref{fig:delta}). Furthermore, there is no way for primary gamma rays to reach Earth from such a distant source, while secondary gamma rays provide a consistent explanation of the PKS 0447-439 spectrum~\\cite{Aharonian:2012fu}. This motivates future observations by ACT of blazars with known large redshifts. Secondary gamma rays with TeV and higher energies can be observed even from some sources located at cosmological ($z\\sim 1$) distances~\\cite{Aharonian:2012fu}. \\begin{wrapfigure}{r}{0.6\\textwidth} \\begin{center} \\includegraphics[height=0.45 \\textwidth,angle=270]{essey_fig2.ps} \\caption{Photon (low energy) and neutrino (high energy) spectra~\\cite{Essey:2009ju} expected from an AGN at $z=0.14$ (such as 1ES0229+200), normalized to HESS data points (shown)~\\cite{Aharonian:2005gh}, for $E_{\\rm max}=10^8$GeV, $10^{10}$GeV, and $10^{11}$GeV shown by the solid, dashed, and dash-dotted lines, respectively. \\label{fig:universality}} \\end{center} \\end{wrapfigure} The spectral slope of protons and the level of EBL do not have a strong effect on the spectrum of secondary photons, as one can see from Fig.~\\ref{fig:universality}. However, for the same photon flux, the neutrino flux varies depending on the maximal energy $E_{\\rm max}$ to which the protons are accelerated. Indeed, there are two competing processes that generate secondary photons: $p\\gamma_{EBL}\\rightarrow p\\pi^0 \\rightarrow p\\gamma\\gamma $ and $p\\gamma_{CMB}\\rightarrow pe^+e^-$. For smaller $E_{\\rm max}$, a larger fraction of photons come from the hadronic channel, which is accompanied by production of neutrinos via $p\\gamma_{EBL}\\rightarrow n\\pi^{\\pm}$ followed by the decays of charged pions and the neutron. Neutrino observations can help determine this parameter~\\cite{Essey:2009ju}. \\begin{figure}[t!] \\begin{center} \\includegraphics[width=0.48 \\textwidth]{fig_ebl.eps} \\hfill \\includegraphics[width=0.45 \\textwidth]{fig_B.eps} \\caption{Sensitivity of secondary gamma-ray spectra to the model of EBL and to the average value of IGMFs. {\\em Fermi} upper limit and tentative detection~\\cite{Orr:2011ha} for blazar 1ES 0229+200 are shown at lower energy, and HESS data points at high energy. The predictions of two models of EBL are shown in the left panel according to Essey et al.~\\cite{Essey:2010er}, and the effects of intergalactic magnetic fields are shown in the right panel~\\cite{Essey:2010nd}. \\label{fig:EBL}} \\end{center} \\end{figure} \\subsection{IGMFs and EBL} The success of this picture lends support to the hypothesis of cosmic ray acceleration in AGN. Furthermore, one can use the spectral gamma-ray data to study EBL and IGMFs. The predicted spectra depend to some extent on the EBL model, as shown in Fig.~\\ref{fig:EBL}, although this dependence is too weak to distinguish between different models\\cite{Essey:2010er}. IGMFs, however, can have a strong effect on the goodness of fit. Based on the spectra of several distant blazars, one can set both upper and lower limits on IGMF\\cite{Essey:2010nd}: $$ 10^{-17} {\\rm G} < B < 3\\times 10^{-14} {\\rm G}. $$ \\subsection{Time variability} An important property of secondary gamma rays is the lack on short-scale time variability\\cite{Prosekin:2012ne}. For $E>1$~TeV and $z>0.15$, one expects the signal to be dominated by secondary photons, and any time variability on short scales should be erased by delays in the propagation of protons and electromagnetic cascades. Fig.~\\ref{fig:delays} shows the time delays as a function of the proton energy. \\begin{wrapfigure}{l}{0.6\\textwidth} \\includegraphics[width=0.6\\textwidth,angle=0]{fig1} \\caption{\\label{fig:delays} Time delays of gamma rays emitted at redshift $z=0.17$ for different proton energies $E_{0}$ in a femtogauss random IGMF with a correlation length of 1~Mpc. } \\end{wrapfigure} The present data are not yet sufficient to probe time variability at the relevant energies and redshifts because the data points are too few. While time variability has been observed for {\\em nearby} TeV blazars at TeV energies, as well as for distant TeV blazars at energies above a few hundred GeV, no variability has been reported so far for {\\em distant} TeV blazars at {\\em TeV energies}. One can infer from Fig.~\\ref{fig:delta} how distant the source has to be for the secondary signals to dominate. It is evident that the secondary component takes over for redshifts beyond 0.15. ", "conclusions": "Based on the recent data, one can make several remarkable inferences about the ultrahigh-energy cosmic rays and magnetic fields inside the Milky Way and in the intergalactic space. Gamma-rays detected from most distant blazars are most likely dominated by the secondary photons produced in line-of-sight interactions of cosmic rays. This interpretation allows one to set both upper and lower bounds on intergalactic magnetic fields, $10^{-17} {\\rm G} < B < 3\\times 10^{-14} {\\rm G}$\\cite{Essey:2010nd}. Furthermore, the energy dependent composition of UHECR, with heavier nuclei at high energy, points to a non-negligible contribution from Galactic sources~\\cite{Calvez:2010uh}. Diffusion in turbulent Galactic magnetic field traps the nuclei more efficiently than protons, leading to an increase in the nuclear fraction up to the energy at which iron escapes ($\\sim 30$~EeV). At higher energies, the extragalactic protons should dominate the flux of UHECR, and their arrival directions should correlate with locations of the known sources. If and when the neutrino telescopes, such as IceCube~\\cite{Halzen:2010yj}, detect point sources, one can learn about the cosmic-ray sources and photon backgrounds by comparing the neutrino flux to the photon flux. Neutrino and gamma-ray observations can help distinguish the local Galactic sources from extragalactic sources of UHE nuclei~\\cite{Murase:2010gj,Murase:2010va,Hooper:2010ze}. These inferences open exciting new opportunities for multi-messenger photon, charged-particle, and neutrino astronomy. This work was supported by DOE Grant DE-FG03-91ER40662." }, "1207/1207.3786_arXiv.txt": { "abstract": "{ We investigate non-singular bounce and cyclic cosmological evolutions in a universe governed by the extended nonlinear massive gravity, in which the graviton mass is promoted to a scalar-field potential. The extra freedom of the theory can lead to certain energy conditions violations and drive cyclicity with two different mechanisms: either with a suitably chosen scalar-field potential under a given St\\\"{u}ckelberg-scalar function, or with a suitably chosen St\\\"{u}ckelberg-scalar function under a given scalar-field potential. Our analysis shows that extended nonlinear massive gravity can alter significantly the evolution of the universe at both early and late times. } ", "introduction": "The question on whether there exits a consistent covariant theory for massive gravity, where the graviton acquires a mass and leads to a modification of General Relativity at large distances, was initiated by Fierz and Pauli a long time ago \\cite{Fierz:1939ix}. However, it was observed that the nonlinear terms required by massive gravity \\cite{Vainshtein:1972sx}, which can give rise to continuity of observables \\cite{vdam, vdam2}, lead inevitably to the existence of the Boulware-Deser (BD) ghost \\cite{Boulware:1973my}, and thus make the theory unstable. Although for the subsequent decades it was believed that there is no consistent way to construct a massive gravity free of ghosts, a nonlinear extension of massive gravity was constructed recently by de Rham, Gabadadze and Tolley \\cite{deRham:2010ik, deRham:2010kj}. In this model, the BD ghost can be removed in the decoupling limit to all orders in perturbation theory through a systematic construction of a covariant nonlinear action (see \\cite{Hinterbichler:2011tt} for a review). Although it is still controversial to prove the absence of BD ghost at the non-perturbative level, the theoretical and phenomenological advantages led to a wide investigation of this theory \\cite{Koyama:2011yg, Hassan:2011hr, deRham:2011rn, CuadrosMelgar:2011yw, D'Amico:2011jj, Hassan:2011zd, Kluson:2011qe, Gumrukcuoglu:2011ew, Volkov:2011an, vonStrauss:2011mq, Comelli:2011zm, Hassan:2011ea, Berezhiani:2011mt, Gumrukcuoglu:2011zh, Khosravi:2011zi, Brihaye:2011aa, Park:2010rp,Park:2010zw, Buchbinder:2012wb, Ahmedov:2012di, Bergshoeff:2012ud, Crisostomi:2012db, Paulos:2012xe, Hassan:2012qv, Comelli:2012vz, Sbisa:2012zk, Kluson:2012wf, Tasinato:2012mf, Morand:2012vx, Cardone:2012qq, Baccetti:2012bk, Gratia:2012wt, Volkov:2012cf, DeFelice:2012mx, Gumrukcuoglu:2012aa, deRham:2012kf, Berg:2012kn, D'Amico:2012pi, Baccetti:2012re, Fasiello:2012rw, D'Amico:2012zv, Baccetti:2012ge, Langlois:2012hk,Gong:2012yv}. Despite the successes of nonlinear massive gravity, it was also noticed that certain cosmological instabilities still exist in the case where the physical and the fiducial metrics have simple homogeneous and isotropic forms \\cite{DeFelice:2012mx}. This behavior motivated researches towards extensions of nonlinear gravity models, namely the construction of nonlinear massive gravity with less symmetric metrics \\cite{D'Amico:2011jj, Gumrukcuoglu:2012aa}. However, in \\cite{Huang:2012pe} a different approach was followed, and nonlinear massive gravity was extended allowing for the graviton mass to vary. This could be realized by introducing an additional scalar field, which coupling to the graviton potentials produces an effective, varying, graviton mass. Moreover, this extension provides a natural way to modify General Relativity not only in the IR but also in the UV, since the graviton mass can be evolved into a large value at the early universe \\cite{Saridakis:2012jy}. Therefore, it is interesting to study the cosmological implications of this scenario at early times, and this is a main topic of the present work. On the other hand, it is well known that cosmological evolution governed by standard Einstein gravity usually suffers from the problem of initial singularity if Null Energy Condition (NEC) is preserved \\cite{Borde:1993xh}. A potential solution to the cosmological singularity problem may be provided by non-singular bouncing cosmologies \\cite{Mukhanov:1991zn, Brandenberger:1993ef, Cai:2012va}. Such scenarios have been constructed within various approaches to modified gravity \\cite{Veneziano:1991ek, Khoury:2001wf, Brustein:1997cv, Kehagias:1999vr, Shtanov:2002mb, Saridakis:2007cf, Cai:2010zma, Cai:2011tc, HLbounce, Cai:2009in,Saridakis:2009bv,Leon:2009rc, Bojowald:2001xe, Martin:2003sf,Saridakis:2010mf,Biswas:2010zk,Leon:2010pu,Biswas:2011ar, Biswas:2012bp}. Generally, a non-singular bouncing model can be acquired by using NEC violating matter \\cite{ Dabrowski:2003jm,Dabrowski:2004hx,Yifu1, Yifu1b, Yifu2, Yifu2b, Cai:2008qw}, by making use of various mechanisms \\cite{Khoury, Creminelli, Chunshan, Tirtho1, Tirtho3, Taotao, Damien, Abramo:2007mp}. Note that in the case of a positive curvature a generic bounce can be obtained by violating Strong Energy Condition (SEC) only \\cite{Starobinskii:1978yy}, however the singularity reappears in the fact that the number of regular bounces is finite \\cite{Page:1984qt}. Furthermore, the perturbation theory of non-singular bounce cosmology and its relation to observables was extensively studied in the literature \\cite{Yifu2, Yifu2b, Cai:2008qw, Cai:2009hc, Cai:2009rd, Cai:2009fn, Liu:2010fm, Cai:2011zx, Cai:2011ci}. The extension of all the above bouncing scenarios is the old idea of cyclic cosmology \\cite{tolman}, in which the universe presents a periodic sequence of contractions and expansions. Cyclic cosmology has attracted a significant interest the last years \\cite{Steinhardt:2001st} since it brings different insights for the origin of the observable universe \\cite{Lidsey:2004ef, Piao:2004hr, Piao:2004me, Xiong:2007cn, Piao:2009ku, Piao:2010cf, Liu:2012gu, Xiong:2008ic, Cai:2009zp, cyclic, cyclic1,Sahni:2012er} (see \\cite{Novello:2008ra, Cai:2011bs} for recent reviews). In the present work we are interested in constructing scenarios of cyclic cosmology in a universe governed by the extended, varying-mass, nonlinear massive gravity. Particularly, we first determine the St\\\"{u}ckelberg-scalar function and we suitably reconstruct the form of the potential of the scalar field which leads to a cyclic universe. Alternatively, we first determine the scalar potential and we reconstruct the St\\\"{u}ckelberg-scalar function in order to obtain cyclicity. Interestingly enough cyclicity is easily acquired in this framework, since extended nonlinear massive gravity can violate certain energy conditions and therefore has fruitful implications to physics of the universe at both early and late times. This paper is organized as follows. In Section \\ref{model} we briefly review the cosmological equations under the extended scenario of nonlinear massive gravity with a scalar field being introduced to evolve the graviton mass. In section \\ref{bouncingsol} we construct scenarios of bouncing and cyclic universe. In particular, in subsection \\ref{caseb} we start with a given St\\\"{u}ckelberg-scalar function and we reconstruct the scalar potential; while in \\ref{caseW} we start from a given scalar potential and we determine the corresponding St\\\"{u}ckelberg-scalar function that leads to cyclicity. Finally, section \\ref{Discussion} is devoted to the summary of our results. ", "conclusions": "\\label{Discussion} In the present work we investigated bouncing and cyclic cosmological behaviors in a universe governed by extended massive gravity, in which the graviton mass has been promoted to a function of an extra scalar field. This model has additional freedom comparing to usual massive gravity, and thus it leads to significantly different and richer behavior \\cite{Huang:2012pe,Saridakis:2012jy}. In particular, although the scenario at late times tends to coincide with standard quintessence, at early and intermediate times the effective graviton mass can be large and thus play a crucial role in the universe evolution. Amongst others, the capability of the scenario to lead to violation of the Null Energy Condition can lead to a bounce or a turnaround, the successive sequence of which can naturally give rise to cyclic cosmology. Extended nonlinear massive gravity has an enhanced freedom (apart from the extra scalar field one can see that the St\\\"{u}ckelberg-scalar function $b(\\tau)$ is not constrained as in usual massive gravity and it can be free), therefore it can drive cyclicity with two different mechanisms. The first is to impose an arbitrary St\\\"{u}ckelberg-scalar function $b(\\tau)$ and suitably choose the usual scalar field potential $W(\\psi)$ in order to obtain a cyclic behavior. It proves that one should use an oscillating $W(\\psi)$ as expected, and the only requirement for this procedure to hold is to obtain a positive $\\dot{\\psi}^2(\\tau)$. Thus, this is not possible for every $b(\\tau)$ form, that is not every $b(\\tau)$ can be consistent with cyclicity. The second mechanism to drive cyclic behavior is exactly the St\\\"{u}ckelberg-scalar function $b(\\tau)$. In particular, imposing an arbitrary scalar-field potential $W(\\psi)$ we suitably choose $b(\\tau)$ in order to obtain a cyclic behavior. It proves that one should use an oscillating $\\dot{b}(\\tau)$ as expected (since $\\dot{b}(\\tau)$ appears in the equations and not $b(\\tau)$). Similarly to the first mechanism above, not all scalar potentials are consistent with cyclicity. A crucial issue in all bouncing and cyclic scenarios is the processing of perturbations through the bounces. A simple check on the stability of this type of models is to look at the generalized Higuchi bound derived in \\cite{Fasiello:2012rw} (see also \\cite{Grisa:2009yy} for a different expression in an earlier model). In particular, as an effective theory, nonlinear massive gravity is reliable only when the scale is well beneath the cut off $\\Lambda_3 = (M_P m_g^2)^{1/3}$, where $m_g^2$ is the graviton mass square, since upon $\\Lambda_3$ helicity 0 mode couples strongly to helicity 1 and helicity 2 modes, and thus effective field theory breaks down. Therefore, in the present extended scenario one should compare\\footnote{We wish to thank the referee for mentioning this point.} the various appeared scales, such are $H$,$\\dot{H}$,$\\ddot{H}$ etc, with $\\Lambda_3 = [M_P V(\\psi)]^{1/3}$. Indeed, in all the above examples $H/\\Lambda_3\\lesssim10^{-3}$, while $\\dot{H}/\\Lambda_3\\lesssim10^{-5}$, $\\ddot{H}/\\Lambda_3\\lesssim10^{-7}$ etc, and thus the scenario is reliable. On the other hand note that $V(\\psi)$ is always much smaller than $M_P^2$ ($V(\\psi)/M_P^2\\lesssim10^{-3}$ and $V_0/M_P^2\\lesssim10^{-1}$), which is an additional requirement for the robustness of the scenario. Therefore, at this level, we can conclude that our model is well behaved when the perturbations pass through the bouncing points. However, we should notice that since a cosmic scalar is introduced to drive the graviton mass varying along background evolution, the stability issue arisen from this scalar field ought to be taken into account in a global analysis. Such a complete perturbation analysis of the extended nonlinear massive gravity lies beyond the scope of the present work and it is left for future investigation. \\vskip .2in \\noindent {\\large{{\\bf" }, "1207/1207.2780_arXiv.txt": { "abstract": "We derive an analytical approximation of nonlinear force-free magnetic field solutions (NLFFF) that can efficiently be used for fast forward-fitting to solar magnetic data, constrained either by observed line-of-sight magnetograms and stereoscopically triangulated coronal loops, or by 3D vector-magnetograph data. The derived NLFFF solutions provide the magnetic field components $B_x({\\bf x})$, $B_y({\\bf x})$, $B_z({\\bf x})$, the force-free parameter $\\alpha({\\bf x})$, the electric current density ${\\bf j}({\\bf x})$, and are accurate to second-order (of the nonlinear force-free $\\alpha$-parameter). The explicit expressions of a force-free field can easily be applied to modeling or forward-fitting of many coronal phenomena. ", "introduction": "The coronal magnetic field can be constrained in a number of ways, such as by extrapolation of photospheric magnetograms or vector-magnetograph data, by radio observations of gyroresonance layers above sunspots, of by coronal seismology of oscillating loops. Before the advent of the STEREO mission, attempts were made to model observed coronal loops with stretched potential field solutions (Gary and Alexander, 1999), to fit a linear force-free model with solar-rotation stereoscopy (Wiegelmann and Neukirch, 2002; Feng \\etal 2007a), by tomographic reconstruction with magnetohydrostatic constraints (Wiegelmann and Inhester, 2003; Ruan \\etal 2008), by magnetic modeling applied to spectropolarimetric loop detections (Wiegelmann \\etal 2005), or by magnetic field supported stereoscopic loop triangulation (Wiegelmann and Inhester, 2006; Conlon and Gallagher, 2010). Recently, stereoscopic triangulation of coronal loops with the STEREO mission became available, which constrains the 3D geometry of coronal magnetic field lines (Aschwanden \\etal 2008; Aschwanden, 2009). The plethora of coronal high-resolution data allows us now to compare different magnetic models and to test whether they are self-consistent. A critical assessment of nonlinear force-free field (NLFFF) codes revealed the disturbing fact that different NLFFF codes yield incompatible results among themselves, and exhibit significant misalignments with stereoscopically triangulated loops (DeRosa \\etal 2009; Sandman \\etal 2009; Aschwanden and Sandman, 2010; Sandman and Aschwanden, 2010; Aschwanden \\etal 2012a,b). The discrepancy was attributed to uncertainties in the boundary conditions as well as to the non-forcefreeness of the photosphere and lower chromosphere. Earlier tests with the virial theorem already indicated that the magnetic fields in the lower chromosphere at altitudes of $h \\lapprox 400$ km are not force-free (Metcalf \\etal 1996). Constraints by coronal tracers thus have become an important criterion to bootstrap a self-consistent magnetic field solution. The misalignment between theoretical extrapolation models and stereoscopically triangulated loops could be minimized by using potential field models with forward-fitted unipolar magnetic charges (Aschwanden and Sandman, 2010) or dipoles (Sandman and Aschwanden, 2011). In this Paper we go a step further by deriving a simple analytical approximation of nonlinear force-free field solutions that is suitable for fast forward-fitting to stereoscopically triangulated loops or to some other coronal observations. While accurate solutions of force-free magnetic fields have been known for special mathematical functions (Low and Lou, 1990) that have been used to reconstruct the local twist of coronal loops (Malanushenko \\etal 2009, 2011), they are not suitable for forward-fitting to entire active regions. In contrast, our theoretical framework entails the representation of a potential or non-potential field by a superposition of a finite number of elementary field components that are associated with buried unipolar magnetic charges at arbitrary locations, each one being divergence-free and force-free to a good approximation, as we test numerically. While this Paper contains the analytical framework of the magnetic field model, the numerical forward-fitting code with applications to observations will be presented in a Paper II (Aschwanden and Malanushenko, 2012), and applications to stereoscopically observed active regions in Aschwanden \\etal (2012a,b). ", "conclusions": "The coronal magnetic field has generally been computed by extrapolation from lower boundary data in form of photospheric magnetograms $B_z(x,y,z=z_{ph})$ or vector-magnetograph data ${\\bf B}(x,y)$, using a numerical extrapolation algorithm that fulfills the conditions of force-freeness ($\\nabla \\cdot {\\bf B}$) and divergence-freeness $\\nabla \\times {\\bf B} = \\alpha({\\bf r}) {\\bf B}$, where $\\alpha({\\bf r})$ is a scalar function in space ${\\bf r}$. These extrapolation algorithms are very computing-intensive, because a good solution requires many iterations on a large computational 3D-grid that has sufficient spatial resolution to resolve the relevant magnetic field gradients. The accuracy of these numerical solutions depends very much on the noise in boundary vector magnetic field data as well as on deviations of photospheric fields from a force-free state. Recent stereoscopic triangulation of coronal loops has demonstrated a considerable mismatch between the extrapolated fields and the actual coronal loops, which cannot easily be reconciled with extrapolation algorithms, since they have only a very limited degree of freedom within the noise of the boundary data. Moreover, since each NLFFF solution is very time-consuming to compute, these algorithms are not suitable for forward-fitting. The forward-fitting of magnetic field solutions to observed data requires a faster algorithm to compute many NLFFF solutions for variable boundary data or for coronal constraints as given by stereoscopic 3D reconstructions. The fastest computational way would be an explicit analytical solution for the coronal field vectors ${\\bf B}({\\bf r})$ as a function of some suitable parameterization of the boundary data or coronal constraints. There exist some analytical solutions of nonlinear force-free fields, such as a class of solutions in terms of Legendre polynomials (Low and Lou, 1990), which is characterized by some spatial symmetry and has been used to test numerical extrapolation algorithms (\\eg DeRosa \\etal 2009; Malanushenko \\etal 2009). However, to our knowledge, the class of analytical NLFFF solutions of Low and Lou (1990) has never been applied to forward-fitting of observed data, such as line-of-sight magnetograms, vector magnetograph 3D data, or to stereoscopically triangulated loops. Moreover, the special class of NLFFF solutions derived in Low and Lou (1990) correspond to harmonics of Legendre polynomials, which have a high degree of symmetry that does not match realistic observations of active regions, and thus is not suitable for forward-fitting to real data. What we need to model observed solar magnetic data with high accuracy is: (1) an explicit formulation of an analytical NLFFF solution; (2) a parameterization of the NLFFF solution with a sufficient large number of free parameters that can be forward-fitted to data and converges close to observations; and (3) a fast computation algorithm that can perform many interations without computing-intensive techniques. Hence, such a project consists of developing a suitable analytical formulation first, and then to implement the analytical solutions into a forward-fitting code. In this paper we have undertaken the first step. We started with a potential-field parameterization in terms of $N_{\\rm m}$ buried magnetic charges, which is defined by $4 N_{\\rm m}$ free parameters that can easily be extracted from an observed line-of-sight magnetogram $B_z(x,y)$ with arbitrary accuracy, as demonstrated in two recent studies (Aschwanden and Sandman, 2010; Aschwanden \\etal 2012a). The key concept of this potential-field representation is that an arbitrary complex 3D magnetic field can be decomposed into a finite number of elementary magnetic field components, where each one simply consists of a quadratically decreasing radial field of a buried magnetic charge. Divergence-freeness is conserved due to the linearity in the superposition of elementary field components. In a next step we extended the elementary potential-field component to a nonpotential-field component by adding a uniform twist that can be parameterized by the force-free $\\alpha$-parameter. Such an elementary nonpotential field component requires five free parameters, consisting of the four potential-field parameters plus the force-free $\\alpha$-parameter. We derived an explicit analytical formulation of the radial $B_r(r,\\theta)$ and azimuthal field vector $B_\\varphi (r, \\theta)$ that represents an approximative solution of the divergence-free and force-free condition to second order ($\\propto \\alpha^2$). This solution is very accurate for weakly non-potential fields and converges to the potential field solution for $\\alpha=0$. In analogy to the potential-field representation, we represent a general non-potential field solution with a superposition of elementary non-potential field components and prove that the divergence-freeness and force-freeness is conserved to second-order accuracy in our NLFFF approximation. We calculated some examples of potential and non-potential fields that mimic an isolated sunspot, a dipolar and a quadrupolar configuration, as well as more complex multi-polar configurations. The examples show that the magnetic field of arbitrary complex active regions can be represented with our parameterization. Increasing the force-free $\\alpha$-parameter distorts circular field lines into helical and sigmoid-shaped geometries. Our parameterization allows one to compute either field lines (starting from arbitrary locations), 3D datacubes of magnetic field vectors, of maps of the force-free $\\alpha$-parameter and electric current $j_z$ (Figure 9). We tested the figures of merit for divergence-freeness and force-freeness, which amount to $L_{\\rm d} \\lapprox 10^{-3}$ and $L_{\\rm f} \\lapprox 10^{-2}$. The examples demonstrate also the computing speed of this algorithm, which amounts to the order of $\\approx 1$ s for a computation grid that encompasses a typical active region with the spatial resolution of MDI. Thus, we envision that a full-fletched forward-fitting code can converge within a few seconds to a few minutes, depending on the number of iterations and number of magnetic field components. Where do we go from here? The next step is the development of a forward-fitting code that uses the magnetic field parameterization described here (see Paper II). We envision the applications to at least three different sets of constraints, requiring three different versions of forward-fitting codes: (i) line-of-sight magnetograms $B_z(x,y)$ and 3D coordinates $[x(s), y(s), z(s)]$ of stereoscopically triangulated loops; (ii) line-of-sight magnetograms $B_z(x,y)$ and 2D coordinates $[(x(s),y(s)]$ of traced loops; and (ii) vector-magnetograph data $[B_x(x,y)$, $B_y(x,y)$, $B_z(x,y)]$. The first application requires STEREO data, while the second one can be obtained from any EUV imager (\\eg AIA/SDO, TRACE, EIT/ SOHO). The third application can be conducted with the new HMI/SDO data and is equivalent to other NLFFF extrapolation codes without coronal constraints, while the first two use coronal tracers and alleviate the force-free assumption of photospheric data. We envision that these three applications will reveal insights into a number of crucial questions in a novel way. There is a large number of physical problems and issues that can be addressed with the anticipated forward-fitting code, such as: (i) The force-freeness of the photosphere; (ii) the accuracy of NLFFF solutions; (iii) the spatial distribution of electric currents in active regions; (iv) the temporal evolution of currents before and during flares; (v) the spatial distribution of current dissipation and coronal heating; (vi) helicity injection; (vii) the 3D geometry of coronal loops which is needed for hydrodynamic modeling; (viii) scaling laws of the volumetric heating function with other physical parameters; (ix) tests of the magnetic field strength inferred from coronal seismology, \\etc There exists hardly a phenomenon in the solar corona that can be modeled without the knowledge of the coronal magnetic field." }, "1207/1207.2749_arXiv.txt": { "abstract": "This is a brief note to comment on some recent papers addressing the Monoceros ring. In our view, nothing new was delivered on the matter: No new evidence or arguments are presented which lead to think that the over-densities in Monoceros must not be due to the flared thick disc of the Milky Way. Again, we restate that extrapolations are easily misleading and that a model of the Galaxy is not the Galaxy. Raising and discussing exciting possibilities is healthy. However, enthusiasm should not overtake and produce strong claims before thoroughly checking simpler and more sensible possibilities within their uncertainties. In particular, claiming that a reported structure, such as the Monoceros Ring, is not Galactic (an exciting scenario) should not be done without rejecting the possibility of being due to the well established warped and flared disc of the Milky Way (simpler). ", "introduction": " ", "conclusions": "This is a brief note to comment on some recent papers addressing the Monoceros ring. As nothing new has been delivered on the matter, there is also nothing new to add. Concretely, the new data produce no evidence requiring the interpretation of the Monoceros over-density as a structure not belonging to the Milky Way. Thus, there is no need to produce a full paper explaining why not. A simple note like the present one should be enough. Since silence gives consent---as the proverb goes---, here we break the silence and express our disagreement. The so called Monoceros ring/stream is a hypothesis for explaining a reported over-density of stars (with respect to some standard models of the Milky Way, such as the Besan\\c con model; Robin et al. 2003) in a large area of the sky approximately parallel to the Galactic plane, in the latitude range $10^{\\circ}<|b|<35^{\\circ}$ and spanning most of the second and third quadrants (e.g., Ibata et al. 2003; Conn et al. 2008). It has been conjectured that this structure would be due to the remnants of a dwarf galaxy cannibalized by the Milky Way. The core of the progenitor would be associated to a further over-density of stars identified in the Canis Major subregion of the Monoceros ring (Martin et al. 2004). However, the over-density of stars in Canis Major was soon explained as an effect of the warped+flared disc of the Milky Way (Momany et al. 2004, 2006; L\\'opez-Corredoira 2006; L\\'opez-Corredoira et al. 2007) with some features in the color-magnitude diagrams due to the warped Norma--Cygnus spiral arm (Carraro et al. 2005, 2007, 2008; Moitinho et al. 2006; Piatti \\& Clari\\'a 2008). Since then, in the last years, the subject of the extragalactic origin of Canis Major was mostly dropped in the literature. It seemed that the purely Galactic explanation of the phenomenon had been mostly accepted. Concerning the Monoceros ring as a whole, Ibata et al. (2003) and Momany et al. (2006) suggested that it can be explained by the flare of the Galactic outer disc, and Hammersley \\& L\\'opez-Corredoira (2011) made explicit calculations showing how a flare in the thick disc (an element not included in models such as the Besan\\c con model) fits approximately the observed over-density in some regions of the anti-centre. The most recent deep surveys clearly show that there are stars out to at least $R$=20~kpc (e.g., Momany et al. 2006, Reyl\\'e et al. 2009, Carraro et al. 2010). That the disc is flared, should not come as a surprise. It is expected (Momany et al. 2006) and not an ad-hoc hypothesis such as the one of an extragalactic stream (or a reported three-fold layer of streams; Li et al. 2012) parallel to the plane. The explanations in terms of the structure of the Milky Way also contemplate the characteristics of the observed populations, including metallicities, kinematics and spatial densities. We find no observations leading to the necessity of considering the reported over-densities not naturally due to the structure the Milky Way. Lately, four recent papers (Sollima et al. 2011; Meisner et al. 2012; Conn et al. 2012; Li et al. 2012) revive the claims of an extra-galactic origin for the Monoceros ring. Below, we address the arguments and conclusions of these articles: \\begin{description} \\item[Density distribution:] Sollima et al. (2011) affirm that no model of the Milky Way is able to explain the density distribution of the Monoceros structure. The statement is puzzling since Hammersley \\& L\\'opez-Corredoira (2011) had previously shown that a simple model of a flared thick disc does explain approximately the density distribution under discussion. The lines of sight considered in both studies were close to each other [one of the lines of sight of Hammersley \\& L\\'opez-Corredoira (2011) is $\\ell=183 ^\\circ$, $b=21^\\circ $, very close to the first field of Sollima et al. (2011) in $\\ell=180^\\circ $, $b=21^\\circ $]. While Sollima et al. (2011) use the same model as Hammersley \\& L\\'opez-Corredoira (2011), Sollima et al. (2011) affirm that the model does not produce a detectable over-density bump. Figure 3 of Sollima et al. (2011) shows synthetic colour-magnitude diagrams with no clear main sequence. However, Hammersley \\& L\\'opez-Corredoira (2011) do reproduce such a main sequence in Monoceros at magnitudes between $g=20$ and 22, for a population of dwarfs with $g-r$ between 0.36 and 0.49. Moreover, other authors supporting the extra-Galactic scenario (e.g., Conn et al. 2012) could also reproduce the density distribution with a flared thick disc. Thus, we are led to impression that the calculations in Sollima et al. (2011) are in error and that a flared model can reproduce the over-density. Conn et al. (2012) also use the model of Hammersley \\& L\\'opez-Corredoira (2011) to fit the morphology of Monoceros over-density and arrive to conclusions roughly similar to those of Hammersley \\& L\\'opez-Corredoira (2011). Although considering different regions than Hammersley \\& L\\'opez-Corredoira (2011) and achieving better fits with some different parameters with respect to Hammersley \\& L\\'opez-Corredoira (2011), Conn et al. (2012) arrive to qualitatively similar results. Interestingly, Conn et al. (2012) find that the over-density with respect to a non-flared model starts at distances of around 5-7 kpc from the Sun, whereas the regions closer to the anti-centre used by Hammersley \\& L\\'opez-Corredoira (2011) indicate a start at distances of 8-10 kpc. The difference is likely due to some lopsidedness of the disc or non-axisymmetry of the flare (L\\'opez-Corredoira \\& Betancort-Rijo 2009). The analysis of the density carried out by Li et al. (2012) is much simpler: the authors simply point out that the number of observed stars is much higher than that expected from a ``canonical'' thick disc, where by ``canonical'' it is meant a non-flared model, but they cannot exclude a flared disc. \\item[Metallicity, stellar populations:] Although observations and discussions on the metallicity of Monoceros were already presented in many previous papers, Meisner et al. (2012), Conn et al. (2012), and Li et al. (2012) revive the theme with new data. These data yield the same results as before: $[Fe/H]\\approx -1.0$ for the first two papers and $[Fe/H]\\approx -0.8$ for Li et al. (2012). Surprisingly, and despite previous work (Momany et al. 2006, Hammersley \\& L\\'opez-Corredoira 2011), these new studies again affirm that such a metallicity is incompatible with the populations of our Galaxy. Conn et al. (2012) argue that the thick disc should have an average metallicity $[Fe/H]\\approx -0.6$ and no radial gradient. The argument is based on analyses of low $R$ and low $|z|$ stars extrapolated to high $R$ and high $|z|$. But extrapolations can easily be inadequate: the observed regions of Monoceros are at $R\\approx 20$ kpc and $z\\approx 4$ kpc for which there are no observations constraining the thick disc which are independent of Monoceros itself. As discussed below, it is not surprising to find in this region stars with a metallicity 0.2-0.4 dex lower than stars at smaller $R$ and $|z|$. The statement by Conn et al. (2012) that there is no radial metallicity gradient in the thick disc is not strictly correct: As an example, in Cheng et al. (2012), which Conn et al. (2012) cite, it is shown that there is a significant negative gradient for stars with $[\\alpha /Fe]<0.2$ (Cheng et al. 2012, Fig. 4), and Monoceros has indeed a significant number of stars with $[\\alpha /Fe]<0.2$ (Meisner et al. 2012). Vertical gradients are also detected in the nearby thick disc (Bilir et al. 2012). In any case, a small average metallicity gradient of $\\sim -0.03$ dex/kpc in the radial and vertical directions is enough to explain $[Fe/H]\\approx -1.0$ and cannot be excluded at the present. Ironically, the Besan\\c con model of the Milky Way --- used in arguing for an extragalactic origin of Monoceros in what concerns predicted stellar densities and colour-magnitude diagrams --- gives that the very outer thick disc should have an average metallicity $[Fe/H]$ very close to -1.0 (L\\'opez-Corredoira et al. 2007, Fig. 3). Li et al. (2012) find that the Monoceros ring population at $g\\approx 20$ has a bluer turn-off than the disc population which is closer at $g=17-18$. This is also expected since most Monoceros stars are presumably thick disc stars (with some small contribution from halo stars), whereas closer stars are a mixture of thin+thick disc with higher metallicity. The turnoff colour of $(g-r)=0.30-0.31$ measured by Li et al. (2012) for Monoceros is indeed very similar to the turnoff colour measured for the thick disc: $(g-r)\\approx 0.33$ (Chen et al. 2001). There is no problem in interpreting this stellar population as belonging to the thick disc. \\item[Radial velocities:] Li et al. (2012) make an interesting point: that the line of sight velocities in the range $150^\\circ <\\ell <190^\\circ $ are significantly higher than those expected from a canonical thick disc. In particular, that the average radial velocity at $\\ell=180^\\circ $ is significantly different from zero. As Li et al. (2012) state, all axisymmetric disc models predict a zero average line-of-sight velocity directly towards the anti-centre, independently of rotation speed and distance to the stars, because at that longitude we are looking perpendicularly to the circular motion of the disc stars. Li et al. (2012) interpret this non-zero detection as proof against a thick disc origin for Monoceros. However, the assumption of perfectly circular orbits for the outer disc is not unquestionable. In fact, many spiral galaxies exhibit some non-axisymmetry/lopsidedness in their outer discs (Jog \\& Combes 2009). There is also the possible phenomenon of stellar migration (Roskar et al. 2008) which displaces stars to different radii in non-circular orbits. A non-axisymmetric outer disc is not only a possibility, but it is also most likely. Thus, a non-zero average radial velocity towards the anti-centre is not enough to discard the thick disc origin of Monoceros. \\end{description} We have shown how and why no new elements on the Monoceros affair are brought by a number of recent papers (Sollima et al. 2011, Conn et al. 2012, Meisner et al. 2012, and Li et al. 2012). We reaffirm that extrapolations are easily misleading and that {\\bf a model of the Galaxy is not the Galaxy}. The Besan\\c con model, despite being a very good model, is not correct in reproducing all the features of the outer disc. In particular, the Besan\\c con model assumes a disc stellar truncation at $R\\sim 14$ kpc. Although some studies argue for a cut-off of the stellar disc at that distance (e.g., Minniti et al. 2012), the deficit of outer in-plane stars is an expected artefact of assuming a non-flared disc. Moreover, stars have been detected at $R\\sim 20$ kpc with higher densities than those expected from models with a closer disc cut-off (e.g., Momany et al. 2006, Reyl\\'e et al. 2009, Carraro et al. 2010). The unnecessary assumption of a close cut-off produces a cascade of further assumptions and loose ends of which the most dramatic is that the majority of stars at far galacto-centric distances are extragalactic. Certainly, the hypothesis of Monoceros as an extragalactic tidal stream is not discarded, but there are no reasons to support it since it can be explained in terms of known features of our Galaxy. Such a strong claim should not be made without verifying the expected features, within the uncertainties, of the Galaxy. This note is a cautionary tale on how models should not be over-interpreted. Li et al. (2012) also discuss structures designated as the``Anti-Center Stream'' and the ``Eastern Banded Structure'' which do not look much better cases than Monoceros, but these will not be discussed here. Nonetheless, there are other tidal streams (e.g., Sagittarius) with solid observational support." }, "1207/1207.0785_arXiv.txt": { "abstract": "\\change{ Hydrodynamic stability has been a longstanding issue for the cloud model of the broad line region in active galactic nuclei. We argue that the clouds \\rvtwo{may be} gravitationally bound to the supermassive black hole. If true, stabilisation by thermal pressure alone becomes even more difficult. We further argue that if magnetic fields should be present in such clouds at a level that could affect the stability properties, they need to be strong enough to compete with the radiation pressure on the cloud. This would imply magnetic field values of \\rvtwo{a few} Gauss for a sample of Active Galactic Nuclei we draw from the literature. } \\change{We then investigate the effect of several magnetic configurations on cloud stability} in axisymmetric magnetohydrodynamic simulations. For a purely azimuthal magnetic field which provides the dominant pressure support, \\change{the cloud} first gets compressed by the opposing radiative and gravitational forces. The pressure inside the cloud then increases, and it expands vertically. Kelvin-Helmholtz and column density instability lead to a filamentary fragmentation of the cloud. This radiative dispersion continues until the cloud is shredded down to the resolution level. For a helical magnetic field configuration, a much more stable cloud core survives with a stationary density histogram which takes the form of a power law. Our simulated clouds develop sub-Alfv\\'enic internal motions on the level of a few hundred~km/s. ", "introduction": "Broad emission lines are produced in the immediate vicinity of optically active super-massive black holes \\citep[SMBH, for reviews see e.g.][]{Ost88,Pet97,Net08}. They may be used to infer the black hole mass in galaxies with such an Active Galactic Nucleus \\citep[AGN, e.g.][]{Benea09}, and are a standard part of optically active AGN. The basic line emitting entity is usually referred to as a cloud. The emission mechanism is photoionisation by the central parts of the accretion disc. Photoionisation models \\change{predict} the clouds to have a typical temperature of order $10^4$~K, number densities of $n_\\mathrm{cl}=10^{10\\pm1}$~cm$^{-3}$, sizes of $R_\\mathrm{cl} = 10^{12\\pm1}$~cm and column densities of \\change{$N_\\mathrm{cl}>2\\times 10^{22}$~cm$^{-2}$} \\citep[e.g.][]{KK81,FE84,RNF89}. \\change{Further important constraints come from reverberation mapping \\citep[e.g.][see below for a comparison of results of these two methods]{Pet88}.} The complete physics of these clouds is however highly complex and involves also pressure, radiative, centrifugal, gravitational, and magnetic forces, probably on a very similar level. We are not aware of any attempt to include all these processes into a single model, but different authors have looked at some particular aspects of the problem: A general assumption for the cloud ensemble is often virial equilibrium. Since the gravitational potential is dominated by the SMBH, this would imply Kepler orbits. This treatment neglects the contribution of radiation pressure to the dynamics. The latter restriction has been relaxed in a recent series of papers (\\citeauthor{Marcea08} \\citeyear{Marcea08}, \\citeyear{Marcea09}; \\citeauthor{Net09} \\citeyear{Net09}, \\citeyear{Net10}; \\citeauthor*{KBS11} \\citeyear{KBS11} (hereafter: paper~I)). The general finding is that the radiation force may contribute significantly, and that in this case, the clouds on bound orbits could be significantly sub-Keplerian. For low angular momentum (strong radiation pressure support), the orbits have to be highly excentric. \\change{The importance of radiation pressure has actually been recognised early on \\citep[e.g.][]{TMcK73,McKT75,BM75,BM79,Mat76}. In fact, many of the earlier articles study the possibility that radiation pressure is dominant, also compared to gravity, and the clouds are unbound. In this case, it would however not be easy to understand why the black hole masses derived from reverberation mapping under the assumption of virial equilibrium may be brought into agreement with the correlation between black hole mass and host galaxy properties with a uniform scaling factor \\citep{Onkea04}. For optically thin clouds, one would expect almost no response to the variability of the ionising continuum. The strong variability of the emission lines, lead by variations in the continuum is therefore evidence for the presence of optically thick clouds \\citep[e.g.][]{SG07}. However, the relative importance of gravity versus radiation pressure is proportional to the optical thickness (compare equation (\\ref{prg_ratio}), below).\\newline This opened up the possibility that the Broad Line Region (BLR) is gravitationally bound and possibly disc-like with a significant total angular momentum. In fact, several pieces of evidence that point in this direction have been found over recent years} (compare paper~I and references therein). One recent piece of evidence comes from spectropolarimetry, where the BLR is spatially resolved by an equatorial scattering region, leading to different polarisation angles in the red and blue wings of emission lines \\citep{Smea05}. Most recently, \\citet{KolZet11} have shown that the shape of the broad emission lines in many objects may be well fit with the assumption of a turbulent thick disc. Cloud stability and confinement is a long-standing issue \\citep[for a review]{OM86}: The clouds should be in rough pressure equilibrium with their environment (e.g. \\citeauthor*{KMKT81} \\citeyear{KMKT81}; \\citeauthor{Krolik88} \\citeyear{Krolik88}). To reach the required pressure of $p\\approx10^{-2}$~dyne~cm$^{-2}$, the inter-cloud medium needs either a high temperature \\citep[$>4\\times10^7$~K, ][]{Krolik88} or a strong magnetic field \\citep{Rees87}. Apart from the confinement issue, the clouds should be hydrodynamically unstable: \\citet{Mat82} has shown that while optically thin clouds may be close to uniformly accelerated by the radiation pressure, there remain internal radial pressure imbalances, which especially in the optically thick case, which is preferred by photoionisation models, lead to lateral expansion of the clouds. He termed the latter state {\\em pancake} clouds. \\citet{Mat86} then assessed the hydrodynamic stability of such pancake clouds with the result that the lateral edges of a pancake cloud are hydrodynamically unstable. Hence, the lateral flows persist and destroy the cloud on a short timescale, comparable to one cloud orbit. Also, the ram pressure by the inter-cloud gas, which is probably pressure supported and moving in a different way than the clouds, may compress the clouds in the direction of relative motion and the increased pressure makes the cloud\\change{s} expand sideways \\citep{Krolik88}, now with regard to the direction of motion. These problems have lead to the idea that clouds must reform or be re-injected steadily. \\citet{Krolik88} develops the idea of clouds forming by the thermal instability from the hot inter-cloud gas. This idea has however been rejected by \\citet{MD90}: It would require very large amplitude fluctuations of unusual type (high density, low temperature). Additionally, a compression by a factor of 10,000 would require an unusually small magnetic field, for the magnetic pressure not to stop the collapse. \\citet{MD90} further argue that the clouds would be destroyed by the radiative shear mechanism before they could contribute to the emission line profiles. We have studied the radiative shearing mechanism for dusty clouds in 2.5D hydrodynamic simulations \\citep{SKB11}. In agreement with \\citet{Mat86} and \\citet{MD90}, we find quick radiative shearing of the cloud. \\change{As mentioned above, the stability considerations were for clouds accelerating due to the dominant radiation pressure. The problem is even more severe for bound clouds where gravity is comparable to the radiation force: The radiation force acts on the illuminated surface of the cloud, while every part of the cloud is uniformly subject to gravity which must dominate by definition for bound clouds. Such clouds are therefore compressed, and must consequently be stabilised by some internal pressure. Yet, radiation pressure and gravity act in one dimension, whereas the thermal pressure is isotropic. It is therefore not possible to stabilise illuminated bound clouds by thermal pressure alone. } In this context, we investigate the stability of magnetised and gravitationally bound BLR clouds -- initially close to an equilibrium orbit as calculated in paper~I. \\change{ Magnetic fields have so far rarely been considered in BLR clouds. We therefore first review literature data for indications on the relative magnitudes of gravitational, magnetic and radiative forces.} (section~\\ref{mclds}). \\change{We then focus on the effect of the magnetic field}, and simulate the evolution of irradiated magneti\\change{cally dominated bound} clouds. We study the two-dimensional, axisymmetric (2.5D), magnetohydrodynamic (MHD) evolution of isolated, initially optically thick clouds, including gravity, rotation and radial radiation pressure via a simple equilibrium photoionisation ansatz. We neglect self-gravity, viscosity and any radiation source other than the central accretion disc. Because of the small size of the clouds, our computational domain is small compared to the full size of the BLR. We describe technical details in section~\\ref{num}, the setup details in section~\\ref{setup}, our results in section~\\ref{res} and discuss our findings in section~\\ref{disc}. We summarise our results in section~\\ref{conc}. ", "conclusions": "\\label{conc} \\change{ Gravitationally bound clouds facing strong radiation pressure are unstable in the purely hydrodynamic case because radiative and gravitational forces compress the clouds radially, whereas the thermal pressure acts isotropic. They may be more stable if they are significantly magnetised. In particular, we find in axisymmetric magnetohydrodynamic simulations with a prescription for the radiation pressure that the magnetic tension force produces more stable clouds, while in a situation where the geometry of the magnetic field is such that its effect is only analogous to an additional pressure, the clouds are similarly unstable as in the hydrodynamic case. In order to be effective, in the BLRs of a discussed sample of AGN, the magnetic field strength should be \\rvtwo{of order a few Gauss, accurate to about an order of magnitude and} constrained by the condition that magnetic, radiative and gravitational forces should be comparable.}" }, "1207/1207.2439_arXiv.txt": { "abstract": "We used the Karl G.~Jansky Very Large Array (VLA) to image one primary beam area at 3~GHz with $8''$ FWHM resolution and $ 1.0 \\,\\mu{\\rm Jy~beam}^{-1}$ rms noise near the pointing center. The $P(D)$ distribution from the central 10 arcmin of this confusion-limited image constrains the count of discrete sources in the $1 < S(\\mu{\\rm Jy}) < 10$ range. At this level the brightness-weighted differential count $S^2 n(S)$ is converging rapidly, as predicted by evolutionary models in which the faintest radio sources are star-forming galaxies; and $\\approx 96$\\% of the background originating in galaxies has been resolved into discrete sources. About 63\\% of the radio background is produced by AGNs, and the remaining 37\\% comes from star-forming galaxies that obey the far-infrared (FIR) / radio correlation and account for most of the FIR background at $\\lambda \\approx 160\\,\\mu$m. Our new data confirm that radio sources powered by AGNs and star formation evolve at about the same rate, a result consistent with AGN feedback and the rough The confusion at centimeter wavelengths is low enough that neither the planned SKA nor its pathfinder ASKAP EMU survey should be confusion limited, and the ultimate source detection limit imposed by ``natural'' confusion is $ \\leq 0.01\\,\\mu$Jy at $\\nu = 1.4$~GHz. If discrete sources dominate the bright extragalactic background reported by ARCADE\\,2 at 3.3~GHz, they cannot be located in or near galaxies and most are $\\leq 0.03\\,\\mu$Jy at 1.4~GHz. ", "introduction": "\\subsection{Counts of Extragalactic Radio Sources} Together the number per steradian $n(S)dS$ of discrete extragalactic radio sources having flux densities $S$ to $S + dS$ and the brightness temperature $T_{\\rm b}$ that sources contribute to the sky background constrain the nature and evolution of all extragalactic radio sources, even those too faint to be counted individually. Counts of radio sources achieved instant fame and notoriety with the announcement by \\citet{ryl55a} and \\citet{ryl55b} that the source count from the 2C survey was dramatically steeper than that expected for a uniformly-filled Euclidean universe. The remarkable \\citet{ryl55b} paper claimed that the majority of ``radio stars'' were extragalactic, that they were so distant and radio-luminous as to be mostly beyond the reach of then-existing optical telescopes, and that the steep count showed dramatic cosmic evolution for individual objects in space density or luminosity. Then revolutionary, it remains a succinct summary of our current understanding of extragalactic radio sources. The claimed evolution sparked a bitter and public dispute, refuting as it did the popular steady-state cosmology \\citep{bon48,hoy48}. Also the 2C radio source count was made before radio astronomers appreciated the importance of confusion and Eddington bias. \\citet{mil57} showed the great majority of the 2C sources were just ``confusion bumps,'' the broad 2C beam often blending two or more faint background sources to appear as a single stronger source. \\citet{mil57} also noted that the count from their higher-resolution survey showed ``no cosmological effects'' with the possible exception of mild clustering. By that time, however, the Cambridge group had fully understood what went wrong, and the aftermath triggered (1) the 3C and 4C Cambridge surveys which were carefully cognizant of confusion, (2) the pioneering $P(D)$ analysis of confusion by \\citet{sch57} which showed how to extract the true source count from a confusion-limited survey, and (3) the first aperture-synthesis images \\citep{ryl62} which revealed a decreased source-count slope (loosely termed ``turnover'') at flux densities below 1 Jy. By 1966 there was general agreement about the counts of sources stronger than $S \\sim 0.1$~Jy at low frequencies. The 4C confusion $P(D)$ \\citep{hew61} and direct \\citep{gow66} counts had mapped their main features: a rise steeper than the static Euclidean slope followed by a decrease to sub-Euclidean slopes at flux densities below 1 Jy. Independent sky surveys, such as the Parkes survey \\citep{bol64}, were in agreement. To explain the source count, several investigators developed mathematical (non-physical) models for the cosmic evolution of extragalactic radio sources. \\citet{lon66} first showed that the relatively rapid change of the count slope could be explained by differential cosmic evolution, in which the more luminous sources underwent more evolution in their comoving space density than their less-luminous counterparts did. Sensitive surveys made with the WSRT and VLA in the 1980s extended the 1.4~GHz source count down to mJy levels and below. \\citet{con84c} and \\citet{win85} discovered a point of inflection near 1 mJy below which the count slope flattens, suggesting the emergence of a new population of radio sources. Spiral galaxies dominate the local radio luminosity function below $L \\sim 10^{23}{\\rm ~W~Hz}^{-1}$. Radio sources powered by star formation in spiral galaxies can account for this new population and for nearly half of the extragalactic sky background if they evolve at about the same rate as the stronger sources powered by active galactic nuclei (AGNs) in elliptical galaxies \\citep{con84a}. Later models \\citep{wal97b,boy98,wil08} find evolution similar to that of the cosmic star-formation rate shown in the Lilly-Madau diagram [see \\citet{hop07} for a modern data-set]. \\citet{mas10} accounted for the inflection with an evolutionary model using the empirical FIR/radio correaltion, FIR counts, and the redshift distributions of star-forming galaxies . A recent compilation of source counts from surveys at frequaencies 150 MHz to 15 GHz is presented by \\citet{dez10}; features and cosmic implications are discussed there in detail. Of particular note in the present context are persistent discrepancies among the most sensitive 1.4~GHz direct counts that far exceed both Poisson and clustering uncertainties \\citep{con07,owe08}. These discrepancies may result from different authors making different corrections for the effects of partial resolution. Resolution corrections can be large and difficult to estimate for an image whose resolution approaches the median angular size $\\langle \\theta_{\\rm s} \\rangle \\sim 1''$ of faint sources. The deep synthesis counts were also extended by confusion $P(D)$ distributions in low-resolution images \\citep{mit85,win93}. We used their approach here to count sources a factor of ten fainter in a low-resolution ($8''$ FWHM) 3~GHz image that does not need significant corrections for partial resolution. \\subsection{The Source Contribution Sky Brightness} The radio source count and the source contribution to the sky brightness at frequency $\\nu$ are connected by the Rayleigh-Jeans approximation \\begin{equation}\\label{dseq} S n(S) dS = {2 k_{\\rm B} d T_{\\rm b} \\nu^2 \\over c^2}~, \\end{equation} where $k_{\\rm B}\\approx 1.38 \\times 10^{-23} {\\rm ~J~K}^{-1}$ is the Boltzmann constant and $dT_{\\rm b}$ is the sky brightness temperature added by the $n(S)$ sources per steradian with flux densities $S$ to $S + dS$. The observed source count now spans eight decades of flux density, so it must be plotted with a logarithmic abscissa. Substituting $d[\\ln(S)] = dS / S$ gives the brightness temperature contribution per decade of flux density \\begin{equation}\\label{dtbeq} \\biggl[{d T_{\\rm b} \\over d \\log(S)}\\biggr] = \\biggl[{\\ln(10) c^2 \\over 2 k_{\\rm B} \\nu^2}\\biggr] S^2 n(S)~. \\end{equation} The cumulative brightness temperature from all sources stronger than any detection limit $S_0$ is \\begin{equation}\\label{tbcumeq} \\Delta T_{\\rm b} (>S_0) = \\biggl[{\\ln(10) c^2 \\over 2 k_{\\rm B}\\nu^2}\\biggr] \\int_{S_0}^\\infty S^2 n(S) d[\\log(S)]~. \\end{equation} If $S_0$ is low enough, $\\Delta T_{\\rm b}$ approaches the total $T_{\\rm b}$ of the source background, and we can say that the background has been resolved into known sources. For example, \\citet{gle10} resolved about half of the far-infrared (FIR) background at $\\lambda = 250,~350$, and $500\\,\\mu{\\rm m}$; and the Herschel/PEP (PACS Evolutionary Probe) survey \\citep{ber11} resolved about 3/4 of the COBE/DIRBE background at $\\lambda = 160\\,\\mu{\\rm m}$. Figure~\\ref{oldtbkndfig} shows the flux-density range covered by published 1.4~GHz source counts. We plotted the weighted count $\\log[S^2 n(S)]$ instead of the traditional Euclidean-weighted count $\\log[S^{5/2}n(S)]$ as a function of $\\log(S)$ because (1) $S^2 n(S)$ is proportional to the source contribution per decade of flux density to the sky temperature (Eq.~\\ref{dtbeq}), (2) the radio universe is neither static nor Euclidean, and (3) the plotted slopes are minimized for easy visual recognition of broad features, unlike the steeply sloped plot of $\\log[n(S)]$ for example. On a plot of $\\log[S^2n(S)]$, the source count must ultimately fall off at both ends to avoid Olbers' paradox. The filled points at $\\log[S\\,{\\rm (Jy)}] > -3$ are from the many surveys referenced in \\citet{con84a}, and the filled data points at lower flux densities are from the \\citet{mit85} confusion-limited VLA survey. The irregular box encloses the range of 1.4~GHz counts consistent with the probablility distribution $P(D)$ of confusion amplitudes $D$ (in brightness units $\\mu$Jy per beam solid angle) in the low-resolution ($17\\,\\farcs5$ FWHM) \\citet{mit85} image. The straight line inside the box indicates the best power-law fit to the $P(D)$ data. The open data points and their power-law fit (upper straight line) indicate the direct count of individual sources from the most sensitive high-resolution 1.4 GHz survey ever made with the original VLA \\citep{owe08}. These two faint-source counts disagree by several times the Poisson counting errors, and nearly all other published faint-source counts are scattered over the wide gap between them \\citep{con07}, as illustrated by Figure~11 in \\citet{owe08}. The sources contributing to the data plotted in Figure~\\ref{oldtbkndfig} show no sign of resolving the extragalactic background at 1.4~GHz because the power-law fits to $S^2 n(S)$ are still rising near the flux density limit $S_0 \\sim 10\\,\\mu$Jy. The solid curve in Figure~\\ref{oldtbkndfig} shows the \\citet{con84b} evolutionary model for the total source count at 1.4\\,GHz. In this model, the stronger radio sources are powered primarily by AGNs of elliptical galaxies (dashed curve) and most of the fainter sources are in star-forming spiral galaxies (dotted curve). The model predicts that (1) $S^2 n(S)$ should converge below $S \\sim 10\\,\\mu$Jy and (2) most of the model's $T_{\\rm b} \\approx 100$\\,mK background should be resolved into sources stronger than $S \\sim 1\\,\\mu$Jy. However, \\citet{fix11} reported that the recent ARCADE\\,2 (Absolute Radiometer for Cosmology, Astrophysics, and Diffuse Emission) balloon measurement of absolute sky brightness leaves a remarkably high extragalactic $T_{\\rm b} = 54 \\pm 6$\\,mK at $\\nu = 3.3$\\,GHz after the Galactic foreground and the $T = 2.731 \\pm 0.004$\\,K cosmic microwave background (CMB) were subtracted. Combining their 3.3~GHz result with low-frequency data from the literature led \\citet{fix11} to fit the ``excess'' extragalactic background with a power-law spectrum \\begin{equation}\\label{arcade2tbeq} T_{\\rm b} = (24.1 \\pm 2.1 {\\rm ~K}) \\times \\biggl( {\\nu \\over 310{\\rm ~MHz}}\\biggr)^{-2.599 \\pm 0.036} \\end{equation} between 22\\,MHz and 10\\,GHz. Equation~\\ref{arcade2tbeq} gives $T_{\\rm b} \\approx 480$\\,mK at 1.4\\,GHz, almost five times the $T_{\\rm b} \\approx 100$\\,mK predicted for the known populations of extragalactic sources. With the goals of (1) determining an accurate source count down to $S \\sim 1 \\,\\mu$Jy, (2) resolving the radio sky background contributed by galaxies, and (3) constraining possible source populations that could produce the high ARCADE\\,2 background at 3.3\\,GHz, we used the Karl G.~Jansky Very Large Array (VLA) to make a very sensitive (rms noise $\\sigma_{\\rm n} \\approx 1\\,\\mu$Jy\\,beam$^{-1}$) low-resolution ($8''$ FWHM) confusion-limited sky image covering one primary beam area at S band (2--4\\,GHz). We pointed at the center of the ``crowded'' \\citet{owe08} field to see if we could confirm the reported high source count there. Our relatively low angular resolution ensures that galaxy-size sources at cosmological distances are unresolved: $8''$ spans at least 40 kpc throughout the redshift range $0.4 < z < 7$ containing most faint radio sources. That low resolution also guarantees that the rms confusion is at least comparable with the rms noise, a necessary condition for using confusion to constrain the sky density of sources as faint as $S_0 \\sim 1\\,\\mu$Jy, which is more than a factor of five below our detection limit for individual sources. \\subsection{Outline} Section~\\ref{obssec} of this paper describes our observations and the production of a confusion-limited sky image. Section~\\ref{pofdsec} presents the 3.02~GHz $P(D)$ distribution from the central region of this image. The 1.4 and 3.02 GHz source counts consistent with our $P(D)$ distribution are derived in Section~\\ref{countsec}, and the contributions of known extragalactic source populations to the sky brightness are discussed in Section~\\ref{tbsec}. In Section~\\ref{arcadesec} we use our narrow $P(D)$ distribution to show that the high ARCADE\\,2 background is too smooth to be produced by, or even spatially associated with, galaxies brighter than $m_{\\rm AB} = +29$. Section~\\ref{summarysec} summarizes our results. ", "conclusions": "\\label{summarysec} We used 21 antennas in the Karl G.~Jansky VLA to observe a single field at S band (2--4~GHz) with a FWHM resolution $\\theta = 8''$ and reached an rms noise $\\sigma_{\\rm n} \\approx 1\\,\\mu{\\rm Jy~beam}^{-1}$ near the image center after 50 hours of integration time. The image is confusion limited with an ``rms'' confusion level $\\sigma_{\\rm c}^* \\approx 1.2\\,\\mu{\\rm Jy~beam}^{-1}$ at $\\nu = 3.02$~GHz. The 3.02~GHz differential source count was derived from the confusion $P(D)$ distribution. For comparison with published source counts, we converted it to 1.4~GHz via the effective spectral index $\\langle \\alpha \\rangle \\approx -0.7$. The power-law approximation $n(S) \\propto S^{-\\gamma}$ to the 1.4~GHz source count has a slope approaching $\\gamma \\approx 1.5$ near $1~\\mu$Jy, significantly lower than the slopes of published counts above $10~\\mu$Jy. If the faintest radio sources have median angular size $\\langle \\theta_{\\rm s} \\rangle \\leq 1''$ as expected for emission coextensive with star-forming regions in distant galaxies, the natural confusion limit for source detection is not more than $5\\sigma_{\\rm c}^* \\approx 0.01\\,\\mu$Jy at 1.4~GHz, and the continuum sensitivity of the planned SKA will not be limited by natural confusion. The observed count is well fit by evolutionary models in which the local radio luminosity functions of all sources associated with both AGNs and star formation evolve at the same rate. This is broadly consistent with the correlation of black hole and stellar bulge masses in massive elliptical galaxies, although this correlation is not yet well established for the lower-mass black holes and bulges in late-type galaxies \\citep{gre10}. The brightness-weighted count $S^2 n(S)$ is clearly converging below $10\\,\\mu$Jy. Our image has resolved about 96\\% of the radio background produced by all galaxies ($T_{\\rm b} \\approx 100$~mK at 1.4~GHz and $T_{\\rm b} \\approx 13$~mK at 3~GHz). Nearly 100\\% of the $\\approx 63$~mK AGN-powered background at 1.4~GHz has been resolved. The remaining $\\approx 37$~mK comes from star-forming galaxies that obey the FIR/radio correlation and account for most of the extragalactic background at $\\lambda = 160\\,\\mu$m. We resolved about 89\\% of the star-forming galaxy contribution. The ARCADE\\,2 balloon experiment indicated a nonthermal excess brightness over the Galaxy, the CMB, and that expected from known populations of radio sources in galaxies. At 3.02 GHz this excess brightness temperature is $52 \\pm 8$~mK. Our narrow 3.02~GHz $P(D)$ distribution implies that the excess background must be very smooth. Any new discrete-source population able to produce such a bright and smooth background is far too numerous to be associated with galaxies brighter than $m_{\\rm AB} = +29$. {\\it Facilities:} \\facility{VLA} \\appendix" }, "1207/1207.1506_arXiv.txt": { "abstract": "{The sets of the synchronous equations are derived from the sets of non-synchronous equations The analytical solutions are given by solving the set of differential equations. The results of the evolutionary tendency of the orbit-spin are that the semi-major axis shrinks gradually with time: the orbital eccentricity dereacses gradually with time until the orbital circularization; the orbital period shortens gradually with time and the rotational angular velocity of primary component speed up with time gradually before the orbit-rotation achieved the circularization The theoretical results are applied to evolution of the orbit and spin of synchronous binary stars Algol A, B on the main sequence phase The circularization time and life time (age) and the evolutional numerical solutions of orbit and spin when circularization time are estimeted for Algol A, B. The results are discussed and concluded. ", "introduction": "The tidal friction plays an important role in the evolution of the orbit and spin of close binary system. Earliest, the author researching this topic is \\cite{Zahn+1965, Zahn+1966a, Zahn+1966b, Zahn+1966c, Zahn+1975}. \\cite{Alexander+1973} firstly studied the dynamical problem of the tidal friction in close binary system by using the method employed by \\cite{Darwin+1879}. Later on, \\cite{Hut+1980, Hut+1981} generalized the method given by \\cite{Alexander+1973}. He studied the stability of tidal equilibrium and tidal evolution in close binary system by using the method of energy and angular momentum. But their research dealt with a few synchronization. The sequential research for the synchronization of rotation are given by \\cite{Zahn+1977, Zahn+1978}. \\cite{Rajamohan+Venkatakrishnan+1981} ever studied synchronization in binary stars. \\cite{Giuricin+etal+1984a} researched synchronization in eclipsing binary stars and \\cite{Giuricin+etal+1984b} also researched synchronization in early-type spectroscopic binary stars, \\cite{Zahn+Bouchet+1989b} studied mainly the orbital circularization of late-type binary stars on the pre-main sequence and the theoretical results are given based on \\cite{Zahn+1989a}. \\cite{Pan+1996} calculated the circularization time scale by using the two mechanisms: one is the equilibrium tidal mechanism given by \\cite{Zahn+1977}, another is purely hydrodynamic mechanissm given by Tassoul~(\\citeyear{Tassoul+1987}). \\cite{Keppens+etal+2000} studied the rotational evolution of binary stars system: synchronization and circularization. \\cite{Huang+Zeng+2000} also researched evolution of non-synchronized binary stars with 9 solar mass and 6 solar mass. \\cite{Meibom+etal+2005} studied obseration tidal synchronization in detached solar-type binary stars and \\cite{Meibom+etal+2006} also researched an observational study of tidal synchronization in solar--type binary starss in open clusters M35 and M34. Although the author \\cite{Li+1998, Li+2004, Li+2009} studied some methods for judging the synchronization of rotation of binary stars, but he does not studied the evolution of orbit-rotation of synchronous binary stars. In the present paper the author examined the evolutional tendency of orbit and spin of synchronous binary stars on the main sequence phase. ", "conclusions": " \\begin{itemize} \\item[(1)] The eccentricity decreases gradually with time until the orbital circularization, i.e until $e$ decreases to the circularization time. \\item[(2)] The semi-major axis shrinks gradually with time or with eccentricity decreases. \\item[(3)] The orbital period shortens gradually with time or with circularization. \\item[(4)] The rotational angular velocity of primary component speeds up with time gradually. \\end{itemize}" }, "1207/1207.3735_arXiv.txt": { "abstract": "We perform relativistic hydrodynamic simulations of the formation and evolution of AGN cocoons produced by very light powerful jets. We calculate the intensity maps of the Sunyaev--Zel'dovich (SZ) effect at high frequencies for the simulated AGN cocoons using the relativistically correct Wright formalism. Our fully relativistic calculations {demonstrate} that the contribution from the high temperature gas ($k_{\\rm{b}} T_{\\rm{e}}\\simeq 100$ keV) to the SZ signal from AGN cocoons at high frequencies is stronger than that from the shocked ambient intercluster medium {owing to the fact that the relativistic spectral functions peak at these temperature values}. We present simulations of the SZ effect from AGN cocoons at various frequencies, and demonstrate that SZ observations at 217 GHz and at higher frequencies, such as 857 GHz, will provide us with knowledge about the dynamically-dominant component of AGN cocoons. ", "introduction": "The standard evolutionary scenario for AGNs suggests that their jets are not in direct contact with the intergalactic medium (IGM), but rather are enveloped in a cocoon. The cocoon around a pair of supersonic, low-density (when compared to the ambient IGM) jets acts as a `wastebasket' for most of the energy deposited by the jets (Scheuer 1974). {A number of cavities in X-ray surface brightness maps were detected by the} {\\it{Chandra}} {X-ray observatory in clusters of galaxies (for a review, see McNamara \\& Nulsen 2007).} {So far the largest} cavities in X-ray surface brightness maps {have been revealed} in the MS 0735+7421 and Hercules A clusters of galaxies (McNamara et al. 2005; Nulsen et al. 2005). These X-ray cavities extend over several hundreds of kpc and are formed by powerful AGN jet activity. The presence of X-ray cavities can be caused by very hot gas embedded in the AGN cocoons. Low-density gas with high temperatures of $k_{\\mathrm{b}} T_{\\mathrm{e}}\\simeq 100$ keV does not significantly emit in the soft X-ray band and, therefore, high temperature regions can be associated with X--ray cavities. Numerical simulations of AGN cocoons in galaxy clusters show that electron temperatures of the order of 100 keV are expected for the plasma in AGN cocoons (see, e.g., Sternberg \\& Soker 2009; Perucho et al. 2011). Inverse Compton (IC) scattering of cosmic microwave background (CMB) photons by free thermal electrons located in an AGN cocoon causes a change of the CMB spectrum (see Pfrommer et al. 2005 and Prokhorov et al. 2010). In this paper, we identify this change of the CMB spectrum with the Sunyaev--Zel'dovich (SZ) effect {(Zel'dovich \\& Sunyaev 1969; for a review, see Birkinshaw 1999)} from AGN cocoons. The advent of new telescopes, such as {\\it{Herschel}}, {\\it{GBT}}, and {\\it{ALMA}} should allow us to measure the CMB distortion towards AGN cocoons at various frequencies with high sensitivity and with high angular resolution (see, Zemcov et al. 2010; Korngut et al. 2011; Yamada et al. 2012, for the first observations of the SZ effect from galaxy clusters by {\\it{Herschel-SPIRE}}, {\\it{GBT-MUSTANG}}, and the recent simulations of SZ maps of {\\it{ALMA}}, respectively). As shown by Prokhorov et al. (2010), the CMB intensity change at a frequency of 217 GHz, which approximately corresponds to the crossover frequency of the SZ effect from a low temperature gas with $k_{\\mathrm{b}} T_{\\mathrm{e}} < 5$ keV (see Sunyaev \\& Zel'dovich 1980), is maximal at the temperature of $\\simeq$100 keV (see Colafrancesco 2005 for a review of the SZ effect from AGN cocoons in non-thermal electron models). The frequency shift of CMB photons owing to IC scattering increases with electron temperature. Therefore, we assume that the CMB distortion caused by IC scattering by a high temperature plasma with $k_{\\mathrm{b}} T_{\\mathrm{e}}\\simeq 100$ keV should also be significant at high frequencies. So far, the SZ effect from AGN cocoons in galaxy clusters has been calculated by using {analytical toy models} for gas pressure and temperature distributions in cocoons (see Colafrancesco 2005; Pfrommer et al. 2005) because of lack of available hydrodynamic simulations of AGN cocoons. {These toy models do not take into account intrinsic properties of the jet--ICM system and, therefore, can lead to oversimplification. To test consistency and feasibility of these toy models, it is necessary to calculate the SZ effect from AGN cocoons which are derived from relativistic hydrodynamic simulations of the jet--ICM system and to confront the results obtained from the analytical toy models with those obtained from more realistic numerical models.} Using the \\textsc{PLUTO} code (Mignone et al. 2007), we performed relativistic hydrodynamic simulations of the formation and evolution of AGN cocoons produced by very light {(i.e. under-dense compared with the ICM)} powerful jets. Our numerical simulations take into account different properties of AGN jets and ICM, such as the power of the jets, {the jet-to-ICM density contrast}, velocity, gas density and temperature distributions, and a dark matter halo profile, all providing us with a more realistic AGN cocoon model. {The performed simulations demonstrate that the internal structure of the region around AGN jets is more complex than that assumed in analytical toy models.} We calculate the CMB spectrum distortion caused by IC scattering of the CMB photons by highly energetic electrons in the framework of the relativistically correct Wright formalism and study the SZ effect from simulated AGN cocoons at various frequencies. In this paper, we show that more realistic plasma models obtained from numerical hydrodynamic simulations {predict: (1) that the morphology of SZ intensity maps changes with frequency and (2) that SZ signals at high frequencies from AGN cocoons are higher than those from the shocked ambient intracluster medium and, therefore, the SZ morphological study can reveal the presence of AGN cocoons on SZ maps}. The fully relativistic approach for calculating the SZ effect from AGN cocoons considered below shows that the cocoon is strongly inhomogeneous and that future measurements of the SZ effect will permit one to study the internal structure of regions formed by AGN activity. {We also show that the tight constraints on the total thermal energy of high temperature electrons which are stored in AGN cocoons can be obtained by taking into account the shape of relativistic SZ spectral functions.} The layout of the paper is as follows. We describe the relativistic hydrodynamic simulations in Sect. 2. We calculate the SZ effect from the simulated AGN cocoons in the relativistically correct Wright formalism {and study the possibility to detect AGN cocoons through SZ observations} in Sect. 3. Our conclusions are presented in Sect. 4. ", "conclusions": "Giant X-ray cavities have been discovered by \\textit{Chandra} in clusters of galaxies, such as MS 0735+7421 at a redshift z=0.22 (McNamara et al. 2005) and Hercules A at z=0.154 (Nulsen et al. 2005). These cavities could be created by powerful AGN outflows with jet kinetic power of $\\gtrsim 10^{46}$ erg s$^{-1}$. Similar X-ray surface brightness depression regions have been observed in the galaxy cluster Abell 2204 at a redshift of 0.152 and the jet's kinetic power of $\\simeq 5\\times10^{46}$ erg s$^{-1}$ should be necessary to create such regions (Sanders et al. 2009). Another example is the intermediate redshift (z=0.29) cluster gas associated with the FR II radio galaxy 3C 438 (see Kraft et al. 2007), the observed surface brightness discontinuity in the gas that extends $\\simeq$600 kpc can be the result of an extremely powerful outburst which is even more powerful than those seen in the nearby clusters MS 0735+7421, Hydra A, and Hercules A. The standard evolutionary scenario for AGNs suggests that their jets are enveloped in a cocoon (Scheuer 1974; Blandford \\& Rees 1974). Modern hydrodynamical simulations led to the conclusion that the gas density is very low and gas temperature is very high, $10^{9}-10^{10}$ K, in AGN cocoons. Therefore, a significant X-ray depression is predicted towards AGN cocoons. X-ray cavities often coincide with the radio lobes of the central radio galaxy, although the non-thermal pressure derived from the equipartition condition for the energy of synchrotron-radiating non-thermal electrons and magnetic fields is a factor of ten smaller than the pressures required to inflate the cavities (e.g. Ito et al. 2008). This implies that most of the energy in the cocoon is carried by an `invisible' component such as, e.g., high energy thermal electrons (Ito et al. 2008). Observations of the SZ effect have been proposed to probe the inferred dynamically-dominant component of plasma bubbles associated with X-ray cavities (see Pfrommer et al. 2005; Colafrancesco 2005; Prokhorov et al. 2010), since the SZ signal is proportional to the electron pressure integrated along the line-of-sight and to the value of the spectral function that depends on electron temperature. In this paper, we present RHD simulations of the formation and evolution of AGN cocoons produced by very light powerful jets. Our simulations have been performed {with} the publicly available \\textsc{PLUTO} computational code. The parameters of the simulations are listed in Table 1. We have constrained our initial conditions based on the X-ray observations of the powerful AGN outburst in the MS 0735+7421 cluster of galaxies. The goal of our simulations is to provide a realistic model of gas pressure and temperature distributions in AGN cocoons. The simulated electron pressure and temperature maps are shown in Figs. 1 and 2 for the two models with different jet lifetimes or `duty cycles'. We have found that the derived temperatures in AGN cocoons, $\\simeq 10^9-10^{10}$ K, are significantly higher than those are in the regions of shocked ambient gas and, therefore, the use of the relativistically correct SZ formalism is necessary to produce the simulated SZ maps of AGN cocoons. We have presented the first SZ intensity maps of AGN cocoons at various frequencies derived from the relativistic hydrodynamical simulations and calculated in the framework of the relativistically correct SZ formalism. This fully relativistic study provides us with more realistic SZ maps of AGN cocoons, more so than those previously obtained by adopting {analytical toy models} of gas pressure and temperature distributions in cocoons. We show that SZ observations at frequencies of 217 GHz and at 857 GHz will provide us with a method to test the presence of very hot gas in AGN cocoons. This result confirms the previous conclusions based on the {analytical toy models} that SZ maps at 217 GHz could be used to probe the dynamically-dominant component of AGN cocoons. Moreover, our study demonstrates that two different regions disturbed by AGN activity are present in the simulated domain, namely an AGN cocoon region and a shocked ambient medium region. We show that by taking into account the presence of these two regions leads us to the conclusion that the AGN cocoon is not clearly seen on the simulated SZ map at a frequency of 90 GHz. Therefore, SZ observations at 217 GHz and at higher frequencies, such as 857 GHz, are more suitable for studying AGN cocoons than those are at lower frequencies. We conclude that fully relativistic simulations of the SZ effect from AGN cocoons are very important, since the spectral properties of the SZ signal should be taken into account in order to produce realistic SZ maps of cocoons. {We have considered the possibility of observing the SZ effect from AGN cocoons by means of modern instruments, such as} {\\it Planck-HFI}, {\\it Herschel-SPIRE}, and {\\it ALMA} {and have derived upper limits on the total thermal energy of high temperature electrons stored in an AGN cocoon that can be obtained with these instruments} (see Table \\ref{tab2}). {We have compared the derived upper limits with the thermal energy stored in the simulated cocoon (Model A), and have found that the upper limits (that can be derived with} {\\it Herschel-SPIRE} and {\\it ALMA}) {on the thermal energy of 100-400 keV electrons are close to the value of the thermal energy of electrons in the AGN cocoon obtained from our simulations. Therefore, AGN cocoons are a suitable target for observations with modern high-resolution SZ instruments.}" }, "1207/1207.4862_arXiv.txt": { "abstract": "On August 9, 2011, there was an X6.9 flare event occurred near the west limb of solar disk. From the observation obtained by the spectrometer of the Chinese Solar Broadband Radio Spectrometer in Huairou (SBRS/Huairou) around the flare, we find that this powerful flare has only a short-duration microwave burst of about only 5 minutes, and during the short-duration microwave burst, there are several kinds of fine structures on the spectrogram. These fine structures include very short-period pulsations, millisecond spike bursts, and type III bursts. The most interesting is that almost all of the pulses of very short-period pulsation (VSP) are structured by clusters of millisecond timescales of spike bursts or type III bursts. And there exists three different kinds of frequency drift rates in the VSP: the frequency drift rates with absolute value of about 55 - 130 MHz s$^{-1}$) in the pulse groups, the frequency drift rates with absolute value of about 2.91 - 16.9 GHz s$^{-1}$) on each individual pulse, and the frequency drift rates with absolute value of about 15 - 25 GHz s$^{-1}$) at each individual spike burst or type III burst. ", "introduction": "Microwave bursts associated with solar flares offer a number of unique diagnostic tools to address long-standing questions about energy release, plasma heating, particle acceleration and propagation in magnetized plasmas (Rosenberg 1970, Aschwanden 1987, Bastian, Benz, \\& Gary, 1998). On 2011 August 9, a most powerful X6.9 solar flare took place in active region NOAA 11263, near the west limb on the solar disk. The X6.9 flare event starts at 08:00 UT, reaches to the maximum at 08:04 UT, and ends at 08:14 UT. It is the largest one in the current solar Schwabe cycle, resulting in a coronal mass ejection (CME). Accompanied with this flare, an extremely powerful microwave burst was observed at a frequency of 2.60 - 3.80 GHz by the Chinese Solar Broadband Radio Spectrometer in Huairou (SBRS/Huairou) (Fu et al 1995, Fu et al, 2004). In this work, we will introduce the main features of the microwave bursts associated with the flare, especially the superfine structures with millisecond timescale. Section 2 is the Observations and Analysis, and the conclusions and some discussions are presented in section 3. ", "conclusions": " (1) Dislike the previous observations of X-class flares which always have long-duration microwave bursts, the above X6.9 flare event has only a short-duration microwave burst of about only 5 minutes. (2) During the short-duration microwave burst, there are several kinds of fine structures on the spectrogram. These fine structures include very short-period pulsations, millisecond spike bursts, and type III bursts. The most interesting is that almost all of the pulses of very short-period pulsation are structured by clusters of millisecond timescales of spike bursts or type III bursts. (3) There exists three different kinds of frequency drift rates in the VSP: the frequency drift rates with absolute value of about 55 - 130 MHz s$^{-1}$) in the pulse groups, the frequency drift rates with absolute value of about 2.91 - 16.9 GHz s$^{-1}$) on each individual pulse, and the frequency drift rates with absolute value of about 15 - 25 GHz s$^{-1}$) at each individual spike burst or type III burst. The flare-associated microwave QPP can provide information on solar flaring regions and give the possible insight into coronal plasma dynamics, such as to remote the microphysics of primary energy releasing regions (Nakariakov \\& Milnikov 2009). The different magnitude of frequency drift rate may reflect the different kinematics of the microwave emission source regions (Kliem et al 2000). ~ \\textbf" }, "1207/1207.0670_arXiv.txt": { "abstract": "Ho{\\v r}ava--Lifshitz gravity models contain higher order operators suppressed by a characteristic scale, which is required to be parametrically smaller than the Planck scale. We show that recomputed synchrotron radiation constraints from the Crab nebula suffice to exclude the possibility that this scale is of the same order of magnitude as the Lorentz breaking scale in the matter sector. This highlights the need for a mechanism that suppresses the percolation of Lorentz violation in the matter sector and is effective for higher order operators as well. ", "introduction": " ", "conclusions": "" }, "1207/1207.5325_arXiv.txt": { "abstract": "The form of depleted sulphur in dense clouds is still unknown. Until now, only two molecules, OCS and SO$_2$, have been detected in interstellar ices but cannot account for the elemental abundance of sulphur observed in diffuse medium. Chemical models suggest that solid H$_2$S is the main form of sulphur in denser sources but observational constraints exist that infirm this hypothesis. We have used the Nautilus gas-grain code in which new chemical reactions have been added, based on recent experiments of H$_2$S ice irradiation with UV photons and high energy protons. In particular, we included the new species S$_n$, H$_2$S$_n$ and C$_2$S. We found that at the low temperature observed in dense clouds, i.e. 10~K, these new molecules are not efficiently produced and our modifications of the network do not change the previous predictions. At slightly higher temperature, 20~K in less dense clouds or in the proximity of protostars, H$_2$S abundance on the surfaces is strongly decreased in favor of the polysulfanes H$_2$S$_3$. Such a result can also be obtained if the diffusion barriers on the grains are less important. In the context of the life cycle of interstellar clouds and the mixing between diffuse and denser parts of the clouds, the depletion of sulphur in the form of polysulfanes or other sulphur polymers, may have occurred in regions where the temperature is slightly higher than the cold inner parts of the clouds. ", "introduction": "Contrary to other elements \\citep[see][]{2009ApJ...700.1299J}, the gas-phase abundance of atomic sulphur, in the diffuse medium, is observed to be constant with cloud density, around its cosmic value of $\\sim 1.5\\times 10^{-5}$ (compared to H) \\citep{1994ApJ...430..650S}. In dense clouds (with densities above $10^4$~cm$^{-3}$), the sum of the abundances of molecules containing sulphur is a small fraction of the cosmic abundance of S \\citep[about $10^{-3}$, see][]{1992IAUS..150..171O,2000ApJ...542..870D}. To reproduce such low S-bearing abundances, one needs to use a very small initial abundance of atomic sulphur of a few $10^{-8}$ \\citep[see for example][]{2008ApJ...680..371W}. This suggests that sulphur depletion occurs very efficiently and rapidly between the diffuse and dense phases of the cloud life leading to an unknown, still unobserved, reservoir of sulphur on grains. When an atom of sulphur sticks to a grain, models predicts that it would be hydrogenated to form H$_2$S \\citep{2007A&A...467.1103G}. H$_2$S has however never been detected in interstellar ices and its abundance has been estimated to be smaller than $5\\times 10^{-8}$ (/H) \\citep{1998ARA&A..36..317V,2011A&A...536A..91J}. Solid OCS has been observed with an abundance of about $5\\times 10^{-8}$ \\citep{1997ApJ...479..839P}. SO$_2$ may also be present with a similar abundance according to \\citet{1997A&A...317..929B} and \\citet{2009ApJ...694..459Z} . Solid H$_2$SO$_4$ has been proposed as a possible reservoir but no detection have confirmed this hypothesis \\citep{2003MNRAS.341..657S}. Finally, iron sulfide FeS has been observed in protoplanetary disks by \\citet{2002Natur.417..148K}. The depletion of sulphur in diffuse clouds does however not follow the one of iron so that the iron left in dense clouds cannot account for such a depletion of sulphur. Laboratory experiments have recently brought some new insights into the sulphur problem. \\citet{garozzo} and \\citet{2011A&A...536A..91J} have independently recently published experimental results of the irradiation of H$_2$S interstellar analogs ices by high energy protons and UV photons respectively. Both studies found that solid H$_2$S was easily destroyed in favor of other species such as OCS, SO$_2$, CS$_2$, H$_2$S$_2$ etc. Based on these results, we have revisited the sulphur depletion problem in dense clouds introducing new mechanisms for the formation of polysulfanes (H$_2$S$_n$), CS$_2$ and sulphur polymers (S$_n$) on the grains. The model, parameters and modifications of the chemical network are described in section~\\ref{model}. The results of our model are presented in section~\\ref{results}. We conclude in the last section. ", "conclusions": "Chemical models that take into account the formation of molecules on interstellar grains predict that sulphur would stick on the grains at high density and would be hydrogenated to form H$_2$S. The non detection of this species in interstellar ices tends to indicate that solid H$_2$S cannot be the reservoir of sulphur. The recent experiments from \\citet{garozzo} and \\citet{2011A&A...536A..91J} showed that H$_2$S is easily dissociated by high energy particles and UV photons, and that other molecules such as CS$_2$, polysulfanes (H$_2$S$_n$) and S$_n$ would be formed on the surfaces. Using a gas-grain model, in which new reactions have been introduced for CS$_2$, H$_2$S$_n$ and S$_n$, we show that we are able to produce large quantities of solid H$_2$S$_3$ and very low abundance of solid H$_2$S but only with temperatures higher than 10~K, which are probably more representative of star forming regions. The key parameter here is the diffusivity of the radical HS on the grains, which is limited at very low temperature. Significant amount of S$_8$ is produced at temperatures between 30 and 50-60~K for densities between $2\\times 10^4$ and $2\\times 10^5$~cm$^{-3}$. The results of our model, as for any chemical model, depends on the chemical reactions considered. More experience on the diffusion of S-bearing species on surfaces at low temperature and the formation of these chains would be a precious ally. The possibility to observe products of the destruction of solid H$_2$S$_n$ or S$_n$ in diffuse or warm medium could bring more stones to this problem. H$_2$S$_2$ is a symmetric molecule and the predicted lines in the millimeter range are weak \\citep{1990JMoSp.141..265B}. Search for these lines, as well as S$_3$ and S$_4$ \\citep{2005JChPh.123e4326T,2005ApJ...619..939G}, in the spectral survey of the low mass protostar IRAS16293-2422 \\citep{2011A&A...532A..23C}, was unsuccessful (E. Caux private communication). The results of this modeling can be put in the context of the larger picture of interstellar dust cycle \\citep{1998ApJ...499..267T} and the fact that the depletion of sulphur into H$_2$S$_n$ may have happened at an early phase of the cloud history when the dust temperature was larger than what is observed at the centre of these cold cores. It is unlikely that simple diffusion mechanisms at the surface of the grains could form pure sulphur polymers (S$_n$). The observation of OCS and SO$_2$ in interstellar ices may indicate however that other mechanisms could take place or that surface reactions are much more efficient than current models assume." }, "1207/1207.4378_arXiv.txt": { "abstract": "A special class of non-trivial topologies of the spherical space ${\\cal S}^3$ is investigated with respect to their cosmic microwave background (CMB) anisotropies. The observed correlations of the anisotropies on the CMB sky possess on large separation angles surprising low amplitudes which might be naturally be explained by models of the Universe having a multiconnected spatial space. We analysed in CQG 29(2012)215005 the CMB properties of prism double-action manifolds that are generated by a binary dihedral group $D^\\star_p$ and a cyclic group $Z_n$ up to a group order of 180. Here we extend the CMB analysis to polyhedral double-action manifolds which are generated by the three binary polyhedral groups ($T^\\star$, $O^\\star$, $I^\\star$) and a cyclic group $Z_n$ up to a group order of 1000. There are 20 such polyhedral double-action manifolds. Some of them turn out to have even lower CMB correlations on large angles than the Poincar\\'e dodecahedron. ", "introduction": "\\label{sec:intro} The $\\Lambda$CDM concordance cosmological model describes nearly all cosmological observations very successfully. Among the few exceptions is the observation of the COBE team \\cite{Hinshaw_et_al_1996} that the fluctuations in the cosmic microwave background (CMB) are nearly uncorrelated on large angular scales $\\vartheta \\gtrsim 60^\\circ$. This surprising result is confirmed by the WMAP team \\cite{Spergel_et_al_2003} and further discussed in \\cite{Aurich_Janzer_Lustig_Steiner_2007,% Copi_Huterer_Schwarz_Starkman_2008,Copi_Huterer_Schwarz_Starkman_2010} with respect to the $\\Lambda$CDM concordance model. In \\cite{Efstathiou_Ma_Hanson_2009} it is argued that there is no significant deviant behaviour from the $\\Lambda$CDM model if the uncertain parts in the CMB map are suitably reconstructed from the less uncertain regions. However, the reconstruction algorithm is analysed by \\cite{Aurich_Lustig_2010,Copi_Huterer_Schwarz_Starkman_2011} showing that this method does not lead to a robust measure of the true CMB sky and the use of masked sky maps is to be preferred. It is concluded in \\cite{Copi_Huterer_Schwarz_Starkman_2011} that the ``lack of large-angle correlation, particularly on the region of the sky outside the Galaxy, remains a matter of serious concern.'' In this paper we try to explain the uncorrelated CMB fluctuations on large scales by relaxing the assumption of the concordance model that the Universe possesses a simply connected spatial topology. Instead, non-trivial topologies are assumed for the spatial 3-manifold, i.\\,e.\\ multiconnected spaces, which can lead to a suppression of CMB correlations on angles corresponding the topological length scale. The simply connected space of the $\\Lambda$CDM concordance model possesses one of the three curvature properties: Euclidean for the ${\\cal E}^3\\equiv\\mathbb{R}^3$, spherical for the ${\\cal S}^3$, or hyperbolic for the ${\\cal H}^3$ depending on the total density $\\Omega_{\\hbox{\\scriptsize tot}}$. These three simply connected spaces are considered as the universal cover which is tessellated by a deck group $\\Gamma$ into cells which are identified. The size of such a cell defines the topological length scale. For an introduction into the topic of cosmic topology, see \\cite{Lachieze-Rey_Luminet_1995,Luminet_Roukema_1999,Levin_2002,% Reboucas_Gomero_2004,Luminet_2008,Fujii_Yoshii_2011}. Below the topological length scale the properties of the concordance model are not altered since the cosmological parameters of the $\\Lambda$CDM concordance model are used, and the local physics is unchanged. For example, possible non-Gaussian features in the CMB are the same as predicted by the $\\Lambda$CDM concordance model \\cite{Monteserin_Barreiro_Sanz_Martinez-Gonzalez_2005}. It is shown in \\cite{Aurich_Janzer_Lustig_Steiner_2010} that the fine structure of the CMB fluctuations for the $\\Lambda$CDM concordance model and for the 3-torus topology cannot be distinguished experimentally due to the same local physics. We investigate the statistical properties of the CMB anisotropies on large separation angles that arise in polyhedral double-action manifolds. These models are not studied in the literature and thus, their CMB properties are unknown. As discussed below, the considered polyhedral double-action manifolds derive from parent manifolds having one of the most severe suppressions of CMB correlations on large scales. This motivates the investigation of polyhedral double-action manifolds since one can hope that they inherit the suppression. These models require a spherical 3-space ${\\cal S}^3$ but we mostly restrict our analysis to almost flat spaces corresponding to a total density $\\Omega_{\\hbox{\\scriptsize tot}}$ in the range $\\Omega_{\\hbox{\\scriptsize tot}}=1.001,\\dots,1.05$. The multiconnected spaces that exist in the spherical 3-space ${\\cal S}^3$ can be classified with respect to three categories of spherical 3-manifolds as described in \\cite{Gausmann_Lehoucq_Luminet_Uzan_Weeks_2001}. The criterion is based on the kind of two subgroups $R$ and $L$ which generate the deck group $\\Gamma$ which in turn defines the spherical 3-manifold. The subgroups $R$ and $L$ act as pure right-handed and left-handed Clifford translations, respectively. The first category consists of the single-action manifolds in which only one of the subgroups $R$ and $L$ acts non-trivially. The double-action manifolds, the second category, require that both subgroups $R$ and $L$ are non-trivial, such that each element of the subgroup $R$ is combined with each element of the subgroup $L$. The third category, the linked-action manifolds, are similar to the second one, except that there are rules specifying which elements of $R$ and $L$ can be combined such that a manifold is obtained instead of an orbifold. For more details on the three categories, see \\cite{Gausmann_Lehoucq_Luminet_Uzan_Weeks_2001}. The single-action manifolds are the simplest with respect to an analysis of the statistical CMB properties, since they are independent of the position of the CMB observer within the manifold. Such manifolds are called homogeneous. This contrasts to the other two categories where the ensemble average of the CMB statistics depends on the observer position, in general, and a much more involved analysis is required for these inhomogeneous manifolds. The aim of this paper is to close a gap that is left by our previous publications \\cite{Aurich_Lustig_2012b,Aurich_Lustig_2012c} which cover some of the possible double-action manifolds. A survey of lens spaces $L(p,q)$ is presented in \\cite{Aurich_Lustig_2012b}. The lens spaces $L(p,q)$ have the amazing property that they have members in all three categories. While the spaces $L(p,1)$ are single-action manifolds, the lens spaces $L(m n,q)$ which are generated by $R=Z_m$ and $L=Z_n$ with $m$ and $n$ relatively prime, are double-action manifolds. The remaining lens spaces belong to the linked-action manifolds so that members of all three categories are studied in \\cite{Aurich_Lustig_2012b}. This study leads to the result that lens spaces $L(p,q)$ with $q\\simeq 0.28 p$ or $q\\simeq 0.38 p$ possess a pronounced suppression of CMB correlations on large angular scales compared to other lens spaces. The prism double-action manifolds, which are generated by a binary dihedral group $R=D^\\star_p$ and a cyclic group $L=Z_n$, are investigated in \\cite{Aurich_Lustig_2012c}, and at least three promising spaces are found. In the notation of \\cite{Aurich_Lustig_2012c}, the prism double-action manifolds are called $DZ(p,n)$ where the letters indicate the subgroups $R$ and $L$, and $p$ and $n$ are the group orders of $D^\\star_p$ and $Z_n$. Three prism double-action manifolds with a remarkable large-scale CMB suppression are $DZ(8,3)$, $DZ(16,3)$, and $DZ(20,3)$. Because of these encouraging results, the question emerges whether there are further interesting double-action manifolds. The double-action manifolds not covered in \\cite{Aurich_Lustig_2012b} and \\cite{Aurich_Lustig_2012c} are those generated by one of the three binary polyhedral groups $R = T^\\star$, $O^\\star$ or $I^\\star$ and a cyclic group $L=Z_n$. For these spaces we introduce the notation $TZ(24,n)$, $OZ(48,n)$, and $IZ(120,n)$. Thus, this paper is devoted to these spaces in order to close the gap with respect to the CMB properties of polyhedral double-action manifolds. We investigate all 20 polyhedral double-action manifolds which exist up to the group order 1000. The polyhedral double-action manifolds can be considered as a dissection of one of the three polyhedral spaces with respect to a cyclic group. The three polyhedral spaces belong to the single-action spaces and are thus homogeneous. They are well studied in several previous papers starting with \\cite{Luminet_Weeks_Riazuelo_Lehoucq_Uzan_2003} which analyses the Poincar\\'e dodecahedral topology that is the binary icosahedral space ${\\cal I}$. A strong suppression of CMB correlations on large angular scales is found for this space at $\\Omega_{\\hbox{\\scriptsize tot}}\\simeq 1.02$. This result is confirmed in \\cite{Aurich_Lustig_Steiner_2004c} by using a much larger set of eigenfunctions for the computation of the CMB statistics. Further studies concerning this model can be found in \\cite{Roukema_et_al_2004,Gundermann_2005,Aurich_Lustig_Steiner_2005a,% Aurich_Lustig_Steiner_2005b,Lustig_2007,% Key_Cornish_Spergel_Starkman_2007,Niarchou_Jaffe_2007,% Lew_Roukema_2008,Roukema_et_al_2008a,Roukema_Kazimierczak_2011}. In \\cite{Gundermann_2005,Aurich_Lustig_Steiner_2005a,Niarchou_Jaffe_2007} the statistical CMB analysis is extended to the binary tetrahedral space ${\\cal T}$ and the binary octahedral space ${\\cal O}$. The central result of \\cite{Aurich_Lustig_Steiner_2005a} is that all three polyhedral spaces lead to a significant suppression of large-scale correlations described by the $S$ statistics of a factor of $\\sim 0.11$ compared to the simply connected spherical 3-space ${\\cal S}^3$. This factor is achieved at $\\Omega_{\\hbox{\\scriptsize tot}}\\simeq 1.07$, $\\Omega_{\\hbox{\\scriptsize tot}}\\simeq 1.04$, and $\\Omega_{\\hbox{\\scriptsize tot}}\\simeq 1.02$ for the spaces ${\\cal T}$, ${\\cal O}$, and ${\\cal I}$, respectively. In the following we analyse the statistical properties on large separation angles $\\vartheta$ of the polyhedral double-action manifolds in order to address the question how strong these spaces suppress the CMB correlations in terms of the $S$ and $I$ statistics defined below in eqs.\\,(\\ref{Eq:S_statistic_60}) and (\\ref{Eq:I_measure}). Since they are based on the three polyhedral spaces with their very low values of the $S$ statistics, they also could yield promising models for the description of our Universe. The polyhedral double-action manifolds are generated by a cyclic subgroup $L=Z_n$ and one of the three binary polyhedral groups $R = T^\\star$, $O^\\star$, and $I^\\star$, where the cyclic groups $Z_n$ have to fulfil $\\hbox{gcd}(24,n)=1$, $\\hbox{gcd}(48,n)=1$, and $\\hbox{gcd}(120,n)=1$, respectively. The generator $g_l=({\\bf 1},g_b)$ of the cyclic group $Z_n$ is given by \\begin{equation} \\label{Def:Z_n} g_b \\; = \\; \\hbox{diag}(e^{+2\\pi\\hbox{\\scriptsize i}/n},e^{-2\\pi\\hbox{\\scriptsize i}/n}) \\hspace{10pt}. \\end{equation} The binary polyhedral groups $R = T^\\star$, $O^\\star$, and $I^\\star$ have two generators $g_{r1}=(g_{a1 },{\\bf 1})$ and $g_{r2}=(g_{a2 },{\\bf 1})$. These two generators can be described by \\begin{equation} \\label{Def:poly_group} g_{ak} \\; = \\; \\left(\\begin{array}{cc} \\cos(\\tau_k)-\\hbox{i}\\sin(\\tau_k)\\cos(\\theta_k) & -\\hbox{i}\\sin(\\tau_k)\\sin(\\theta_k)e^{-\\hbox{\\scriptsize i}\\phi_k}\\\\ -\\hbox{i}\\sin(\\tau_k)\\sin(\\theta_k)e^{\\hbox{\\scriptsize i}\\phi_k}& \\cos(\\tau_k)+\\hbox{i}\\sin(\\tau_k)\\cos(\\theta_k) \\end{array}\\right) \\hspace{6pt} \\end{equation} using the spherical coordinates $(\\tau_k, \\theta_k, \\phi_k)$, $k=1,2$. The values of $\\tau_k$, $\\theta_k$ and $\\phi_k$ given in table \\ref{Tab:generators_poly_spherical} determine the representation of the groups $T^\\star$, $O^\\star$, and $I^\\star$. \\begin{table}[!htbp] \\centering \\begin{tabular}{|c|c|c|} \\hline group $R$ & ($\\tau_1$, $\\theta_1$, $\\phi_1$) & ($\\tau_2$, $\\theta_2$, $\\phi_2$)\\\\ \\hline $T^\\star$ & $(\\frac{\\pi}{3}, 0, 0)$&$(\\frac{\\pi}{3}, \\arccos\\big(\\frac{1}{3}\\big), 0)$\\\\ \\hline $O^\\star$ & $(\\frac{\\pi}{4}, 0, 0)$&$(\\frac{\\pi}{3}, \\arccos\\big(\\frac{1}{\\sqrt 3}\\big), 0)$\\\\ \\hline $I^\\star$ & $(\\frac{\\pi}{5}, 0, 0)$&$(\\frac{\\pi}{5}, \\arccos\\big(\\frac{1}{\\sqrt 5}\\big), 0)$\\\\ \\hline \\end{tabular} \\caption{\\label{Tab:generators_poly_spherical} These values of $(\\tau_1, \\theta_1, \\phi_1)$ and $(\\tau_2, \\theta_2, \\phi_2)$ determine the two generators in eq.\\,(\\ref{Def:poly_group}) for the binary polyhedral groups $T^\\star$, $O^\\star$, and $I^\\star$. } \\end{table} Although the central topic of this paper concerns the correlation of the CMB fluctuations on large angular scales, some remarks on the circles-in-the-sky (CITS) signature are in order which serves as a topological test \\cite{Cornish_Spergel_Starkman_1998b}. The CITS test requires a full CMB sky survey and has been applied to different sky maps derived from the WMAP mission. The first year CMB data are analysed with respect to nearly back-to-back circle pairs by \\cite{Cornish_Spergel_Starkman_Komatsu_2003,Key_Cornish_Spergel_Starkman_2007} and no significant signature was found, whereas a search for the Poincar\\' e dodecahedral space, being a single action manifold, yields a tentative signal \\cite{Roukema_et_al_2004}. It is shown in \\cite{Aurich_Janzer_Lustig_Steiner_2007} that the error in the CMB signal has to be significantly lower than 50$\\mu\\hbox{K}$ in order to get a CITS signal. It is hard to obtain a statement about the size of the error in the heavily processed WMAP data leading to the maps used for the CITS searches. The constraint to nearly back-to-back circle pairs is investigated in \\cite{Mota_Reboucas_Tavakol_2010,Mota_Reboucas_Tavakol_2011} where the probability for the deviation from the back-to-back orientation is studied. The seven year WMAP data are analysed by \\cite{Bielewicz_Banday_2011} again for the special case of back-to-back circles, and no topological signature is detected. A complete CITS search without the back-to-back restriction is carried out in \\cite{Vaudrevange_Starkman_Cornish_Spergel_2012} using the WMAP seven year data. Several signatures are found, but they are all ascribed to foreground sources, so that the paper concludes that no hint for a non-trivial topology is found. Since no statement on the accuracy of the CMB signal is made, one cannot exclude the possibility that a possible CITS signal is swamped by foreground sources which can even produce spurious signals. In order to reduce the computer time, the analysis of \\cite{Vaudrevange_Starkman_Cornish_Spergel_2012} uses a search grid for the screening of circle pairs that is coarser than that of the CMB map. Our preliminary investigations show that the probability for missing circle pairs increases by such an algorithm. For this reason topologies with few circle pairs have a high probability to get missed in this way. Since these results are devoted to a future publication, we turn to the CMB correlations now. ", "conclusions": "This paper analyses the large-scale correlations in the CMB sky for the polyhedral double-action manifolds. With this analysis, the CMB correlations are finally investigated for all double-action manifolds since those belonging to the lens spaces and to the prism double-action manifolds are already studied in \\cite{Aurich_Lustig_2012b} and \\cite{Aurich_Lustig_2012c}. The large-scale correlation measure (\\ref{Eq:S_min}) is used in order the search for spaces with a significant suppression of correlations in the CMB anisotropy on scales above $\\vartheta>60^\\circ$. This quantity is normalised to the simply connected spherical space ${\\cal S}^3$. The lens spaces $L(p,q)$ can lead to a suppression relative to ${\\cal S}^3$ by a factor of about $\\sim 0.5$ \\cite{Aurich_Lustig_2012b}. The lens spaces with such a large suppression have lenticular fundamental cells whose two faces have to be rotated by a relative angle of $\\sim 101^\\circ$ or $\\sim 137^\\circ$ before the faces are identified. Among the prism double-action manifolds $DZ(p,n)$, there are spaces with even smaller large-scale correlations with suppression factors in the range $0.3\\dots 0.4$. The three best candidates are $DZ(8,3)$, $DZ(16,3)$, and $DZ(20,3)$ \\cite{Aurich_Lustig_2012c}. Although this CMB suppression is remarkable, it is less pronounced than in the cases of the three binary polyhedral spaces ${\\cal T}$, ${\\cal O}$, and ${\\cal I}$ where the suppression factor is of the order of $0.11$. The three binary polyhedral spaces ${\\cal T}$, ${\\cal O}$, and ${\\cal I}$ lead to the three classes $TZ(24,n)$, $OZ(48,n)$, and $IZ(120,n)$ of polyhedral double-action manifolds. The analysis of this paper shows that several polyhedral double-action manifolds can possess even stronger suppressions than those found in the three binary polyhedral spaces (see figure \\ref{Fig:s_statistic_min}). From these spaces, the octahedral double-action manifolds $OZ(48,n)$ with $n=7$, 11, 13, 17, and 19 have suppression factors below 0.11 for $\\Omega_{\\hbox{\\scriptsize tot}}$ in the range $\\Omega_{\\hbox{\\scriptsize tot}}=1.03\\dots 1.04$. With the constraint $\\Omega_{\\hbox{\\scriptsize tot}}\\leq 1.02$, the best octahedral double-action manifold is the space $OZ(48,5)$ with a suppression factor 0.2. In addition, three further octahedral spaces with $n=7$, 11, and 13 possess suppression factors between 0.3 and 0.4 in that $\\Omega_{\\hbox{\\scriptsize tot}}$ range. The icosahedral double-action manifold $IZ(120,7)$ also reveals an interesting behaviour with a suppression factor below 0.11 close to $\\Omega_{\\hbox{\\scriptsize tot}}=1.02$. Remarkably, insisting on the constraint $\\Omega_{\\hbox{\\scriptsize tot}}\\leq 1.01$, the space $IZ(120,7)$ has the largest suppression of all investigated spherical manifolds. The minimum in the correlation measure is obtained at $\\Omega_{\\hbox{\\scriptsize tot}}=1.007$ with a suppression factor of 0.27. The tetrahedral double-action manifolds $TZ(24,n)$ do not provide comparable candidates to explain the low correlations on the CMB sky at large angles. They possess such small suppression factors only for significantly larger values of the total density $\\Omega_{\\hbox{\\scriptsize tot}}$ which are beyond the range considered in this paper. Some $TZ(24,n)$ spaces have nevertheless CMB suppressions comparable to the prism double-action manifolds $DZ(p,n)$ mentioned above. The ensemble averages of the correlation functions $C(\\vartheta)$ of the polyhedral double-action manifolds are also compared to the observed $C^{\\hbox{\\scriptsize obs}}(\\vartheta)$ using the $I$ statistics. This analysis confirms the result obtained from the $S$ statistics. Concluding, there are five octahedral double-action manifolds and one icosahedral double-action manifold with a group order below 1\\,000 with a pronounced suppression of CMB correlation on large angular scales which deserve further investigations. \\appendix" }, "1207/1207.3244_arXiv.txt": { "abstract": "{Open clusters are ideal test particles for studying the formation and evolution of the Galactic disk. However, the number of clusters with information about their radial velocities and chemical compositions remains largely insufficient.} {We attempt to increase the number of open clusters with determinations of radial velocities and metallicities from spectroscopy.} {We acquired medium-resolution spectra (R$\\sim$8000) in the region of the infrared \\ion{Ca}{II} triplet lines ($\\sim$8500\\AA) for several stars in four open clusters with the long-slit spectrograph IDS at the 2.5m Isaac Newton Telescope, Roque de los Muchachos Observatory, Spain. Radial velocities were obtained by cross-correlating the observed spectra with those of two template stars. We used the relationships available in the literature between the strength of infrared \\ion{Ca}{II} lines and metallicities to derive the metal content of each cluster.} {We provide the first spectroscopic determinations of radial velocities and metallicities for the open clusters Berkeley~26, Berkeley~70, NGC~1798, and NGC~2266. We obtain $\\langle V_r \\rangle$=68$\\pm$12, -15$\\pm$7, 2$\\pm$10, and -16$\\pm$15 km s$^{-1}$ for Berkeley~26, Berkeley~70, NGC~1798, and NGC~2266, respectively. For Berkeley~26 we derive a metallicity of [Fe/H]=-0.35$\\pm$0.17 dex. Berkeley~70 has a solar metallicity of [Fe/H]=-0.01$\\pm$0.14 dex, while NGC~1798 has a slightly lower metal content of [Fe/H]=-0.12$\\pm$0.07 dex. Finally, we derive a metallicity of [Fe/H]=-0.38$\\pm$0.06 dex for NGC~2266.} {} ", "introduction": "Open clusters (OCs) have been widely used to investigate the existence of trends in the Galactic disk such as radial and vertical metallicity gradients or an age-metallicity relationship \\citepads[e.g.][]{1979ApJS...39..135J,1980A&A....81..375P,1997AJ....114.2556T, 2002AJ....124.2693F,2003AJ....125.1397C,2010A&A...511A..56P,2011A&A...535A..30C} . However, our knowledge of OC properties is still far from complete. Age and distance estimations, mainly obtained from isochrone fitting, are available for about $\\sim$70\\% of the more than 2100 OCs known in our Galaxy \\citepads{2002A&A...389..871D}\\footnote{The updated version of this catalog can be found at {\\tt http://www.astro.iag.usp.br/~wilton/}.}. However, radial velocities have been determined only for a 24\\% of them. The picture is even worse in the case of chemical compositions, which have been obtained for only 9\\%, mainly by means of different studies in the Washington \\citepads[e.g.][]{2005MNRAS.363.1247P}, DDO \\citepads[e.g.][]{1977AJ.....82...35J,1999A&AS..134..301C}, Str\\\"omgren \\citepads[e.g.][]{2003A&A...403..937P}, UBV \\citepads[e.g.][]{1998A&AS..128..131C}, IR \\citepads[e.g.][]{1999A&A...343..825C,2000A&A...353..147V}, and Vilnius \\citepads[e.g.][]{2011BaltA..20...27B,2011BaltA..20....1Z} photometric systems, often giving rise to considerable differences among them and to those obtained from spectroscopy. Reliable information about the chemical composition can be retrieved only from spectroscopy. However, the acquisition of high-resolution spectra (R$\\geq$20000), which is the only way to derive detailed abundances, needs large amounts of telescope time. This together with the complexity of the associated analysis explain why this kind of study has been performed for only 4\\% of known OCs \\citepads{2011A&A...535A..30C}. The alternative is to perform low- and medium-resolution spectroscopy, although these data provide information only information about the metallicities and radial velocities of stars \\citepads[e.g.][]{1993A&A...267...75F,2002AJ....124.2693F,2007AJ....134.1298C, 2009MNRAS.393..272W}. \\begin{table*}[htb] \\caption{Observing logs and program star information.} \\label{obsstars} \\centering \\renewcommand{\\footnoterule}{} \\begin{scriptsize} \\begin{tabular}{l c c c c c c c c c c c c c c} \\hline\\hline Cluster & Star\\footnote{Identification taken from WEBDA database.} & $\\alpha_{2000}$ & $\\delta_{2000}$ & V & V-I & B-V & V$_{r}$ & EW$_{8498}$ & EW$_{8542}$ & EW$_{8662}$ & $\\Sigma$ Ca & t$_{exp}$ & S/N$^{tot}$ & Note\\\\ & & (hrs) & (deg) & (mag) & (mag) &(mag) & (km s$^{-1}$) & (\\AA) & (\\AA) & (\\AA) & (\\AA) & (sec) & &\\\\ \\hline Be~26 & 1038 & 06:50:02.9 & +05:44:59 & 18.845 & 1.283 & --- & 66.1$\\pm$11.9 & --- & --- & --- & --- & 2$\\times$950 & 25 & 2 \\\\ & 1155 & 06:50:09.5 & +05:44:28 & 15.754 & 1.729 & --- & 71.0$\\pm$7.9 & 1.17$\\pm$0.13 & 3.39$\\pm$0.12 & 2.74$\\pm$0.12 & 7.30$\\pm$0.21 & 2$\\times$950 & 29 & 1 \\\\ & 1231 & 06:50:08.6 & +05:44:07 & 16.702 & 1.819 & --- & 66.1$\\pm$7.7 & 1.13$\\pm$0.14 & 3.21$\\pm$0.21 & 2.77$\\pm$0.10 & 7.11$\\pm$0.27 & 2$\\times$1500 & 25 & 1 \\\\ & 1288 & 06:50:05.7 & +05:43:53 & 15.500 & 1.735 & --- & 68.3$\\pm$7.2 & 1.27$\\pm$0.11 & 3.51$\\pm$0.12 & 2.93$\\pm$0.12 & 7.71$\\pm$0.20 & 2$\\times$850 & 29 & 1 \\\\ & 1421 & 06:50:18.4 & +05:43:22 & 15.556 & 2.063 & --- & 68.3$\\pm$6.7 & 1.45$\\pm$0.08 & 3.94$\\pm$0.09 & 2.62$\\pm$0.07 & 8.01$\\pm$0.14 & 2$\\times$900 & 40 & 1 \\\\ & 1650 & 06:50:14.9 & +05:42:17 & 14.974 & 2.092 & --- & 97.1$\\pm$7.8 & --- & --- & --- & --- & 2$\\times$850 & 45 & 9 \\\\ Be~70 & 1105 & 05:25:44.8 & +41:56:44 & 12.651 & 2.035 & 2.014 & -13.8$\\pm$7.6 & 1.94$\\pm$0.05 & 4.66$\\pm$0.05 & 3.50$\\pm$0.06 & 10.10$\\pm$0.09 & 2$\\times$700 & 186 & 1 \\\\ & 1088 & 05:25:45.0 & +41:55:55 & 14.097 & 1.994 & 1.849 & -4.3$\\pm$7.5 & 1.61$\\pm$0.02 & 4.32$\\pm$0.02 & 3.20$\\pm$0.03 & 9.13$\\pm$0.04 & 2$\\times$800 & 98 & 1 \\\\ & 0820 & 05:25:49.9 & +41:56:51 & 14.482 & 1.944 & 1.816 & -17.3$\\pm$7.9 & 1.53$\\pm$0.04 & 4.18$\\pm$0.03 & 2.99$\\pm$0.04 & 8.70$\\pm$0.06 & 2$\\times$800 & 69 & 1 \\\\ & 0609 & 05:25:54.6 & +41:58:22 & 14.608 & 1.734 & 1.585 & -11.7$\\pm$7.2 & 1.56$\\pm$0.04 & 3.69$\\pm$0.05 & 2.80$\\pm$0.06 & 8.05$\\pm$0.09 & 2$\\times$850 & 44 & 1 \\\\ & 0811 & 05:25:50.1 & +41:57:46 & 14.774 & 1.796 & 1.645 & -19.7$\\pm$7.7 & 1.46$\\pm$0.08 & 3.87$\\pm$0.07 & 2.94$\\pm$0.08 & 8.27$\\pm$0.12 & 2$\\times$850 & 58 & 1 \\\\ & 0690 & 05:25:52.7 & +41:57:13 & 14.923 & 1.617 & 1.457 & -20.4$\\pm$7.3 & 1.49$\\pm$0.05 & 3.71$\\pm$0.07 & 2.71$\\pm$0.06 & 7.91$\\pm$0.11 & 2$\\times$950 & 58 & 1 \\\\ & 0735 & 05:25:51.9 & +41:57:35 & 15.142 & 1.762 & 1.624 & -19.7$\\pm$7.9 & 1.46$\\pm$0.08 & 3.89$\\pm$0.08 & 2.80$\\pm$0.10 & 8.15$\\pm$0.15 & 2$\\times$950 & 40 & 1 \\\\ & 0171 & 05:26:03.3 & +41:58:54 & 15.304 & 1.716 & 1.547 & -12.9$\\pm$8.9 & 1.35$\\pm$0.12 & 4.01$\\pm$0.13 & 2.35$\\pm$0.13 & 7.71$\\pm$0.22 & 2$\\times$900 & 30 & 1 \\\\ NGC~1798 & 0005 & 05:11:36.7 & +47 41 49 & 14.380 & 1.633 & 1.483 & 4.3$\\pm$7.3 & 1.54$\\pm$0.06 & 3.79$\\pm$0.07 & 2.91$\\pm$0.07 & 8.24$\\pm$0.12 & 2$\\times$800 & 56 & 1 \\\\ & 0009 & 05:11:42.8 & +47:41:32 & 14.777 & 1.653 & 1.486 & 0.2$\\pm$7.3 & 1.31$\\pm$0.10 & 3.45$\\pm$0.10 & 2.60$\\pm$0.08 & 7.36$\\pm$0.16 & 2$\\times$750 & 40 & 1 \\\\ & 0011 & 05:11:37.0 & +47:40:15 & 15.305 & 1.666 & 1.529 & -3.4$\\pm$7.0 & 1.40$\\pm$0.09 & 3.85$\\pm$0.11 & 2.71$\\pm$0.08 & 7.96$\\pm$0.16 & 2$\\times$800 & 41 & 1 \\\\ & 0013 & 05:11:41.2 & +47:40:41 & 15.325 & 1.660 & 1.524 & -5.6$\\pm$7.2 & 1.38$\\pm$0.07 & 3.52$\\pm$0.10 & 2.64$\\pm$0.13 & 7.54$\\pm$0.18 & 2$\\times$850 & 40 & 1 \\\\ & 0043 & 05:11:47.9 & +47:40:26 & 12.875 & 2.252 & 1.970 & 11.2$\\pm$7.8 & 1.89$\\pm$0.02 & 4.68$\\pm$0.03 & 3.44$\\pm$0.03 & 10.01$\\pm$0.05 & 2$\\times$700 & 110 & 1 \\\\ & 0608 & 05:11:41.0 & +47:40:43 & 16.653 & 0.909 & 0.816 & 12.8$\\pm$11.0 & --- & --- & --- & --- & 2$\\times$850 & 25 & 2 \\\\ NGC~2266 & 0034 & 06:43:16.0 & +26 57 45 & 12.454 & --- & 0.957 & 21.0$\\pm$7.7 & --- & --- & --- & --- & 2$\\times$350 & 55 & 9 \\\\ & 0067 & 06:43:20.2 & +26:57:32 & 10.538 & --- & 1.288 & 30.2$\\pm$7.2 & --- & --- & --- & --- & 2$\\times$50 & 55 & 9 \\\\ & 0075 & 06:43:12.9 & +26:57:07 & 15.242 & --- & 0.341 & -18.0$\\pm$9.2 & --- & --- & --- & --- & 2$\\times$350 & 24 & 2\\\\ & 0096 & 06:43:28.0 & +26:58:10 & 12.287 & --- & 1.253 & -8.9$\\pm$7.5 & 1.47$\\pm$0.05 & 3.84$\\pm$0.05 & 2.65$\\pm$0.06 & 7.96$\\pm$0.09 & 2$\\times$250 & 58 & 1 \\\\ & 0101 & 06:43:25.3 & +26:57:49 & 12.978 & --- & 1.187 & -19.1$\\pm$7.8 & 1.46$\\pm$0.06 & 3.71$\\pm$0.06 & 2.25$\\pm$0.10 & 7.42$\\pm$0.13 & 2$\\times$500 & 50 & 1 \\\\ & 1011 & 06:43:25.3 & +26:57:50 & 12.117 & --- & 1.273 & -18.5$\\pm$7.1 & 1.43$\\pm$0.06 & 4.13$\\pm$0.05 & 2.97$\\pm$0.06 & 8.53$\\pm$0.10 & 2$\\times$500 & 58 & 1\\\\ \\hline\\hline \\end{tabular} \\tablefoot{ (1) member star used both for radial velocity and metallicity determination; (2) member star used only for radial velocity determination; (9) non-member star. } \\end{scriptsize} \\end{table*} The goal of this paper is to increase the number of clusters with radial velocities and metallicities determined from spectroscopy. For this purpose, we have selected four OCs, Berkeley~ 26, Berkeley~70, NGC~1798, and NGC~2266, that had not been studied spectroscopically before. These clusters are located towards the Galactic anticenter, at relatively large distances (R$_{GC}\\geq$10 kpc) where the trend of the radial metallicity seems to flatten \\citepads[e.g.][]{2011A&A...535A..30C}. Medium-resolution spectra in the region of the near-infrared \\ion{Ca}{II} triplet (CaT) at $\\sim$8500\\AA~were obtained for several red giant branch (RGB) stars in each cluster. The observations and data reduction are described in Sect.~\\ref{sec2}, radial velocities and metallicities are obtained in Sects.~\\ref{sec3}--\\ref{sec4}, the results for each cluster are discussed and compared with the literature in Sect.~\\ref{sec5}, and the comparison of the results obtained here with the trends observed in the disk is performed in Sec.~\\ref{sec6}. Finally, our main conclusions are summarised in Sect.~\\ref{sec7}. ", "conclusions": "\\label{sec7} We have analyzed medium-resolution spectra (R$\\sim$8000) in the infrared CaT region ($\\sim$8500 \\AA) of several stars in four OCs: Berkeley~26, Berkeley~70, NGC~1798, and NGC~2266. To our knowledge, this is the first time that these clusters have been studied spectroscopically. Our main results can be summarised as follows: \\begin{itemize} \\item For Berkeley~26, we derived a mean radial velocity of $\\langle V_r \\rangle$=68$\\pm$12 km s$^{-1}$ based on six stars. One of them is a main-sequence star that was not used in the metallicity determinations. Using $V$ and $I$ magnitudes, we determined a metallicities of [Fe/H]=-0.43$\\pm$0.20 and -0.35$\\pm$0.17 dex, respectively. \\item The eight observed stars in Berkeley~70 were confirmed as members based on their radial velocity. For them, we derived a mean radial velocity of $\\langle V_r \\rangle$=-15$\\pm$7 km s$^{-1}$ and metallicities of [Fe/H]=-0.02$\\pm$0.18 and -0.01$\\pm$0.14 dex based on $V$ and $I$ magnitude data, respectively. \\item We derived a mean radial velocity for NGC~1798 of $\\langle V_r \\rangle$=2$\\pm$10 km s$^{-1}$ from six stars, although one object is on the main sequence. For the RGB stars, we obtained a metallicity of [Fe/H]=-0.16$\\pm$0.10 and -0.12$\\pm$0.07 dex in $V$ and $I$ bandpasses, respectively. \\item In the case of NGC~2266, its mean radial velocity, $\\langle V_r \\rangle$=-16$\\pm$15 km s$^{-1}$ was obtained for four objects. A metallicity of [Fe/H]=-0.38$\\pm$0.07 dex was derived from the $V$ magnitudes of the three RGB stars observed in this cluster. \\end{itemize} Finally, we investigated how the four analyzed clusters fit the trends defined by other well-studied OCs. This comparison is motivated by our clusters being situated at distances where other investigations have observed a change in the slope of the metallicity gradient. In general, Berkeley~26 and NGC~1798 follow the trends described by other coeval systems situated at the same distance. In the case of NGC~2266, its height above the Galactic plane can explain its low metallicity compared to other clusters of the same age. In contrast, Berkeley~70 seems to be more metal-rich than other coeval clusters situated at similar distances. We suggest that this cluster may have formed at relatively small Galactocentric distances and has migrated outwards with time. Nevertheless, more information is needed to confirm or discard this hypothesis." }, "1207/1207.2845_arXiv.txt": { "abstract": "We present a tool for measuring the equivalent width (\\ew ) in high-resolution spectra. The Tool for Automatic Measurement of Equivalent width (\\TAME) provides the \\ews\\ of spectral lines by profile fitting in the automatic or the interactive mode, which can yield a more precise result through the adjustment of the local continuum and fitting parameters. The automatic \\ew\\ results of \\TAME\\ have been verified by comparing them with the manual \\ew\\ measurements by IRAF {\\tt splot} task using the high-resolution spectrum of the Sun, and measuring \\ews\\ in the synthetic spectra with different spectral resolutions and S/N ratios. The \\ews\\ measured by \\TAME\\ agree well with manually measured values, with a dispersion of less than 2 m\\AA. By comparing the input \\ews\\ for synthetic spectra and \\ews\\ measured by \\TAME, we conclude that it is reliable for measuring the \\ews\\ in a spectrum with a spectral resolution, R $\\gtrsim$ 20000 and find that the errors in \\ews\\ is less than 1 m\\AA\\ for a S/N ratio $\\gtrsim$ 100. ", "introduction": "The measurement of equivalent width (\\ew) for spectral absorption lines is essential in a spectral analysis, particularly for determining the atmospheric parameters and chemical abundances of stars. In the study of stellar spectroscopy, it is critical to determine the atmospheric parameters of stars, such as the effective temperature (\\teff), surface gravity (\\logg), metallicity ([Fe/H]), and micro-turbulence ($\\xi_t$), because atmospheric parameters are fundamental to understand spectroscopic properties and construct the model atmosphere for an abundance analysis. For the atmospheric parameters, however, the most common method is to analyse the abundances that can be obtained from \\ew\\ measurements of neutral and singly ionized lines. Additionally, the chemical abundances are also estimated by measuring the \\ews\\ of atomic lines \\citep[e.g.,][]{bensby03, santos04, bond06, gilli06, sousa06, kang11}. The \\ew\\ measurement, therefore, is undoubtedly the most important task in spectroscopic studies. The \\ews\\ of spectral lines have generally been measured by using the {\\tt splot} task in IRAF\\footnote{IRAF is the Image Reduction and Analysis Facility software. It is written and supported by the IRAF programming group at the National Optical Astronomy Observatories (NOAO) that is operated by the Association of Universities for Research in Astronomy (AURA), Inc., under cooperative agreement with the National Science Foundation} {\\tt echelle} package, which makes it possible to manually estimate the \\ew\\ of each line. Although this method guarantees a high degree of accuracy for \\ew\\ measurement, it requires a disciplined expert in the field of stellar spectroscopy and the result depends on the personal bias. For an abundance analysis, it is necessary to measure the \\ews\\ for many lines for each star, which is a tedious and time-consuming task. Therefore, a uniform and fast method for \\ew\\ measurement is required in stellar abundance studies using a large number of high-resolution spectra. \\citet{ARES} presented a new C++ code, called ARES (Automatic Routine for line Equivalent widths in stellar Spectra), which can automatically and simultaneously measure the \\ews\\ of spectral lines in stellar spectra. ARES provides quick measurement results for \\ews, without manual operation, from high-resolution spectra. However, ARES code focuses on the performance of the code, and hence deprives a user of the interactive operation that can be used to control an environment for each line. Further, a Fortran code for \\ew\\ measurement, called DAOSPEC, was recently presented by \\citet{DAOSPEC}. In order to achieve more accurate measurement, DAOSPEC offers the enhanced interactive mode for detailed manipulative tasks, such as the adjustment of the local continuum and the deblending of nearby lines. Unfortunately, ARES and DAOSPEC were written in C++ and Fortran, respectively. Therefore, the installation of ARES and DAOSPEC depends on the platform OS, and it would be difficult and inconvenient to compile and run these codes coherently, because of the required libraries (e.g., cfitsio\\footnote{http://heasarc.nasa.gov/fitsio/fitsio.html}, GSL\\footnote{http://www.gnu.org/s/gsl/}, SuperMongo\\footnote{http://www.astro.princeton.edu/$\\sim$rhl/sm/sm.html}, IRAF). To avoid these practical difficulties, we have developed the Tool for Automatic Measurement of Equivalent width (\\TAME)\\footnote{\\TAME\\ can be downloaded from http://astro.snu.ac.kr/$\\sim$wskang/tame/}, which is written in IDL\\footnote{The Interactive Data Language (IDL) is a cross-platform software, providing support for Microsoft Windows\\textsuperscript{\\textregistered}, Mac OS X, Linux, and Solaris (http://www.exelisvis.com)} and uses a graphical user interface (GUI). \\TAME\\ can be used with any platform OS on which IDL has been installed, and it contains various features required to adjust the environment of \\ew\\ measurement such as the local continuum and radial velocity of a star. Its semi-automatic mode ($hereafter$, the interactive mode) offers more flexible measurement of the \\ew\\ as similar to DAOSPEC. And its fully automatic mode ($hereafter$, the automatic mode) can simultaneously measure the \\ews\\ for a large set of lines. \\TAME\\ produces a formatted text file containing the \\ew\\ result which can be used directly in the abundance analysis code MOOG \\citep{MOOG}, in addition to a graphical output file with the fitting results of the local continuum and line profile. In this work, we describe the procedure by which \\TAME\\ measures the \\ews\\ of spectral lines and examine the results of \\ews\\ obtained by \\TAME. In Sect. 2, we introduce the user interface and input parameters of \\TAME. In Sect. 3, we explain the automatic processes used to measure the \\ew\\ with \\TAME, such as determining the local continuum, searching for blended lines, and fitting the lines with a Gaussian/Voigt profile. In Sect. 4, we present the comparison of the manual \\ew\\ measurements obtained using IRAF and those obtained using \\TAME\\ for the high-resolution spectra of the Sun, whose atmospheric parameters are well known. We also discuss the difference between the \\ew\\ estimated by \\TAME\\ and the input \\ew\\ for a synthetic spectrum having different spectral resolutions and S/N ratios. In Sect. 5, we summarize the advantages of using \\TAME\\ along with its performance results. ", "conclusions": "We have developed a new software tool for automatic \\ew\\ measurement called \\TAME\\ for measuring \\ews\\ in a high-resolution spectrum. It has the following features: \\begin{itemize} \\item \\TAME\\ can automatically measure \\ews\\ for a large set of lines in a spectrum simultaneously. \\item \\TAME\\ offers an interactive mode, in which a user can adjust the local continuum level precisely and change parameters such as the SMOOTHER, radial velocity, and type of fitting profile (Gaussian/Voigt). \\item \\TAME\\ provides a text file including the \\ews\\ with a format suited for MOOG code and a graphical post-script file for confirming the \\ew\\ results when performing abundance analysis. \\end{itemize} We verified \\TAME\\ in two ways. By using the solar spectrum, we measured solar \\ews\\ by \\TAME\\ and compared them with those obtained by the traditional method with IRAF {\\tt splot} task. The \\ews\\ measured by \\TAME\\ showed a good agreement with the precise manual measurements made using IRAF, with a standard deviation of only 1.76 m\\AA, and the atmospheric parameters of the Sun were determined to \\teff\\ $=$ 5791 K, \\logg\\ $=$ 4.54 dex, [Fe/H] = 0.03 dex, and $\\xi_t = $ 0.81 \\kms\\ from the \\ew\\ result of \\TAME. In order to examine the effect of the S/N ratio and spectral resolution on \\ew\\ measurement, we performed \\ew\\ measurement for different synthetic spectra by using \\TAME\\ in the fully automatic mode without any manual interactions. From the test results obtained for the synthetic spectra, we concluded that the \\ew\\ measurements obtained by \\TAME\\ are reliable for high-resolution spectra with R $\\gtrsim$ 20000 and found that the errors in \\ews\\ could be expected to be less than 1 m\\AA\\ for a S/N ratio $\\gtrsim$ 100." }, "1207/1207.2768_arXiv.txt": { "abstract": "We distinguish between Local Group field galaxies which may have passed through the virial volume of the Milky Way, and those which have not, via a statistical comparison against populations of dark matter haloes in the Via Lactea II (VLII) simulation with known orbital histories. Analysis of VLII provides expectations for this escaped population: they contribute 13 per cent of the galactic population between 300 and 1500 kpc from the Milky Way, and hence we anticipate that about 7 of the 54 known Local Group galaxies in that distance range are likely to be Milky Way escapees. These objects can be of any mass below that of the Milky Way, and they are expected to have positive radial velocities with respect to the Milky Way. Comparison of the radius-velocity distributions of VLII populations and measurements of Local Group galaxies presents a strong likelihood that Tucana, Cetus, NGC3109, SextansA, SextansB, Antlia, NGC6822, Phoenix, LeoT, and NGC185 have passed through the Milky Way. Most of these dwarfs have a lower HI mass fraction than the majority of dwarfs lying at similar distances to either the Milky Way or M31. Indeed, several of these galaxies -- especially those with lower masses -- contain signatures in their morphology, star formation history and/or gas content indicative of evolution seen in simulations of satellite/parent galactic interactions. Our results offer strong support for scenarios in which dwarfs of different types form a sequence in morphology and gas content, with evolution along the sequence being driven by interaction history. ", "introduction": "Dwarfs within the approximate 300 kpc virial radii of the Milky Way and M31 are preferentially small, gas-poor spheroids, compared to their field counterparts which are typically larger, gaseous, and irregularly shaped \\citep[e.g.][]{Grebel03,Grcevich09,Weisz11,vandenBergh94}. This position-morphology relationship, first noted by \\citet{Einasto74}, appears universal, as it is found in other galaxy groupings as well \\citep[e.g.][]{Skillman03b,Bouchard09}. The position-morphology relationship is attributed to a transformation of gas-rich dwarf irregular galaxies into gas-poor dwarf spheroidals via environmental effects. That the cumulative environmental effects encountered during a passage through a larger potential are sufficient to transform the morphology of a dwarf is very well motivated by simulations \\citep[e.g.][]{Mayer01a,Mayer01b,Kravtsov04,Mayer06}. Environmental effects each leave a multitude of signatures on a galaxy. Tidal stirring has been shown to convert stellar components from disks to bars and finally to pressure supported spheroidal systems \\citep[e.g.][]{Klimentowski09}. Shocking and ram-pressure stripping of gas \\citep{Sofue94,Grebel03,Mayer10} leaves signatures in the satellite's star formation history, either as starbursts \\citep{Hernquist89, Barnes96,Mihos96} or as starvation and quenching of the star formation (see \\citet{Kawata08} for a low mass group). Tidal shock heating is known to disrupt or destroy star clusters \\citep{Kruijssen11}. Although initially it appeared that these effects might only be highly effective within 50 kpc of a Milky Way-size object \\citep{Sofue94,Grebel03}, recent studies (including other effects e.g. tidal effects with UV background \\citet{Mayer06}, resonant stripping \\citet{D'Onghia09}) show that such a close passage may not be necessary for a morphological transformation. There are objects that do not fit the rough distance-morphology relationship, because they exist outside the virial radius of the nearest large galaxy, but nevertheless exhibit a morphology that suggests strong interactions (e.g. Tucana). However, interaction with a Milky-Way-size object is not the only way to affect changes in dwarfs: dwarf-dwarf interactions (or even mergers) have been shown to stimulate bursts of star formation, and to create irregular morphologies \\citep{Mendez99,Bekki08, Besla12}; interactions between dark satellites and dwarf galaxies can also trigger starbursts or a transformation to a spheroidal morphology \\citep{Helmi12}; episodic star formation \\citep{Gerola80} of the bursty \\citep[e.g.][]{Davies88} or quiescent variety \\citep[e.g.][]{Tosi92} has been shown to reduce high gas content and lower metallicity through the interaction of stellar feedback and the interstellar medium; and small galaxies can ionize and blow out (via stellar feedback, and including supernova feedback) enough gas to shut off a star formation episode \\citep[e.g.][]{deYoung94, Brinks98}. Knowledge of the past orbit of a dwarf would be helpful in determining whether prior interaction with the Milky Way is sufficient to explain the properties of objects like Tucana or whether alternative explanations (such as dwarf-dwarf encounters or internal effects) need to be invoked. Unfortunately, drawing direct, clear connections between the current morphology of an observed object and its past orbit is limited by our observational perspective. It is difficult or impossible to measure more than the angular position, distance and line-of-sight velocity for field dwarfs, and these quantities have been shown to be insufficient to determine a complete, accurate, orbital history for objects in the Local Group \\citep{Lux10}. However, there is precedence for using distance and velocity measurements to draw a connection between morphology and rough orbital history on the larger scale of galaxy clusters. These clusters exhibit a high incidence of so-called ``backsplash galaxies'', defined to be objects on extreme orbits that have taken them through the inner 0.5 \\Rvir of a larger potential and subsequently carried them back outside \\Rvir. \\citet{2005MNRAS.356.1327G} demonstrated in simulations how a population of backsplash galaxies might be probabilistically separated from those infalling to the cluster for the first time using their observed velocities. Subsequent observations demonstrate that galaxies selected using this approach indeed exhibit unusual or unique morphologies \\citep{2010arXiv1012.3114M,2006ApJ...647..946S,2002ApJ...580..164S,2002AJ....124.2440S,2000ApJ...540..113B}. Owing to the approximately self-similar clustering of dark matter, the research done on clusters provokes questions about the existence and nature of backsplash galaxies on a smaller scale, specifically in the Local Group. Theoretical work on these scales suggests the existence of satellites on extreme orbits around potentials about the size of the Milky Way. Around galaxy potentials, \\citet{Sales07b} identifies an ``associated'' population of haloes which have at some point passed through the virial volume of the main halo. Of these, $\\sim$6 per cent have apocentric radii greater than 50 per cent of their turnaround radius, and a few have been ejected as far as 2.5 \\Rvir. \\citep[Similar populations have also been seen in simulations analysed by][.]{Warnick08,Wang09,Ludlow09,Knebe11a} Data samples which further inform the extent to which morphology and gas content can be related to dynamical history are growing rapidly. The study of Local Group objects has recently been invigorated by an influx of new members: SDSS enabled an expansion in the volume probed by star count surveys, which resulted in the discovery of numerous new dwarf satellite galaxies of both the Milky Way and M31 \\citep[e.g.][]{Willman05,Belokurov06,Irwin07,Zucker04}. Moreover, new observational surveys, such as DES \\citep{Bernstein11}, SkyMapper \\citep{Keller07}, Pan-STARRS \\citep{Kaiser02}, and LSST \\citep{LSST_ScienceBook_2009,Ivezic08}, will be even more sensitive to faint magnitude and low surface brightness objects, and are expected to reveal even lower surface brightness objects over even larger volumes of space \\citep{Tollerud08}. Motivated by this confluence of theoretical analyses, recent observational discoveries and promising new surveys, this paper makes connections between dynamically distinct histories for subhaloes seen in a cosmological simulation of structure formation (Via Lactea II, hereafter VLII), and properties of Local Group dwarf galaxies. More specifically, we establish that it is possible to distinguish field populations which may have passed within the Milky Way-like halo of VLII from those which have not, using observable properties at z=0 (radial distance, line-of-sight velocity and mass). The $z=0$ distributions of these observable properties for haloes in VLII are given in Section \\ref{sec.theory}. The simulated populations can be used to categorise the orbital histories of Local Group field objects (Section \\ref{sec.obs}). Assuming that morphology is a result of environmental changes over time, we can connect morphology to orbit. Finally, we discuss whether this rough orbital characterisation provides insight into the morphologies and gas content of nearby field objects in the Local Group (Section \\ref{sec.summ}). The methods we employ, and details of the VLII simulation itself, are described in Section \\ref{sec.method}. ", "conclusions": "\\label{sec.summ} We demonstrate that with just the line-of-sight distance and velocity, we can obtain a rough interaction history for field objects in the Local Group via comparison with VLII populations. We separate field haloes in VLII into categories: associated haloes have been within the virial radius of the main Milky Way-like halo, unassociated haloes have not. We find $\\sim$13 per cent of field haloes in the simulations to have passed through the virial volume of the Milky Way-like halo at some point during their histories. These associated haloes could be found out to 5 \\Rvir. This suggests that, for the Local Group, of the 54 known galaxies within this distance range, we expect at least 7 to have interacted with the Milky Way. Further analysis of VLII suggest that these associated objects are likely to have positive radial velocities with respect to the Milky Way of order or greater than the Hubble Flow, which will make them distinguishable from the unassociated populations. From our analysis we do not expect a mass-distance bias in the associated dwarfs around the Milky Way. About 4 per cent of the MW-associated haloes may have become renegade haloes bound to M31. The separation between the associated and unassociated populations in the distance-velocity plane in VLII was applied in the Local Group to identify field dwarfs that may be associated with the Milky Way: Tucana, Cetus, Antlia, NGC3109, SextansA, SextansB, NGC6822, Phoenix, LeoT and NGC185. Several of these objects have signatures in their morphology, gas content, or stellar populations that could be the result of their passage through the Milky Way. This possibility should be considered when analyzing transformative internal and external effects for these objects. Overall we conclude that our simple test provides strong support for scenarios in which the gas-poor, dwarf spheroidal objects in the field result from the transformation of gas-rich irregulars during past interactions with Milky Way or Andromeda." }, "1207/1207.3839_arXiv.txt": { "abstract": "We have analysed all the good quality \\xmm\\ data publicly available for the bright ULXs Holmberg\\,IX X-1 and NGC\\,1313 X-1, with the aim of searching for discrete emission or absorption features in the Fe K band that could provide observational evidence for the massive outflows predicted if these sources are accreting at substantially super-Eddington rates. We do not find statistically compelling evidence for any atomic lines, and the limits that are obtained have interesting consequences. Any features in the immediate Fe K energy band (6--7\\,\\kev) must have equivalent widths weaker than $\\sim$30\\,eV for Holmberg\\,IX X-1, and weaker than $\\sim$50\\,eV for NGC\\,1313 X-1 (at 99 per cent confidence). In comparison to the sub-Eddington outflows observed in GRS\\,1915+105, which imprint iron absorption features with equivalent widths of $\\sim$30\\,eV, the limits obtained here appear quite stringent, particularly when Holmberg\\,IX X-1 and NGC\\,1313 X-1 must be expelling at least 5--10 times as much material if they host black holes of similar masses. The difficulty in reconciling these observational limits with the presence of strong line-of-sight outflows suggests that either these sources are not launching such outflows, or that they must be directed away from our viewing angle. ", "introduction": "Ultraluminous X-ray Sources (ULXs) are extra-nuclear point sources in external galaxies observed to be more luminous in X-rays than the Eddington luminosity \\le\\ for a stellar mass ($\\sim$10\\msun) black hole (\\lx\\,$\\gtrsim$\\,$10^{39}$\\,\\ergps); on rare occasions ULX luminosities have been observed to exceed $10^{41}$\\,\\ergps\\ (\\citealt{WaltonULXcat, Sutton12}) and even to reach as high as $\\sim10^{42}$\\,\\ergps\\ (\\citealt{Farrell09}). This combination has led to extended debate over the nature of these sources. The majority of early theories were based around either the presence of intermediate mass black holes (IMBHs: $10^{2}$ $\\lesssim$ \\mbh\\ $\\lesssim$ $10^{5}$\\,\\msun; \\citealt{Colbert99}), super-Eddington accretion states for standard mass binary systems (\\eg \\citealt{Poutanen07}, \\citealt{Finke07}) or the inferred luminosities being artificially high due to anisotropic emission (\\citealt{King01}), although observations of excited emission lines, typically He {\\small II} $\\lambda$4686, from either the accretion disc of the ULX or its surrounding nebula appear to rule out highly beamed emission via photon counting arguments (\\eg \\citealt{Pakull02}, \\citealt{Kaaret04}, and \\citealt{Berghea10}). More recently, proposed interpretations for bright ULXs (\\lx\\,$\\sim$\\,$10^{40}$ \\ergps) tend to incorporate elements from each of these three ideas, invoking black holes only slightly larger than those seen in Galactic binary systems (\\mbh\\ $\\lesssim$ 100\\,\\msun) with mildly super-Eddington accretion rates (luminosities up to a few times \\le), such that the accretion disc is geometrically thicker than the standard thin discs predicted for more moderate accretion rates (\\citealt{Shakura73}) and forces some mild anisotropy (`slim' discs; see \\eg \\citealt{Abram80}). For recent reviews on the observational status and the potential nature of ULXs see \\cite{Roberts07rev} and \\cite{Feng11rev}. \\begin{table*} \\caption{Basic details of the two bright ULXs considered in this work, Holmberg\\,IX X-1 and NGC\\,1313 X-1, and the \\xmm\\ observations analysed. Distances are taken from \\citet{Paturel02} and \\citet{TULLY} respectively.} \\begin{center} \\begin{tabular}{c c c c c c c c c} \\hline \\hline \\\\[-0.3cm] Source & RA & DEC & $D$ & $z$ & OBSID & Obs. Date & Duration & Good Exposure \\\\ & (h:m:s) & (d:m:s) & (Mpc) & & & & (ks) & (\\epicpn; ks) \\\\ \\\\[-0.3cm] \\hline \\hline \\\\[-0.25cm] Holmberg\\,IX X-1 & 09:57:53.2 & +69:03:48.3 & 3.55 & 0.000153 & 0112521001 & 10/04/2002 & 11 & 7 \\\\ & & & & & 0112521101 & 16/04/2002 & 12 & 8 \\\\ & & & & & 0200980101 & 26/09/2004 & 119 & 88 \\\\ \\\\[-0.15cm] NGC\\,1313 X-1 & 03:18:20.0 & -66:29:11.0 & 3.7 & 0.001568 & 0106860101 & 17/10/2000 & 42 & 27\\\\ & & & & & 0150280601 & 08/01/2004 & 55 & 10 \\\\ & & & & & 0150181101 & 16/01/2004 & 9 & 6 \\\\ & & & & & 0205230201 & 01/05/2004 & 13 & 3 \\\\ & & & & & 0205230301 & 05/06/2004 & 12 & 9 \\\\ & & & & & 0205230601 & 07/02/2005 & 14 & 10 \\\\ & & & & & 0405090101 & 15/10/2006 & 123 & 99 \\\\ \\\\[-0.25cm] \\hline \\hline \\end{tabular} \\label{tab_obs} \\end{center} \\end{table*} One of the observational results in recent X-ray studies of ULXs that has sparked significant interest is that many of these sources display curvature in their X-ray continuum above $\\sim$3\\,\\kev\\ (\\citealt{Stobbart06, Gladstone09, Walton4517}). Although some of these are lower luminosity sources that display (hot) thermal disc-like spectra, in which high energy curvature would naturally be expected, even the brighter ULXs with two apparently distinct emission components, potentially analogous to the disc and Comptonised corona observed in Galactic black hole binaries (BHBs), also appear to show curvature in their high energy components. The hard, Comptonised emission rarely displays similar curvature in the standard sub-Eddington accretion states of BHBs (for a recent review see \\citealt{Remillard06rev}), and if this is an intrinsic difference in the high energy continuum, it could imply that there are fundamental differences between the accretion onto ULXs and their lower luminosity BHB cousins. Given that the basic accretion geometry is not expected to evolve with black hole mass, this might indirectly suggest we are not viewing sub-Eddington accretion. In this high-Eddington framework, \\cite{Gladstone09} propose that the high energy curvature could be due to Comptonisation in an optically thick corona, which shrouds the inner accretion disc. The effect of the corona causes the observed temperature of the disc (identified here with the soft component) to appear artificially low as only the outer disc is observed directly. Low disc temperatures have frequently been inferred for bright ULXs, which if otherwise associated with the inner disc would imply the presence of IMBHs (see \\eg \\citealt{MilFab04}). This interpretation has been slightly modified by \\cite{Middleton11a}, who propose that the cool, quasi-thermal emission previously associated with the disc is instead thermalised emission from a photosphere associated with the base of an outflowing wind, as strong outflows are ubiquitously predicted by models for high and particularly super-Eddington accretion (see \\citealt{Ohsuga11}, and references therein). Galactic BHBs frequently display evidence for mass outflows, which at moderately high accretion rates (during the thermal dominated states) take the form of equatorial disc winds with outflow velocities $v_{\\rm out} \\lesssim 1000$\\,\\kms. Prominent examples are the BHBs GRO\\,J1655-40 (\\citealt{Miller06a, DiazTrigo07}) and GRS\\,1915+105 (\\citealt{Kotani00, Lee02, Ueda09, Neilsen09}). When these outflows impinge upon our line of sight to the central source, their observational consequence is to imprint absorption features onto the intrinsic X-ray continuum, typically in the form of narrow absorption lines from highly ionised species, with the most prominent features arising from the \\ka\\ transitions of highly ionised iron (Fe {\\small XXV} and/or {\\small XXVI}). In the case of GRS\\,1915+105, perhaps the best studied BHB outflow, \\cite{Neilsen09} show that the strength of the absorption features generally increases as the luminosity of the source, and hence its Eddington ratio increases. The mass of this source is relatively well known (\\mbh\\ = $14\\pm4$\\,\\msun; \\citealt{Greiner01}), and it is observed to radiate up to, and possibly slightly in excess of its Eddington limit (\\citealt{Vierdayanti10grs}). While the geometry of the BHB disc winds observed at moderate accretion rates are largely equatorial, the scale height of these winds is generally expected to increase with Eddington ratio (see \\eg \\citealt{Abramowicz05}, \\citealt{King09}), so at higher accretion rates the outflows should probably cover a larger solid angle. If instead, super-Eddington outflows are present in ULXs but, despite their potential large scale heights, do not occur along our line of sight, rather than viewing absorption features we should instead observe reprocessed emission from the material in the outflow. Indeed, reprocessed emission lines consistent with neutral (or moderately ionised) iron are seen ubiquitously in the spectra of Galactic high mass X-ray binaries (HMXBs; see {\\citealt{Torrejon10}), which are expected to accrete largely from the stellar winds launched by their massive companions. Furthermore, the majority of ULXs themselves are widely expected to be HMXB analogues, owing to the their apparent correlation with recent star formation (\\citealt{Swartz09}). Recent population studies of ULXs appear to support this picture, as their luminosity function appears consistent with being a smooth extrapolation of the HMXB luminosity function (at least for sources in spiral/irregular type galaxies; \\citealt{Grimm03}, \\citealt{WaltonULXcat}, \\citealt{Swartz11}). Here, we present an analysis of the Fe K band for two bright ($L_{\\rm X} \\sim 10^{40}$\\,\\ergps) ULXs with hard X-ray spectra, Holmberg\\,IX X-1 (also known as M\\,81 X-9) and NGC\\,1313 X-1, in which we search for evidence of atomic iron features that could be associated with outflowing material. These sources have been selected as they represent the datasets with the highest quality data at high energies for such bright, isolated ULXs. The paper is structured as follows: section \\ref{sec_obs} describes our data reduction prodecure, and section \\ref{sec_analysis} describes the analysis performed. We discuss our results in section \\ref{sec_discussion} and summarise our conclusions in section \\ref{sec_conc}. ", "conclusions": "\\label{sec_conc} Making use of all the good quality \\xmm\\ data publically available for two bright ULXs, Holmberg\\,IX X-1 and NGC\\,1313 X-1, we have searched for (persistent) discrete atomic features in their high energy spectra that could be associated with either iron emission or absorption, and provide observational evidence for the massive outflows predicted if these sources are accreting at substantially super-Eddington rates. We do not find any statistically compelling evidence for any features, either in absorption or in emission. However, we do find that the data currently available for these sources is of sufficient quality that we can place fairly stringent limits on the strengths of any features intrinsically present these spectra that remain undetected. Any features present (absorption or emission) in the immediate Fe K energy band (6--7\\,\\kev) must have equivalent widths weaker than $\\sim$30\\,eV for Holmberg\\,IX X-1, and weaker than $\\sim$50\\,eV for NGC\\,1313 X-1. In comparison to the strongest sub-Eddington outflows observed in GRS\\,1915+105, which imprint iron absorption features with equivalent widths of $\\sim$30\\,eV, the limits obtained for the ULXs considered here appear quite restrictive, particularly when these sources must be radiating at $\\sim$5--10 times their Eddington limits if they host black holes of similar masses, and should therefore be expelling at least 5--10 times as much material as GRS\\,1915+105. The difficulty in trying to reconcile these observational constraints with the presence of strong line-of-sight outflows leads us to conclude that either these sources do not launch such outflows, which would strongly argue against a highly super-Eddington interpretation, or that they much be launched away from our viewing angle. Additional deep observations of bright ULXs with current instrumentation, and in particular with the micro-calorimeter due for launch on board \\textit{Astro-H} in the near future, will be essential for detecting any weak iron features in such sources, and hence in placing stronger constraints on the nature of any outflows present, and the ultimate nature of these sources." }, "1207/1207.6288_arXiv.txt": { "abstract": "{We present the \\grid (Astro-rivelatore Gamma a Immagini LEggero -- \\textit{Gamma-Ray Imaging Detector}) monitoring of \\mbox{Cygnus X-3}, during the period between November 2007 and July 2009. We report here the whole \\grid monitoring of \\mbox{Cygnus X-3} in the AGILE ``pointing'' mode data-taking, to confirm that the \\gray activity coincides with the same repetitive pattern of multiwavelength emission and analyze in depth the overall \\gray spectrum by assuming both leptonic and hadronic scenarios. Seven intense \\gray events were detected in this period, with a typical event lasting one or two days. These durations are longer than the likely cooling times of the \\gray emitting particles, implying we see continuous acceleration rather than the result of an impulsive event such as the ejection of a single plasmoid that then cools as it propagates outwards. Cross-correlating the \\grid light curve with both X-ray and radio monitoring data, we find that the main events of \\gray activity were detected while the system was in soft spectral X-ray states (\\textit{RXTE}/ASM (\\textit{Rossi X-ray Timing Explorer}/All-Sky Monitor) count rate in the 3-5 keV band $\\gtrsim 3~\\mathrm{counts~s^{-1}}$), that coincide with local and often sharp minima of the hard X-ray flux (\\textit{Swift}/BAT (Burst Alert Telescope) count rate $\\lesssim 0.02~\\mathrm{counts}$ $\\mathrm{cm^{-2}~s^{-1}}$), a few days before intense radio outbursts. This repetitive temporal coincidence between the \\gray transient emission and spectral state changes of the source turns out to be the \\textit{spectral signature} of \\gray activity from this microquasar. These \\gray events may thus reflect a sharp transition in the structure of the accretion disk and its corona, which leads to a rebirth of the microquasar jet and subsequent enhanced activity in the radio band. The \\gray differential spectrum of \\mbox{Cygnus X-3} (100 MeV -- 3 GeV), which was obtained by averaging the data collected by the \\grid during the \\gray events, is consistent with a power law of photon index $\\alpha=2.0~\\pm~0.2$. Finally, we examine leptonic and hadronic emission models for the \\gray events and find that both scenarios are valid. In the leptonic model -- based on inverse Compton scatterings of mildly relativistic electrons on soft photons from both the Wolf-Rayet companion star and the accretion disk -- the emitting particles may also contribute to the overall hard X-ray spectrum, possibly explaining the hard non-thermal power-law tail seen during special soft X-ray states in \\mbox{Cygnus X-3}. } ", "introduction": "\\mbox{Cygnus X-3} is the brightest radio source among all known microquasars and was discovered, as an X-ray source, in 1966 \\citep{giacconi_67}. It is a high-mass X-ray binary, whose companion star is a Wolf-Rayet (WR) star \\citep{vankerk_92} with a strong helium stellar wind \\citep{szo_zdzia_08}. The system is located at a distance of about 7-10 kpc \\citep{bonnet_88,ling_09}. The orbital period is 4.8 hours, as inferred from infrared \\citep{becklin_73}, X-ray \\citep{parsignault_72}, and \\gray \\citep{abdo_09} observations. Owing to its very tight orbit (orbital distance $d \\approx 3 \\times 10^{11}$ cm), the compact object is totally enshrouded in the wind of the companion star\\footnote{The observational evidence of this strong wind can be found in the prominent attenuation of the \\mbox{Cygnus X-3} power density spectrum (PDS) for frequencies above 0.1 Hz \\citep{axelsson_09, koljonen_11}.}. The nature of the compact object is still uncertain\\footnote{Published results suggest either a neutron star of 1.4 $M_{\\odot}$ \\citep{stark_03} or a black hole with a mass $\\lesssim 10~M_{\\odot}$ \\citep{hanson_00,shrader_10}.} \\citep{vilhu_09}, although a black hole scenario is favored \\citep{szo_zdzia_08, szostek_08}. In the radio band, the system shows strong flares (``\\textit{major radio flares}'') reaching up to few tens of Jy. Radio observations at milliarcsec scales confirm emissions (at cm wavelengths) from both a core and a one-sided relativistic jet ($v \\sim 0.81c$), with an inclination to the line-of-sight of $\\lesssim14^{\\circ}$ \\citep{mioduszewski_01}. The radiation from the jet dominates the radio emission from the core during (and soon after) the major flares \\citep{tudose_10}.\\\\ \\mbox{Cygnus X-3} exhibits a clear, repetitive pattern of (anti)correlations between radio and X-ray emission, and an overall anticorrelation between soft and hard X-ray fluxes \\citep{mccollough_99,szostek_08}. The most important pattern of correlations found by \\citet{szostek_08} is related to the connection between radio (8.3 GHz band, GBI) and soft X-ray emissions (3-5 keV band of the \\textit{Rossi X-ray Timing Explorer}/All-Sky Monitor (\\textit{RXTE}/ASM)). When the soft X-ray flux is above the transition level (3 counts/s), the source can be found in different states, depending on the level of the radio flux density. In particular, the \\textit{quenched state} is characterized by a radio flux density $\\leqslant$ 30 mJy and followed by a \\textit{major-flaring state} with values of radio flux density $\\geqslant$ 1 Jy. It is very important to emphasize that all major radio flares have been observed after a quenched state, and in almost all cases the quenched state is followed by a major flare. After a major flare, a ``\\textit{hysteresis}'' in the radio/soft-X-ray plane is found, because the decline in the radio flux density never occurs by means of a quenched state. Firm detections of high-energy \\grays (HE \\grays: $>$100 MeV) from \\mbox{Cygnus X-3}\\footnote{\\gray detections of \\mbox{Cygnus X-3} were reported in both the 1970s and 1980s at TeV \\citep{vladimirsky_73,danaher_81,lamb_82} and PeV energies \\citep{samorski_83,bhat_86}. However, subsequent observations by more sensitive ground-based telescopes did not confirm TeV and PeV emission from this source \\citep{oflaherty_92}. Furthermore, the \\textit{COS-B} satellite could not find any clear emission from \\mbox{Cygnus X-3} at MeV-GeV energies \\citep{hermsen_87}, and both \\textit{CGRO}/EGRET observations of the Cygnus region (1991-1994) and the first-year analysis of AGILE observations could not demonstrate that there was a solid association with the microquasar, although they confirmed a \\gray detection above 100 MeV in a region including \\mbox{Cygnus X-3} \\citep{mori_97, pittori_09}.} were published at the end of 2009: the AGILE (Astro-rivelatore Gamma a Immagini LEggero) team found evidence that strong \\gray transient emission above 100 MeV coincided with special X-ray/radio spectral states \\citep{tavani_09a}, and the \\lat (Large Area Telescope) collaboration announced the detection of \\gray orbital modulation \\citep{abdo_09}. The peak \\gray isotropic luminosity detected above 100 MeV is $L_{\\gamma}\\sim10^{36}~\\mathrm{erg~s^{-1}}$ (for a distance of 7-10 kpc). {The \\gray emission is most likely associated with a relativistic jet \\citep{tavani_09a, abdo_09, dubus_10, cerutti_11, zdziarski_12}, but the radiative process (leptonic or hadronic) is uncertain}. A possible leptonic scenario for \\gray emission in \\mbox{Cygnus X-3} was proposed by \\citealp{dubus_10}: stellar ultraviolet (UV) photons are Compton upscattered to HE by relativistic electrons accelerated in the jet. The particle acceleration could take place in a shock where the jet interacts with the dense stellar wind of the WR star. The emerging picture is that of a jet with moderate bulk relativistic speed and oriented not too far from the line-of-sight.\\\\ The \\gray modulation -- coherent with the orbital period -- suggests that the emitting region is located at distances of between $\\sim$$10^{10}$ cm and $\\sim$$3 \\times 10^{12}$ cm ($10d$) from the compact object \\citep{dubus_10,cerutti_11}. The lack of modulation at radio wavelengths and the delay ($\\sim$5 days, \\citealp{abdo_09}) between the onset of \\gray activity and the radio flare suggest that different emission regions are linked by the collimated jet. The \\gray emission -- related to inverse Compton (IC) scatterings -- most likely occurs close to the compact object, while the radio emission -- assumed to be synchrotron in origin -- occurs farther out in the jet, at an angular distance from the core of a few tens of milli-arcseconds (e.g., \\citealp{tudose_07, tudose_10}), corresponding to $\\sim10^{15}$--$10^{16}$ cm. The \\gray modulation is due to the anisotropic efficiency of the IC scattering \\citep{aharonian_81}. Thus, the \\gray maximum occurs at the superior conjuction (where the compact object is behind the WR star), when relativistic electrons of the jet, moving towards the Earth, have head-on collisions with stellar UV photons. This orbital phase corresponds to the minimum of the X-ray modulation, produced in turn by the maximum of absorption/scattering by the companion's wind \\citep{abdo_09,dubus_10,zdziarski_12}. A hadronic scenario accounting for \\gray emission in microquasars was discussed by \\citet{romero_03,romero_05}. Their model is based on the interaction of a mildly relativistic jet with the dense wind of the companion star, and the \\gray emission is due to the decay of neutral pions ($\\pi^{0}$) produced by $pp$ collisions. Furthermore, TeV emission from relativistic jet in microquasars has been predicted by several models (e.g., see \\citealp{atoyan_99}). A search for very-high-energy (VHE) \\grays from the microquasar GRS 1915+105 with H.E.S.S. (High Energy Stereoscopic System) was carried out, but no significant detection was found in the direction of the source \\citep{hess_09}. On the other hand, hints of VHE \\grays were found in Cygnus X-1 \\citep{albert_07}. The Major Atmospheric Gamma-ray Imaging Cherenkov Telescope (MAGIC) observed \\mbox{Cygnus X-3} several times between March 2006 and August 2009, during both its hard and soft states\\footnote{The MAGIC telescope was also pointed at \\mbox{Cygnus X-3} after two \\gray alerts from the \\grid team (the first one after the \\gray event of 16-17 April 2008, and the second after the event of 13-14 July 2009, see Appendix \\ref{agile_data}). In both cases, they found a $2\\sigma$ upper limit, for energies above 250 GeV, of $\\sim$$10^{-11}$ \\flx \\citep{aleksic_10}.}, but no evidence of clear VHE \\gray emission from the microquasar was found: an overall $2\\sigma$ upper limit to the integral flux was set at $2.2 \\times 10^{-12}$ \\flx for energies above 250 GeV \\citep{aleksic_10}. Here we present a comprehensive and homogeneous analysis of \\mbox{Cygnus X-3} that takes into account \\gray events found in the data between 2007 November 2 and 2009 July 29, during the AGILE ``pointing'' mode data-taking. We analyzed a dataset previously published by \\citet{tavani_09a} and \\citet{bulgarelli_12a}. We report here the whole \\grid monitoring of \\mbox{Cygnus X-3} during the ``pointing'' mode, to confirm that the \\gray activity coincides with the same repetitive pattern of multiwavelength emission and to analyze in depth the overall \\gray spectrum by assuming both leptonic and hadronic scenarios. ", "conclusions": "Several events of \\gray activity were detected by the \\grid from \\mbox{Cygnus X-3} while the system was in a special radio/X-ray spectral state: intense \\gray activity was detected during prominent minima of the hard X-ray light curve (corresponding to strong soft X-ray emission), a few days before intense radio outbursts (major radio flares). This temporal repetitive coincidence turned out to be the spectral signature of \\gray activity from this puzzling microquasar, which might open new areas to study the interplay between the accretion disk, the corona, and the formation of relativistic jets. The simultaneous strong soft X-ray emission from the disk and \\gray emission from the jet preceding the intense radio outbursts are consistent with a scenario in which the hot thermal corona ``dissolves'' and the accretion power from the disk directly charges the jet, emitting \\grays and, subsequently, radio outbursts (via synchrotron processes) far from the compact object. The \\gray detections of \\mbox{Cygnus X-3} provide new constraints on emission models for this powerful X-ray binary, indicating that hybrid-Comptonization mechanisms \\citep{coppi_99} alone cannot account for the \\gray fluxes detected by AGILE and Fermi above 100 MeV, unless we assume unrealistic physical parameters \\citep{cerutti_11}. This implies that the corona cannot be the site of the \\gray emission. We found that the innermost part of the jet (distances $\\lesssim 10^{10}$ cm from the compact object) could provide a strong contribution to the hard X-rays at $\\sim$100 keV during the \\gray emitting interval, while the farthest part (distances $\\gtrsim 10^{10}$ cm from the compact object) produces the bulk of the \\gray emission above 100 MeV. We found that the \\gray spectrum of \\mbox{Cygnus X-3} detected by the \\grid is significantly harder than the time-averaged spectrum obtained by \\lat for the ``\\gray active periods'' of the microquasar, lasting $\\sim$4 months (see Figure \\ref{cyg_x3_agile_fermi_only}). Although both the AGILE main \\gray events and the Fermi \\gray active periods are both likely related to the presence of an active jet, the spectral difference may imply that there was a fast hardening of the spectrum during the peak \\gray events, lasting $\\sim$1-2 days. We have demonstrated that both a leptonic model based on inverse Compton emission from a relativistic plasmoid injected into the jet and a hadronic model based on $\\pi^0$-decays, might account for the \\gray emission observed by the \\grid. Both of these models require the introduction of a new component (``IC bump'' or ``$\\pi^0$-bump'') into the SED of the system. In both the leptonic and hadronic pictures, the inclination of the jet to the line of sight is assumed to be $i = 14^{\\circ}$.\\\\ A leptonic scenario seems to be more likely than a hadronic one: the \\gray modulation, the spectral link between hard X-ray and \\gray spectra, and the temporal link between \\gray events and radio flares could be interpreted in a natural way by assuming that the electrons are the main emitters. According to our results, the HE \\gray emission occurs at distances up to $\\sim$$ 10^{12}$ cm from the compact object. If we were to interpret the $\\sim$4-day delay between the onset of \\gray and radio flaring emission as the propagation time of the relativistic jet ($v = \\sqrt{5}c/3$), the radio burst would occur at a distance of $\\sim$$8\\times10^{15}$ cm.\\\\ Our hadronic model, with the assumption of a standard WR wind, would require a jet kinetic power of $L_{kin,~p} \\approx 1.5 \\times 10^{38}$ $\\mathrm{erg~s^{-1}}$ to explain the \\gray emission detected by AGILE. This value is of the same order of magnitude as the bolometric luminosity of the disk/corona during the hypersoft spectral state, and lower than the Eddington accretion limit for a black hole with a mass of $M_x \\approx 10 M_{\\odot}$ ($L_{Edd} \\approx 10^{39}$ $\\mathrm{erg~s^{-1}}$). Thus, a hadronic picture is physically reasonable and not energetically less likely than a leptonic one. At present, there is no strong evidence that one of these hypotheses can be excluded, and it remains an open question whether the dominant process for \\gray emission in microquasars is either hadronic or leptonic \\citep{mirabel_12}. The firm discovery of \\gray emission from this microquasar represents the experimental proof that these astrophysical objects are capable of accelerating particles up to relativistic energies, through a mechanism -- related to the disk-corona dynamics -- that leads to jet formation." }, "1207/1207.4186.txt": { "abstract": "We characterize the dust in NGC\\,628 and NGC\\,6946, two nearby spiral galaxies in the KINGFISH sample. With data from 3.6\\um\\ to 500\\um, dust models are strongly constrained. Using the \\citet{Draine+Li_2007} dust model, (amorphous silicate and carbonaceous grains), for each pixel in each galaxy we estimate (1) dust mass surface density, (2) dust mass fraction contributed by polycyclic aromatic hydrocarbons (PAH)s, (3) distribution of starlight intensities heating the dust, (4) total infrared (IR) luminosity emitted by the dust, and (5) IR luminosity originating in regions with high starlight intensity. We obtain maps for the dust properties, which trace the spiral structure of the galaxies. The dust models successfully reproduce the observed global and resolved spectral energy distributions (SEDs). The overall dust/H mass ratio is estimated to be $0.0082\\pm0.0017$ for NGC~628, and $0.0063\\pm0.0009$ for NGC~6946, consistent with what is expected for galaxies of near-solar metallicity. Our derived dust masses are larger (by up to a factor 3) than estimates based on single-temperature modified blackbody fits. We show that the SED fits are significantly improved if the starlight intensity distribution includes a (single intensity) ``delta function'' component. We find no evidence for significant masses of cold dust $(T\\lesssim12\\rm{K})$. Discrepancies between PACS and MIPS photometry in both low and high surface brightness areas result in large uncertainties when the modeling is done at PACS resolutions, in which case SPIRE, MIPS70 and MIPS160 data cannot be used. We recommend against attempting to model dust at the angular resolution of PACS. ", "introduction": "Introduction} Interstellar dust affects the appearance of galaxies, by attenuating short wavelength radiation from stars and ionized gas, and contributing IR, submm, mm, and microwave emission [for a review, see \\citet{Draine_2003a}]. Dust is also an important agent in the fluid dynamics, chemistry, heating, cooling, and even ionization balance in some interstellar regions, with a major role in the process of star formation. Despite the importance of dust, determination of the physical properties of interstellar dust grains has been a challenging task -- even the overall amount of dust in other galaxies has often been very uncertain. Many previous studies have used far-infrared photometry to estimate the dust properties of galaxies. For example, \\citet{Draine+Dale+Bendo+etal_2007} used global photometry of 65 galaxies in the ``Spitzer Infrared Nearby Galaxies Survey'' (SINGS) galaxies to estimate the total dust mass and PAH abundance in each galaxy, and to characterize the intensity of the starlight heating the dust. For most of these galaxies the photometry extended only to 160$\\micron$, although ground-based global photometry at 850$\\micron$ was also available for 17 of the 65 galaxies.\\footnote{Photometry at 850$\\micron$ was available for 26 galaxies in the SINGS sample, but the data were only reliable (and used) for 17 galaxies.} \\citet{Munoz-Mateos+Gil_de_Paz+Boissier+etal_2009} used images of the SINGS galaxies to examine the radial distribution of the dust surface density. \\citet{Sandstrom+Bolatto+Draine+etal_2010} studied the PAHs in the Small Magellanic Cloud SMC on a pixel-by-pixel basis with a very similar dust model as the present work. Very recently, \\citet{Totani+Takeuchi+Nagashima+etal_2011} used global photometry in 6 bands -- 9, 18, 65, 90, 140, and 160$\\micron$ -- to estimate dust masses for a sample of more than 1600 galaxies in the Akari All Sky Survey \\citep{Ishihara+Onaka+Kataza+etal_2010, Yamamura+Makiuti+Ikeda+etal_2010}. In the present work, we develop state-of-the-art image processing and dust modeling techniques aiming to reliably determine the dust properties in resolved studies. The present study makes use of combined imaging by the Spitzer Space Telescope and Herschel Space Observatory, covering wavelengths from 3.6$\\micron$ to 500$\\micron$, to produce well-resolved maps of the dust in two nearby galaxies. The present study is focused on two well-resolved spiral galaxies, NGC~628 and NGC~6946, as examples to illustrate the methodology. Subsequent work %\\citep{Aniano+Draine+KINGFISH_2012b} (Aniano et al.\\ 2012, in preparation) will extend this analysis to all 61 galaxies in the KINGFISH sample \\citep{Kennicutt+Calzetti+Aniano+etal_2011}. The paper is organized as follows. In \\S \\ref{sec:datasources} we discuss the data sources. Background subtraction and data processing are discussed in \\S \\ref{sec:dataproc}. The dust model is summarized in \\S \\ref{sec:dustmodel}, and the fitting procedure is described in \\S \\ref{sec:SEDFIT}. Results for NGC\\,628 and NGC\\,6946 are given in \\S \\ref{sec:results}. The sensitivity of the derived parameters to the set of cameras used to constrain the dust models is explored in \\S \\ref{sec:importance of SPIRE}, where we compare dust mass estimates obtained at high spatial resolution (without using MIPS160 or the longest-wavelength SPIRE bands) with estimates made at lower spatial resolution (using all the cameras available). We also investigate the reliability of the photometry by comparing MIPS70 and MIPS160 images with PACS70 and PACS160 images. In \\S \\ref{sec:maps vs global} we compare dust mass estimates based on spatially-resolved images with dust mass estimates based on global photometry, as would apply to distant, unresolved galaxies. The ``goodness of fit'' of different dust models is discussed in \\S 9. The principal results are discussed in \\S 10 and summarized in \\S \\ref{sec:summary}. Appendices A-D describe the method for image segmentation (i.e., galaxy and background recognition), background subtraction, estimation of photometric uncertainties in the images and estimation of dust modeling uncertainties. Appendix E is a comparison of PACS and MIPS photometry. %%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%% %%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%% %%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%% %%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%% %%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%% %%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%% %%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%% %%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%% %%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%% %%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%% %%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%% %%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%% ", "conclusions": "Multi-pixel and single-pixel best-fit model parameters at MIPS160 resolution.} \\begin{tabular}{|c|cc|cc|} \\hline & \\multicolumn{2}{|c|}{ NGC~628} & \\multicolumn{2}{|c|}{ NGC~6946}\\\\ Parameter & Resolved$^a$ & Global$^b$ & Resolved$^a$ & Global$^b$\\\\ \\hline $M_\\dust$ ($\\Msol$) & 2.87$\\times10^7$ & 2.33$\\times10^7$ & 6.74$\\times10^7$ & 6.62$\\times10^7$ \\\\ $L_\\dust $ ($\\Lsol$) & 6.83$\\times10^9$ & 7.08$\\times10^9$ & 3.31$\\times10^{10}$ & 3.20$\\times10^{10}$ \\\\ $\\Umin$ & 1.46 & 2.00 & 3.05 & 3.00 \\\\ $\\overline{U}$ & 1.75 & 2.23 & 3.61 & 3.55 \\\\ %$\\alpha$ & 1.78 & 1.70 & 1.75 & 1.70 \\\\ $100 \\times \\fpdr$ & 11.6 & 10.2 & 14.2 & 15.3 \\\\ $100 \\times \\qpah$ & 3.67 & 3.70 & 3.55 & 3.30 \\\\ %$100 \\times \\gamma$ & 13.66 & 0.055 & 3.51 & 0.081 \\\\ $100\\times M_\\dust/M_{\\rm H}$$^c$ & 0.82 & 0.66 & 0.63 & 0.61 \\\\ \\hline \\multicolumn{5}{l}{$^a$ Total or mean values are defined in equations (\\ref{eq:mean_mdust} - \\ref{eq:mean_fpdr}).}\\\\ \\multicolumn{5}{l}{$^b$ Values are from a model fit to the global photometry. }\\\\ \\multicolumn{5}{l}{$^c$ For $\\XCOxx=4$.}\\\\ \\end{tabular} \\end{center} \\end{table} %%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%% %%%%%%%%%%%%%%% TAB 04 %%%%%%%%%%%%%%% %%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%% %%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%% %%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%% %%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%% %%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%% %%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%% %%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%% %%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%% %%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%% %%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%% %%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%% %%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%% %%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%% Discussion} \\subsection{Dust Mass Estimation} Dust mass estimates are, of course, model-dependent. Above we have estimated the dust mass using a specific dust model, the DL07 silicate-graphite-PAH model, with an assumed parametric form for the starlight intensities heating the dust. The physical dust model and the ansatz for the distribution of starlight intensities together appear to successfully reproduce the observed SEDs, and to give reliable estimates of the dust masses. As discussed in \\S \\ref{sec:dustmodel}, the far-infrared opacity of the amorphous silicate originally put forward by DL84 was subsequently adjusted slightly \\citep{Li+Draine_2001b} to improve agreement with the far-infrared and submm emission observed for the local high-latitude dust. Thus the DL07 model uses dust properties that can reproduce the observed FIR-submm emission from the Milky Way (MW) cirrus with a single starlight heating intensity -- no ``cold dust'' is needed for the MW cirrus. In this dust model, the opacities are assumed to be independent of the dust temperatures. Here we see that the same dust model, with suitable adjustment of the starlight assumed to be heating the dust, is able to reproduce the emission from NGC\\,628 and NGC\\,6946 out to 500\\um, without introducing any ``cold dust'' component. The modeling performed at MIPS160 resolution, using all the IRAC, MIPS, PACS, and SPIRE cameras, in addition to being able to reproduce the observed SED, gives dust masses that are in line with the expected dust/H mass ratios for these galaxies: the dust/H mass ratio images in Figures \\ref{fig:ngc0628-2} and \\ref{fig:ngc6946-2}, at MIPS160 resolution, are smooth, and the global dust/H mass ratios are $0.0082\\pm0.0017$ and $0.0063\\pm0.0009$ for NGC~628 and NGC~6946, respectively. Observed depletions in the local Milky Way indicate a dust/H mass ratio $M_{\\rm dust}/M_{\\Ha} = 0.0091\\pm0.0006$ \\citep[][Table 23.1]{Draine_2011a}. Thus, a galaxy with a similar fraction of interstellar heavy elements in dust would be expected to have \\beq \\label{eq:Md/MH expected} \\frac{M_{\\rm dust}}{M_{\\Ha}} \\approx 0.0091 \\times 10^{[{\\rm O}]}~~~, \\eeq where $[{\\rm O}]\\equiv \\log_{10}[({\\rm O/H)}/({\\rm O/H})_\\odot]$. Thus galaxies with heavy element abundances (and depletions) similar to the local Milky Way should have $M_{\\rm dust}/M_\\Ha\\approx 0.009$. NGC~628 and NGC~6946 appear to be mature star-forming galaxies that would be expected to have interstellar metallicities similar to the local Milky Way. If so, then the values of $M_{\\rm dust}/M_{\\rm H}$ found by fitting a dust model to the observations appear to be in excellent agreement with expectations. We note that the H\\,II region oxygen abundances found by \\citet{Moustakas+Kennicutt+Tremonti+etal_2010} together with $(A_{\\rm O})_\\odot=12+\\log_{10}({\\rm O/H})_\\odot=8.73$ \\citep{Asplund+Grevesse+Sauval+Scott_2009} yield [O]$\\,=-0.46$ and $-0.36$ for NGC~628 and NGC~6946, and Eq.\\ (\\ref{eq:Md/MH expected}) would predict $M_\\dust/M_\\Ha=0.0032$ and 0.0040 for NGC~628 and NGC~6946, respectively. However, we suspect that the PT05 oxygen abundance estimates may be biased low: \\citet{Moustakas+Kennicutt+Tremonti+etal_2010} list PT05 HII-region oxygen abundances for 38 galaxies; the highest oxygen abundance is $A_{\\rm O}=8.59\\pm0.11$ for NGC~4826. It seems unlikely that none of the galaxies in their sample have oxygen abundances that are solar or supersolar. \\subsection{Dust Opacity in Molecular Gas and the Value of $\\XCO$} Both NGC~628 and NGC~6946 are rich in molecular gas, and the estimated gas mass depends on the adopted value of $\\XCO$. In the present study we have adopted $\\XCOxx=4$, which, as discussed in Section \\ref{sec:gas}, is significantly larger than the value $\\XCOxx\\approx2$ found for resolved CO clouds in the Milky Way. The larger value of $\\XCO$ adopted here may reflect the presence of so-called ``dark gas'', diffuse $\\HH$ with very low CO abundances, which does not radiate effectively in either \\ion{H}{1} 21-cm or CO$\\,J = 2 \\rightarrow1$ \\citep{Wolfire+Hollenbach+McKee_2010,Leroy+Bolatto+Gordon+etal_2011}. However, if the dust opacity in molecular regions is actually larger than the opacity in \\ion{H}{1} gas, then the present approach -- where we favor a value of $\\XCO$ that minimizes small-scale structure in maps of dust optical depth/H surface density -- will overestimate $\\XCO$. \\citet{Planck_molecular_clouds_2011}, using measurements of submm emission by Planck, and $\\NH$ inferred from NIR reddening of stars \\citep{Pineda+Goldsmith+Chapman+etal_2010}, conclude that the dust opacity per H nucleus in the Taurus molecular cloud is larger than in the local diffuse \\ion{H}{1} by a factor $R\\approx 2.0\\pm0.4$. Such an enhancement in the far-infrared and submm opacity might be a consequence of coagulation, which is expected to increase the far-infrared and submm opacity \\citep{Ossenkopf+Henning_1994,Stognienko+Henning+Ossenkopf_1995}. However, grain coagulation could also flatten the NIR extinction curve, so that the value of $\\NH$ inferred from the $(J-H)$ and $(H-K)$ stellar colors might be an underestimate. It should also be noted that studies of the Corona Austrina molecular cloud, using MIPS160/LABOCA870$\\micron$ ratios to determine the dust temperature, found a normal ratio of 870$\\micron$ optical depth to visual extinction \\citep{Juvela+Pelkonen+Porceddu_2009}. The dust model used here has been calibrated on dust in \\ion{H}{1} regions. If the actual dust opacity per H nucleon in molecular clouds is larger than in \\ion{H}{1} clouds by a factor $R$, then the actual value of $\\XCO$ in NGC~628 and NGC~6946 would be $\\XCOxx=4/R$. A value of $R\\approx2$ would then bring us into agreement with the $\\XCOxx\\approx2$ inferred from other estimators (virial mass estimates, $\\gamma$-ray emission) of molecular mass. \\subsection{Dust Mass Estimates from Single-Temperature Fits} \\citet{Skibba+Engelbracht+Dale+etal_2011} used a single dust temperature $T_\\dust$ with an assumed $\\kappa_\\nu\\propto\\lambda^{-1.5}$ opacity to fit the 70--500$\\micron$ photometry from MIPS and SPIRE. The opacity at $500\\micron$ was taken to be that of the DL07 dust model. They estimated $T_\\dust=24.0\\K$, $M_\\dust=10^{7.03\\pm0.08}\\Msol$ for NGC\\,628, and $T_\\dust=26.0\\K$, $M_\\dust=10^{7.47\\pm0.08}\\Msol$ for NGC\\,6946. The dust masses estimated by \\citet{Skibba+Engelbracht+Dale+etal_2011} for these two galaxies are smaller than the dust masses found here (see Table \\ref{tab:galaxy properties}) by factors of 3.3 and 3.6, respectively. This is the result of using a single dust temperature to try to reproduce emission from 70--500\\um. With the more realistic assumption of a distribution of dust temperatures, a small amount of warmer dust can provide much of the 70\\um emission, thus requiring an increased mass of ``normal'' temperature dust to account for the emission at $\\lambda\\gtsim 160\\micron$. A range of dust temperatures is of course expected from both spatial variations in the starlight intensity heating the dust, and the fact that a grain model with more than one dust composition, and a broad range of grain sizes, will have a distribution of temperatures even in a single radiation field. Single-temperature dust fits, if constrained by emission at 70\\um, will not provide a reliable estimate of the dust mass. In a separate work (Aniano \\& Draine 2012, in prep.) we discuss the bias introduced when the DL07 SED is approximated by a (single or dual) temperature blackbody multiplied by a power law opacity. \\subsection{Radial Gradients in Dust and Starlight} In both NGC~628 and NGC~6946, we find that the dust/H mass ratios outside the nucleus vary slowly, with little indication of a decrease in the dust/H ratio as one moves outward (see Figures \\ref{fig:ngc0628-1}l and \\ref{fig:ngc6946-1}l). In NGC~6946, however, the dust/H mass ratio appears to have a pronounced minimum at the center. We interpret this as due to overestimation of the gas mass in this region: we employ a single value of $\\XCOxx=4$ for this galaxy, but \\citet{Donovan_Meyer+Koda+Momose+etal_2011}, using virial mass estimates, find $\\XCOxx=1.2$ for the giant molecular clouds (GMCs) in the central 5 kpc. Because the molecular gas dominates near the center, our use of a higher value of $\\XCO$ implies an overestimate of the gas mass, resulting in an underestimate of the dust/H ratio. We therefore suspect that the central minimum in the dust/H ratio in Figure \\ref{fig:ngc6946-2} is entirely an artifact of using a value of $\\XCO$ that is too large for the central region. If $\\XCOxx=1.2$ is appropriate in the center of NGC~6946, the gas surface density will be lowered by a factor $\\sim 1.2/4$, and the dust/H mass ratio increased by a factor $\\sim3$, from the central value $\\sim0.004$ in Figure 7 to the expected value $\\sim 0.012$. The present study of the dust surface density therefore supports the finding by \\citet{Donovan_Meyer+Koda+Momose+etal_2011} of a low value of $\\XCO$ in the center of NGC~6946. In both NGC~628 and NGC~6946, the PAH abundance parameter $\\qpah$ appears to be quite uniform, with little evidence for a radial decline (see Figures \\ref{fig:ngc0628-2} and \\ref{fig:ngc6946-2}). Previous studies have shown that low-metallicity galaxies have low values of $\\qpah$ \\citep[e.g.,][]{Engelbracht+Gordon+Rieke+etal_2005, Draine+Dale+Bendo+etal_2007}; a galaxy with a negative radial gradient in metallicity might then show a radial decline in $\\qpah$. \\citet{Munoz-Mateos+Gil_de_Paz+Boissier+etal_2009} found radial declines in $\\qpah$ for a number of SINGS galaxies. For both NGC~628 and NGC~6946, $\\qpah$ appeared to be quite uniform out to a radius of $\\sim$$10\\kpc$, in qualitative agreement with our findings here. Future papers will examine radial trends in the KINGFISH sample. The starlight intensity parameter $\\Umin$ is presumed to characterize the average starlight intensity in the diffuse interstellar medium. The maps of $\\Umin$ (Figs.\\ \\ref{fig:ngc0628-2}ghi and \\ref{fig:ngc6946-2}ghi) have a peak $\\Umin\\approx 4$ and $6$ in the central regions of NGC~628 and NGC~6946, respectively, declining to $\\Umin\\approx 0.3$ near the edge of the galaxy mask ($\\sim$10\\,kpc from the center). It is gratifying that $\\Umin\\approx 1$ at $\\sim$8\\,kpc from the center, just as for our location in the Galaxy. The parameter $\\fpdr$ is the fraction of the dust heating that is produced by starlight intensities $U>100$. Averaged over the full galaxy, $\\langle f_\\PDR\\rangle =0.12$ and 0.14 for NGC~628 and NGC~6946, respectively, but both galaxies have hot spots where $f_\\PDR$ is much higher, with peak values of $\\sim$$0.30$ (see Figs.\\ \\ref{fig:ngc0628-2}jkl and \\ref{fig:ngc6946-2}jkl). With the $\\sim$635\\,pc resolution of the SPIRE250 camera at the 7.2 Mpc distance of NGC~628, these hot spots presumably correspond to regions of very active star formation. \\subsection{Interpretation of the Starlight Heating Parameters $\\alpha$ and $\\gamma$} %%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%% %%%%%%%%%%%%%%% FIGURE 19 %%%%%%%%%%%%%%% %%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%% \\renewcommand \\RoneCone {NGC6946_P160_100_SSS_000_Gal_Parameter_MW_dUm_pl_fitAl_Umm_0oo700_Alpha.eps} \\renewcommand \\RoneCtwo {NGC6946_S250_100_SSS_100_Gal_Parameter_MW_dUm_pl_fitAl_Umm_0oo070_Alpha.eps} \\renewcommand \\RoneCthree {NGC6946_M160_111_SSS_111_Gal_Parameter_MW_dUm_pl_fitAl_Umm_0oo010_Alpha.eps} \\renewcommand \\RtwoCone {NGC6946_P160_100_SSS_000_Gal_Parameter_MW_dUm_pl_fitAl_Umm_0oo700_Gamma.eps} \\renewcommand \\RtwoCtwo {NGC6946_S250_100_SSS_100_Gal_Parameter_MW_dUm_pl_fitAl_Umm_0oo070_Gamma.eps} \\renewcommand \\RtwoCthree {NGC6946_M160_111_SSS_111_Gal_Parameter_MW_dUm_pl_fitAl_Umm_0oo010_Gamma.eps} \\renewcommand \\RthreeCone {NGC0628_P160_100_SSS_000_Gal_Parameter_MW_dUm_pl_fitAl_Umm_0oo700_Alpha.eps} \\renewcommand \\RthreeCtwo {NGC0628_S250_100_SSS_100_Gal_Parameter_MW_dUm_pl_fitAl_Umm_0oo070_Alpha.eps} \\renewcommand \\RthreeCthree {NGC0628_M160_111_SSS_111_Gal_Parameter_MW_dUm_pl_fitAl_Umm_0oo010_Alpha.eps} \\renewcommand \\RfourCone {NGC0628_P160_100_SSS_000_Gal_Parameter_MW_dUm_pl_fitAl_Umm_0oo700_Gamma.eps} \\renewcommand \\RfourCtwo {NGC0628_S250_100_SSS_100_Gal_Parameter_MW_dUm_pl_fitAl_Umm_0oo070_Gamma.eps} \\renewcommand \\RfourCthree {NGC0628_M160_111_SSS_111_Gal_Parameter_MW_dUm_pl_fitAl_Umm_0oo010_Gamma.eps} \\ifthenelse{\\boolean{make_heavy}}{ } { \\renewcommand \\RoneCone {No_image.eps} \\renewcommand \\RtwoCone {No_image.eps} \\renewcommand \\RthreeCone {No_image.eps} \\renewcommand \\RfourCone {No_image.eps} \\renewcommand \\RoneCtwo {No_image.eps} \\renewcommand \\RtwoCtwo {No_image.eps} \\renewcommand \\RthreeCtwo {No_image.eps} \\renewcommand \\RfourCtwo {No_image.eps}} \\begin{figure} \\centering \\begin{tabular}{c@{$\\,$}c@{$\\,$}c} \\footnotesize PACS160 PSF & \\footnotesize SPIRE250 PSF & \\footnotesize MIPS160 PSF \\\\ \\FirstNormal \\SecondNormal \\ThirdNormal \\FourthLast \\end{tabular} \\vspace*{-0.5cm} \\caption{\\footnotesize\\label{fig:alpha-gamma} Characterization of the starlight power-law component, at the resolution of PACS160 (left), SPIRE250 (center), and MIPS160 (right). Row 1 and 3: power-law exponent $\\alpha$ map for NGC~628 and NGC~6946 respectively. Row 2 and 4: dust mass fraction in the power-law component $\\gamma$ map for NGC~628 and NGC~6946 respectively.} \\end{figure} %%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%% %%%%%%%%%%%%%%% FIGURE 19 %%%%%%%%%%%%%%% %%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%% As discussed in \\S \\ref{DustHeating}, a fraction $\\gamma$ of the dust mass is taken to be heated by a power-law distribution of starlight intensities between $\\Umin$ and $\\Umax$ with $dM/dU\\propto U^{-\\alpha}$, and the remaining fraction $(1-\\gamma)$ of the dust mass is heated by a radiation field with intensity $\\Umin$. Dust heated by this simple parameterization of the starlight intensities is quite successful in reproducing the observed SED (see the SEDs in Figures \\ref{fig:ngc0628-4}c,f; \\ref{fig:ngc0628-5}c,f,i,l; \\ref{fig:ngc6946-4}c,f; and \\ref{fig:ngc6946-5}c,f,i,l). \\citet{Dale+Helou+Contursi+etal_2001} discussed two ideal cases that have analytical predictions for the value of $\\alpha$: a single star in a homogeneous diffuse medium, and a dark cloud. The first case, with $U \\propto r^{-2}$, has $dM_{\\dust}/{dU} \\propto U^{-2.5}$, i.e., $\\alpha = 2.5$. In this case, $\\Umax$ would be very large, $\\gtsim 10^{10}$, the heating rate at which dust grains would vaporize. In the second case, a slab where the heating intensity is primarily attenuated by dust absorption, one would get $dM_{\\dust}/{dU} \\propto U^{-1}$, i.e., $\\alpha = 1$, and $\\Umax$ would be set by the starlight intensity at the edge of the slab. This case might correspond to a photodissociation region, the edge of a dark cloud. Different clouds in the pixel would have different values of $\\Umax$, so their superposition could be approximated as a power-law with a slightly larger value of $\\alpha$. We expect dust in a real galaxy to have $1 \\lesssim \\alpha \\lesssim 2.5$. Consider for the moment the fraction $\\gamma$ of dust with $dM/dU\\propto U^{-\\alpha}$. The power radiated is $dL\\,\\propto\\, U\\, dM\\, \\propto\\, U^{2-\\alpha}\\, d\\log U$. A value of $\\alpha = 2$ would have uniform dust luminosity per unit interval in $\\log U$. A value $\\alpha<2$ would concentrate the dust luminosity $L$ in the high radiation field regions $(U\\approx \\Umax)$, and a value $\\alpha >2$ would have $L$ dominated by dust with $U\\approx \\Umin$. Figure \\ref{fig:alpha-gamma} shows the starlight power-law component parameters $\\alpha$ and $\\gamma$. The power-law index $\\alpha$ has a preferred value $\\alpha \\approx 1.6$ in the bright regions. In these regions, the power-law distribution is representing the heating of dust in high-$U$ regions, e.g., photodissociation regions near OB stars. There are also regions with $\\alpha \\approx 2.4$. In these regions, most of the dust luminosity is concentrated in regions with $U \\approx \\Umin$, and the power law component is essentially making the $U=\\Umin$ component broader, i.e., is allowing for variations in the starlight intensities in the diffuse ISM. We do not attach great physical significance to the parameters $\\gamma$ and $\\alpha$. The more physically meaningful quantities are the mass-weighted mean starlight intensity $\\Ubar$, and $f_\\PDR$, the fraction of the dust luminosity that originates in regions with $U>10^2$. In Figures \\ref{fig:ngc0628-2} and \\ref{fig:ngc6946-2}, one sees that $f_\\PDR > 0.05$ over most of the galaxy mask for both galaxies. %%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%% %%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%% %%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%% %%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%% %%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%% %%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%% %%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%% %%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%% %%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%% %%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%% %%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%% %%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%% Summary} Using Spitzer and Herschel data, we perform a resolved study of the dust physical parameters, and of the starlight heating the dust, in the nearby galaxies NGC~628 and NGC~6946. We employ the DL07 dust model, having amorphous silicate grains and carbonaceous grains, including PAHs (see \\S \\ref{sec:dustmodel}). The model has a distribution of grain sizes, and we allow for a distribution of starlight intensities heating the dust within each pixel. The model assumes a frequency-dependent opacity that reproduces the observed emission spectrum of high-latitude dust. In order to perform resolved studies of the galaxies, it is important to convolve all the images into a common PSF. This is achieved using the convolution kernels generated by \\citet{Aniano+Draine+Gordon+Sandstrom_2011}. We perform our dust modeling using all appropriate combinations of PSFs and cameras. Table \\ref{tab:resolutions} lists some of the resolutions and cameras used in this work. Our principal findings are as follows: \\begin{enumerate} \\item The dust model is quite successful in reproducing the observed SED over the full wavelength range 6--500$\\micron$ where dust emission dominates (see the SEDs in Figs.\\ \\ref{fig:ngc0628-4}c,f; \\ref{fig:ngc0628-5}c,f,i,l; \\ref{fig:ngc6946-4}c,f; and \\ref{fig:ngc6946-5}c,f,i,l; and the model/observed maps in Figs.\\ \\ref{fig:ngc0628-3}c,f and \\ref{fig:ngc6946-3}c,f). The dust model reproduces the emission out to 500$\\micron$ without introduction of a ``cold dust'' component, and with no allowance for possible temperature-dependent dust opacities. The DL07 dust opacities therefore appear to be consistent with both the observed emission from local ``cirrus'' at high galactic latitudes, and the observed SED from 500\\,pc-sized regions of the ISM in NGC~628 and 6946. There is no indication that $T$-dependent opacities are needed to reproduce these data. \\item Maps of the dust/H mass ratio show it to be relatively uniform over both galaxies, outside of the nucleus (see Figs.\\ \\ref{fig:ngc0628-1}l and \\ref{fig:ngc6946-1}l). The derived dust/H mass ratio for these galaxies is consistent with that expected if the interstellar abundances in NGC~628 and NGC~6946 are close to solar. As near-solar abundances seem likely, this is strong support for the quantitative accuracy of the dust masses obtained by modeling the IR and FIR emission. \\item Figure \\ref{fig:dustmass} shows how dust mass estimates depend on the camera and PSF used. The ``gold standard'' is taken to be dust modeling done at MIPS160 resolution, using all cameras, with Scanamorphos \\citep{Roussel_2012} processing of the PACS data. Relative to this standard, mass estimates done with smaller pixels (giving up the lowest resolution cameras) tend to be biased slightly high. At SPIRE250 resolution, the bias is $\\sim$$38\\%$. Working at SPIRE350 resolution provides a good compromise between resolution and accuracy, with $\\sim$$30\\%$ errors for the total dust mass estimated for NGC~628 and NGC~6946. \\item Even though the dust modeling process is non-linear, resolved modeling is compatible with modeling using the global photometry of the galaxy. Differences in the inferred parameters are in most cases under 20\\%. \\item In NGC~6946 the dust/H mass ratio calculated with a single value of $\\XCOxx=4$ appears to have a minimum at the center (see Fig.\\ \\ref{fig:ngc6946-2}a,b,c). This minimum is interpreted as an artifact due to overestimation of the molecular mass. The dust surface densities found here therefore support the finding by \\citet{Donovan_Meyer+Koda+Momose+etal_2011} of $\\XCOxx\\approx 1.2$ in the center of NGC~6946. \\item The present fitting procedures employ six adjustable parameters for each pixel: $\\Omega_\\star$, $\\Mdust$, $\\qpah$, $\\Umin$, $\\alpha$, and $\\gamma$ (we fix $\\Umax=10^7$). For NGC~628 and NGC~6946 we find that this approach provides substantially better fits than the approach advocated by \\citet{Galliano+Hony+Bernard+etal_2011} who treat $\\Omega_\\star$, $\\Mdust$, $\\qpah$, $\\Umin$, $\\alpha$, and $\\Umax$, as adjustable parameters and fix $\\gamma=1$: compare the $\\chi^2$ maps in Fig.\\ \\ref{fig:chicomp}a with \\ref{fig:chicomp}b, and \\ref{fig:chicomp}d with \\ref{fig:chicomp}e. \\item The starlight heating the dust within a single pixel is characterized by three adjustable parameters: $\\Umin$, $\\gamma$, and $\\alpha$ (see eq. \\ref{eq:dMd/dU} and \\ref{eq:dMd/dU1}). $\\Umin$ is interpreted as the intensity of the diffuse starlight, responsible for the bulk of the dust heating. Maps of $\\Umin$ (see Figs.\\ 2 and 7) show significant structure: there is a tendency of $\\Umin$ to be higher in spiral arms, as well as a significant radial decline in $\\Umin$. The overall mass-weighted mean value of $\\Umin=1.5$ for NGC~628 and $\\Umin=3.1$ for NGC~6946, but $\\Umin$ declines to values of 0.3 or lower in the outer regions of both galaxies. \\item The parameter $\\fpdr$ is the fraction of the dust heating power that is contributed by starlight intensities $U > 100$. The overall value is $\\langle \\fpdr\\rangle =0.12$ for NGC~628 and $\\langle \\fpdr\\rangle =0.14$ for NGC~6946, but in both galaxies $\\fpdr$ peaks at local maxima of the dust surface brightness, which are associated with star-forming regions. At SPIRE250 resolution, $\\fpdr$ reaches values as high as 0.25 in both galaxies. \\item Current PACS photometry disagrees with MIPS photometry (see Fig.\\ \\ref{fig:color_0628} and \\ref{fig:color_6946}). PACS70 photometry is up to 80\\% higher than MIPS70, and PACS160 is up to 50\\% higher than MIPS70 photometry in the bright areas of the galaxies. \\item The PACS-MIPS photometry disagreement induces a bias when trying to model dust at high spatial resolutions (PACS160 PSF), where MIPS70 and MIPS160 cannot be used. We do not recommend modeling dust at PACS160 resolution in low surface brightness areas. Modeling done at SPIRE250 resolution is more reliable than modeling at PACS160, but the inferred dust masses still disagree with our ``gold standard'' (i.e., modeling done at MIPS160 resolution using all the IRAC, MIPS, PACS and SPIRE cameras) by up to $\\approx 30\\%$. \\end{enumerate}" }, "1207/1207.5074_arXiv.txt": { "abstract": "We have conducted a kinematic study of 165 young M dwarfs with ages of $\\lesssim$300 Myr. Our sample is composed of stars and brown dwarfs with spectral types ranging from K7 to L0, detected by \\rosat\\ and with photometric distances of $\\lesssim$25 pc assuming the stars are single and on the main-sequence. In order to find stars kinematically linked to known young moving groups (YMGs), we measured radial velocities for the complete sample with Keck and CFHT optical spectroscopy and trigonometric parallaxes for 75 of the M dwarfs with the CAPSCam instrument on the du Pont 2.5-m Telescope. Due to their youthful overluminosity and unresolved binarity, the original photometric distances for our sample underestimated the distances by 70\\% on average, excluding two extremely young ($\\lesssim$3 Myr) objects found to have distances beyond a few hundred parsecs. We searched for kinematic matches to 14 reported YMGs and identified 9 new members of the AB~Dor YMG and 2 of the Ursa Majoris group. Additional possible candidates include 6 Castor, 4 Ursa Majoris, 2 AB Dor members, and 1 member each of the Her-Lyr and $\\beta$~Pic groups. Our sample also contains 27 young low-mass stars and 4 brown dwarfs with ages $\\lesssim$150 Myr which are not associated with any known YMG. We identified an additional 15 stars which are kinematic matches to one of the YMGs, but the ages from spectroscopic diagnostics and/or the positions on the sky do not match. These warn against grouping stars together based only on kinematics and that a confluence of evidence is required to claim that a group of stars originated from the same star-forming event. ", "introduction": "Planet-formation studies require a range of stellar host mass and age in order to test models of planet and disk evolution. The vast majority of confirmed exoplanets are in systems greater than 1 Gyr old, and efforts to find the youngest planets, those just forming in their disks around T Tauri stars (1 -- 3 Myr old), are just beginning \\citep{prat08,croc11,krau11}. Young stars with ages spanning 10 to 100 Myr fill a particularly interesting, and relatively unexplored, gap as this time scale coincides with the end of giant planet formation (e.g.~\\citealt{boss11}) and the onset of active terrestrial planet formation. From an observational perspective, this time range is critical for direct imaging searches of planets and disks (e.g.~\\citealt{maro08,liu10}). Planets cool and fade dramatically during this period, e.g.~the luminosity of a 5-M$_{Jup}$ planet drops by 2 -- 3 orders of magnitude in this time span and its effective temperature decreases from 1300 K to 400 K (\\citealt{bara03}). Also in this time period, the fraction of debris disks around AFGK main-sequence stars decreases significantly (e.g.~\\citealt{riek05,hill08}). These attributes have made young, nearby stars prime targets at direct imaging searches for exoplanets and circumstellar disks, as supported by the imaged planets around the 10 -- 30-Myr stars $\\beta$ Pic and HR 8799 \\citep{lagr10,maro08}. M dwarfs offer an additional observational advantage. Their intrinsic faintness provides a more favorable contrast in star-planet flux, aiding the detection of faint planetary-mass companions. M~dwarfs also dominate the stellar mass function by number$\\colon$ roughly 3 out of 4 stars in a volume-limited sample of the solar neighborhood are M~dwarfs (\\citealt{reid95,boch10}). So in principle low-mass stars could represent the most common and nearest hosts of planetary systems, providing a much larger population of targets for direct imaging of exoplanets than has been studied to date. Searches for young low-mass stars have been carried out using X-ray (e.g.~\\citealt{jeff95,webb99,mont01,shko09b}, SLR09 hereafter) and UV (\\citealt{shko11,rodr11}) surveys to identify strong coronal and chromospheric emission. SLR09 and \\cite{shko11} estimated that M dwarfs with fractional X-ray and UV luminosities of $log(F_X/F_J) > -2.5$ and $log(F_{NUV}/F_J) > -4$ are younger than $\\sim$300 Myr, based on comparisons with Pleiades (120~Myr) and Hyades (625 Myr) members. However, determining a more accurate age for an individual star is more difficult. Over the past decade, many dispersed young stars have been kinematically linked to coeval moving groups (e.g., \\citealt{zuck04,torr08}, and references therein), for which more accurate ages are available through comparison of the bulk group properties with stellar isochrones and lithium-depletion models (e.g.~\\citealt{sode10}). Thus, pinning an X-ray- or UV- bright M dwarf to one of the known young moving groups (YMGs) provides a more accurate means to determine its age. The current census of YMG members is mostly limited to AFGK-type stars due to the past reliance on optical catalogs for distances and proper motions, e.g., the Hipparcos and Tycho catalogs \\citep{perr97}, which exclude most of the fainter nearby M dwarfs. If the initial mass function of YMGs follows the field, we can expect that there are many unidentified low-mass members of the known YMGs, e.g.~$\\approx$60 M dwarfs missing from the $\\beta$~Pic YMG and $\\approx$20 mid-Ms from the TW Hydrae Association \\citep{shko11}. Recent searches for these low-mass members using large photometric and proper motion catalogs have found many candidates awaiting confirmation with follow-up parallaxes and radial velocities (e.g., \\citealt{lepi09,schl12}). This paper presents the kinematic analysis of the 165 young M dwarfs first described in SLR09,\\footnote{There are 10 additional stars in this paper compared to SLR09, including 8 common proper-motion companions found during the CAPSCam astrometric observations and 2 young M dwarfs observed in support of the Gemini NICI Planet-Finding Campaign \\citep{liu10}.} including radial velocities (RVs) for all the targets, trigonometric parallaxes for half the sample, and three-dimensional space motions with the goal of providing more accurate ages by linking these stars to known YMGs. ", "conclusions": "\\label{summary} We present a kinematic analysis of the 165 young ($\\lesssim$300 Myr), nearby and X-ray-bright M dwarfs compiled by SLR09. We have further characterized this ``25-pc sample'' with radial velocities and distances, showing that nearly half are outside of 25 pc as expected if they are indeed young. We acquired distances for half the sample from trigonometric parallaxes while we estimated photometric distances for the remainder of the sample using models and correcting for visual binarity and youth determined from spectroscopic age indicators. Our astrometric survey showed that for this sample of young stars, photometric distances underestimate the true distances by an average of roughly 70\\%, and range from 5 to 500 pc. This implies that the $UVW$s of young stars based on photometric distances that assume the stars lie on the main-sequence are most likely unreliable. By combining our RVs and distances with proper motions, we calculated the targets' $UVW$ velocities in search of new YMG members. We identified 21 likely members of YMGs, which have accurate $UVW$s determined from trigonometric parallaxes, positional coincidence with known members, and age agreement using the independent spectroscopic and X-ray age diagnostics presented in SLR09. Of these, 10 were previously known AB Dor, Hyades, or $\\beta$~Pic members. Newly proposed members include 9 AB Dor members and 2 UMa members. Additional possible candidates include 6 Castor members, 4 Ursa Majoris members, 2 AB Dor members, and 1 member each of the Her-Lyr and $\\beta$~Pic groups. In addition, we identified three stars (two of which form a binary pair) with identical kinematics and age ranges, potentially forming their own young association. We also found kinematic matches to YMGs for 40 stars, but either their positions on the sky or ages do not agree with previously known members. For 12 stars that are kinematic and age matches but are positionally offset to the other group members, chemical analysis would be one way to confirm that the stars did indeed originate from a common star-forming event. However, metallicity indicators of M dwarfs may not yet be accurate or precise enough to discriminate group members from the field population (e.g.~\\citealt{neve11} and references therein). Lastly, our sample also contains 31 young M dwarfs, including 4 brown dwarfs, with ages $\\lesssim$150 Myr, which are not associated with any known YMG. Even relaxing the requirements of spectroscopic and positional age agreement, we did not recover the expected numbers of M dwarf members to the YMGs. After accounting for our observing limitations, the likeliest explanation is that the \\rosat\\ All-Sky Survey is not sensitive enough to detect YMG members lying beyond 25--30 pc. In order to understand this further, as well as identify even more members, we are completing a larger survey of young M dwarfs with photometric distances out to 100 pc using the more sensitive \\galex\\ archive." }, "1207/1207.2148_arXiv.txt": { "abstract": "{The X--ray source 1RXS J180431.1$-$273932 has been proposed as a new member of the symbiotic X--ray binary (SyXB) class of systems, which are composed of a late-type giant that loses matter to an extremely compact object, most likely a neutron star. In this paper, we present an optical campaign of imaging plus spectroscopy on selected candidate counterparts of this object. We also reanalyzed the available archival X--ray data collected with {\\it XMM-Newton}. We find that the brightest optical source inside the 90\\% X--ray positional error circle is spectroscopically identified as a magnetic cataclysmic variable (CV), most likely of intermediate polar type, through the detection of prominent Balmer, He {\\sc i}, He {\\sc ii}, and Bowen blend emissions. On either spectroscopic or statistical grounds, we discard as counterparts of the X--ray source the other optical objects in the {\\it XMM-Newton} error circle. A red giant star of spectral type M5\\,III is found lying just outside the X--ray position: we consider this latter object as a fore-/background one and likewise rule it out as a counterpart of 1RXS J180431.1$-$273932. The description of the X--ray spectrum of the source using a bremsstrahlung plus black-body model gives temperatures of $kT_{\\rm br} \\sim$ 40 keV and $kT_{\\rm bb} \\sim$ 0.1 keV for these two components. We estimate a distance of $d \\sim$ 450 pc and a 0.2--10 keV X--ray luminosity of L$_{\\rm X} \\sim$ 1.7$\\times$10$^{32}$ erg s$^{-1}$ for this system and, using the information obtained from the X--ray spectral analysis, a mass $M_{\\rm WD} \\sim$ 0.8 $M_\\odot$ for the accreting white dwarf (WD). We also confirm an X--ray periodicity of 494 s for this source, which we interpret as the spin period of the WD. In summary, 1RXS J180431.1$-$273932 is identified as a magnetic CV and its SyXB nature is excluded.} ", "introduction": "Symbiotic X--ray binaries (SyXBs; see e.g. Masetti et al. 2006a) form a minor class of low mass X--ray binaries in which the compact accretor, most likely a neutron star (NS), receives matter from a red giant rather than from a late-type companion star on the main sequence (or possibly slightly evolved) and with mass of generally $\\la$1 $M_\\odot$. These objects are defined SyXBs by analogy with symbiotic binary systems, which are formed by an evolved late-type star and a white dwarf (WD). There are currently only seven confirmed objects of this type known in the Galaxy: six cases listed in Masetti et al. (2007), Nespoli et al. (2010), and references therein, to which a newly-identified one, XTE J1743$-$363, was recently added (Smith et al. 2012). It is therefore equally important to investigate possible new candidates (cf. Masetti et al. 2011) and to spectroscopically confirm the known candidates. In addressing the latter issue, Masetti et al. (2012) found by using optical spectroscopy that the SyXB candidate 2XMM J174016.0$-$290337 (also known as AX J1740.2$-$2903) proposed by Farrell et al. (2010) is actually a cataclysmic variable (CV) of dwarf nova type. Likewise, the availability of (sub)arcsec-sized X--ray positions allows one to determine possible optical counterpart misidentifications, especially in extremely crowded fields. This occurred in the case of IGR J16393$-$4643, for which a {\\it Chandra} snapshot (Bodaghee et al. 2012) pinpointed the correct near-infrared counterpart and dismissed the one proposed by Nespoli et al. (2010) as a SyXB. With the aim of confirming (or disproving) the nature of yet another SyXB candidate, we performed an optical imaging and spectroscopic campaign on two possible counterparts of the X--ray source 1RXS J180431.1$-$273932 (Nucita et al. 2007); we also took this opportunity to reanalyze the X--ray data presented by those authors. The X--ray object 1RXS J180431.1$-$273932, first detected in the {\\it ROSAT} bright source survey (Voges et al. 1999), was subsequently observed on October 2005 with {\\it XMM-Newton}. The main results of this observation, reported by Nucita et al. (2007), are: (i) the detection of an X--ray period of 494 s, most likely due to the spin of the compact accretor; (ii) the description of its X--ray spectrum in the 0.2-7 keV range in the form of a power-law with index $\\Gamma \\sim$ 1 plus a Gaussian emission line at $\\sim$6.6 keV; and (iii) the detection, with the Optical Monitor (OM) onboard {\\it XMM-Newton}, of an object with magnitude $v \\sim$ 17.2 at a position consistent with the $\\sim$2$''$-radius (1$\\sigma$, corresponding to 3$\\farcs$3 at the 90\\% confidence level) X--ray error circle of the source. Concerning the last point, Nucita et al. (2007) found that the OGLE catalog (Wray et al. 2004) reports a red optical object at $\\sim$5$''$ from the X--ray position of the source that has a periodicity of about 20.5 days in its $I$-band light curve. On the basis of the optical and near-infrared magnitudes of this object (assuming that the OM and OGLE sources are one and the same), Nucita et al. (2007) concluded that its colors are compatible with those of a red giant star of type M6\\,III, thus making 1RXS J180431.1$-$273932 a viable SyXB candidate. However, the non-negligible (albeit small) difference in the positions of the X--ray and the OGLE objects, together with the lack of optical spectroscopy for the latter, calls for an in-depth investigation of the properties of this source in the optical bands. In addition, the location of the source towards the Galactic center ($l$ = 3$\\fdg$2; $b$ = $-$2$\\fdg$9) suggests that the field crowdedness might produce source confusion when the positional uncertainty of an object is as large as a few arcsec. We therefore started an optical imaging and spectroscopic campaign to clarify the nature of 1RXS J180431.1$-$273932 using the Italian Telescopio Nazionale Galileo. We also decided to reanalyze here the {\\it XMM-Newton} data first reported in Nucita et al. (2007) using updated software and response matrices and a more physical model to describe the X--ray spectrum. The outline of the present paper is as follows: in Sect. 2, we describe our optical and X--ray observations, while Sect. 3 reports our results and Sect. 4 discusses them. Finally, in Sect. 5 we present our conclusions for this source. \\begin{figure}[!t] \\vspace{-3.6cm} \\hspace{-3.1cm} \\psfig{figure=19334f1.ps,width=15cm,angle=0} \\vspace{-2cm} \\caption[]{{\\it Upper panel:} TNG+DOLORES white (open) filter image of the field of 1RXS J180431.1$-$273932 with a superimposed 3$\\farcs$3-radius 90\\% confidence level {\\it XMM-Newton} X--ray error circle. The field size is about 3$'$$\\times$3$'$. {\\it Lower panel}: zoomed image of a 30$''$$\\times$30$''$ box centered on the {\\it XMM-Newton} position. The actual counterpart (see text) is the object readily recognizable within the circle; the red giant star mentioned by Nucita et al. (2007) is the bright source just outside the circle on the right. In both panels, north is at top and east is to the left.} \\end{figure} ", "conclusions": "Our multiwavelength optical/X--ray study of 1RXS J180431.1$-$273932 has allowed us to identify its actual counterpart and pinpoint its real nature. We have found that this object is a magnetic CV, most likely of IP type; the SyXB hypothesis, put forward by Nucita et al. (2007), is thus ruled out. This misidentification was likely induced by the presence of a red giant lying along the line of sight and just outside the border of the X--ray error box. We also exclude any connection of the X--ray source with other, fainter optical objects within its positional uncertainty obtained from {\\it XMM-Newton} data. We have confirmed the X--ray periodicity of 494 s first detected by Nucita et al. (2007) and interpreted it as the spin period of the accreting WD hosted in this system. We could also successfully model the X--ray spectrum of 1RXS J180431.1$-$273932 using a bremsstrahlung plus black-body emission model, as typically found in magnetic CVs. We encourage the acquisition of follow-up observations in the optical and X--rays to determine the orbital period and the main physical characteristics of this CV, to also confirm the IP nature proposed here. This research moreover stresses that it is of paramount importance to have very precise (smaller than a few arcsec) X--ray localizations, especially in cases of crowded fields and particularly for objects concentrated towards the Galactic bulge, in order to determine the optical counterpart." }, "1207/1207.7071_arXiv.txt": { "abstract": "{Although it has important ramifications for both the formation of star clusters and their subsequent dynamical evolution, rotation remains a largely unexplored characteristic of young star clusters (few Myr). Using multi-epoch spectroscopic data of the inner regions of 30~Doradus in the Large Magellanic Cloud (LMC) obtained as part of the VLT-FLAMES Tarantula Survey, we search for rotation of the young massive cluster R136. From the radial velocities of 36 apparently single O-type stars within a projected radius of 10\\,pc from the centre of the cluster, we find evidence, at the 95\\% confidence level, for rotation of the cluster as a whole. We use a maximum likelihood method to fit simple rotation curves to our data and find a typical rotational velocity of $\\sim$3\\,$\\kms$. When compared to the low velocity dispersion of R136, our result suggests that star clusters may form with at least $\\sim$20\\% of the kinetic energy in rotation.} ", "introduction": "Despite their spherical shape, Milky Way globular clusters (GCs) rotate with amplitudes up to half the 1D velocity dispersion \\citep[$0\\lesssim\\vrotsigma\\lesssim0.5$, e.g.][]{1997A&ARv...8....1M}, so their amount of rotational energy is typically not dominant but also not negligible. Numerical studies have shown that this rotation has an important influence on star clusters by accelerating their dynamical evolution, for example by speeding up the collapse of the core through the gravogyro instability or by significantly increasing the escape rate for clusters in a strong tidal field \\citep[e.g.][]{1999MNRAS.302...81E, 2002MNRAS.334..310K, ernst2007}. Most of the rotation signatures are found through radial velocity (RV) studies, but rotation has also been confirmed in the plane of the sky for $\\omega$ Centauri and 47 Tucanae \\citep[][respectively]{2000A&A...360..472V, 2010ApJ...710.1032A}. RV studies are now able to measure rotational amplitudes in GCs below 1 $\\kms$ and despite these precise measurements rotation has not been detected in some clusters \\citep[e.g.][]{2010MNRAS.406.2732L}, although this could also be an inclination effect. It is unclear what the origin of the rotation is in some of these old clusters. It could be the result of the merging of two clusters \\citep{2003ApJ...589L..25B} or imprinted during the formation process. Observational input is now getting sufficiently abundant to look for correlations between rotational amplitude and other cluster properties. \\citet{2012A&A...538A..18B} report a correlation between horizontal branch (HB) morphology and $\\vrotsigma$ in a sample of 20 globular clusters. Given that metallicity is the first parameter determining HB morphology, this in turn suggests a correlation between $\\vrotsigma$ and metallicity, such that clusters with higher metallicity have greater fractions of energy in rotation. Since a higher metallicity in a gas implies a higher efficiency of energy dissipation by atomic transitions, this then hints at a significant role of dissipation in the process of cluster formation \\citep[e.g.][]{2010ApJ...724L..99B}. Rotation may therefore be intimately linked to the formation of clusters. Little is known about rotation in young clusters, partially because it is very challenging to measure accurately $\\sigma_{\\rm 1D}$ given the high multiplicity fraction of massive stars \\citep[e.g.][]{vhb2012a}. \\citet{1982MNRAS.199..565F} found an age-ellipticity relation for clusters in the Large Magellanic Cloud (LMC) with older clusters presenting less elongated shapes, which was interpreted as internal evolution erasing any asymmetry stemming from the violent relaxation process of the formation \\citep[although see][]{goodwin1997}. However, rotation and ellipticity are not necessarily equivalent. Ellipticity can be due to rotation \\citep[e.g. $\\omega$ Cen;][]{meylan1986} but also to velocity anisotropy \\citep{stephens2006, Henon1973}, and rotating clusters can be spherical \\citep{LB60, meza2002}. Marginal evidence for rotation was found for the young (few 100 Myrs) Galactic cluster GLIMPSE-C01 with an amplitude of $\\vrotsigma\\simeq0.2$ \\citep{2011MNRAS.411.1386D}. A rotational signal in the RVs was also detected in the $\\sim$100 Myr cluster NGC\\,1866 \\citep{fischer1992} and in the $\\sim$50\\,Myr binary cluster NGC\\,1850 \\citep{fischer1993}, both in the LMC. To really confirm whether clusters form with a significant amount of angular momentum, we need to look for an even younger cluster. The young massive cluster (YMC) R136 in the 30 Doradus star forming region in the LMC is an ideal target to establish this. With an estimated mass of about $10^5\\,\\msun$ \\citep{2009ApJ...707.1347A} and its sub solar metallicity, it may at some stage resemble a typical metal-rich GC as we find them in the Milky Way Bulge. With an age of less than 2\\,Myr \\citep{koterheap, massey1998, Crowther2010}, it is so young that any rotation needs to be attributed to the formation process, be it from merging of sub-clusters or directly from the angular momentum of the progenitor cloud. A rough estimate of the half-mass relaxation time ($t_{\\rm rh}$) of R136 can be obtained by assuming $N=10^5$ stars and a half-mass radius of 2.27\\,pc, which is found from multiplying the half-light radius of 1.7\\,pc \\citep[e.g.][]{vhb2012a} by 4/3 \\citep{spitzer:1987}. Following the formula of \\citet{spitzerhart1971}, we obtain $t_{\\rm rh}\\simeq366$\\,Myr, so relaxation would not have had time to erase the original signature of rotation. Here we report on evidence for rotation of R136 deduced from RV measurements of massive stars obtained as part of the VLT-FLAMES Tarantula Survey \\citep[VFTS;][]{evans:2011}. We briefly present the data in Sect.~\\ref{data} and describe our analysis of the rotational signature in Sect.~\\ref{analysis}. We discuss the implications of the rotation of R136 for cluster evolution in Sect.~\\ref{discuss}, and present our conclusions in Sect.~\\ref{conc}. \\begin{figure}[!h] \\centering \\includegraphics[width=8.5cm]{rv_map_new.eps} \\caption{Illustration of the positions and RVs of the stars considered in this study. Symbol sizes denote the magnitude of the stellar velocities with respect to the average cluster velocity. The solid, dotted, and dashed lines correspond to the optimal rotation axis determined for models with a constant rotational velocity, a constant rotation rate, and a more realistic rotation curve (see Sect. \\ref{analysis}), respectively.} \\label{rot_circles} \\end{figure} ", "conclusions": "} We presented evidence that the young massive cluster R136 is rotating with a rotational velocity amplitude of about 3\\,km\\,s$^{-1}$, which implies that at least $\\sim$20\\% of its total kinetic energy is in rotation. Obviously, RV measurements of more stars in this cluster would be desirable to better populate the rotation curve and confirm the rotational signal with a confidence level higher than the current 95\\%. Given the young age of R136, our results suggests that star clusters may form with a significant amount of angular momentum. This will place useful constraints on models of cluster formation. We finally argued that the rotation of globular clusters could originate from their formation, but this is clearly a topic where more detailed numerical investigations are welcome." }, "1207/1207.0651_arXiv.txt": { "abstract": "{The study of stellar multiple systems provides us with important information about the stellar formation processes and can help us to estimate the multiplicity fraction in the Galaxy. 65~UMa belongs to a rather small group of stellar systems of higher multiplicity, whose inner and outer orbits are well-known. This allows us to study the long-term stability and evolution of the orbits in these systems.} {We obtained new photometric and spectroscopic data that when combined with interferometric data enables us to analyze the system 65~UMa and determine its basic physical properties.} {We perform a combined analysis of the light and radial velocity curves, as well as the period variation by studying the times of the minima and the interferometric orbit. A disentangling technique is used to perform the spectra decomposition. This combined approach allows us to study the long-term period changes in the system for the first time, identifying the period variation due to the motion on the visual orbit, in addition to some short-term modulation.} {We find that the system contains one more component, hence we tread it as a sextuple hierarchical system. The most inner pair of components consists of an eclipsing binary orbiting around a barycenter on a circular orbit, both components being almost identical of spectral type about A7. This pair orbits on an eccentric orbit around a barycenter, and the third component orbits with a period of about 640~days. This motion is reflected in the period variation in the minima times of the eclipsing pair, as well as in the radial velocities of the primary, secondary, and tertiary components. Moreover, this system orbits around a barycenter with the distant component resolved interferometrically, whose period is of about 118~years. Two more distant components (4\\arcs and 63\\arcs) are also probably gravitationally bound to the system. The nodal period of the eclipsing-pair orbit is on the order of only a few centuries, which makes this system even more interesting for a future prospective detection of changing the depths of minima.} {We identify a unique solution of the system 65~UMa, decomposing the individual components and even shifting the system to higher multiplicity. The study of this kind of multiple can help us to understand the origin of stellar systems. Besides 65~UMa, only another 11 sextuple systems have been studied. } ", "introduction": "As members of more complex multiple systems, the eclipsing binaries can provide us important information about their physical properties, as derived from different methods. This is the case for 65~UMa, a system whose the close components form an eclipsing binary, and additional components found to be gravitationally bounded to this pair \\citep{2004A&A...424..727P}. Thanks to the combined analysis, we have been able to derive the radii, masses, and evolutionary statuses of the close components, in addition to some properties of the distant ones. These systems are still very rare and mostly lie relatively close to the solar system. Only 39 such systems are known where a close eclipsing binary is a member of a wide visual binary and we know both orbits, their mutual inclinations, ratio of periods, etc. For instance, the ratio of periods can tell us something about the long-term stability of the system. These unique systems are the most suitable for studies of dynamical effects, the short and long-term evolution of the orbits, etc. (see e.g. \\citealt{1975A&A....42..229S}). The study of systems of higher multiplicity is still relatively undeveloped yet, and can provide insight into their formation. Moreover, \\cite{2005A&A...439..565G} found that the majority of the early-type stars are found in multiple systems. Star-forming theories are still based on many ad hoc assumptions and the physical characteristics of the multiple systems can provide strong constraints on some of them. These can be e.g. the mass ratios of the inner and outer pairs, the ratio of periods, and inclinations, see for instance (\\citealt{2007prpl.conf..133G} and \\citealt{2008MNRAS.389..925T}). In addition, the multiplicity fraction is one of the most crucial parameters in theoretical models and nowadays we know of only 20 quintuples, 11 sextuples, and 2 septuple systems \\citep{2008MNRAS.389..869E}. ", "conclusions": "\\label{Discussion} We have performed our first attempt to perform a detailed combined solution of all available data for 65~UMa, namely photometry, spectroscopy, and interferometry, obtaining quite a reliable picture of this unusual sextuple hierarchical system (see Fig. \\ref{Picture}). \\begin{figure} \\centering \\includegraphics[width=89mm]{TABUL.eps} \\caption{Schematic structure of the whole system 65~UMa.} \\label{Picture} \\end{figure} The inner close eclipsing pair consists of two almost-identical stars of A7 spectral type. This finding is consistent with the photometric indices $B-V$, $V-R$, and $R-I$ being constant for the whole phase of the eclipsing binary at a level of 0.005~mag. The stars move on circular orbits with periods of about 1.73043~d, both being located on the main sequence. Thanks to the combined analysis, we were also able to compute its distance as $d = 234 \\pm 29$~pc, independently of the Hipparcos satellite data. The 640 day orbit was confirmed by both the minima times and the RV variations. Applying the spectral disentangling and rough estimation of the mass ratios from this analysis, one can estimate the masses of the outer components. The Ab component is probably of A1 spectral type and has a mass of about 2.4~$\\pm$~0.4~\\Mo, from the total mass of the 118 yr visual orbit, we can estimate the mass of the fourth component (in WDS named B), to be about 2.4~$\\pm$~2.0~\\Mo. If we assume both these masses of Ab and B components, we can estimate the magnitude difference. \\cite{1999AAS..134..545A} found this value to be 1.9~$\\pm$~0.1~mag, while here we derive 0.7~$\\pm$~4.5~mag. The very large error is due to the large uncertainty in the mass of the fourth body. Another approach is to use the standard mass-luminosity relation and derive the individual luminosities of the components. Using this approach, we have plotted Fig. \\ref{Color-mag}, where all components of the system 65~UMa are placed in the color-magnitude diagram. As one can see, the two eclipsing binaries are slightly under-luminous, while the D component seems to be over-luminous. The same finding about its higher luminosity was found elsewhere, e.g. \\cite{2010MNRAS.401.1299J} or \\cite{2007A&A...475.1053A}. However, some properties of the Aa-Ab orbit remain unclear, such as the inclination angle between the orbits. We can do some rough estimation of this quantity. The {\\sc KOREL} $K_{Aa}$ value and the predicted amplitude of radial-velocity variations from the LITE$_{Aa-Ab}$ are connected via $\\sin{i_{Aa-Ab}}$. Hence, we obtain $i_{Aa-Ab} \\approx 47^\\circ,$ which lies well between the inclinations $i$ of the eclipsing pair and the $i_{A-B}$ of the visual orbit. Nevertheless, its error is large but this is still only an estimation. We can also compute the predicted minimal angular separation of the Aa-Ab pair for a prospective interferometric detection. This resulted in about 11~mas, which is very favorable for modern stellar interferometers, because the magnitude difference between the Aa and Ab components should also be rather low. On the other hand, the angular separation of the eclipsing pair components is still rather low, at about only 0.18~mas. \\begin{figure} \\centering \\includegraphics[width=89mm]{New.eps} \\caption{Color-magnitude diagram for all components of the system. Their position is compared with the Hipparcos stars (small black dots).} \\label{Color-mag} \\end{figure} Dealing with a multiple system, we should also consider the nodal period of the close pair and the 640 day orbit, hence the change in the inclination of the eclipsing binary \\citep{1975A&A....42..229S}. The most crucial here is the ratio of periods ${p_{Aa-Ab}}^2/P$, which implies that the nodal period was about 650~years, a duration that should be practical to observe. Unfortunately, we do not have a complete set of orbital parameters of the Aa-Ab pair, so this is only first rough estimation. However, this nodal period is not too long and potentially detectable. Further observations would help us to detect the change in the eclipse depths. Despite these being rather shallow, this effect was detected in only nine other cases, hence it would be interesting to reattempt detections, especially with the modern ultra-precise satellite photometry. Nevertheless, 65~UMa is a rather unusual system, we presently know of only 11 other sextuple systems (see \\citealt{2008MNRAS.389..869E}). The mass ratio of close to unity for the inner pair seems to agree with some theoretical models of star formation, e.g. \\cite{2003MNRAS.342..926D}. Moreover, some studies (e.g. \\citealt{2002A&A...382..118T}) indicate that about one-third of all multiples are higher-order systems. \\cite{2007prpl.conf..133G} discussed a finding that there is a difference between the number of observed and expected higher-order multiples (quadruples and higher). Perhaps the discovery of other systems similar to 65~UMa would diminish this discrepancy." }, "1207/1207.7137_arXiv.txt": { "abstract": "The Sloan Digital Sky Survey III (\\mbox{SDSS-III}) presents the first spectroscopic data from the Baryon Oscillation Spectroscopic Survey (BOSS). This ninth data release (DR9) of the SDSS project includes 535,995 new galaxy spectra (median $z\\sim 0.52$), 102,100 new quasar spectra (median $z\\sim 2.32$), and 90,897 new stellar spectra, along with the data presented in previous data releases. These spectra were obtained with the new BOSS spectrograph and were taken between 2009 December and 2011 July. In addition, the stellar parameters pipeline, which determines radial velocities, surface temperatures, surface gravities, and metallicities of stars, has been updated and refined with improvements in temperature estimates for stars with $T_{\\rm eff}<5000$~K and in metallicity estimates for stars with $\\rm [Fe/H]>-0.5$. DR9 includes new stellar parameters for all stars presented in DR8, including stars from SDSS-I and II, as well as those observed as part of the \\mbox{SDSS-III} Sloan Extension for Galactic Understanding and Exploration-2 (SEGUE-2). The astrometry error introduced in the DR8 imaging catalogs has been corrected in the DR9 data products. The next data release for \\mbox{SDSS-III} will be in Summer 2013, which will present the first data from the Apache Point Observatory Galactic Evolution Experiment (APOGEE) along with another year of data from BOSS, followed by the final \\mbox{SDSS-III} data release in December 2014. ", "introduction": "\\label{sec:introduction} The Sloan Digital Sky Survey III (\\mbox{SDSS-III}; \\citealt{eisenstein11}) is an extension of the SDSS-I and II projects \\citep{york00}. It uses the dedicated 2.5-meter wide-field Sloan Foundation Telescope \\citep{gunn06} at Apache Point Observatory (APO), and fiber-fed multi-object spectrographs to carry out four surveys to study dark energy through observations of distant galaxies and quasars (the Baryon Oscillation Sky Survey; BOSS), to understand the structure of the Milky Way Galaxy (the Sloan Extension for Galaxy Understanding and Exploration; SEGUE-2, and the APO Galactic Evolution Experiment; APOGEE), and to search for extrasolar planets (the Multi-object APO Radial Velocity Exoplanet Large-area Survey; MARVELS). \\mbox{SDSS-III} commenced in Fall 2008, and will carry out observations for six years through Summer 2014. The first data release of this phase of SDSS (and the eighth release overall; DR8; \\citealt{DR8}) was made public in Winter 2011. In addition to all the data from SDSS-I and II \\citep{DR7}, DR8 included additional five-band imaging data over 2500 deg$^2$ over the Southern Galactic Cap, as well as stellar spectra from SEGUE-2. This paper presents the ninth data release (DR9) from SDSS, including all survey-quality data from BOSS gathered through 2011 July. BOSS \\citep{dawson12} uses new spectrographs \\citep{smee12} to obtain spectra of galaxies with $0.15 < z < 0.8$ and quasars with $2.15 < z < 3.5$ to measure the scale of the baryon oscillation peak in the correlation function of matter in order to probe the geometry and dynamics of the universe. DR9 includes the first year of BOSS data, and this paper describes the characteristics of these data (summarized in \\S\\ref{sec:scope}), with a particular emphasis on how it differs from the spectroscopy carried out in SDSS-I and SDSS-II (\\S\\ref{sec:BOSS}). The erratum to the DR8 paper \\citep{DR8_erratum} describes a systematic error in the astrometry in the imaging catalogs in DR8. This has now been fixed, as we describe in \\S\\ref{sec:astrometry}. The SEGUE Stellar Parameters Pipeline (SSPP) fits detailed models to the spectrum of each star, to determine surface temperatures, metallicities, and gravities. It has been continuously improved since its introduction in the sixth data release (DR6, \\citealt{DR6}; see also \\citealt{lee08a}). In \\S\\ref{sec:SSPP}, we describe the improvements since DR8 that are incorporated into the DR9 outputs. Section~\\ref{sec:distribution} describes how one can access the DR9 data, and we conclude and outline the planned future data releases in \\S\\ref{sec:conclusions}. ", "conclusions": "\\label{sec:conclusions} The \\mbox{SDSS-III} Data Release 9 presents the first data from the BOSS survey, with $\\sim$102,000 new quasar spectra, $\\sim$91,000 new stellar spectra and $\\sim$536,000 new galaxy spectra. The astrometry has been improved since DR8, and the stellar properties for SEGUE-I/II and \\mbox{SDSS-I/II} stars have been updated. These data are already sufficient for cosmological analyses of large-scale structure, investigations of the structure of the Milky Way, measurements of quasar physics, clustering, and demographics, and countless other science investigations. We invite the larger scientific community to investigate and explore this new data set. The \\mbox{SDSS-III} project will present two more public data releases. DR10, in summer 2013, will include the first data from the APOGEE survey and another year of BOSS data. DR11 will be an internal release only, as a public release would occur only six months before the final public data release for \\mbox{SDSS-III}, DR12, which will be released in December 2014 and will contain all of the data taken during the six years of the project." }, "1207/1207.0521_arXiv.txt": { "abstract": "We describe a search for infra-red excess emission from dusty circumstellar material around 180,000 stars observed by the \\emph{Kepler} and WISE missions. This study is motivated by i) the potential to find bright warm disks around planet host stars, ii) a need to characterise the distribution of rare warm disks, and iii) the possible identification of candidates for discovering transiting dust concentrations. We find about 8,000 stars that have excess emission, mostly at 12$\\mu$m. The positions of these stars correlate with the 100$\\mu$m background level so most of the flux measurements associated with these excesses are spurious. We identify 271 stars with plausible excesses by making a 5MJy/sr cut in the IRAS 100$\\mu$m emission. The number counts of these excesses, at both 12 and 22$\\mu$m, have the same distribution as extra-Galactic number counts. Thus, although some excesses may be circumstellar, most can be explained as chance alignments with background galaxies. The one exception is a 22$\\mu$m excess associated with a relatively nearby A-type star that we were able to confirm because the disk occurrence rate is independent of stellar distance. This detection implies a disk occurrence rate consistent with that found for nearby A-stars. Despite our low detection rate, these results place valuable upper limits on the distribution of large mid-infrared excesses; e.g. fewer than 1:1000 stars have 12$\\mu$m excesses ($F_{\\rm obs}/F_\\star$) larger than a factor of five. In contrast to previous studies, we find no evidence for disks around 1790 stars with candidate planets (we attribute one significant 12$\\mu$m excess to a background galaxy), and no evidence that the disk distribution around planet hosts is different to the bulk population. Higher resolution imaging of stars with excesses is the best way to rule out galaxy confusion and identify more reliable disk candidates among \\emph{Kepler} stars. A similar survey to ours that focusses on nearby stars would be well suited to finding the distribution of rare warm disks. ", "introduction": "\\label{s:intro} \\emph{Kepler} \\citep{2003SPIE.4854..129B} is revolutionising our perspective on extra-Solar planets \\citep[e.g.][]{2010Sci...330...51H,2011Natur.470...53L,2011ApJ...729...27B,2011Sci...333.1602D,2011arXiv1112.2165H,2012ApJ...745..120B} and will likely yield many Earth-sized planets in the terrestrial zones of their host stars. Like the Solar System, these planetary systems will comprise not only planets, but also smaller objects that for one reason or another did not grow larger. In the Solar System these make up the Asteroid and Kuiper belts, along with other populations such as the Oort cloud, and Trojan and irregular satellites. Characterisation of these populations has been critical to building our understanding of how the Solar System formed. For example, one of the primary validation methods of the so-called ``Nice'' model for the origin of the outer Solar System's architecture has been the reproduction of these minor body populations and their properties \\citep[e.g.][]{2005Natur.435..462M,2008Icar..196..258L,2007AJ....133.1962N}. The exquisite detail with which the Solar System minor body populations are characterised is made clear when they are contrasted with their extra-Solar analogues, collectively known as ``debris disks.'' First discovered around Vega \\citep{1984ApJ...278L..23A}, they are almost always detected by unresolved infra-red (IR) emission, visible as an excess above the stellar photosphere. Detection of an excess at multiple wavelengths yields the dust temperature, and thus the approximate radial distance from the star (to within a factor of a few). The radial location can be refined further when spectral features are present \\citep[e.g.][]{2007ApJ...658..584L}. However, because the temperature of a dust or ice grain depends on size, the true radial location (and any radial, azimuthal, or vertical structure) can generally only be found by resolved imaging \\citep[e.g.][]{1984Sci...226.1421S,2005Natur.435.1067K} or interferometry \\citep[e.g.][]{2006A&A...452..237A,2009A&A...503..265S}. It is therefore difficult to draw links between the regions of planetary systems occupied by planets and small bodies, and if and how they interact. The best examples of extrasolar systems where known dust and planets are likely to interact are $\\beta$ Pictoris, Fomalhaut, and HR 8799 \\citep{1995AAS...187.3205B,1997MNRAS.292..896M,2005Natur.435.1067K,2008Sci...322.1345K,2008Sci...322.1348M,2009ApJ...705..314S,2010ApJ...717.1123M}. In these cases, the spatial dust distribution is fairly well known because the disk is resolved, but the orbits of the planets, which were only discovered recently with direct imaging, are not. These are rare cases however, and typically the search for links between the major and minor body components of extra-Solar planetary systems means asking whether the presence of one makes the presence of the other more or less likely. So far no statistically significant correlation between the presence of planets and debris has been found \\citep{2009ApJ...700L..73K,2009ApJ...705.1226B,2011AJ....141...11D}. However, there is new tentative evidence that nearby stars with low-mass planetary systems are more likely to harbour debris than those with no planet detections (Wyatt et al, in press), an exciting possibility that has only been achievable recently with better sensitivity to such planetary systems around nearby stars. One of the key limiting factors in the search for links between debris and planets is the small number of stars known to host both. Two recent \\emph{Spitzer} surveys observed about 150 planet host stars, of which about 10\\% were found to have disks \\citep{2009ApJ...705.1226B,2011AJ....141...11D}. The small number of disk detections is therefore the product of the number of nearby stars known to host planets that could be observed with \\emph{Spitzer}, and the $\\sim$10\\% disk detection rate (for both planet and non-planet host stars). One way to sidestep this problem is therefore to look for disks around a much larger sample of planet host stars; the \\emph{Kepler} planet host candidates \\citep{2011ApJ...736...19B,2012arXiv1202.5852B}. The method we use to look for disks in this study is to find infra-red (IR) excess emission above that expected from the stellar photosphere. An IR excess is usually interpreted as being thermal emission from an Asteroid or Kuiper-belt analogue, which is heated by the star it orbits. We use photometry from the Wide-field Infrared Survey Explorer (WISE) mission's \\citep{2010AJ....140.1868W} all-sky catalogue, which is most sensitive to dust in the terrestrial region of Sun-like stars. Three properties of warm dust at these relatively close radial distances provide motivation. First, this warm dust, if discovered, is located in the vicinity of the planets being discovered with \\emph{Kepler}. Currently, only one system, HD 69830, is known to host both a planetary system in the terrestrial region and warm dust \\citep{2005ApJ...626.1061B,2006Natur.441..305L}. The origin of this dust is unclear, but given the proximity to the planetary system is plausibly related \\citep{2007ApJ...658..584L,2011ApJ...743...85B}. Through discovery of similar systems the links between planets and warm dust can be better understood. For planets discovered by \\emph{Kepler}, the knowledge that a transiting planetary system is almost exactly edge-on provides the second motivational aspect. If planets pass in front of the host star, so will coplanar minor body populations. Indeed, the discovery of systems where multiple planets transit their stars \\citep[e.g.][]{2010Sci...330...51H,2011Natur.470...53L,2011ApJS..197....8L} provides striking evidence that the Solar System's near-coplanar configuration is probably typical. While transits of individual small bodes will be impossible to detect, it may be possible to detect concentrated populations that arise from a recent collision \\citep{2005AJ....130..269K} or perturbations by planets \\citep{2011AJ....142..123S}. The dust must reside on a fairly close orbit---within a few AU---to allow multiple transits within the mission lifetime. Thus, the WISE sensitivity to terrestrial dust, and likely difficulties in discerning dust transits from other instrumental and real effects, mean that the odds of finding dust transits might be maximised by the prior identification of dusty systems. Finally, but most importantly, detections of terrestrial dust are rare \\citep[e.g.][]{1991ApJ...368..264A,2006ApJ...638.1070H,2006ApJ...636.1098B,2006ApJ...639.1166B}. Because only a few such systems are known, their occurrence rate is poorly constrained. More discoveries are therefore needed to add to our understanding of the processes that create it. The collision rate in a debris disk is proportional to the orbital period, so warm terrestrial debris disks decay to undetectable levels rapidly, hence their rarity. Indeed, the few that are known are usually thought to be the result of recent collisions, and thus transient phenomena \\citep[e.g. HD 69830, HD 172555, BD +20 307, $\\eta$ Corvi, HD 165014, HD 169666, HD 15407A][]{2005ApJ...626.1061B,2005Natur.436..363S,2007ApJ...658..569W,2009ApJ...701.2019L,2009ApJ...700L..25M,2010ApJ...714L.152F,2011arXiv1110.4172L,2012ApJ...749L..29F}. Possible scenarios include objects thrown into the inner regions of a planetary system from an outer reservoir \\citep{1999ApJ...510L.131G,2007ApJ...658..569W,2009MNRAS.399..385B,2011A&A...530A..62R,2012MNRAS.420.2990B}, or the remnant dust from a single catastrophic collision \\citep{2005Natur.436..363S,2005ApJ...626.1061B,2011ApJ...726...72W}. Clearly, there are reasons that discovery of debris in the terrestrial regions of known planetary systems is important. However, because WISE is sensitive to the rarest and brightest disks around \\emph{Kepler} stars, the third point above is of key importance. As stars become more distant, they and their debris disks become fainter, and the number of background galaxies at these fainter flux levels increases. Thus, the bulk of the stars in the \\emph{Kepler} field, which lie at distances of hundreds to thousands of parsecs may not be well suited to debris disk discovery. Practically, the importance of contamination depends on the galaxy contamination frequency relative to the disk frequency (i.e. only if disks are too rare will they be overwhelmed by contamination). \\emph{Therefore, because the occurrence rate of the rare disks that WISE is sensitive to is unknown, whether \\emph{Kepler} stars are a good sample for disk detection with WISE is also unknown.} Characterising the occurrence rate of rare bright disks is therefore the main goal of this study, because this very distribution sets what can be discovered. While the sample of \\emph{Kepler} stars is not specifically needed for this goal, there is the possibility that the disk occurrence rate is higher for stars that host low-mass planets. Such a trend could make this particular sub-sample robust to confusion, even if the general population is not. An additional potential issue specific to the \\emph{Kepler} field is the importance of the Galactic background. High background regions are sometimes avoided by debris disk observations because they make flux measurement difficult, and can even mask the presence of otherwise detectable emission. Unfortunately this issue cannot be avoided for the present study, as the \\emph{Kepler} field is necessarily located near the Galactic plane to maximise the stellar density on the sky. In what follows, we describe our search for warm excesses around $\\sim$180,000 stars observed by \\emph{Kepler} using the WISE catalogue. We first outline the data used in this study in \\S \\ref{s:cat} and in \\S \\ref{s:xs} describe our SED fitting method for finding excesses and the various issues encountered. We discuss the interpretation of these excesses in \\S \\ref{s:interp}, and place our findings in the context of disks around nearby stars in \\S \\ref{s:nearby}. We discuss the disk-planet relation, rarity of warm bright excesses, and some future prospects in \\S \\ref{s:disc} and conclude in \\S \\ref{s:sum}. Readers only interested in the outcome of this search may wish to skip the details described in \\S\\S \\ref{s:cat}-\\ref{s:xs}. ", "conclusions": "\\label{s:disc} One of several goals for this study was to test for a correlation between the existence of debris disks and planets discovered by \\emph{Kepler}. However, the distribution of the rare bright excesses that WISE is sensitive to around \\emph{Kepler} stars was not known at the outset, so whether this goal was possible was not known either. We noted that even if bright disks were too rare among the bulk population, that a possible correlation between disks and low-mass planets may allow robust disks detections among this subset. Only one \\emph{Kepler} planet candidate host (of 348 KOIs that were not excluded by the 100$\\mu$m background cut) was found to have an excess, so this possibility appears unlikely. In addition, Figure \\ref{fig:counts} shows that this detection rate is close to that expected from galaxy confusion. Thus, for the bright warm excesses that WISE is sensitive to, there is no evidence that planet host candidates have a disk occurrence rate that is different from the bulk population. Similarly, excesses around the remaining \\emph{Kepler} stars are also consistent with arising from chance alignments with background galaxies, with the exception of a single A-type star. However, the possibility that a small number of the excesses are true debris disks means that the chance of detecting transiting dust concentrations is at least as good as for \\emph{Kepler} stars without excesses, and may be higher (the 271 Kepler stars with W3 or W4 excesses are listed in Table \\ref{tab:xs}). Discovery of many such dust transits that preferentially occur around stars with excesses would argue that at least some excesses are debris disks, though this method of verification seems unlikely. We have therefore set new limits on the distribution of warm excesses. The range of flux ratios for which we have set limits for Sun-like stars is 2-20 at 12$\\mu$m (Fig. \\ref{fig:cumxsw3}) and 10-300 at 22$\\mu$m (left panel of Fig. \\ref{fig:cumxs}). For such large 12-22$\\mu$m excesses to arise from steady-state processes the planetesimal belts would have to be either around very young stars or relatively distant from their central star \\citep{2007ApJ...658..569W}, which in turn requires fractional luminosities $\\gtrsim$1\\% (see Fig. \\ref{fig:sens}). Detecting large warm excesses around main-sequence stars is very unlikely because collisional evolution depletes belts near the central star to undetectable levels rapidly, so the conclusion is that such mid-IR excesses are most likely transient. Two main processes seem to be plausible causes of such excesses. The first, delivery of material from an outer reservoir \\citep{2005ApJ...626.1061B,2007ApJ...658..569W}, is appealing because short-lived warm dust can be replenished using material from a long-lived outer belt. Alternatively, because we are here interested in large excesses, the debris from a giant impact between large bodies is a possibility \\citep[i.e. perhaps similar to the Earth-Moon forming event,][]{2012arXiv1206.4190J}. Several possibilities exist for the delivery of objects from an outer belt to terrestrial regions. A system of sufficiently many planets on stable orbits can pass objects inwards from an outer belt \\citep{2012MNRAS.420.2990B}, or a planetary system instability can severely disturb a planetesimal population, some of which end up in the terrestrial zone \\citep{2005Natur.435..466G}. Such possibilities have been suggested as mechanisms to generate the warm dust component observed around $\\eta$ Corvi \\citep{2009MNRAS.399..385B,2011arXiv1110.4172L}. Because at least 15\\% of Sun-like stars have cool outer planetesimal belts \\citep[e.g.][]{2008ApJ...674.1086T}, our limits of 0.01-0.1\\% for warm belts (for the flux ratios noted above for 12 and 22$\\mu$m) mean that fewer than 1 in 150-1500 can be generating large levels of warm dust from cool outer belts at any given time. This fraction could in fact be larger because the 15\\% only represents disks down to a particular detection limit, and cool disks too faint to detect could still have enough material to produce large warm dust levels \\citep{2007ApJ...658..569W}. \\citet{2009MNRAS.399..385B} placed similar limits on the number of systems that could be caught in the act of an instability that delivers large amounts of debris to the terrestrial region, estimating that less than about 0.2\\% (i.e. 1/500) of Sun-like stars might be observed undergoing an instability at 24$\\mu$m. Whether such instabilities do produce very large excesses is another question. In studying the dust emission generated in their model of the Solar System's proposed planetary instability, \\citet{2009MNRAS.399..385B} find that while the relative changes can be very large, the flux ratios are near unity at 12$\\mu$m and of order 10 at 24$\\mu$m. However, these ratios may be underestimated because they do not incude emission that could arise from the sublimation of comets within 1AU. It is therefore hard to say whether these results are representative, since they will also depend on the specific system architecture. The $\\eta$ Corvi system has been suggested as a possible candidate currently undergoing such an instability, and shows a 12$\\mu$m flux ratio of 1.3. If typical, these results suggest that instabilities may not produce the larger excesses considered here. In contrast, the giant impact scenario can produce extremely large excesses \\citep{2012arXiv1206.4190J}. The relatively nearby star BD+20 307 (at 96pc), which has a 10$\\mu$m flux ratio of about 100, is a good candidate for such an event \\citep{2005Natur.436..363S,2011ApJ...726...72W}. While such events would generally be expected to be associated with young systems, where the final $\\sim$10-100Myr chaotic period of giant impacts and terrestrial planet formation is winding down \\citep[e.g.][]{1998Icar..136..304C,2001Icar..152..205C}, BD+20 307 is a $\\gtrsim$Gyr old main-sequence binary \\citep{2008ApJ...688.1345Z}. The excess may therefore be indicative of a recent instability that has greatly increased the chance of collisions within the terrestrial zone, and is unrelated to planet formation \\citep{2008ApJ...688.1345Z}. Clearly, age estimates for the host stars are important for understanding the origin of dust in such systems. While the WISE mission might appear to permit near-unlimited sample sizes to help detect the aftermath of the rarest collision events, we have shown that their detection among \\emph{Kepler} stars is fundamentally limited. This limit arises because the occurrence rate of excesses that can be detected is too low, so the disks are overwhelmed by galaxy contamination. Because \\emph{Kepler} stars represent a sample that will remain unique for the forseeable future, it is desirable to find ways to overcome this issue. Based on the findings of \\S \\ref{ss:w34}, one option is to create sub-samples that maximise the chance of disk detection, because higher disk occurrence rates are more robust to galaxy contamination. Younger stars tend to have larger excesses that are also more frequent \\citep[e.g.][]{2005ApJ...620.1010R,2007ApJ...654..580S,2009ApJS..181..197C}, so a sub-sample of young stars will be more robust to confusion. The long-term monitoring of \\emph{Kepler} stars may provide some help if accurate stellar ages can be derived, for example if rotation periods can be derived to yield age estimates via gyrochronology \\citep{1972ApJ...171..565S,2007ApJ...669.1167B,2008ApJ...687.1264M}. Another way to split the sample is by spectral type, because earlier-type stars are both brighter and have higher disk occurrence rates (for fixed sensitivity). This approach is less appealing for studying the links between disks and planets however, because the bulk of stars observed by \\emph{Kepler} are Sun-like. If we allow for the possibility of observing \\emph{Kepler} stars with WISE excesses with other instruments, there is a potential gain with better resolution. A galaxy that is unresolved with WISE might be resolved with \\emph{Spitzer}'s IRAC instrument, or using ground-based mid-IR observations on 8m-class telescopes for example. Assuming that it could be detected, the high ($\\sim$0\\farcs5) resolution of such ground-based observations would have over a 99\\% chance of detecting a galaxy that was not resolved with the WISE beam at 12$\\mu$m. Therefore, detection of fewer galaxies than expected in a sample of targets (e.g. significantly fewer than 99 out of 100) would be evidence that the excesses do not randomly lie within the WISE beam and that some are therefore due to excesses centered on the star (i.e. are debris disks). Ultimately, we found that searching for debris disks around stars in the \\emph{Kepler} field with WISE is limited by the high background level and galaxy contamination. While high background regions can be avoided, background galaxies will always be an issue for such distant stars. Though it means being unable to study the planet-disk connection with such a large planet-host sample, nearby stars should be the focus of studies that aim to better define the distribution of warm excesses. Characterising this distribution is very important, particularly for estimating the possible impact of terrestrial-zone dust on the search for extrasolar Earth analogues \\citep[e.g.][]{2006ApJ...652.1674B,2012arXiv1204.0025R}. For example, extending the distribution to the faintest possible level available with photometry (calibration limited to a 3$\\sigma$ level of $\\sim$5\\%) yields a starting point to make predictions for instruments that aim to detect faint ``exozodi'' with smaller levels of excess. Because bright warm debris disks must decay to (and be observable at) fainter levels, the distribution will also provide constraints on models that aim to explain the frequency and origin of warm dust." }, "1207/1207.2524_arXiv.txt": { "abstract": "We produce the most comprehensive public void catalog to date using the Sloan Digital Sky Survey Data Release 7 main sample out to redshift $z=0.2$ and the luminous red galaxy sample out to $z=0.44$. Using a modified version of the parameter-free void finder {\\tt ZOBOV}, we fully take into account the presence of the survey boundary and masks. Our strategy for finding voids is thus appropriate for any survey configuration. We produce two distinct catalogs: a complete catalog including voids near any masks, which would be appropriate for void galaxy surveys, and a bias-free catalog of voids away from any masks, which is necessary for analyses that require a fair sampling of void shapes and alignments. Our discovered voids have effective radii from 5 to 135~\\hmpc. We discuss basic catalog statistics such as number counts and redshift distributions and describe some additional data products derived from our catalog, such as radial density profiles and projected density maps. We find that radial profiles of stacked voids show a qualitatively similar behavior across nearly two decades of void radii and throughout the full redshift range. ", "introduction": "\\label{sec:introduction} The hierarchical clustering of matter in the universe naturally leads to large underdense regions, called voids. Indeed, the presence of voids in the large-scale distribution of galaxies was one of the early predictions of cold dark matter cosmological theories~\\citep{Hausman1983}, and the discovery of voids in some of the first galaxy redshift surveys~\\citep{Gregory1978,Kirshner1981} quickly provided a rich source of interest. Today, galaxy surveys, such as the Void Galaxy Survey~\\citep{VandeWeygaertR.2011} and the Las Companas Redshift Survey~\\citep{Muller2000} routinely find and characterize both voids themselves and their contents for useful astrophysical and cosmological information (see~\\citealt{Thompson2011} for a review). A combination of observations and simulations gives a coherent picture of void properties. In the cosmological constant plus cold dark matter ($\\Lambda$CDM) picture of cosmic evolution, voids --- usually defined to have densities of 10-20\\% the cosmic mean --- have characteristic radii of 10-40~\\hmpc, with the smallest identifiable voids in the local universe having radii $\\sim7$~\\hmpc ~\\citep{Tikhonov2006}. Early structure formation simulations successfully reconstructed observed voids~\\citep{Hoffman1982, White1987}. Later studies of voids from surveys such as IRAS~\\citep{Plionis2002}, the 2dF Galaxy Redshift Survey~\\citep{Hoyle2004}, and the Sloan Digital Sky Survey~\\citep{Pan2011} confirmed these properties. Semi-analytic models~\\citep{Benson2003, Tinker2009} and large-scale \\emph{ab initio} simulations~\\citep{Dubinski1993, Colberg2005, Ceccarelli2006, Park2007, Kreckel2011} have further elucidated the evolution, internal structure, and distribution of voids. Since voids are nearly empty, their dynamics are dominated by dark energy. Thus, they may provide crucial probes of primordial density fluctuations~\\citep{Sahni1994}, fifth forces~\\citep{Li2009}, and $F(r)$ gravity models~\\citep{Li2012}. The Alcock-Paczynski test~\\citep{Alcock1979} can be applied to measurements of void ellipticities, directly probing the expansion history of the universe~\\citep{Ryden1995, Ryden1996, Park2007, Biswas2010, LavauxGuilhem2011}. The internal structure of voids behaves as a universe in miniature, allowing for probes of the history of dark energy~\\citep{Gottlober2003, Goldberg2004}. The ellipticity distribution of voids can provide insights into the growth of structure and the correctness of General Relativity~\\citep{Shoji,Lavaux2010}. Void orientation and spin statistics reveal information on large-scale tidal fields~\\citep{Lee2006, Platen2008}. Understanding the locations and sizes of voids is also crucial for cosmic microwave background (CMB) missions, since they affect the CMB signal via the integrated Sachs-Wolfe effect~\\citep{Thompson1987, Vadas1998, Cruz2008, Gurzadyan2009, Granett2008}. The reliability of the above conclusions and predictions rests on the ability to robustly produce void catalogs from galaxy surveys. While void finders are well-studied in the context of the large-scale dark matter simulations~\\citep[e.g.,][]{Colberg2005, Colberg2008}, few are applied directly to large-scale redshift surveys. The largest void catalogs previous to this work use void finders that rely on overlapping spheres of underdensities~\\citep{Hoyle2004, Pan2011}. While simple to apply, this approach fails to capture the full geometry of the voids and relies on finely-tuned parameters to correctly capture them. Additionally, previous works ignore the presence of survey boundaries and masks and do not extend to the full redshift range of the available surveys. In this work, we describe our techniques for accounting for biases due to the presence of a survey boundary and masks. We employ these techniques along with a modified version of the parameter-free void finder {\\tt ZOBOV}~\\citep{Neyrinck2008, LavauxGuilhem2011} to produce a catalog of voids from both the main sample~\\citep{Strauss2002} and luminous red galaxy (LRG)~\\citep{Eisenstein2001} sample of the Sloan Digital Sky Survey (SDSS) Data Release 7~\\citep{Abazajian2009}. The void catalog we produce will be useful for many void-based astrophysical and cosmological studies, as already noted. This void catalog extends to higher redshifts than other catalogs~\\citep[e.g.,][]{Plionis2002,Hoyle2004,Pan2011,Sousbie2011,VandeWeygaertR.2011} and is the first to include not only main sample galaxies but also LRGs. While the voids we will identify in the LRG sample are topologically consistent (based on the tessellation and watershed procedures in {\\tt ZOBOV}), they may not fully correspond to underdensities in the cosmological sense due to undersampling of the density field and galaxy biasing effects. We will return to this discussion in Section~\\ref{sec:conclusions}. We begin in Section~\\ref{sec:samples} with a presentation of our selection of data samples from the SDSS catalog. In Section~\\ref{sec:finding} we describe our modifications to {\\tt ZOBOV} to handle the survey boundary and masks as well as our process for eliminating alignment biases in the void catalog. We characterize the demographics of our void catalog, including redshift-dependent number counts and radial density profiles, in Section~\\ref{sec:properties}. We provide two examples of derived data products from the void catalog: radial profiles of stacked voids in Section~\\ref{sec:radial} and projected density maps of stacked voids in Section~\\ref{sec:projected}. We discuss the potential for future applications in Section~\\ref{sec:conclusions} and provide details of the layout of the public void catalog in the Appendix. ", "conclusions": "\\label{sec:conclusions} We have modified the parameter-free void finding algorithm {\\tt ZOBOV} to account for the survey boundaries and internal masks in observational data sets. This prevents voids from growing past the survey boundary or into any internal masks. Thus our approach is more generally applicable to any given survey and mask. To demonstrate our technique we have constructed the first public void catalog using the full extent of the SDSS DR7 spectroscopic survey, including the LRGs. We combined multiple volume-limited samples of the SDSS galaxy catalogs to maximize the number of discovered voids. We have produced two catalogs: one catalog that includes all discovered voids, including truncated voids near the survey boundaries, and a central catalog, which removes voids with questionable shapes and alignments. The general statistics of our void catalog, such as number counts as a function of redshift and size distributions, broadly agree with --- but significantly extend --- both past analyses of observational data~\\citep[e.g.,][]{Muller2000,VandeWeygaertR.2011,Pan2011,Patiri2012} and results from simulations~\\citep[e.g.,][]{Dubinski1993, Park2007}. In addition, radial profiles and projections of stacked voids show a qualitatively similar shape across the entire sample and agree well with previous efforts. Due to the relatively poor sampling and the high redshift of the LRG samples, the topological voids we identify there may not be strict cosmological features understood as underdensities bounded by filaments and walls. We may also be overestimating the size of these voids and possibly miscalculating their centers. However, the largest voids found in the main sample ($\\sim 50-60$~\\hmpc) overlap with the size distribution of voids from the \\emph{lrgdim} sample, indicating that there is at least some correspondence between the void populations in these samples. Also, our radial profiles show a qualitatively universal shape in all volume-limited subsamples (excepting the \\emph{lrgbright} sample), which again is a point of evidence that these are truly cosmic voids (note especially the similarity in shape in the $50-55$~\\hmpc~bin of Figure~\\ref{fig:profile_samples}). In either case, these structures are useful for many kinds of analysis ~\\citep[e.g.,][]{Granett2008}. Our catalogs are useful for many pursuits, including studies of the ellipticity distribution of voids, correlations of void positions with CMB fluctuations, Alcock-Paczynski tests using the shapes of voids in redshift space, and studies of the properties of galaxies within voids. We have constructed useful data sets to enable these studies, such as catalogs of void galaxies, void stacks of various radial sizes, and two-dimensional projections of void densities. We have constructed these data sets using both all discovered voids and a central void catalog free from survey edge effects. We are making our catalogs and data products publicly available at~\\url{http://www.cosmicvoids.net}." }, "1207/1207.7301_arXiv.txt": { "abstract": "We present results on the age and metallicity estimates of the astonishingly unstudied SMC cluster ESO\\,51-SC09, from CCD $BVI$ photometry obtained at the ESO NTT with the EMMI attached. ESO\\,51-SC09 turns out to be a relatively small cluster (FWHM = (10 $\\pm$ 1) pc) located $\\sim$ 4$\\degr$ northward from the galaxy center. We report for the first time a mean cluster age of (7.0 $\\pm$ 1.3) Gyr and a mean cluster metallicity of [Fe/H] = (-1.00 $\\pm$ 0.15) dex, concluding that ESO\\,51-SC09 belongs to the group of the oldest SMC clusters. We found that the cluster is projected onto a dominant field stellar population older (age $\\sim$ 10-13 Gyr) and more metal-poor ([Fe/H] = -1.3$\\pm$0.2 dex), so that the cluster could reach its current location because of its orbital motion. ", "introduction": "As far as we are aware, ESO\\,51-SC09 (RA = 00$^h$ 58$^m$ 57$^s$.96, Dec. = -68$\\degr$ 54$\\arcmin$ 55$\\arcsec$.7, J2000) is a cataloged star cluster of the Small Magellanic Cloud \\citep[SMC,][]{bietal08} which has remained unstudied until the present. It is located in the outer disk of the SMC, at $\\sim$ 4$\\degr$ northward from the galaxy center as supposed to be at RA = 00$^h$ 52$^m$ 45$^s$, Dec. = \u221272$\\degr$ 49$\\arcmin$ 43$\\arcsec$ (J2000) \\citep{cetal01}. Its position should facilitate astrophysical studies, since reddening effects are at a minimum regime and the unavoidable field contamination does not represent a real constraint. On the other hand, bearing in mind the enormous interest in identifying new relatively old/old clusters in the SMC \\citep{detal10} and the appearance in the sky of ESO\\,51-SC09 like a candidate old cluster \\citep{l82}, it is astonishing that it has not been mentioned in the literature as a valuable target. According to the hierarchical star-formation scenario found by \\citet{bb10}, the outer SMC disk would appear to be a genuine reservoir of old clusters. In fact, the oldest known cluster \\citep[NGC\\,121, age = 10.6$\\pm$0.5 Gyr,][]{detal01} is also placed in the outer disk, whereas Lindsay\\,32 and 38 - two relatively old/old clusters \\citep{petal01} - are in the ESO\\,51-SC09's zone, at distances smaller than $\\sim$ 1$\\degr$. In this Letter we present for the first time age and metallicity estimates for ESO\\,51-SC09. The results show that this cluster belongs to the handful of oldest clusters in the SMC (age $\\ge$ 5 Gyr). The impact of this finding would appear to be twofold: firstly, we actually found a new SMC cluster in the relatively old/old range. Note that different campaigns have carried out until the present searching for old star clusters in the SMC and, unfortunately, new candidates have not been identified. These results would appear not only to show that the task of finding more old star clusters in the SMC is arduous, but also it would appear a venture hardly to reach success. The amazing scarce amount of old SMC star clusters results even more noticeable when comparing it with the 456 star clusters cataloged in the SMC \\citep{bb10}, thus representing $\\approx$ 1\\% of the SMC star cluster population. Secondly, taking into account that \\citet{p11a} predicted that we should expect to identify only one outer disk relatively old/old cluster not studied yet within those cataloged by \\citet{bietal08}, and ESO\\,51-SC09 is uncovered to be a relatively old/old cluster, the SMC outer disk would appear not to be populated by any other unstudied old cluster. This Letter is organized as follows: In Section 2 we describe the data collected, the reduction procedures performed, and the subsequent photometry standardization. In Section 3 we deal with the infamous cleaning process of the decontamination of field stars in the cluster Color-Magnitude Diagram, while Section 4 is devoted to the estimation of the cluster age and metallicity. Finally, Section 5 summarizes our results. ", "conclusions": "In this study we present for the first time CCD $BVI$ photometry of stars in the field of the unstudied SMC cluster ESO\\,51-SC09. The data were obtained at the ESO NTT with the EMMI attached under high quality photometric conditions. We are confident that the photometric data yield accurate morphology and position of the main cluster features in the CMD. To disentangle the cluster features from those belonging to its surrounding field, we applied a subtraction procedure to statistically clean the cluster CMD from field star contamination. The method has shown to be able to eliminate stochastic effects in the cluster CMDs caused by the presence of isolated bright stars, as well as, to make a finer cleaning in the most populous CMD regions. The FWHM of the genuine cluster stellar density radial profile turned out to be (0.58 $\\pm$ 0.06)$\\arcmin$, which converts into (10 $\\pm$ 1) pc if a SMC distance of 60 kpc is adopted. Using the cleaned cluster ($V$,$B-V$) diagram, we estimated its age and metallicity using the $\\delta V$ and $\\delta T_1$ indices and fitting theoretical isochrones. The three different age estimates are in excellent agreement, resulting in a mean value of (7.0 $\\pm$ 1.3) Gyr. A metallicity of [Fe/H] = (-1.00 $\\pm$ 0.15) dex was estimated from the fitting of theoretical isochrones to the cluster CMDs. We thus report that ESO\\,51-SC09 belongs to the group of the oldest SMC clusters, only younger (mean values) than NGC\\,121, HW\\,42, NGC\\,361, and Linday\\,1. The cluster is placed in a region of the SMC where probably it has not been born, since the mean age and metallicity of the dominant field stellar population is remarkable older (age $\\sim$ 10-13 Gyr) and more metal poor ([FeH] = -1.3$\\pm$0.2 dex). The cluster could reach its current location because of its orbital motion." }, "1207/1207.2689_arXiv.txt": { "abstract": "{The abundances of many observed compounds in interstellar molecular clouds still lack an explanation, despite extensive research that includes both gas and solid (dust-grain surface) phase reactions.} {We aim to qualitatively prove the idea that a hydrogen-poor subsurface chemistry on interstellar grains is responsible for at least some of these chemical ``anomalies''. This chemistry develops in the icy mantles when photodissociation reactions in the mantle release free hydrogen, which escapes the mantle via diffusion. This results in serious alterations of the chemical composition of the mantle because pores in the mantle provide surfaces for reactions in the new, hydrogen-poor environment.} {We present a simple kinetic model, using existing astrochemical reaction databases. Gas phase, surface and subsurface pore reactions are included, as are physical transformations of molecules.} {Our model produces significantly higher abundances for various oxidized species than most other models. We also obtain quite good results for some individual species that have adequate reaction network. Thus, we consider that the hydrogen-poor mantle chemistry may indeed play a role in the chemical evolution of molecular clouds.} {The significance of outward hydrogen diffusion has to be proved by further research. A huge number of solid phase reactions between many oxidized species is essential to obtain good, quantitative modeling results for a comparison with observations. We speculate that a variety of unobservable hydrogen-poor sulfur oxoacid derivatives may be responsible for the ``disappearance'' of sulfur in dark cloud cores.} ", "introduction": "\\label{intro} It is widely accepted that molecular hydrogen and many other interstellar molecules form on interstellar dust grains. There has been wide research in the field of gas-grain chemistry occurring in the dark, dense cores of interstellar molecular clouds, including observations and calculations. Various desorption mechanisms and chemical reactions on the surfaces of the grains are investigated. Several molecules in the interstellar medium at least are known whose abundances are not easily explained by gas-phase and grain surface-phase chemistry. These include OCS, HCN, SO, cyanopolyynes, and others. In the warming star-formation regions there appear highly oxidized organic compounds that indicate that a different chemistry occurs in these regions or, as we believe, the ejection of heavily processed interstellar grain mantles into the gas phase. The grain surface reactions in interstellar clouds by definition are subjected to heavy hydrogenation. The proper production of heavier and hydrogen-poor species is difficult to reproduce by calculations at least in some cases (Hasegawa \\& Herbst \\cite{11}, van Weeren et al. \\cite{01}, Hatchell et al. \\cite{35}). We present an alternative explanation of this problem by considering the possibility that within the grain mantles below the surface layer there is a hydrogen-poor, chemically active material. The aim of this article is to qualitatively answer this question through existing knowledge and means of astrochemical problem solving. There are only few such models that take into account more than one layer of the accreted species. There are some serious researches done which insist that the chemistry of the species frozen onto grains does not end with the formation of the next accreted layer. Shalabiea \\& Greenberg (\\cite{26}) examine photon-induced processes inside the mantle. They conclude that photoprocessing of grain mantles is the start of the synthesis of many species. Hasegawa \\& Herbst (\\cite{10}) use a 3-phase model to examine the formation and composition of the inner mantle. They found that radicals like OH, $\\mathrm{CH_{3}}$ etc. are absorbed into mantles in large numbers. We argue that the chemistry below the outer surface of the icy mantle has to be common and is an important path in interstellar molecule synthesis. Besides, they also note that grain surface reactions tend to overproduce hydrogenated species, which is the main problem we attempt to tackle in this work. Schutte \\& Greenberg (\\cite{27}) examine the possibility of molecule desorption from the grains by chemical explosions within the grain mantle. Naturally, these reactions can be expected to alter the chemical composition of the mantle itself. Freund \\& Freund (\\cite{38}) present a dust grain model based on the principle of solid solutions, producing results that explain important features in the molecular cloud composition. We present a model of the processes on the surface and inside the ``frozen'' grain mantles with a basic concept that chemical reactions occur on the surface of pores (or cracks) inside the mantles. The point that makes the difference between the surface and mantle reactions is that the mantle is not directly exposed to the ocean of hydrogen in a nebula, and thus a significantly different chemistry may develop. This gives an opportunity to present an explanation for some of the astrochemical mysteries. We do not use advanced calculation techniques or new reactions; we evaluate the importance of hydrogen diffusion through the grain mantle on the chemical composition of the mantles in dark molecular cloud cores. Thus, the model includes various physiochemical processes but otherwise keeps a rather conservative approach. ", "conclusions": "\\label{conclusions} \\begin{enumerate} \\item According to our model, which includes several basic assumptions (see Sect.~\\ref{model}), it is highly possible that chemical processes below the outer surface of interstellar grain mantles play an important role in the chemistry of dark molecular cloud cores. \\item An important factor that should be further investigated is the hydrogen diffusion through the grain mantles. This, combined with the continued dissociation of molecules by cosmic ray induced photons, leads to an overall outward flux of hydrogen from the mantle. According to our model results, the more dense the grain mantle, the more efficient is the outward diffusion of hydrogen. A more diverse, H-poor chemistry is encouraged below the surface, explaining the abundance of at least some species observed in dark molecular clouds and hot molecular cores. \\item A model based on the concept of pore surface reactions can at least partially describe the transformations occurring within the icy mantle. \\item Fe nuclei of cosmic-rays can be a cause of physical alteration of the mantle structure, but other possibilities, like light cosmic-ray particles, CR induced UV photons, and slow thermal diffusion may provide alteration with generally similar chemical consequences. \\item A combination of outer and inner mantle surface chemistry is able to produce a wide set of mantle species. It may ultimately lead to more accurate calculations of the composition of molecular clouds. \\item Further research is required to clarify many factors that are only approximately estimated. These include photodissociation and desorption yields, H and $\\mathrm{H_{2}}$ diffusion rate in amorphous dirty ices, the thickness of the mantle, the properties and number of the pores, reaction mechanism inside the grain mantles, effectiveness of selective desorption mechanisms, etc. \\item A chemically active and hydrogen poor environment in the mantle may explain the difficulties of observing sulfur in dark molecular cores. In grain mantles the rich chemistry of sulfur permits the formation of many various S molecules (mostly oxoacid derivatives) low on hydrogen, most of them with abundances too low to be observed. The large abundance of sulfur oxides in hot star-forming cores may be a direct consequence of these compounds being ejected into the gas phase. \\end{enumerate}" }, "1207/1207.3646_arXiv.txt": { "abstract": "In this work is presented the software OGCOSMO. This program was written using high level design methodology (HLDM), that is based on the use of very high level (VHL) programing language as main, and the use of the intermediate level (IL) language only for the critical processing time. The languages used are PYTHON (VHL) and FORTRAN (IL). The core of OGCOSMO is a package called OGC{\\_}lib. This package contains a group of modules for the study of cosmological and astrophysical processes, such as: comoving distance, relation between redshift and time, cosmic star formation rate, number density of dark matter haloes and mass function of supermassive black holes (SMBHs). The software is under development and some new features will be implemented for the research of stochastic background of gravitational waves (GWs) generated by: stellar collapse to form black holes, binary systems of SMBHs. Even more, we show that the use of HLDM with PYTHON and FORTRAN is a powerful tool for producing astrophysical softwares. \\bigskip {\\footnotesize {\\bf Keywords}:Computational Physics, Cosmology, Gravitational Waves, Black Hole, Structure Formation.} ", "introduction": " ", "conclusions": "" }, "1207/1207.6914_arXiv.txt": { "abstract": "{ We introduce a novel approach, a Dense Shell Method (DSM), for measuring distances for cosmology. It is based on original Baade idea to relate absolute difference of photospheric radii with photospheric velocity. We demonstrate that this idea works: the new method does not rely on the Cosmic Distance Ladder and gives satisfactory results for the most luminous Type IIn Supernovae. This allows one to make them good primary distance indicators for cosmology. Fixing correction factors for illustration, we obtain with this method the median distance of $\\approx 68^{+19}_{-15}$(68\\%CL)~Mpc to SN~2006gy and median Hubble parameter $79^{+23}_{-17}$(68\\%CL)~km/s/Mpc. } ", "introduction": "Supernovae are among the most luminous phenomena in the Universe, and they can serve as cosmological distance indicators. In some cases one can use a standard candle method. Nobel prize 2011 in physics is given ``for the discovery of the accelerating expansion of the Universe through observations of distant supernovae''. Actually, Type Ia supernovae have been used for this. Although SNe~Ia are not uniform in luminosity, they can be standardised. The standardisation is based on statistical correlations found for nearby events \\cite{Pskovsky77, Phillips93}. Thus they are \\emph{secondary} distance indicators, see reviews, e.g. \\cite{Leibundgut01,Phillips05}. Type II supernovae, on the other hand, have a much larger variance in luminosity and therefore cannot provide an accurate distance by photometry alone. Nevertheless, their great advantage is the possibility of direct measurement of distance, e.g. by Expanding Photosphere Method (EPM) \\cite{KK1974} when applied to SNe~IIP. The development of EPM is the spectral-fitting expanding atmosphere method (SEAM) \\cite{BaronSEAM}. Thus, Type II supernovae are interesting because there are ways to make them \\emph{primary} distance indicators. A standard candle assumption and its calibration is not needed for direct methods. Applications of SNe~IIP in cosmography do not depend upon the steps of Cosmic Distance Ladder avoiding their systematic and statistical errors. Due to absolute weakness of SNe~IIP they cannot be used at large cosmological distances. In this paper we introduce a novel approach to measuring distances for cosmology with the help of the most luminous Type IIn Supernovae. The method is based on the formation of an expanding dense shell in SN~IIn and allows one to find a linear size of a supernova shell in absolute units and distance to it. This Dense Shell Method (DSM) is partly based on ideas introduced in EPM and SEAM, and partly in Expanding Shock Front Method (ESM) \\cite{Bartel2007} used for SNR~1993J. ", "conclusions": "Now, we can summarise essential features of the new method, DSM (Dense Shell Method), for finding cosmological distances with the help of SNe~IIn. The method is based on the following steps: \\begin{itemize} \\item Measurement of \\emph{wide} emission components of lines and determination of the velocity at photosphere level $v_{\\rm m} = v_{\\rm ph}$ (with highest possible accuracy). \\item Measurement of \\emph{narrow} components of spectral lines for estimating properties (density, velocity) of circumstellar envelope. One does not need a very high accuracy of measurements and modelling here. \\item Building of a set of \\emph{best fitting models} (``suitable'') for broad band photometry and speed $v_{\\rm ph}$, for a set of trial distances satisfying the constraints for the circumstellar envelope found from narrow lines. \\item Although the free expansion assumption $v = {r}/{t}$ is not applicable, $v_{\\rm m}$ now measures a true velocity of the photospheric radius (not only the matter flow speed, as in type IIP). \\item Now the original Baade's idea works for measuring the radius $R_{\\rm ph}$ by integrating $dR_{\\rm ph}=v_{\\rm ph} dt$ (of course, with due account of scattering, limb darkening/brightening etc. in a time-dependent modelling). \\item The observed flux then gives the {\\it distance} $D$ from the system~(\\ref{sqrtDistAv}). \\end{itemize} The constraining of cosmological parameters and our understanding of Dark Energy depend strongly on accurate measurements of distances in Universe. SNe~IIn may be used for cosmology as \\emph{primary distance indicators} with the new DSM method. Application of EPM and SEAM requires crafting a best fitting hydro model for each individual SN. This procedure is in principle simpler in DSM. The case of SN~2006gy shows that the DSM distance agrees well with other most reliable techniques when the correct model is used, without the assumption on free expansion which is needed for EPM and SEAM." }, "1207/1207.6125_arXiv.txt": { "abstract": "Studying CMEs in coronagraph data can be challenging due to their diffuse structure and transient nature, and user-specific biases may be introduced through visual inspection of the images. The large amount of data available from the SOHO, STEREO, and future coronagraph missions, also makes manual cataloguing of CMEs tedious, and so a robust method of detection and analysis is required. This has led to the development of automated CME detection and cataloguing packages such as CACTus, SEEDS and ARTEMIS. Here we present the development of a new CORIMP (coronal image processing) CME detection and tracking technique that overcomes many of the drawbacks of current catalogues. It works by first employing the dynamic CME separation technique outlined in a companion paper, and then characterising CME structure via a multiscale edge-detection algorithm. The detections are chained through time to determine the CME kinematics and morphological changes as it propagates across the plane-of-sky. The effectiveness of the method is demonstrated by its application to a selection of SOHO/LASCO and STEREO/SECCHI images, as well as to synthetic coronagraph images created from a model corona with a variety of CMEs. The algorithms described in this article are being applied to the whole LASCO and SECCHI datasets, and a catalogue of results will soon be available to the public. ", "introduction": "Coronal mass ejections (CMEs) are large-scale eruptions of plasma and magnetic field from the Sun into interplanetary space, and have been studied extensively since they were first discovered four decades ago \\citep{1972BAAS....4R.394T}. They propagate with velocities ranging from $\\sim$\\,20~km~s$^{-1}$ to over 2000~km~s$^{-1}$ \\citep{2004JGRA..10907105Y, 2001JGR...10629219G}, and with masses of 10$^{14}$\\,--\\,10$^{17}$~g \\citep{1985SoPh..100..563J, 1996ApJ...470..629H, 1992ApJ...390L..37G}, and are a significant driver of space weather in the near-Earth environment and throughout the heliosphere \\citep{2010heliophysics, 2005AnGeo..23.1033S}. Traveling through space with average magnetic field strengths of 13~nT \\citep{2003SoPh..212..425L} and energies of $\\sim$\\,10$^{25}$~J \\citep{2004JGRA..10910104E}, they can cause geomagnetic storms upon impacting Earth's magnetosphere, possibly damaging satellites, inducing ground currents, and increasing the radiation risk for astronauts \\citep{2007A&G....48f..11L}. Thus, models of CMEs and the forces acting on them during their eruption and propagation through the corona remain an active area of research \\citep[see reviews by][]{2011LRSP....8....1C, 2012LRSP_Webb}. The Large Angle Spectrometric Coronagraph suite \\citep[LASCO;][]{1995SoPh..162..357B} onboard the Solar and Heliospheric Observatory \\citep[SOHO;][]{1995SoPh..162....1D} has observed thousands of CMEs from 1995 to present; and since 2006 the Sun-Earth Connection Coronal and Heliospheric Imaging suite \\citep[SECCHI;][]{2008SSRv..136...67H} onboard the Solar Terrestrial Relations Observatory \\citep[STEREO;][]{2008SSRv..136....5K} has provided twin-viewpoint observations of the Sun and CMEs from off the Sun-Earth line. Defined as an outwardly moving, bright, white-light feature, CMEs appear in a variety of geometrical shapes and sizes, typically exhibiting a three-part structure of a bright leading front, dark cavity, and bright core \\citep{1985JGR....90..275I}. Their geometry is attributed to the underlying magnetic field, generally believed to have a flux-rope configuration. The eruption of the CME is triggered by a loss of stability and its subsequent outward motion is governed by the interplay of magnetic and gas pressure forces in the low plasma-$\\beta$ environment of the solar corona. CMEs are commonly linked to filament/prominence eruptions and solar flares \\citep{2002ApJ...581..694M, 2002ApJ...566L.117Z}, or labelled `stealth CMEs' if they cannot be associated with any on-disk activity \\citep{2009ApJ...701..283R}, but knowledge about their specific driver mechanisms remains elusive. Several theoretical models have been developed in order to describe the forces responsible for the observed characteristics of CME initiation and propagation. The most favoured models explain CMEs in the context of tether straining and release, whereby the outward magnetic pressure increases due to flux injection or field shearing, and overcomes the magnetic tension of the overlying field \\citep{2001AGUGM.125..143K}. Different approaches to such models provide different force-balance interpretations, that lead to a variety of predictions on the kinematic and morphological evolution of CMEs \\citep[e.g.][]{2002A&ARv..10..313P, 2003JGRA..108.1410C, 2006PhRvL..96y5002K, 2008ApJ...683.1192L}. To this end, there is a motivation to resolve the observations of CMEs with robust, high-accuracy methods, in order to determine their kinematics and morphology with the greatest possible precision. From the large number of CMEs observed to date, many exhibit a general multiphased kinematic evolution. This often consists of an initiation phase, an acceleration phase, and a propagation phase which can show positive or negative residual acceleration as the CME speed equalizes to that of the local solar wind \\citep{2006ApJ...649.1100Z, 2009SoPh..256..149M}. Statistical analyses can provide a general indication of CME properties \\citep[e.g.][]{2000GeoRL..27..145G, 2003AdSpR..32.2637D, 2005AnGeo..23.1033S}, but it remains true that individual CMEs must be studied with rigour in order to satisfactorily derive the kinematics and morphology to be compared with theoretical models. The CME catalogue hosted at the Coordinated Data Analysis Workshop (CDAW\\footnote{http://cdaw.gsfc.nasa.gov/CME\\_list}) Data Center grew out of a necessity to record a simple but effective description and analysis of each event observed by LASCO \\citep{2009EM&P..104..295G}, but its manual operation is both tedious and subject to user biases. Ideally an automated method of CME detection should be applied to the whole LASCO and SECCHI datasets in order to glean the most information possible from the available statistics. A number of catalogues have therefore been developed in an effort to do this, namely the Computer Aided CME Tracking catalogue \\citep[CACTus\\footnote{http://sidc.oma.be/cactus/};][]{2004A&A...425.1097R}, the Solar Eruptive Event Detection System \\citep[SEEDS\\footnote{http://spaceweather.gmu.edu/seeds/};][]{2008SoPh..248..485O} and the Automatic Recognition of Transient Events and Marseille Inventory from Synoptic maps \\citep[ARTEMIS\\footnote{http://www.oamp.fr/lasco/};][]{2009SoPh..257..125B}. However, these automated catalogues have their limitations. For example CACTus imposes a zero acceleration, while SEEDS and ARTEMIS employ only LASCO/C2 data. The motivation thus exists to develop a new automated CME detection catalogue that overcomes such drawbacks, and indeed methods of multiscale analysis have shown excellent promise for achieving this \\citep{2009A&A...495..325B}. In this paper we discuss a new coronal image processing (CORIMP) technique for detecting and tracking CMEs. We outline our application of an automated multiscale filtering technique, to remove small scale noise/features and enhance the larger scale CME in single coronagraph frames. This allows the CME structure to be detected with increased accuracy for deriving the event kinematics and morphology. A companion paper \\citep[][hereafter referred to as Paper~\\RNum{1}]{2012ApJ...752..144M} outlines the steps used in preprocessing the coronagraph data with a normalizing radial graded filter \\citep[NRGF;][]{2006SoPh..236..263M} and deconvolution technique for removing the quiescent background features, leading to a very clean input for the automatic CME detection algorithm. These image processing steps are based on ideas first developed by \\citet{2010ApJ...711..631M}, where a more rudimentary approach was taken to isolate the dynamic component of coronagraph images. The new methods of Paper~\\RNum{1}, in conjunction with those outlined here, have led to a significant improvement in our ability to automatically detect and track CMEs in coronagraph data, such that a wealth of information on their structure and evolution may be obtained. In Section~\\ref{sect_automation} we outline the multiscale filtering techniques employed, and our method of automatically detecting and tracking CMEs in coronagraph images. In Section~\\ref{sect_data} the effectiveness of the CORIMP algorithms is demonstrated through their application to sample cases of LASCO and SECCHI data, and to a model corona with CMEs of known morphology. In Section~\\ref{sect_conclusions} we discuss the results and conclusions from the techniques. ", "conclusions": "\\label{sect_conclusions} The main objective in implementing an automated detection and tracking routine is to output reproducible, robust, accurate CME measurements (height, width, position angle, etc.). Current methods of CME detection have their limitations, mostly since these diffuse objects have been difficult to identify using traditional image processing techniques. These difficulties arise from the transient nature of the CME morphology, the scattering effects and non-linear intensity profile of the surrounding corona, the presence of coronal streamers, and the addition of noise due to cosmic rays and solar energetic particles (SEPs) that impact the coronagraph detector, along with instrumental effects of stray light, the limitations imposed by low cadence observations, and data corruption or dropouts. In the introduction to this paper, the drawbacks of current cataloguing procedures for investigating CME dynamics (CDAW, CACTus, SEEDS, ARTEMIS) were highlighted as the motivation for establishing a new catalogue. However, given the highly variable nature of CME phenomena and the coronal atmosphere they traverse, there are certain limitations that can never be overcome but only minimized; and it is exactly such a minimizing of current limitations that these new CORIMP methods achieve. The methods are completely automated, making them robust and reproducible - important for back-dating the full LASCO dataset and inspecting the statistics across thousands of events. The automated detection has been extended through both the LASCO/C2 and C3 fields-of-view without any need for differencing, thus minimizing the issues of under-sampling events and of the uncertainty involved when subtracting and scaling images. The multiscale filtering technique reveals the CME structure and so minimizes the uncertainty in determining their often complex geometry. The number of scales in the multiscale decomposition also allows a strength of detection to be assigned through both the magnitude and angular information, thus minimizing the chances that a CME, or parts thereof, go undetected. Furthermore, the spread of measurements available for inspection of the CME kinematics minimizes the uncertainty involved when deriving velocity and acceleration profiles, which is important for comparing with physical theory of CME propagation. Indeed, the overall CORIMP method of automatically detecting, tracking, and deriving CME parameters has been described and demonstrated here on a number of well-conceived models, and real data, with excellent results." }, "1207/1207.6639_arXiv.txt": { "abstract": "We present the Survey for High-z Absorption Red and Dead Sources (SHARDS), an ESO/GTC Large Program carried out with the OSIRIS instrument on the 10.4m Gran Telescopio Canarias (GTC). SHARDS is an ultra-deep optical spectro-photometric survey of the GOODS-N field covering 130~arcmin$^2$ at wavelengths between 500 and 950~nm with 24 contiguous medium-band filters (providing a spectral resolution R$\\sim$50). The data reach an AB magnitude of 26.5 (at least at a 3$\\sigma$ level) with sub-arcsec seeing in all bands. SHARDS main goal is obtaining accurate physical properties of intermediate and high-z galaxies using well-sampled optical SEDs with sufficient spectral resolution to measure absorption and emission features, whose analysis will provide reliable stellar population and AGN parameters. Among the different populations of high-z galaxies, SHARDS principal targets are massive quiescent galaxies at z$>$1, whose existence is one of the major challenges of current hierarchical models of galaxy formation. In this paper, we outline the observational strategy and include a detailed discussion of the special reduction and calibration procedures which should be applied to the GTC/OSIRIS data. An assessment of the SHARDS data quality is also performed. We present science demonstration results about the detection and study of emission-line galaxies (star-forming objects and AGN) at z$=$0--5. We also analyze the SEDs for a sample of 27 quiescent massive galaxies with spectroscopic redshifts in the range 1.0$<$z$\\lesssim$1.4. We discuss on the improvements introduced by the SHARDS dataset in the analysis of their star formation history and stellar properties. We discuss the systematics arising from the use of different stellar population libraries, typical in this kind of studies. Averaging the results from the different libraries, we find that the UV-to-MIR SEDs of the massive quiescent galaxies at z$=$1.0--1.4 are well described by an exponentially decaying star formation history with scale $\\tau$$=$100--200~Myr, age around 1.5-2.0~Gyr, solar or slightly sub-solar metallicity, and moderate extinction, A(V)$\\sim$0.5~mag. We also find that galaxies with masses above M$^\\ast$ are typically older than lighter galaxies, as expected in a downsizing scenario of galaxy formation. This trend is, however, model dependent, i.e., it is significantly more evident in the results obtained with some stellar population synthesis libraries and almost absent in others. ", "introduction": "\\label{sect:intro} The current paradigm of galaxy formation establishes that the baryons closely follow the evolution of the Cold Dark Matter (CDM) halos, which cluster and grow hierarchically as shown in cosmological simulations and semi-analytical models (such as those in \\citealt{2005Natur.435..629S}; see also \\citealt{1998ApJ...498..504B,1999MNRAS.310.1087S,2000MNRAS.319..168C, 2008MNRAS.391..481S,2010A&A...518A..14R}). In this scenario, star formation started within the cooling gas clouds in merging dark matter halos after a relatively slow early collapse regulated by feedback processes. This early star formation produced relatively small disk systems that later merged and generated larger (i.e., more massive) spheroidal systems \\citep[see, e.g.,][]{1993MNRAS.264..201K,2000A&G....41b..10E,2001ApJ...560L.119E} The global picture about the co-evolution of matter in the Universe (including all graviting components: CDM and baryons) is self-consistent and has been successful in reproducing and even predicting many observables about galaxy evolution, especially at low redshift. Among the most relevant successes, we find the good comparison of models with the observed power spectrum of the Cosmic Microwave Background \\citep{2007ApJS..170..377S,2011ApJS..192...18K} or the Large Scale Structure of the Universe \\citep{2001MNRAS.327.1297P,2001MNRAS.328.1039C}. Also very convincing is the link between observations and theoretical expectations such as the existence and properties of the acoustic baryonic oscillations \\citep{2005ApJ...633..560E,2009MNRAS.399.1663G,2010MNRAS.401.2148P}. In addition, the hierarchical scenario for galaxy formation is also supported by the observations of galaxy mergers at different cosmological distances \\citep[e.g.,][]{1999ApJ...520L..95V,2005AJ....130.2647V,2006ApJ...650L..29M}, and the increase of the fraction of galaxies undergoing mergers as we move to higher redshifts \\citep[among others,][]{1993MNRAS.262..627L,2000MNRAS.311..565L,2003AJ....126.1183C, 2006ApJ...652..270B,2008ApJ...672..177L,2009A&A...501..505L}. However, the hierarchical picture contrasts with several observational evidences, especially at high redshift (z$>$1--2), where a more classical monolithic collapse is favored. This formation path was proposed 50 years ago to explain the origin of bulges such as the Milky Way's and spheroidal galaxies. This would be done through a free-fall rapid collapse causing the formation of the bulk of the stars in these systems in a short period of time. Later, the star formation is shut off by some {\\it quenching} phenomena, and the galaxy henceforth evolves passively \\citep{1962ApJ...136..748E,1974MNRAS.166..585L}. This theory was largely abandoned due, first, to the compilation of evidence supporting that spheroidal galaxies suffer merging episodes \\citep{1972ApJ...178..623T}. Furthermore, globular clusters and the general stellar population in the MW present a relatively wide range of ages \\citep[e.g.,][]{1978ApJ...225..357S}, directly pointing out to the hierarchical scenario. Eventually, the hierarchical picture was adopted instead of the monolithic collapse due to the high degree of success of the $\\Lambda$CDM models and semi-analytic models mentioned above. Nevertheless, rapid early episodes of intense star formation are indeed consistent (although not uniquely) with observational facts in nearby galaxies, such as the dominant old stellar populations in bulges and ellipticals, their metallicity and $\\alpha$-elements enhancement, and the dynamics and shape of these systems \\citep[e.g., ][]{1997ApJS..111..203V,1997AJ....114.1771F,2000AJ....119.1645T, 2000AJ....120..165T}. In addition, hierarchical models still present severe drawbacks in several aspects. The most challenging observational facts for hierarchical models refer to the lightest and heaviest galaxies. Indeed, hierarchical models typically present a ``missing satellite problem'', i.e., they predict many more low-mass galaxies than what is actually observed \\citep[see][]{1993MNRAS.264..201K,1999MNRAS.303..188K,1999ApJ...522...82K, 2004ApJ...609..482K,2009MNRAS.398.2177L,2012ApJ...752L..19Q}. On the bright end, models tend also to overpredict the number of massive galaxies observed in the local Universe, although they are getting closer to the observations after taking into account quenching mechanisms \\citep{2006MNRAS.365...11C,2007MNRAS.375....2D,2008MNRAS.391..481S, 2010A&A...518A..14R,2010MNRAS.404.1111G}. The discrepancies between the predictions of current galaxy formation models based on the $\\Lambda$CDM paradigm and the data are more obvious as we move to higher redshifts. In the last 15 years, a wide variety of papers using very heterogeneous data and methods have presented compelling evidence that the formation of galaxies follows a so-called {\\it downsizing} scenario \\citep{1996AJ....112..839C,2004Natur.428..625H,2004Natur.430..181G, 2005ApJ...621L..89B,2005ApJ...630...82P,2007A&A...476..137A,2008ApJ...675..234P}. In this theory, the most massive galaxies formed first in the history of the Universe, thus having the oldest stellar populations seen today. The formation of less massive systems continued at lower redshifts. Downsizing implies that the bulk of the star formation in the most massive galaxies happened quick and stopped for some reason in early times. This also means that there should be massive passively evolving galaxies at high redshift. This kind of objects have indeed been detected at redshifts around z$\\sim$1--3 with a variety of techniques \\citep{2000AJ....120..575Y2,2003ApJ...587L..79F, 2004ApJ...617..746D,2006ApJ...640...92P,2008A&A...482...21C}. The finding of massive galaxies at z$=$1--3, some of them already evolving passively, is indeed extremely challenging for current models of galaxy formation based on the $\\Lambda$CDM paradigm. Indeed, models predict many less massive systems at high-z than observed \\citep[see, e.g.,][]{2007MNRAS.381..962C,2009ApJ...701.1765M,2012MNRAS.421.2904H, 2012ApJ...744..159L}. In contrast, the downsizing scenario contradicts, at least at first sight, the predictions of a hierarchical assembly of the stellar mass in galaxies, i.e., the most massive galaxies do not seem to be the result of multiple mergers occurring in a extended period along the Hubble time \\citep{1998ApJ...498..504B,2000MNRAS.319..168C, 2007ApJ...665..265F}. Still, a hierarchical assembly with (maybe multiple) mergers occurring at high redshift between gas-rich systems (a process close in nature to a monolithic collapse) would be consistent with both the evidences for downsizing and the properties of the dominant stellar populations seen in nearby spheroidal systems \\citep[e.g.][]{2009Natur.457..451D}. From the observational point of view, our understanding of the processes involved in the early (z$>$1) assembly of galaxies (and also the evolution from the early Universe to the present) is still hampered by the significant (often systematic) uncertainties in our estimations of their physical properties. Our global picture of galaxy formation will only improve if we are able to get more robust estimations of some key properties of galaxies, such as the stellar masses, SFRs, and extinctions. Jointly with those, we of course need better estimations of the distances to galaxies based on spectroscopic or photometric redshifts, which can be used to relate the mentioned galaxy properties with other relevant parameters such as the environment. The improvements in the determination of stellar masses and SFRs/extinctions will also mean a better estimation of the age of the stellar population and the Star Formation Histories, SFH \\citep[see, e.g,][]{2001ApJ...559..620P,2006A&A...459..745F, 2008ApJ...677..219K,2008A&A...477..503E,2012MNRAS.tmp.2944P,2012MNRAS.421.2002P}. Jointly with this observational effort, models should also be improved, including better physics. For example, models are still to provide more certain emissivities of the stellar populations in the rest-frame NIR, now affected by strong uncertainties due to limitations of about knowledge about the properties and importance of stellar evolutionary phases such as the thermally-pulsating TP-AGB phase \\citep[see][]{2005MNRAS.362..799M,2010ApJ...722L..64K}. The task of obtaining more robust physical parameters of galaxies at cosmological distances is even more interesting for those massive galaxies which have already reached a quiescent state and are evolving passively at high-z, whose number densities and properties are the most demanding challenges for current galaxy evolution models. The cosmological importance of these systems is very high, since they most probably represent the early formation phases of present-day early-type galaxies. In this paper, we present the basics of the {\\it Survey for High-z Absorption Red and Dead Sources} (SHARDS), an ESO/GTC Large Program awarded 180 hours of GTC/OSIRIS time during 2010-2013. This project consists of an ultra-deep (m$<$26.5~AB mag) imaging survey in 24 medium-band filters covering the wavelength range between 500 and 950~nm and targeting the GOODS-N field. The observations carried out by SHARDS allow to accurately determine the main properties of the stellar populations present in these galaxies through spectro-photometric data with a resolution R$\\sim$50, sufficient to measure absorption indices such as the D(4000) \\citep[e.g.,][]{1983ApJ...273..105B,1999ApJ...527...54B, 2003MNRAS.341...33K,2011ApJ...743..168K} or \\Mg\\, index \\citep{1997ApJ...484..581S,2004ApJ...614L...9M,2005MNRAS.357L..40S, 2005ApJ...626..680D,2008A&A...482...21C}. The analysis of these spectral features is a powerful method to constrain the solutions of stellar population synthesis models and to improve our estimations of parameters such as the age, SFH, mass, and extinction of galaxies at cosmological distances. SHARDS inherits the observational strategy of past and on-going optical surveys such as COMBO17 \\citep{2001A&A...377..442W,2003A&A...408..499W}, the COSMOS medium-band survey \\citep{2009ApJ...690.1236I}, ALHAMBRA \\citep{2008AJ....136.1325M}, and PAU/J-PAS \\citep{2009ApJ...691..241B,2011arXiv1108.2657A}. These projects have demonstrated the impact of large photometric datasets on our understanding of the formation of galaxies (see, among many papers, \\citealt{2004ApJ...608..752B,2007ApJ...665..265F,2004A&A...421..913W, 2004ApJS..152..163R,2006A&A...453..869B,2006ApJ...637..727C, 2007ApJS..172..150S,2009ApJ...692L...5B,2010ApJS..189..270C,2011ApJ...735...86W}). SHARDS intends to be a step forward from these surveys in terms of depth, spectral resolution, and data quality. Our survey prioritizes the detailed study of the faintest galaxies at the highest redshifts over the analysis of closer galaxy populations and the Large Scale Structure at intermediate redshift, and thus focuses on a smaller area than the surveys mentioned above. Indeed, SHARDS was planned to reach up to 3 mag fainter than those surveys, uses typically twice the number of filters in the same wavelength range (i.e., our spectral resolution is better), and was obtained in excellent (sub-arcsec) seeing conditions with a 10m class telescope. In contrast, it covers a fraction of the area surveyed by other projects. In this paper, we present the main technical characteristics of the survey in Section~\\ref{sect:shards_description}, including a thorough discussion of the reduction and calibration procedures in Section~\\ref{sect:shards_data}. Next, we present our science verification results about emission-line and absorption systems. In Sections~\\ref{sect:elgs} and \\ref{sect:laes}, we discuss about the ability of the SHARDS data to select and study emission-line sources (star-forming galaxies and AGN) at intermediate (z$<$1) and high redshifts (up to z$\\sim$5 and beyond). In Section~\\ref{sect:redanddead}, we present detailed spectral energy distributions of massive quiescent galaxies at z$>$1, and demonstrate the power of our spectro-photometric data to analyze the stellar populations in this kind of object through a detailed comparison with stellar population synthesis models. Throughout this paper we use AB magnitudes. We adopt the cosmology $H_{0}=70$ km~s$^{-1}$Mpc$^{-1}$, $\\Omega_{m}=0.3$, and $\\Omega_{\\lambda}=0.7$. ", "conclusions": "" }, "1207/1207.7229_arXiv.txt": { "abstract": "{} {Hot Jupiters are thought to belong to single-planet systems. Somewhat surprisingly, some hot Jupiters have been reported to exhibit transit timing variations (TTVs). The aim of this paper is to identify the origin of these observations, identify possible periodic biases leading to false TTV detections, and refine the sample to a few candidates with likely dynamical TTVs.} {We present TTV frequencies and amplitudes of hot Jupiters in $Kepler$ Q0--6 data with Fourier analysis and a frequency-dependent bootstrap calculation to assess the false alarm probability levels of the detections.} {We identified 36 systems with TTV above four standard deviation confidence, about half of them exhibiting multiple TTV frequencies. Fifteen of these objects (\\object{HAT-P-7b}, \\object{KOI-13}, 127, 183, 188, 190, 196, 225, 254, 428, 607, 609, 684, 774, 1176) probably show TTVs due to a systematic observational effect: long cadence data sampling is regularly shifted transit-by-transit, interacting with the transit light curves, introducing a periodic bias, and leading to a stroboscopic period. For other systems, the activity and rotation of the host star can modulate light curves and explain the observed TTVs. By excluding the systems that were inadequately sampled, showed TTV periods related to the stellar rotation, or turned out to be false positives or suspects, we ended up with seven systems. Three of them (KOI-186, 897, 977) show the weakest stellar rotation features, and these are our best candidates for dynamically induced TTV variations.} {Those systems with periodic TTVs that we cannot explain with systematics from observation, stellar rotation, activity, or inadequate sampling may be multiple systems or even exomoon hosts.} ", "introduction": "Transit timing tariation (TTV) is a major diagnostics of various system parameters of extrasolar planets \\citep{holman05,agol05}. In multiplanet systems, planets perturb each other, leading to correlated TTVs of them \\citep{holman10, lissauer11}. TTVs have also uncovered the presence of further non-transiting planets in planetary systems \\citep{ballard11, ford12b}. According to our current view, hot Jupiters occur as single planets, since they have been not detected in multiplanet systems. This picture suggests that hot Jupiters occupy unperturbed orbits, hence their orbital motion is Keplerian, and they exhibit strictly periodic transit times. In contrast to this picture, current literature reports a considerable number of hot Jupiters with TTV, which are often periodic \\citep{steffen12a,ford12a}. In this work we publish TTV periods, TTV amplitudes, and significance levels for hot-Jupiter candidates in $Kepler$ data. Transit times covering Q0--Q6 were published in a catalog by \\cite{ford12a}\\footnote{http://www.astro.ufl.edu/$\\sim$eford/data/kepler/} (see \\cite{steffen12b} for more details), whose data set has become the basis of several TTV studies). The main aims of this study are to critically revise the \\cite{ford12a} catalog and to look for possible nondynamical processes that may cause virtual TTVs. To this end, we analyzed the stellar variations of the systems in the original $Kepler$ data up to Q9 besides exploring the catalog itself. The main conclusions of this study are the following. \\begin{enumerate} \\item{} Timing data of some already published systems with suspected TTV can be satisfactorily explained by nondynamical reasons, such as stroboscopic period due to even sampling or light curve distortions due to stellar activity; \\item{} About 2\\%{} of the Jupiter-size candidates passed all tests and show a TTV that presently has an unknown origin, and obviously needs further studies with follow-ups. \\item{} A fraction of systems with TTV signals tend to exhibit multiple TTV periods (confirmed by \\cite{mazeh13}) that are incompatible with sampling or stellar rotation effects. \\end{enumerate} We briefly introduce the most exciting systems and discuss the possible sources of TTVs for these planets. \\begin{figure*} \\centering\\includegraphics[bb=75 450 740 640,width=12cm]{modulation1.ps} \\centering\\includegraphics[bb=75 475 740 640,width=12cm]{modulation2.ps} \\caption{Illustration of the detection of a super-Nyquist modulation in the orbital motion with Fourier analysis. Top panel: one quarter (90-d) long detail from a 600-d long simulation. Dots represent sampling by transits. Bottom panel: Fourier series of the 600 d long data set, extended beyond the Nyquist frequency. The actual frequency and the sub-Nyquist detections are highlighted. This specific model had an orbital (sampling) and modulation period of 11.21 and 4.76 days (which were taken from the ZIP code and altitude-above-sea-level of the Budapest station of Konkoly Observatory, as uncorrelated random numbers).} \\label{Nyquist} \\end{figure*} ", "conclusions": "" }, "1207/1207.5530_arXiv.txt": { "abstract": "We report on Chandra, RXTE, Swift/BAT and MAXI observations of a $\\sim$1 day X-ray flare and subsequent outburst of a transient X-ray source observed in October--November 2011 in the globular cluster Terzan~5. We show that the source is the same as the transient that was active in 2000, i.e., the neutron star low-mass X-ray binary EXO~1745--248. For the X-ray flare we estimate a 6--11 hr exponential decay time and a radiated energy of $2-9 \\times 10^{42}$ erg. These properties, together with strong evidence of decreasing blackbody temperature during the flare decay, are fully consistent with what is expected for a thermonuclear superburst. We use the most recent superburst models and estimate an ignition column depth of $\\approx 10^{12}$ g cm$^{-2}$ and an energy release between $0.1-2 \\times 10^{18}$ erg g$^{-1}$, also consistent with expected superburst values. We conclude therefore that the flare was most probably a superburst. We discuss our results in the context of theoretical models and find that even when assuming a few days of low level accretion before the superburst onset (which is more than what is suggested by the data), the observations of this superburst are very challenging for current superburst ignition models. ", "introduction": "\\label{sec:intro} Thermonuclear Type-I X-ray bursts are caused by unstable burning of a several meters thick layer of accreted H/He on the surface of neutron stars (NSs) in low-mass X-ray binary (LMXB) systems \\cite[e.g.][]{Lewin93}. Manifesting themselves as a sudden (seconds) increase in the X-ray luminosity and reaching levels that can be many times brighter than the persistent (accretion) luminosity, typical bursts emit about $10^{39}-10^{40}$ ergs, last seconds to minutes, and have light curves that are well described by a fast-rise exponential-decay profile. Their spectra are generally consistent with a blackbody temperature $T_{\\rm bb}=2$--3~keV, where $T_{\\rm bb}$ increases until the burst peak, and then decreases exponentially. This is naturally interpreted as heating resulting from the initial fuel ignition, followed by cooling of the ashes \\citep[and additional hydrogen burning through a series of rapid proton captures and $\\beta$-decays, e.g.,][]{Schatz01} once the main available fuel is exhausted. Type-I X-ray bursts are a common phenomenon in NS-LMXBs. They have been observed in about 100 sources and, depending on the conditions (e.g, accretion rate, composition of the fuel, etc.) can have recurrence times between minutes and weeks \\cite[e.g.][]{Galloway08,Linares12}. Superbursts are a class of extremely long-duration bursts which are attributed to the unstable thermonuclear burning of a $\\sim$100 meter thick carbon-rich layer, formed from the ashes of normal Type-I X-ray bursts \\citep{Cumming01a}. Superbursts tend to quench the regular Type-I bursts for weeks afterwards, probably because the cooling flux from the superburst temporarily stabilizes the H/He burning \\citep{Cumming01a,Cumming04,Keek12}. The difference in fuel composition between Type-I X-ray bursts and superbursts leads to a clear difference in time scales, recurrence times and energetics, where superbursts last for a few hours, recur every one-to-few years and emit $10^{41}-10^{42}$ ergs. With such long recurrence times superbursts are difficult to catch. While thousands of Type-I X-ray bursts have been observed \\citep[e.g.][]{Galloway08}, to date only about 22 (candidate) superbursts have been observed from 13 sources \\citep[see, e.g., ][and references therein]{Wijnands01, Kuulkers04, Keek11, Altamirano11f, Chenevez11, Mihara11, Asada11}. Terzan 5 is a globular cluster containing ~50 known X-ray sources, of which $\\sim$12 are likely LMXBs containing neutron stars \\citep[e.g.][]{Heinke06}. During 2011 we monitored Terzan 5 on a weekly basis with \\textit{Rossi X-ray Timing Explorer} (RXTE) observations to search for transient X-ray flares and/or outbursts. At 4:57 UT, October 26th 2011, an RXTE pointed observation measured a 2--16 keV intensity of $\\sim$8 mCrab, significantly above the typical quiescent intensity of $\\sim$2 mCrab \\citep{Altamirano11e}. Approximately 8 hours earlier, INTEGRAL monitoring observations of Terzan 5 did not detect any enhanced activity, with a $5\\sigma$ upper limit of 6 mCrab in the 3--10 keV energy band \\citep{Vovk11}. The RXTE detection was confirmed by the Swift/BAT daily-averaged flux measurements \\citep{Altamirano11e}, as well as by a Swift/XRT pointed observation performed $\\sim$11 hours after the RXTE one \\citep{Altamirano11f}. The position of the source from these Swift/XRT data was consistent \\citep{Altamirano11f,Evans11} with that of the transient NS-LMXB that was active in 2000 \\citep[which we refer to as EXO 1745--248, though it is not necessarily the EXOSAT source; see][]{Markwardt00, Wijnands05a}. This result was later confirmed by a preliminary analysis of a pointed Chandra observation \\citep{Pooley11}. Just before the INTEGRAL non-detection, MAXI and Swift/BAT light curves of Terzan 5 revealed an X-ray flare that lasted less than a day. We identified this flare as a possible superburst based on its duration, shape of its light curve and estimated radiated energy of $\\sim$10$^{42}$ ergs \\citep{Altamirano11f}. Our speculations were supported by the results of \\citet{Mihara11} who used the MAXI data and showed that (i) the spectra of the flare were well modeled with a blackbody component at $\\sim$2--3~keV and that (ii) there was an apparent decrease of the black-body temperature, which is usually interpreted as the cooling of the neutron star surface after a thermonuclear burst \\citep[see, e.g.,][]{Lewin96}. Very recently, \\citet{Serino12} have presented a detailed analysis of the MAXI data supporting the superburst identification. The occurrence of a superburst in the transient NS-LMXB 4U~1608--522 after 55 days of low ($\\lesssim10$\\% Eddington) accretion rate has challenged superburst theory, as it is difficult to explain carbon ignition at the observed depths when the NS surface is still cool \\citep{Keek08}, i.e., when accretion has not yet been able to ``warm up'' the NS. The superburst candidate in Terzan 5 is even more challenging for theoretical models, as the NS is very cool \\citep{Wijnands05a, Degenaar12} and the superburst onset was coincident with a period of only low-level accretion, or no accretion at all. \\begin{figure} \\centering \\resizebox{1\\columnwidth}{!}{\\rotatebox{0}{\\includegraphics{./f1.eps}}} \\caption{39\" x 52\" Chandra images of Terzan 5 from different epochs show the 2011 outburst source is EXO 1745--248. The upper-left panel shows the combined image of two observations (July 24 and 29, 2000, ObsIDs 655 and 644, respectively) for a total time of 47 ksec \\citep[see][]{Heinke03a}. The upper-right panel shows a 35-ksec observation on July 13th, 2003 \\citep[ObsID 3798; see][]{Wijnands05a,Heinke06}. The lower left panel is a 10 ksec observation on Oct. 24, 2010 \\citep[ObsID 11051;][]{Pooley10}, and the lower right panel is our 9.8 ksec observation on Nov. 3, 2011 \\citep[ObsID 12454;][]{Pooley11}. All images were extracted in the 1-3 keV energy range (chosen to try to maximize S/N in the 2011 image). Ten X-ray sources from \\citet{Heinke06} are marked with red circles (2-pixel -- 0.984\" radius). Diamonds mark the position of EXO~1745--248 as detected in its 2000 outburst. The active source in the 2010 observation \\citep{Pooley10} was the 11~Hz pulsar IGR~J17480--2446 \\citep{Strohmayer10a,Papitto11} }\\label{fig:chandra} \\end{figure} ", "conclusions": "\\label{sec:discussion} We present Chandra, RXTE, Swift/BAT and MAXI data of the X-ray flare and subsequent outburst of EXO~1745--322 in Terzan~5. We show that the active source is the same as that active in 2000 and that the characteristics of the flare are consistent with what is expected for a superburst. We also show that the outburst may have started just before the superburst onset, although our results are not conclusive due to systematics in the data. The Swift/BAT peak in the superburst flux was delayed by about 0.5 days compared to the flux peak on the MAXI data. Similar delays between soft and hard energy bands have already been seen in Type-I X-ray bursts \\citep[order of seconds, see, e.g.,][]{Lewin93,Falanga08,Chelovekov11} and in at least one superburst \\citep[about $\\sim$1000 sec in the LMXB 4U 1820--30, see, e.g.,][]{Intzand10}. These delays have been interpreted as due to photospheric-radius expansion (PRE) bursts, where the X-ray intensity first peaks in the low-energy band and later X-rays become visible at higher energies \\citep[see, e.g., ][]{Lewin93,Intzand10}. The $\\sim1000$ sec duration of the PRE phase in the superburst observed in the LMXB 4U 1820-30 is already at the limit of what current superburst models can explain. Irrespective of the mechanism, the delay is by far the largest. The fact that it is so much larger, may raise the question of its origin being the same as that proposed for Type-I PRE X-ray bursts. In Section~\\ref{sec:lc} we show marginal evidence that EXO 1745--248 may have been detected before the peak of the superburst. In the rest of this section we will discuss the implications of our results on superburst theory taking into account both the possibilities that the outburst started a few days before, or approximately a day after the peak of the superburst. For a discussion on how the superburst emission may have affected the accretion disk to trigger the subsequent outburst, we refer the reader to \\citet{Serino12}. We note that if the pre-superburst detections of the source are real, then the superburst most probably momentarily affected the normal outburst evolution \\citep[see, e.g., ][for the study of the evolution of the accretion disk around the NS system 4U 1820--30 during a superburst]{Ballantyne04}. \\begin{table} \\begin{centering} \\begin{tabular}{lll} \\hline & EXO~1745--248 & 4U~1254-690\\tabularnewline \\hline $\\tau_{\\mathrm{exp}}(\\mathrm{hr})$ & $6-11$ & $6\\pm0.3$\\tabularnewline $E_{\\mathrm{b}}(10^{42}\\,\\mathrm{erg})$ & $2-9$ & $0.8\\pm0.2$\\tabularnewline $\\log(y/(\\mathrm{g\\, cm^{-2}}))$ & $12.0\\pm0.3$ & $12.4$\\tabularnewline $E_{18}(10^{18}\\,\\mathrm{erg\\, g^{-1}})$ & $>0.1$ & $0.15$\\tabularnewline \\hline \\end{tabular} \\par\\end{centering}\\caption{Comparison to the superburst of 4U~1254-690 \\citep{Intzand03,Cumming06a}.}\\label{table:table} \\end{table} \\subsection{Comparison to other superbursts and theoretical implications} Previously, the longest and most energetic superburst known from a hydrogen accreting source was from 4U 1254--690 \\citep{Intzand03a}. Unlike for that superburst, we did not observe the start of the superburst from EXO~1745--248, resulting in large uncertainties in the superburst properties. The superburst of EXO~1745--248 is \\emph{at least} of equal duration, and twice as energetic (Table~\\ref{table:table}). The largest values of the bolometric radiated energy, $E_{\\mathrm{b}}$, consistent with the observations, are close to the predicted maximum radiated energy for a superburst set by neutrino emission (\\citealt{Keek11}; see also \\citealt{Cumming06a}). The decay time, $\\tau_{\\mathrm{exp}}$, depends on the thickness of the cooling layer, and, therefore, on the ignition column depth, $y_{\\mathrm{ign}}$. For 4U~1254--690 the depth was determined using the instantaneous burning model, yielding a depth comparable to the larger values in the range we derive for EXO~1745--248, which are also favored by our fits with the same model \\citep{Cumming06a}. This suggests that the ignition depths and decay times of the two superbursts likely have similar values. The larger $E_{\\mathrm{b}}$ for EXO~1745--248 can be explained by the burning of more carbon-rich material, which is accommodated by the larger values in the range found for the specific energy release, $E_{18}$. Therefore, this is the most energetic and possibly the longest superburst observed to date. Most superbursting sources, including 4U~1254--690, are observed to accrete continuously at a high rate of around $10\\%$ of the Eddington limited rate $\\dot{M}_{\\mathrm{Edd}}=2\\times10^{-8}\\, M_{\\odot}\\mathrm{yr}^{-1}$ \\citep[for solar composition and a $10\\,\\mathrm{km}$ radius; e.g.,][]{Keek08a}. The high rate ensures a hot outer crust, forces unstable ignition of the carbon, and may be necessary for the production of a mixture of carbon and heavy isotopes that is thought to be the fuel for superbursts (\\citealt{Cumming01a}; see also \\citealt{Cooper09}). Both sufficient heat and carbon are required for superburst ignition. This scenario was challenged by the observation of the superburst from the NS-transient 4U~1608--522, which occurred only 55 days after the onset of an accretion outburst. \\citet{Keek08} showed that the neutron star envelope does not heat up quickly enough to explain the ignition of runaway carbon burning. In the past year three more superbursts have been detected from transient sources, including the one discussed in this paper \\citep[for the other detections see][]{Chenevez11,Asada11}. Even if we assume that EXO~1745--248 started accreting at an increased rate 0.5 days or even a few days before the superburst (but at levels undetected by MAXI and Swift/BAT), the time is much too short for the envelope at the derived ignition depths to heat up from either thermonuclear burning in the envelope or from nuclear processes in the inner crust. Therefore, sufficient heat must have been generated at the superburst ignition depth within this short time interval. Currently there is no known process that could provide this. The case for a substantial additional heat source in the outer crust (close to the superburst ignition depth) has also been made from observations of crustal cooling after outbursts in long-duration transients \\citep{Brown09}. The 0.5 days time scale that we find, however, puts strong constraints on the immediacy with which this heating process must take place. \\subsection{On the Carbon production} But where does the carbon-fuel necessary for a superburst come from? Hydrogen-accreting superbursters display a high ratio of the persistent fluence between two (Type-I X-ray) bursts to the burst fluence \\citep{Intzand03a}, indicating that apart from during Type-I X-ray bursts, a substantial fraction of the accreted hydrogen and helium burns in a stable manner. This is thought to be a required process to produce the carbon fuel for superbursters \\citep{Schatz03}, and it is observed to occur close to an accretion rate of $10\\%\\,\\dot{M}_{\\mathrm{Edd}}$, i.e., the rate inferred for most superbursters. During the outburst in 2000, EXO~1745--248 accreted at a comparable rate of on average $17\\%\\,\\dot{M}_{\\mathrm{Edd}}$ for two months \\citep{Degenaar12}, during which there were bursts as well as periods without bursts. In fact, because the quiescent luminosity is over a factor $10^{4-5}$ lower \\citep[$L_x$ in quiescence is $5-7 \\times 10^{32}$ erg s$^{-1}$, see][]{Degenaar12}, effectively all of the superburst fuel must have been created in such short outbursts. This conclusion is still valid even if we consider that at the above level of quiescent emission, EXO~1745--248's luminosity might vary by a factor of a few on timescales of hours-years \\citep[which may indicate that the accretion does not fully switch off in quiescence, but continues at a very low rates, see, e.g.,][]{Wijnands05a}. During the outburst in 2000, about $8\\%$ of the inferred $y_{\\mathrm{ign}}=1.0\\times10^{12}\\,\\mathrm{g\\, cm^{-2}}$ was accreted. Using the shortest suggested outburst recurrence time of $11\\,\\mathrm{yr}$ \\citep{Degenaar12}, a superburst recurrence time of $186\\,\\mathrm{yr}$ is inferred, but it may very well be longer (unless the outburst recurrence time is shorter, which would translate in shorter superburst recurrence times). Because of the low average luminosity, the neutron star is relatively cool, which reduced the carbon burning rate at the bottom of the accreted pile to allow for sufficient carbon to remain to trigger a thermonuclear runaway after such a long recurrence time. Of course, we cannot exclude the possibility that the superburst ignition conditions had been almost reached during the previous outburst, such that only a short accretion episode of a few days was required to set it off. Although not impossible, we find it improbable as the outer crust is expected to have reached a higher temperature by heating during the two-month outburst in 2000 (i.e. conditions more favorable for ignition), than after a few days in 2011. A more plausible scenario could be that the carbon-fuel necessary for a superburst was created mostly during outburst, and then concentrated during the long period of quiescence, as after accretion ceases, there is time for the light and heavy elements to separate out from each other \\citep[see, e.g.,][]{Brown02}. An additional potentially important process is chemical separation by freezing at the interface of the ocean and the outer crust \\citep[see][and references therein]{Medin11}. After the previous outburst there was plenty of time for carbon to separate out from iron and heavier isotopes, and so substantially increasing the carbon fraction at the bottom of the accreted column. If this scenario is correct, then it is possible to explain the $y_{\\mathrm{ign}}$ necessary in cases where superburst ignition occurs at early times of the outburst. However, it could be problematic for models, as pure carbon layers have a higher thermal conductivity and will remain colder than an impure carbon layers \\citep[see][]{Cumming01a}; moreover, upwards transport of carbon could make it harder for the carbon to reach ignition depth. In any case, still unexplained is how the neutron star temperature can rise so quickly at the start of the outburst to be able to ignite the superburst. The difficulties faced by superburst models that invoke carbon ignition may point to a different fuel for superbursts. In the analysis of bursts from the likely ultra compact X-ray binary 4U~0614+091, \\citet{Kuulkers10a} pointed out that in principle helium ignition could explain many of the observed column depths of superbursts. This would require accumulation of a deep and cold layer of helium on the star. Further theoretical work on this scenario is needed, but the fact that the superburst from EXO 1745--248 appears to have one of the largest ignition column depths of known superbursts may place it too deep for helium ignition. \\vspace{1cm} \\textbf{Acknowledgment:} We thank H. Krimm for his help on the Swift/BAT data and J.J.M. in't Zand for insightful discussions. RW acknowledges support from a European Research Council (ERC) starting grant. LK is supported by the Joint Institute for Nuclear Astrophysics (JINA; grant PHY08-22648), a National Science Foundation Physics Frontier Center. ND is supported by NASA through Hubble Postdoctoral Fellowship grant number HST-HF-51287.01-A from the Space Telescope Science Institute. AC, GRS, and COH are supported by Discovery Grants from the Natural Sciences and Engineering Research Council of Canada (NSERC) . COH is also supported by an Alberta Ingenuity New Faculty Award. JH acknowledges financial support through Chandra award GO1-12055B. AC and LK are members of an International Team in Space Science on type I X-ray bursts sponsored by the International Space Science Institute (ISSI) in Bern. DP gratefully acknowledges support provided by the National Aeronautics and Space Administration through Chandra Award Number GO1-12055A issued by the Chandra X-ray Observatory Center, which is operated by the Smithsonian Astrophysical Observatory for and on behalf of the National Aeronautics Space Administration under contract NAS8-03060. We thank the CXC, RXTE and Swift staff for their quick response to our ToO requests. This research has made use of the MAXI data provided by RIKEN, JAXA and the MAXI team." }, "1207/1207.2256_arXiv.txt": { "abstract": "Most stars form in a cluster environment. These stars are initially surrounded by discs from which potentially planetary systems form. Of all cluster environments starburst clusters are probably the most hostile for planetary systems in our Galaxy. The intense stellar radiation and extreme density favour rapid destruction of circumstellar discs via photoevaporation and stellar encounters. Evolving a virialized model of the Arches cluster in the Galactic tidal field we investigate the effect of stellar encounters on circumstellar discs in a prototypical starburst cluster. Despite its proximity to the deep gravitational potential of the Galactic centre only a moderate fraction of members escapes to form an extended pair of tidal tails. Our simulations show that encounters destroy one third of the circumstellar discs in the cluster core within the first 2.5\\,Myr of evolution, preferentially affecting the least and most massive stars. A small fraction of these events causes rapid ejection and the formation of a weaker second pair of tidal tails that is overpopulated by disc-poor stars. Two predictions arise from our study: (i)~If not destroyed by photoevaporation protoplanetary discs of massive late B- and early O-type stars represent the most likely hosts of planet formation in starburst clusters. (ii)~Multi-epoch $K$- and $L$-band photometry of the Arches cluster would provide the kinematically selected membership sample required to detect the additional pair of disc-poor tidal tails. ", "introduction": "\\label{sec:introduction} Observations in the past decade have shown that most young stars do not form in isolation but as part of a cluster environment \\citep[e.g.][]{2003ARA&A..41...57L,2003AJ....126.1916P,2009ApJS..181..321E}. The accretion discs of these stars are thus exposed to environmental effects that could affect their evolution \\citep{1998A&A...340..508R,2000prpl.conf..401H,2001MNRAS.325..449S,2004ApJ...611..360A,2006ApJ...642.1140O,2006ApJ...652L.129P,2006A&A...454..811P,2007A&A...462..193P,2007MNRAS.376.1350C,2010A&A...509A..63O}. As these discs are the prerequisites for the formation of planetary systems this process might be influenced by the cluster environment. However, so far theoretical investigations about the effect of irradiation by massive stars or strong gravitational interactions have concentrated on low- and intermediate mass star clusters like IC348 or the Orion Nebula Cluster (ONC). Much more massive systems like NGC~3603, the Arches cluster or Westerlund~1 - known as starburst clusters - that are expected to trigger the strongest effects have not been treated. Considering the huge amount of massive stars and extreme densities, i.e. more than 100 O-stars and a core density $\\gtrsim$$10^5\\,\\Msun\\,\\pcdens$ in the Arches cluster \\citep{1999ApJ...525..750F}, an extrapolation of the previous investigations towards starburst clusters would suggest that the lifetime of circumstellar discs could be shortened dramatically and so planet formation in such an environment hindered significantly. However, recent observations of the Arches cluster in the near-infrared $JKL'$ bands by \\citet{2010ApJ...718..810S} have revealed emission from circumstellar matter around at least a few percent of stars in the mass range $2-20\\,\\Msun$. This detection of discs in the Arches B-star population was surprising for two reasons. First, it is expected that the stars' UV radiation causes depletion of their own inner disc. Considering that the characteristic UV evaporation timescale of a primordial disc around Herbig Be stars is less than 1\\,Myr \\citep{2009A&A...497..117A}, a disc lifetime of 2.5\\,Myr for B-type stars implies that the self-photoerosion of discs is probably less efficient than various models suggest \\citep[e.g.][]{2000prpl.conf..401H,2008NewAR..52...60A}. Second, extrapolating from the above mentioned numerical studies of less massive clusters, in a starburst cluster environment disc destruction is expected to be much increased by external processes. The extreme UV radiation from numerous O-stars and gravitational interactions induced by the high stellar density could boost the removal of external disc material. In the present investigation we focus on the mechanism of encounter-induced disc-destruction to constrain its contribution to the overall disc life time. We use simulations of the Arches cluster preformed by the $\\starlab$ simulation package \\citep[][]{1996ASPC...90..413M,2001MNRAS.321..199P,2003IAUS..208..331H} and carry out the analysis of the disc-mass loss analogous to our previous publications referenced above. Hence this study of a massive cluster complements previous investigations of encounter-induced disc evolution in low- and intermediate-mass clusters \\citep{2001MNRAS.325..449S,2004ApJ...611..360A,2006ApJ...642.1140O,2006ApJ...652L.129P,2006A&A...454..811P,2007A&A...462..193P,2007MNRAS.376.1350C,2010A&A...509A..63O}. In contrast to the low- and intermediate-mass clusters in the solar neighbourhood the location of the massive Arches cluster near the Galactic centre exposes it to strong tidal fields. Therefore another difference to previous simulations is that our model includes the contribution of the Galactic tidal field. Stars escaping from a star cluster in the gravitational field of a galaxy form extended tidal tails known from various observations in the Milky Way \\citep{1995AJ....109.2553G,1997A&AS..121..439K,2000A&A...359..907L,2006ApJ...637L..29B} and numerical models \\citep{1999A&A...352..149C,2002ApJ...565..265P,2005AJ....129.1906C,2007MNRAS.380..749F}. Except for high-speed escapers created in few-body encounters, stars escape as a result of two-body encounters passing at slow speed close to the saddle points of the effective potential, known as Lagrange points \\citep{2008MNRAS.387.1248K}. With our study we aim to investigate the evolution of the encounter induced-disc mass loss in a tidally distorted starburst cluster and its tidal arms. Of particular relevance is the work of \\citet{2009MNRAS.392..969J} who use the epicycle theory for a quantitative analysis of the tidal tail structure of star clusters moving on a circular orbit in the Galactic disc. They find that the radial offset of a star's epicyclic motion relative to the orbit of the star cluster, $\\Delta{R_0}$, is first order in the angular momentum, $\\Delta{L}$, while its radial amplitude~$r_\\mathrm{m}$ depends on the energy excess, $\\Delta{E}$, which is of second order in~$\\Delta{L}$: \\begin{equation} \\label{eq:just09} \\begin{aligned} \\Delta{R_0} &\\propto \\Delta{L} \\,,\\\\ r_\\mathrm{m} &\\propto \\sqrt{\\Delta{E}} \\,. \\end{aligned} \\end{equation} \\citet{2009MNRAS.392..969J} conclude that the circumstances are more complicated near the Galactic Centre but to lowest order the theory is still applicable. Throughout this work we assume that initially all stars are surrounded by protoplanetary discs. This is justified by observations that reveal disc fractions of nearly 100\\,\\% in very young star clusters \\citep[e.g.][]{2000AJ....120.1396H,2000AJ....120.3162L,2001ApJ...553L.153H,2005astro.ph.11083H}. The typical disc diameter is a few hundred AU for low- and intermediate-mass stars \\citep{1996AJ....111.1977M,2008Ap&SS.313..119A}, though discs of more than several thousand AU surrounding massive stars have been observed \\citep[see][and references therein]{2005IAUS..227..135Z}. However, it remains unclear whether a clear correlation of disc extension and stellar mass does exist \\citep[see e.g.][]{2005A&A...441..195V}. In Section~\\ref{sec:observations_onc} we outline the observationally determined basic properties of the Arches cluster that serves as a reference for our cluster models. The computational method and the properties of the numerical models are described in Section~\\ref{sec:numerical_method}. Afterwards we present results from our numerical simulations in Section~\\ref{sec:numerical_results}. The conclusion and discussion mark the last section of this paper. ", "conclusions": "\\label{sec:conclusions} We have investigated the effect of stellar encounters on the evolution of protoplanetry discs using a numerical model of the Arches cluster, one of the most extreme Galactic star clusters observed so far. Our numerical method is based on a hybrid approach where we combine simulations of pure star cluster dynamics and a parameter study of isolated star-disc encounters.\\footnote{``star-disc encounters'' refer to the interaction of a star-disc system and a disc-less star and serve as a numerical simplification of the disc-disc encounters that occur in reality.} This approach involves necessarily some assumptions and simplifications that potentially mildly overestimate the effect of the encounter-induced perturbations of circumstellar material. Nonetheless, our work represents the most detailed study of the effect of stellar dynamics on the evolution of protoplanetary discs in a starburst cluster so far. We find that stellar interactions in the Arches cluster are dominated by hyperbolic fly-bys, with more than 70\\,\\% having eccentricities $\\varepsilon > 10$. This is much different in sparser clusters where nearly parabolic encounters are by far most frequent \\citep{2010A&A...509A..63O}. According to Eq.~\\eqref{eq:fit_function_mass_loss_vs_eccentricity} the normalized disc-mass loss in hyperbolic interactions with $\\varepsilon > 10$ is reduced by a factor $<$\\,0.38 compared to the parabolic case. Consequently, encounters drive the destruction of only about one third of the stellar discs in the cluster core until its present-day age of $\\sim$2.5\\,Myr. The process continues with reduced efficiency such that after 6\\,Myr half of the core population becomes disc-less. The extended period of disc destruction in a starburst cluster is again in stark contrast to sparser clusters in which it ceases after roughly 1\\,Myr due to significant cluster expansion \\citep{2010A&A...509A..63O}. Note that by neglecting primordial binaries in our models we underestimate the efficiency of disc destruction for three reasons: (i) the larger number of stars increases the probability for encounters, (ii) binaries increase the probability for strong few-body interactions, and (iii) discs are truncated in tight binaries. The inclusion of a Galactic tidal field in our model of the Arches cluster induces the growth of an extended pair of tidal tails. The fraction of disc-less stars in the tidal tails is nearly constant over time at $\\sim$5\\,\\% and hence much lower than in the cluster at ages later than 1\\,Myr. Only about half of the disc-less tidal tail stars have escaped slowly via the Lagrange points while the other half has been ejected at high speed from the cluster centre directly after disc destruction. These stars are the dominant population of a weaker and shorter second pair of ``disc-poor'' tidal arms at the present age of the cluster. With time the disc-poor tails mix with the dominant ``classical'' tails and form new features: a concentration of disc-less stars at the inner edge of the trailing arm and the outer edge of the leading arm. In fact, we expect that dynamical features arising from strong encounters -- such as the disc-poor tidal arms -- could be even more prominent than shown by our model. For one, as outlined in Section~\\ref{sec:numerical_results:cluster_dynamics}, our model of the Arches cluster in the Galactic tidal field is slightly underdense such that close encounters are underrepresented in the simulation. Second, we model a population of single stars only and ignore the dynamically important effect of binaries. Assuming the binary fraction in starburst clusters is potentially as high as in clustered star forming regions of lower mass \\citep[$\\sim$50\\,\\%: e.g.][]{2003ApJ...583..358B,2006A&A...458..461K}, the probability for strong few-body encounters would increase significantly \\citep{1975AJ.....80..809H,1975MNRAS.173..729H,1986Ap&SS.124..217A}, boosting the production of high-speed escapers and disc-less stars. However, one could argue that taking into account additional disc destruction mechanisms that are independent of cluster dynamics could mask the difference between the classical and disc-poor tidal arms. We discuss here shortly the potential contribution from two other important processes. For one, grain growth \\citep{2003A&A...400L..21V,2006ApJ...644L..71S,2008A&A...480..859B,2008ARA&A..46...21B,2010A&A...513A..79B} would hinder detection of protoplanetary discs in the infrared, the only wavelength range available to detect signatures of individual circumstellar discs in the distant Arches cluster. Whether grain growth in its disc population is expected to be as efficient as in the less massive star-forming regions in our solar neighborhood depends on the environmental conditions. Recently, \\citet{2011ApJ...731...95O,2011ApJ...731...96O} have shown that dust growth in discs can be suppressed by significant ionization, a probable scenario given the strong X-ray irradiation of the numerous massive stars in the Arches cluster \\citep{2000ApJ...536..896C,2005MNRAS.361..679O}. Second, if the radiation field of UV and X-ray photons emitted by massive stars is strong enough to heat the discs efficiently, photoevaporation will drive additional serious disc-mass loss \\citep{1998ApJ...499..758J,2004ApJ...611..360A,2007MNRAS.376.1350C,2009ApJ...705.1237G}. However, \\citet{2010ApJ...718..810S} have found a fraction of a few percent of A5V to B0V stars (corresponding to a mass range of $2-20\\,\\Msun$) in the Arches cluster that show indisputable features of protoplanetary discs. Regarding its current age of 2.5\\,Myr and enormous radiation field fed by more than 100~OB stars it seems that photoevaporation was not acting entirely destructive over this period. So we expect that characteristic features of encounter-induced disc-mass loss -- such as the disc-poor tidal arms -- could be still apparent at the present age of the Arches cluster. As the velocities of tidal tail stars do not differ by more than $2\\sigma^O_{3D} \\approx 20\\,\\kms$ from the large space motion of the cluster of $\\sim$200$\\,\\kms$ (see also Fig.~\\ref{fig:tidal_diagnostics}), kinematically they can be well separated from the field star population \\citep[e.g.][]{2008ApJ...675.1278S}. This makes the Arches cluster a highly interesting observational target to study the combined and distinct effect of photoevaporation and encounters on the evolution of protoplanetary discs in an extreme environment. Note that observations of the stellar population and its circumstellar discs in this dense and distant cluster require instruments with very high angular resolution $\\theta \\lesssim 0.1''$. Unlike successfully performed for the nearby, less dense young star clusters in the solar neighbourhood \\citep[e.g.][]{1999ARA&A..37..363F,2005ApJS..160..511K,2005ApJS..160..401P,2007A&A...464..211A,2009ApJ...696...47W}, members of the young Arches cluster cannot be separated from the older field stars using the Chandra X-ray Observatory because of its limited resolution $\\theta \\gtrsim 0.5''$. Thus at present the only viable strategy to discern cluster members from field stars is kinematic selection based on long-term proper motion studies. A natural extension of such a study would include the $\\sim$4\\,Myr old, more extended Quintuplet cluster, that has been considered to be the Arches cluster's ``older borther'' \\citep{1999ApJ...525..750F}. Indeed, core collapse of our cluster model at 2.5\\,Myr induces significant expansion (see Fig.~\\ref{fig:dynamics__isolated_vs_orbit}), naturally explaining the two starburst clusters as different evolutionary stages emerging from the same initial conditions. Hence the Quintuplet cluster represents the ideal observational target to verify the future evolution of the Arches cluster as predicted by our simulation. The characteristic concentration of disc-less stars at the edges of the extended tidal arms provides a critical test. We have thus initiated a multi-epoch $K$-band campaign and $L$-band photometry of the Arches and the Quintuplet cluster to obtain a kinematically selected membership sample required to observationally detect the characteristic disc-poor tidal features. However, we note that near-infrared surveys are not ideal to trace the predicted spatial gradients in disk fractions. This is because this wavelength regime is dominated by emission from the inner disc regions (at $\\lesssim$1\\,AU from the central star) while the encounter-induced disc-mass loss studied here mostly removes material from the outer disc parts (at $\\sim$100\\,AU). Hence, ideally, observations would be carried out at millimeter wavelengths. The instrument of choice for this purpose is ALMA as it provides the high spatial resolution ($\\theta \\approx 0.1''$) required to resolve individual sources in the starburst clusters at the Galactic Centre." }, "1207/1207.2583_arXiv.txt": { "abstract": "{Cool, evolved stars undergo copious mass loss but the detailed mechanisms and the form in which the matter is returned to the ISM are still under debate.}{We investigated the structure and evolution of the wind at 5 to 50 stellar radii from Asymptotic Giant Branch and Red Supergiant stars.} {22-GHz water masers around seven evolved stars were imaged using MERLIN, at sub-AU resolution. Each source was observed at between 2 and 7 epochs, covering several stellar periods. We compared our results with long-term single dish monitoring provided by the Pushchino radio telescope.} {The 22-GHz emission is located in approximately spherical, thick, unevenly filled shells. The outflow velocity increases twofold or more between the inner and outer shell limits. Water maser clumps could be matched at successive epochs separated by less than two years for AGB stars, or at least 5 years for RSG. This is much shorter than the decades taken for the wind to cross the maser shell, and comparison with spectral monitoring shows that some features fade and reappear. In five sources, most of the matched maser features brighten or dim in concert from one epoch to the next. A number of individual maser features show idiosyncratic behaviour, including one cloud in W Hya caught in the act of passing in front of a background cloud leading to 50-fold, transient amplification. The masing clouds are one or two orders of magnitude denser than the wind average and contain a substantial fraction of the mass loss in this region, with a filling factor $<1\\%$. The RSG clouds are about ten times bigger than those round the AGB stars. }{ Proper motions are dominated by expansion, with no systematic rotation. The maser clouds presumably survive for decades (the shell crossing time) but the masers are not always beamed in our direction. Only radiative effects can explain changes in flux density throughout the maser shells on short timescales. The size of the clouds is proportional to that of the parent star, being of a similar radius to the star once the clumps reach the 22-GHz maser shell. Stellar properties such as convection cells must determine the clumping scale.} ", "introduction": "Stars of less than eight Solar mass evolve onto the Asymptotic Giant Branch (AGB) when they have exhausted much of the available H. They become very large ($R_{\\star} > 1$ AU) and cool ($T_{\\star} \\approx 2500$ K) and lose mass at the rate of up to an Earth mass per year. Red Supergiants (RSG) had main sequence masses $>8$ M$_{\\odot}$ and swell to ten times the size of AGB stars, with ten times the mass loss rate. Although the internal stellar processes are mass-dependent, the winds are similar, being cool, dusty and molecular. O-rich stars produce SiO, H$_2$O and OH masers with spatially very compact, spectrally narrow lines which can be imaged at milli-arcsec resolution using radio interferometry. A combination of pulsations, radiation pressure and possibly other e.g. magnetic forces drives material away from the stellar surface. The exact mechanisms are still under debate (e.g. \\citealt{Woitke06}) and at a few stellar radii ($R_{\\star}$), SiO masers show both infall and outflow (e.g. \\citealt{Assaf11}). Dust nucleation is likely to start close to the star \\citep{Wittkowski07} and by 4--6 $R_{\\star}$, radiation pressure on dust is the main force driving the wind. The H$_2$O masers around these stars are found in approximately spherical shells extending from about 5-50 $R_{\\star}$. They provide a very high-resolution probe of the structure, kinematics and conditions in the circumstellar envelope (CSE) in the region where the outflow is most strongly accelerated and passes through the escape velocity, see e.g. \\citet{Bowers94}, \\citet{Yates94} and also \\citet{Habing96} for a general review. MERLIN\\footnote{the UK radio interferometer, operated by the University of Manchester on behalf of STFC} has been used to monitor H$_2$O and OH masers around a number of evolved stars with no known companions. We analyse 2--7 epochs of images of 22-GHz H$_2$O masers around seven objects, listed in Table~\\ref{tab:stars}. The observations of the RSG VX Sgr and S Per (first epoch) were reported in \\citet{Murakawa03} (M03) and \\citet{Richards99} (R99), respectively. The first epochs of the AGB stars IK Tau, U Her, U Ori and RT Vir were reported in \\citet{Bains03} (B03). W Hya MERLIN 22-GHz observations have not been published previously. These stars are close enough and unobscured enough to have AAVSO optical monitoring data, \\emph{Hipparcos} or other parallax measurements and in many cases optical or IR interferometric meaurements of the stellar diameter. They are among the objects which have been monitored by the Pushchino radio telescope at 22 GHz for many decades. Component fitting allows the size of AU-scale clouds to be measured from MERLIN data. The H$_2$O maser properties of the sources above declination 0\\degr\\/ were investigated in detail by \\citet{Richards11} (R11), who found that they are predominantly unsaturated. The maser appearance is mostly characteristic of spherical clouds but, in some objects, amplification appears to be occuring along the long axis of flattened clouds. \\begin{table*} \\begin{tabular}{lcrrrccccccc} \\hline Star &Position &Type&\\multicolumn{1}{c}{$V_{\\star}$}&\\multicolumn{1}{c}{Distance}&\\multicolumn{1}{c}{$D$} & $R_{\\star}$&$P$&$\\dot{M}$& Previous imaging \\\\ &(J2000)& &(km s$^{-1}$) & \\multicolumn{1}{c}{(pc)}& \\multicolumn{1}{c}{(pc)} & (AU)&(d) &(M$_{\\odot}$ yr$^{-1}$)& \\\\ \\hline VX Sgr &18 08 04.05 --22 13 26.6&RSG&--5.3&$1570\\pm^{270}_{270}$ &1700 &$7.4\\pm0.7$ &732& $7.2\\times10^{-5}$ &C86,B93,M98,M03,V05 \\\\ S Per &02 22 51.71 +58 35 11.4& RSG& --38.5 & $2312\\pm^{65}_{32}$ &2300&$8\\pm3$ &822& $3.8\\times10^{-5}$ &D87,Y94.R99,M98,V01,\\\\ &&&&&&&&&A10 \\\\ U Ori &05 55 49.17 +20 10 30.7&Mira& --39.5 & $260\\pm^{50}_{50}$ &266& $1.5\\pm0.1$ &368& $2.3\\times10^{-7}$ &B88,Y94,B94,B03,V05 \\\\ U Her &16 25 47.47 +18 53 32.9&Mira&--14.5 & $266\\pm^{32}_{28}$ &266& $1.3\\pm0.1$&406& $3.4\\times10^{-7}$&Y94,B94,M98,C00,B03,\\\\ &&&&&&&&&V02,V05\\\\ IK Tau &03 53 28.89 +11 24 21.9&Mira& +34.0 & $250\\pm^{20}_{20}$ &266& $2.8\\pm0.3$ &470& $2.6\\times10^{-6}$&L87,B93,Y94,M98,B03 \\\\ RT Vir &13 02 37.98 +05 11 08.4&SRb& +18.2 & $135\\pm^{15}_{15}$ &133& $0.8\\pm0.04$ &158& $1.3\\times10^{-7}$ &Y94,B93,B94,B03,I03\\\\ W Hya &13 49 02.00 --28 22 03.4&SRb&+40.6 & $98\\pm^{30}_{18}$&98&$2.1\\pm0.2$ &375& $2.3\\times10^{-7}$&R90,B93 \\\\ \\multicolumn{6}{l}{References}\\\\ VX Sgr &vL07&&C86 &\\,\\,C07 &&\\,\\,M04 &S11&D10 \\\\ S Per &vL07& &D87 & \\,\\,M08 & & \\,\\,H94, L05 &S11&vL05\\\\ U Ori &vL07& &C91 & \\,\\,C91 & &\\,\\,R06 &S11&K98\\\\ U Her &vL07& &C94 & \\,\\,V07 &&\\,\\,R06 &S11&Y95\\\\ IK Tau &C06& &K87 & \\,\\,098 &&M04, R06 &S11&O98, B00\\\\ RT Vir &vL07& &N86 & \\,\\,vL07 & &\\,\\,M04 &S11&K98, K99\\\\ W Hya &vL07& &N96 & \\,\\,V03 & & \\,\\,Z11 &S11&K98\\\\ \\hline \\end{tabular} \\caption{Properties of the sample stars. The stellar velocity $V_{\\star}$ is given in the Local Standard of Rest (LSR) convention. The positions given are for epoch 2000, see Section~\\ref{sec:align} for more on astrometry. The most recent distances are given along with those used in our calculations, $D$, in order to remain consistent with M03, R99 and B03. The values of $R_{\\star}$ for the AGB stars were measured using IR interferometry at H and K bands. The radius of S Per is the average of values deduced from spectral fits. The stellar period $P$ is complex, changeable and/or uncertain for some objects, see Figs.~\\ref{VXSgrSPer_AAVSO.png} to~\\ref{IKTauRTVirWHya_AAVSO.png}. The mass loss rates $\\dot{M}$ are adjusted to our adopted distances. The final column gives some of the previous interferometric 22-GHz H$_2$O maser images which have been published. \\newline References in the table are: A10~\\citet{Asaki10}; B88~\\citet{Bowers88}; B93~\\citet{Bowers93}; B94~\\citet{Bowers94}; B00~\\citet{Bieging00};B03~\\citet{Bains03}; C86~\\citet{Chapman86}; C91~\\citet{Chapman91}; C94~\\citet{Chapman94}; C00~\\citet{Colomer00}; C06~\\citet{Carlsberg06}; C07~\\citet{Chen07}; D87~\\citet{Diamond87}; I03~\\citet{Imai03}; K87~\\citet{Kirrane87}; K98~\\citet{Knapp98}; K99~\\citet{Kerschbaum99}; L87~\\citet{Lane87}; L05~\\citet{Levesque05}; M98~\\citet{Marvel98}; M03~\\citet{Murakawa03} M04~\\citet{Monnier04}; M08~\\citet{Mayne08}; N96~\\citet{Neufeld96}; N86~\\citet{Nyman86}; O98~\\citet{Olofsson98}; R90~\\citet{Reid90}; R99~\\citet{Richards99}; R06~\\citet{Ragland06}; S11~\\citet{Samus11} (GCVS); V01~\\citet{Vlemmings01}; V02~\\citet{Vlemmings02}; V03~\\citet{Vlemmings03}; V05~\\citet{Vlemmings05}; V07~\\citet{Vlemmings07}; vL05~\\citet{vanLoon05}; vL07~\\citet{vanLeeuwen07}; Y94~\\citet{Yates94}; Z11~\\citet{Zhao-Geisler11}} \\label{tab:stars} \\end{table*} The H$_2$O masers have also been imaged using VLBI and the VLA, see Table~\\ref{tab:stars}. However, a large fraction of the flux (half is not unusual) is resolved-out by VLBI, whilst the VLA cannot resolve individual maser clumps. Allowing for this, our observational results are consistent with other images. Comparisons with previous interpretations are included in the relevant sections of this paper. For completeness, we note that we observed several other objects with MERLIN at 22 GHz. The first epochs of NML Cyg and VY CMa were published by \\citet{Richards96} and \\citet{Richards98v}; further monitoring of these two exceptional RSG will be presented by Yates et al., as will obervations of the heavily-obscured SRb, R Crt. We also observed R Cas (20000404, 20010511, 20020405) and R Leo (20000429) but these were non-detections at a conservative upper limit of 200 mJy. This is consistent with non-detections for several years around this period (upper limit 10 Jy) using Pushchino (\\citet{Pashchenko04} R Cas; \\citet{Esipov99} R Leo). This paper presents a large-scale analysis of multiple observations (between two and seven epochs per star, spanning up to 7 years). Data acquisition and analysis is summarised in Section~\\ref{sec:data}, including estimating the centre of expansion, aligning epochs and assessing the astrometric accuracy. We examine proper motions and cloud survival in Section~\\ref{sec:pm} and compare images with Pushchino single dish monitoring in Section~\\ref{sec:Pushchino}, to investigate the survival of masers with respect to their parent clouds. We analyse the kinematics and estimate the mass concentration in maser clouds in Section~\\ref{sec:mass}. Local and global maser variability is discussed in Section~\\ref{sec:variability}. The size of the maser clumps and their relationships with stellar properties is discussed in Section~\\ref{sec:cloudsize} and Section~\\ref{sec:conclusions} presents the summary and future work. ", "conclusions": "\\label{sec:conclusions} We have resolved the detailed structure of the approximately spherical, thick H$_{2}$O maser shells around two RSG (VX Sgr, S Per), three Miras (U Ori, U Her, IK Tau) and two SRb stars (RT Vir, W Hya), at multiple epochs. MERLIN detects all the 22-GHz maser emission (or possibly almost all, in the case of the closest object), at high enough resolution to measure the sizes and proper motions of individual clouds. We compare these images with single dish monitoring by the Pushchino radio telescope. Our results provide new insights into the development of the stellar wind in the region where the escape velocity is exceeded. \\paragraph{\\bf Survival} The H$_{2}$O maser shells are 10--250 AU thick with crossing times from a few tens of years for AGB stars to nearly a century for RSG, but most individual maser features survive less than 1--2 yr and 1--2 decades, in the AGB stars and RSG, respectively. The disappearance and reappearance of individual masers identified by comparing imaging with single dish monitoring shows that this contradiction is resolved if the masers emanate from long-lived clouds. These undergo subsonic or mildly supersonic turbulence of a few km s$^{-1}$, which causes detectable masers to be alternately boosted and supressed without destroying the clouds themselves. \\paragraph{\\bf Expansion proper motions} The proper motions of VX Sgr, S Per, U Her, IK Tau, RT Vir and most epochs of W Hya are unambiguously dominated by spheroidal expansion. There is a hint of rotation in the IK Tau maser proper motions. The small tangential component could be caused by random motions, which is also likely to be a significant factor affecting U Ori. The radial proper motion velocities (with respect to the assumed stellar position) are generally consistent with $V_{\\mathrm{LSR}}$ measurements, but the former have larger measurement errors and scatter, especially for the low declination source W Hya. Discrepancies between velocities derived from proper motion and $V_{\\mathrm{LSR}}$ measurements could also be due to asymmetry, as seen for VX Sgr (M03). \\paragraph{\\bf Acceleration} The 22-GHz masers lie outside the dust formation zone, where the wind is driven by radiation pressure on grains. The outflow velocity increases twofold or more across the maser shells of all sources. Moreover, comparison of different sources shows that the larger the inner and outer shell radii, the greater the expansion velocities at the outer rims, although the acceleration is more gradual. This is likely to be due to changes in the dust optical properties so a higher proportion of radiation is absorbed, and/or increasing optical depth so that the stellar photons are re-emitted making radiation pressure more efficient. This is consistent with \\emph{Herschel} and other results which show that terminal velocity is not reached until hundreds of $R_{\\star}$. \\paragraph{\\bf Dense clumps} The wind at $\\sim5-50 R_{\\star}$ is concentrated in clumps which are 40--110 times denser than the wind average. They comprise between a fifth and almost all of the mass in the stellar wind in the maser shell, although accurate comparisons are difficult due to uncertainties in measuring the total mass loss rate. The volume filling factor is less than 1\\%. At least two to six clouds are formed per stellar period. \\paragraph{\\bf Flare due to cloud overlap} Single-dish monitoring of H$_2$O masers reveals various forms of variability; in some instances a single spectral feature brightens by more than one order of magnitude for between a few weeks and a year or two. One such flare, in W Hya, was imaged at 3 epochs from 2000--2002, bracketing the 2500-Jy maximum captured by Pushchino. The images show that the variable cloud appears to pass in front of another. Two sets of components are spatially differentiated, at slightly different $V_{\\mathrm{LSR}}$, both with distinct Gaussian spectral profiles. The flaring feature is initially to the S of the other cloud, both $<10$ Jy. Two years later it has exceeded 400 Jy and moved to the NW of the fainter feature (in a direction consistent with outflow). This is a dramatic confirmation of the predictions of \\citet{Kartje99} for maser amplification by overlapping clouds. \\paragraph{\\bf Variability throughout the maser shell} The majority of matched features in each of VX Sgr, U Her and IK Tau decrease in flux density between the two imaging epochs regardless of their position in the maser shell; the majority of U Ori matched features increase. The intervals between epochs are of order 1--2 stellar periods, but there is no obvious connection between optical phase and brightening or dimming. Features throughout the RT Vir maser shell, imaged six times during the decline of optical brightness, rise and then slightly decline again in 22-GHz flux density. S Per and W Hya matched features show an apparently random mixture of brightening and dimming. The coordinated behaviour of features in 5 CSEs suggests that radiative effects play a major role in determining maser brightness, probably linked to the IR stellar phase (which lags the optical). The 22-GHz maser shells are 10--200 AU thick, so the light travel time is less than a few days but a shock at 5--10 km S$^{-1}$ would take many years to cross even the thinnest shell. Nonetheless, the H$_2$O masers may be affected by shocks close to the inner rim of the shell. \\paragraph{\\bf Cloud size depends on star size} The average radius of water maser clouds is 0.5--2 AU for the AGB stars and 6--9 AU for the RSG, showing a close dependence on $R_{\\star}$. Assuming the clouds expand steadily in the outflow, this corresponds to birth radii $\\sim5-10\\% R_{\\star}$. Previous CO and dust imaging have hinted at similar clumping scales and overdensity, but only MERLIN H$_{2}$O maser observations have well-resolved all the emission from such clumps, close to the star. The dependence on stellar size suggests that stellar surface phenomena such as convection cells determine the scale of clumps. \\paragraph{\\bf Future work} We are in the process of analysing multi-epoch MERLIN and VLBI observations of OH masers which probe different conditions in these objects. A future paper will compare the H$_2$O morphology with the distribution of OH masers to investigate asymmetry and inhomogeneity. There are very few estimates published for the sizes of SiO or OH maser clouds in our sources, but we will seek available data where the clouds are likely to be resolved but not resolved-out, to test whether the sizes do increase with distance from the star. The advent of ALMA, \\emph{e}-MERLIN, upgraded VLBI and other interferometers will allow the unanswered questions to be tackled. The current baselines of ALMA, up to 1 km, will locate sub-mm H$_2$O masers with respect to the 22 GHz transition and the dust formation zone, mapping changes in excitation conditions. Eventual longer baselines and full sensitivity will resolve both dust and molecular species in clumps, showing whether the 22-GHz clouds are indeed dustier and denser. ALMA will also resolve the stars, as will \\emph{e}-MERLIN and the EVLA at its highest frequencies. This will reveal any convective disurbances and timely VLBA SiO maser monitoring will show whether these lead directly to clumpy mass loss, and if the SiO clumps in turn become dusty water maser clouds, resolved by ALMA and \\emph{e}-MERLIN in multiple transitions. Simultaneous imaging of the star and masers at 22 GHz (only previously achieved for a very few objects e.g. \\citealt{Reid90}; \\citealt{Reid97}) will solve the current astrometric ambiguities." }, "1207/1207.2898_arXiv.txt": { "abstract": "\\vspace{1cm} \\centerline{\\bf ABSTRACT}\\vspace{2mm} The fate of our universe is an unceasing topic of cosmology and the human being. The discovery of the current accelerated expansion of the universe significantly changed our view of the fate of the universe. Recently, some interesting scenarios concerning the fate of the universe attracted much attention in the community, namely the so-called ``Little Rip'' and ``Pseudo-Rip''. It is worth noting that all the Big Rip, Little Rip and Pseudo-Rip arise from the assumption that the dark energy density $\\rho(a)$ is monotonically increasing. In the present work, we are interested to investigate what will happen if this assumption is broken, and then propose a so-called ``Quasi-Rip'' scenario, which is driven by a type of quintom dark energy. In this work, we consider an explicit model of Quasi-Rip in detail. We show that Quasi-Rip has an unique feature different from Big Rip, Little Rip and Pseudo-Rip. Our universe has a chance to be rebuilt from the ashes after the terrible rip. This might be the last hope in the ``hopeless'' rip. ", "introduction": "\\label{sec1} Since its discovery in 1998, the current accelerated expansion of our universe~\\cite{r1} has been one of the most active fields in modern cosmology. As is well known, it could be due to an unknown energy component (dark energy) or a modification to general relativity (modified gravity)~\\cite{r1,r2}. Before 1998, it was commonly believed that in the future our universe will either expand forever or contract again into a final Big Crunch. However, the discovery of the current accelerated expansion of the universe significantly changed our view of the fate of the universe. In fact, many novel possibilities are under the active consideration in the community nowadays. Today, there are many dark energy candidates in the market. Among them, the famous phantom dark energy~\\cite{r3,r4} is very interesting. Its equation-of-state parameter (EoS) is smaller than $-1$. Although phantom dark energy is consistent with the current observational data~\\cite{r1,r3}, it violates all the energy conditions. One of the consequences is that our universe will encounter a singularity at a finite time, namely the so-called Big Rip~\\cite{r4}. At this singularity, the scale factor $a$, energy density and pressure of our universe are all divergent. In fact, besides the traditional Big Bang, Big Crunch, and the Big Rip, many novel singularities have been considered in the literature, such as Sudden singularities, Generalized sudden singularities, Quiescent singularities, Big Boost, Big Brake, Big Freeze, $w$ singularities, Inaccessible singularities, Directional singularities (see e.g.~\\cite{r5,r6} and references therein for some brief reviews). These singularities arise at the price of violation of one or several energy conditions. As is well known, in~\\cite{r6} (see also e.g.~\\cite{r5,r7,r8}) the future singularities have been classified into four types, namely \\begin{itemize} \\item Type I (Big Rip): $a\\to\\infty$, $\\rho\\to\\infty$, $H\\to\\infty$, $|p|\\to\\infty$, when $t\\to t_s<\\infty\\,$; \\item Type II (Sudden singularity): $a\\to a_s$, $\\rho\\to\\rho_s$, $H\\to H_s$, $|p|\\to\\infty$, $\\dot{H}\\to\\infty$, when $t\\to t_s<\\infty\\,$; \\item Type III (Big Freeze): $a\\to a_s$, $\\rho\\to\\infty$, $H\\to\\infty$, $|p|\\to\\infty$, when $t\\to t_s<\\infty\\,$; \\item Type IV (Generalized sudden singularity): $a\\to a_s$, $\\rho\\to\\rho_s$, $H\\to H_s$, $|p|\\to p_s$, $\\dot{H}\\to\\dot{H}_s$, and higher derivatives of $H$ diverge, when $t\\to t_s<\\infty\\,$, \\end{itemize} where $t_s$, $a_s$, $\\rho_s$, $H_s$, $p_s$, $\\dot{H}_s$ are all finite constants ($a_s\\not=0$); $\\rho$ and $p$ are energy density and pressure respectively; $H\\equiv\\dot{a}/a$ is the Hubble parameter, and a dot denotes a derivative with respect to the cosmic time $t$. In fact, the above four types can include almost all known future singularities. Of course, singularities usually are not desirable in physics. Therefore, other possible fates of our universe are also considered in the literature, such as the cyclic/oscillatory cosmology. Recently, some interesting scenarios concerning the fate of the universe attracted much attention in the community, namely the so-called ``Little Rip''~\\cite{r9} and ``Pseudo-Rip''~\\cite{r10}. In~\\cite{r9,r10,r11}, their authors showed that if the cosmic energy density will remain constant or monotonically increase in the future, then all the possible fates of our universe can be divided into four categories based on the time asymptotics of the Hubble parameter $H(t)$~\\cite{r10}, namely \\begin{itemize} \\item Big Rip: $H(t)\\to\\infty$, when $t\\to t_{rip}<\\infty\\,$; \\item Little Rip: $H(t)\\to\\infty$, when $t\\to\\infty\\,$; \\item Cosmological Constant: $H(t)=const.\\,$; \\item Pseudo-Rip: $H(t)\\to H_\\infty <\\infty$, when $t\\to\\infty\\,$, \\end{itemize} where $H_\\infty$ is a constant. Obviously, the Big Rip singularity is not the only fate of our universe with the phantom-like dark energy. Both Little Rip and Pseudo-Rip are non-singular, and hence fall outside of the four categories in~\\cite{r6}. Similar to the Big Rip, the Little Rip dissociates {\\em all} bound structures, but the strength of dark energy is not enough to rip apart spacetime (unlike the Big Rip)~\\cite{r9,r10}. On the other hand, the Pseudo-Rip dissociates the bound structures which are held together by a binding force at or below a particular threshold, and hence it is possible that only {\\em some} bound structures are dissociated while the others are {\\em not} dissociated (depending on the model parameters)~\\cite{r10}. In fact, the Little Rip is an intermediate case between the cosmological constant and the Big Rip~\\cite{r9}, while the Pseudo-Rip is an intermediate case between the cosmological constant and the Little Rip~\\cite{r10}. It is worth noting that all the Big Rip, Little Rip and Pseudo-Rip arise from the assumption that the dark energy density $\\rho(a)$ is monotonically increasing~\\cite{r9,r10,r11}, i.e., the dark energy is phantom-like (its EoS $w<-1$). In the present work, we are interested to investigate what will happen if this assumption is broken. Obviously, in the case of the dark energy density $\\rho(a)$ is monotonically decreasing (i.e., the dark energy is quintessence-like with an EoS $w>-1$), no rip will happen and no bound structures will be dissociated. On the other hand, in the case of the dark energy density $\\rho(a)$ monotonically decreases (namely $w>-1$) in the first stage and then monotonically increases (namely $w<-1$) in the second stage (this is the case of the so-called ``quintom A'' dark energy in terminology of e.g.~\\cite{r12} and references therein), the fate of our universe is the Big Rip, which is trivial in some sense. The third case is that the dark energy density $\\rho(a)$ monotonically increases (namely $w<-1$) in the first stage and then monotonically decreases (namely $w>-1$) in the second stage (this is the case of the so-called ``quintom~B'' dark energy in terminology of e.g.~\\cite{r12} and references therein). It can be expected that in the first stage some or all bound structures will be dissociated (similar to the case of Pseudo-Rip), but then the disintegration process will stop, and the already disintegrated structures have the possibility to be recombined in the second stage. We dub it ``Quasi-Rip'', which is the subject of the present work (here we temporarily do not consider the case of oscillatory quintom dark energy). Since the quintom-like dark energy~\\cite{r13} (whose EoS can cross the so-called phantom divide $w=-1$) is slightly favored by the observational data~\\cite{r1} (see e.g.~\\cite{r12} for a comprehensive review), we note that the Quasi-Rip is well-motivated in fact. This paper is organized as follow. In Sec.~\\ref{sec2}, we discuss the disintegration of bound structures. In Sec.~\\ref{sec3}, we present an explicit model of Quasi-Rip. We constrain this model with the observational data, and then clearly show the Quasi-Rip in this model. In Sec.~\\ref{sec4}, some concluding remarks are given. ", "conclusions": "\\label{sec4} The fate of our universe is an unceasing topic of cosmology and the human being. The discovery of the current accelerated expansion of the universe significantly changed our view of the fate of the universe. Recently, some interesting scenarios concerning the fate of the universe attracted much attention in the community, namely the so-called ``Little Rip'' and ``Pseudo-Rip''. It is worth noting that all the Big Rip, Little Rip and Pseudo-Rip arise from the assumption that the dark energy density $\\rho(a)$ is monotonically increasing. In the present work, we are interested to investigate what will happen if this assumption is broken, and then propose a so-called ``Quasi-Rip'' scenario, which is driven by a type of quintom dark energy. In this work, we consider an explicit model of Quasi-Rip in detail. We show that Quasi-Rip has an unique feature different from Big Rip, Little Rip and Pseudo-Rip. Our universe has a chance to be rebuilt from the ashes after the terrible rip. This might be the last hope in the ``hopeless'' rip. Some remarks are in order. Firstly, as is shown in Sec.~\\ref{sec3}, our Quasi-Rip model is well consistent with the current observational data. However, even in the $1\\sigma$ C.L. region of $\\alpha-\\beta$ parameter space, the future behavior of our universe can be different enough (depending on the particular model parameters $\\alpha$ and~$\\beta$). In fact, as is well known, the current observational data can be consistent with all the phantom-like, quintessence-like and quintom-like dark energy models. Therefore, the current observational data cannot tightly tell what is the true fate of our universe. Most of the possibilities (including Big Rip, Little Rip, Pseudo-Rip, Quasi-Rip, de~Sitter expansion, other future singularities and so on) are still living. Secondly, the explicit model of Quasi-Rip considered in the present work is the simplest case. One can construct other more complicated $\\rho(a)$ to implement the Quasi-Rip. For example, one might construct an EoS $w(a)$ as a function of scale factor $a$, which is smaller than $-1$ when $aa_t$. Then the corresponding $\\rho(a)$ can be found from the energy conservation equation $\\dot{\\rho}+3H\\rho(1+w(a))=0$. Of course, other smart methods to construct the desirable $\\rho(a)$ are awaiting us. Thirdly, as mentioned in~\\cite{r10}, in its second Pseudo-Rip model, the reduced inertial force $f(a)$ can also have a peak, similar to our Quasi-Rip model. However, we note that in the Pseudo-Rip model, after the peak, the reduced inertial force $f(a)\\to const.$ which is still higher than the corresponding threshold to disintegrate the bound structure. Therefore, the already disintegrated structures have {\\em no} possibility to be recombined in the Pseudo-Rip models. On the contrary, in the Quasi-Rip models, the reduced inertial force $f(a)$ monotonically decreases in the second stage. Eventually, it will become lower than all the thresholds to disintegrate the bound structure. Therefore, the already disintegrated structures have the possibility to be recombined in the second stage. Fourthly, it is well known that phantom is unstable at quantum level and hence the perturbations grow large. Noting that in the present work our discussions are at classical level instead, this problem could be set aside. In fact, the quantum stability of a phantom phase has been considered in~\\cite{r24}. The authors of~\\cite{r24} studied the perturbations in the quantum-corrected effective field equation at one- and two-loop order, and they found that the system is stable. On the other hand, it is claimed in~\\cite{r25} that scalar perturbations can grow during a phantom phase if EoS $w<-5/3$. However, from Eq.~(\\ref{eq15}) and Fig.~\\ref{fig1}, it is easy to see that the corresponding $w$ is only slightly smaller than $-1$ in the phantom phase for our particular $\\alpha$ and $\\beta$ (n.b. Fig.~\\ref{fig1}). So, $w>-5/3$ instead and hence our Quasi-Rip model can avoid the problem raised in~\\cite{r25}. Further, in fact the dark energy considered in this work is not necessarily a scalar field. It could even be an effective dark energy from modified gravity, namely the so-called ``geometric dark energy''. So, it might avoid the corresponding problems in the phantom phase. Finally, in the present work, we consider only the quintom dark energy which crosses the phantom divide $w=-1$ once. In fact, we can also consider the quintom dark energy which can cross the phantom divide for many times. The most attractive Quasi-Rip model might be the one driven by the oscillatory quintom dark energy. In this oscillatory Quasi-Rip model, our universe will be destroyed and then be rebuilt again and again." }, "1207/1207.1063_arXiv.txt": { "abstract": "{ We consider the possibility that the primordial curvature perturbation was generated through the curvaton mechanism from a scalar field with an electric charge, or precisely the Standard Model U(1) weak hypercharge. This links the dynamics of the very early universe concretely to the Standard Model of particle physics, and because the coupling strength is known, it reduces the number of free parameters in the curvaton model. The gauge coupling also introduces several new physical effects. Charge fluctuations are generated during inflation, but they are screened by electron-positron pairs therefore do not violate observational constraints. After inflation, the curvaton interacts with thermal radiation which destroys the curvaton condensate and prevents the generation of curvature perturbations, unless the inflaton dynamics satisfy strong constraints. The curvaton also experiences a period of parametric resonance with the U(1) gauge field. Using the standard perturbative approach, we find that the model can generate the observed density perturbation for Hubble rate $H_* \\gtrsim 10^8 \\GeV$ and curvaton mass $m \\gtrsim 10^{-2}H_*$, but with a level of non-Gaussianity ($f_{NL} \\gtrsim 130$) that violates observational constraints. However, previous studies have shown that the parametric resonance changes the predicted perturbations significantly, and therefore fully non-linear numerical field theory simulations are required. } ", "introduction": "The inflationary paradigm is commonly accepted as the origin of the primordial fluctuations that seeded structure in the Universe. Models of inflation are usually based on quantum field theory, with fields unrelated to the Standard Model of particle physics. However, some interaction between the inflaton and the Standard Model fields is required in order to reheat the Universe, which means that a single theoretical framework for both inflation and particle physics is ultimately required. If the inflaton field is responsible for generating the primordial curvature perturbations, its coupling to other fields has to be extremely weak in order to avoid radiative corrections that would spoil the flatness of the potential. This makes it difficult to couple the inflaton to Standard Model gauge fields, which all interact relatively strongly. A less constrained alternative is the curvaton model \\cite{Linde:1996gt,Enqvist:2001zp,Lyth:2001nq,Moroi:2001ct}, in which the perturbations are generated by a separate scalar field known as the curvaton. The curvaton is light and subdominant during inflation and gains an isocurvature perturbation. After inflation has ended, the energy density of the curvaton grows relative to the background radiation. When the curvaton finally decays, its isocurvature perturbation is converted to an adiabatic perturbation that can seed the structure in the Universe. The purpose of this paper is to investigate whether the curvaton field could be charged under a Standard Model gauge group. This would have the attractive feature that the properties of these interactions are known, in contrast with typical curvaton models, which have so much freedom that it is hard to make definite predictions. Of the three Standard Model gauge groups, the most promising is the U(1) weak hypercharge, because the SU(2) and SU(3) groups are confining and would give more complicated physics. In contrast, a U(1) charge would essentially mean that the curvaton has an electric charge, and it is relatively easy to investigate the various constraints that arise from this. To be specific, we consider the Lagrangian \\be{lag} {\\cal L} = - m^2 \\sigma^\\dagger \\sigma - \\lambda (\\sigma^\\dagger \\sigma)^2- \\frac{1}{4}F_{\\mu\\nu}F^{\\mu\\nu} + \\left| (i\\partial_\\mu - g' A_\\mu)\\sigma \\right|^2, \\ee where $\\sigma$ is the curvaton field, $g'$ is the Standard Model U(1) gauge coupling, $A_\\mu$ is the Standard Model U(1) gauge field and $F_{\\mu\\nu}$ is the corresponding field strength tensor. The coupling strength is known~\\cite{pdg}, $g' \\approx 0.36$, and thus the number of free parameters is reduced.\\footnote{The coupling will run to larger values --- we quote the weak scale value.} For concreteness, we assume that the curvaton carries one unit of hypercharge, $Y=1$. The hypercharge has to be an integer to allow the curvaton to decay into Standard Model particles, and a higher value would not change our conclusions significantly. For $Y=2$, the curvaton could have a Yukawa coupling to right-handed electrons, but for other values the only other renormalisable term is a bilinear coupling $\\sigma^\\dagger\\sigma \\Phi^\\dagger\\Phi$ to the Higgs field. For simplicity, we assume that this coupling is negligible. The opposite case would essentially correspond to the model discussed in Refs.~\\cite{Enqvist:2008be,Chambers:2009ki}. To determine whether such a model is viable, we apply both theoretical and observational constraints, for example requiring a stable vacuum and the correct amplitude of the curvature perturbation $\\zeta$. In this model, three mechanisms affect the curvaton's dynamics after inflation: non-perturbative parametric resonance into photons, interactions with the thermal bath, and perturbative decay. We find that a perturbative estimate of the non-Gaussianity ($f_{\\rm NL}$) appears to rule out the model at 95\\% confidence level, but this ignores the parametric resonance which is known to have a significant effect on perturbations~\\cite{Chambers:2009ki} and therefore it cannot be relied on. Instead, non-linear field theory simulations are ultimately needed to determine the perturbations. In \\sect{over}, we give a brief overview of dynamics of the model. In \\sect{charge}, we discuss the charge fluctuations produced during inflation and show that, because of electric screening, they are compatible with observations. Then in \\sect{curvaton}, we give the conditions that the model must satisfy to have the qualitative behaviour of a curvaton model. These are conditions such as requiring a light curvaton during inflation and no false vacuum in the curvaton potential. In \\sect{decay}, we then expand the discussion of the dynamics after inflation, analytically calculating timescales for each process. In \\sect{results} we discuss the generation of curvature perturbations and production of transient cosmic string-like structures. We conclude in \\sect{conclusions}. ", "conclusions": "\\label{conclusions} We have demonstrated that it is possible to have a consistent model of the early Universe in which a scalar field charged under the Standard Model $U(1)$ weak hypercharge plays the role of the curvaton, generating a nearly scale-invariant spectrum of curvature perturbations with the observed amplitude. Besides curvature perturbations, the curvaton charge gives several other potentially observable effects. In principle, field fluctuations during inflation generate significant electric charge fluctuations on superhorizon scales, but we found that these are suppressed by the screening provided by Schwinger pair creation of other charged particles such as electron-positron pairs during inflation. The standard calculation for this model predicts relatively a high Hubble rate during inflation, $H_* \\gtrsim 10^8\\GeV$, and significant non-Gaussianity ($f_{\\rm NL} \\gtrsim 130$), which would rule it out at 95\\% confidence level by WMAP. However, this ignores a period of parametric resonance between the curvaton and photon fields, whose effect on the curvature perturbation can only be calculated with numerical lattice field simulations. There is also freedom to relax some of the constraints in \\sect{curvaton}, such as allowing a meta-stable vacuum or permitting a quartic term to initially dominate the curvaton potential. Inflaton dynamics after inflation are also strongly constrained to avoid thermal photons destroying the curvaton condensate. In practice, this means that inflaton reheating should be substantially delayed, and that the inflaton should oscillate in a quartic potential after inflation (the inflaton should behave like radiation so that the relative magnitude of the curvaton can grow). The existence of an extra U(1) charged field, with a non-zero value, could also play a significant role in the electroweak phase transition. However, investigating these possibilities is beyond the scope of this paper, and we simply raise them here as potential directions for future research. In summary, this paper presents an interesting model that concretely links the dynamics of the Standard Model of particle physics with the very early Universe. In the era of new data from both LHC and Planck, the mechanism presented in this paper deserves further investigation." }, "1207/1207.1639_arXiv.txt": { "abstract": "In this talk, we report results of our recent studies to delineate effects of the tensor force on the density dependence of nuclear symmetry energy within phenomenological models. The tensor force active in the isosinglet neutron-proton interaction channel leads to appreciable depletion/population of nucleons below/above the Fermi surface in the single-nucleon momentum distribution in cold symmetric nuclear matter (SNM). We found that as a consequence of the high momentum tail in SNM the kinetic part of the symmetry energy $E^{kin}_{sym}(\\rho)$ is significantly below the well-known Fermi gas model prediction of approximately $12.5 (\\rho/\\rho_0)^{2/3}$. With about $15\\%$ nucleons in the high momentum tail as indicated by the recent experiments at J-Lab by the CLAS Collaboration, the $E^{kin}_{sym}(\\rho)$ is negligibly small. It even becomes negative when more nucleons are in the high momentum tail in SNM. These features have recently been confirmed by three independent studies based on the state-of-the-art microscopic nuclear many-body theories. In addition, we also estimate the second-order tensor force contribution to the potential part of the symmetry energy. Implications of these findings in extracting information about nuclear symmetry energy from nuclear reactions are discussed briefly. ", "introduction": "Nuclear symmetry energy $E_{sym}(\\rho)$, which encodes the energy related to neutron-proton asymmetry in the nuclear matter Equation of State (EOS), is a vital ingredient in the theoretical description of neutron stars and of the structure of neutron-rich nuclei and reactions involving them. Since the density-dependence of $E_{sym}(\\rho)$ is still the most uncertain part of the EOS of neutron-rich nucleonic matter especially at supra-saturation densities, to better determine the $E_{sym}(\\rho)$ has become a major goal of both nuclear physics and astrophysics \\cite{ireview98,ibook01,dan,bar,li1,Sum94,Lat04,Ste05a,Xuli10a,Xuli10b,LWC05,tsa,Cen09,xia}. While significant progress has been made recently in narrowing down the symmetry energy near normal nuclear matter density $\\rho_0$, see, e.g., \\cite{Nat10,Tsang11,Newton11,War11,Dut11,BALI11,Dong12,ts12,St12,Lat12}, much more efforts are needed to pin down the $E_{sym}(\\rho)$ at both sub- and supra-saturation densities. Moreover, it is now broadly recognized that essentially all of the constraints extracted from experimental data are model dependent. Thus, to make further progress in the field, it is imperative to identify clearly the key physics ingredients determining the density dependence of nuclear symmetry energy in each model \\cite{Xuli10a,Far12}. Besides different techniques used in various nuclear many-body theories, several ingredients, such as, the spin-isospin dependence of the three-body force, tensor force induced high momentum tail in the single-nucleon momentum distribution of symmetric nuclear matter (SNM) and the associated isospin-dependence of short-range two-nucleon correlations, the isospin-dependence of nucleon pairing and clustering at low densities, are particularly known to affect significantly the $E_{sym}(\\rho)$. Of course, these ingredients may be approximately equally important and interfere strongly at some densities but individually dominate at other densities in models where they are all considered. In reality, however, they are rarely all taken into account simultaneously in a given model. Also, among these ingredients effects of the tensor force are least known so far. For instance, in most of the Relativistic Mean-Field (RMF) models, the $E_{sym}(\\rho)$ are determined by the coupling schemes and properties of the $\\rho$ and $\\delta$ mesons. Generally, no tensor coupling is considered. In phenomenological models, such as the Skyrme and/or Gogny Hartree-Fock approaches, the spin-isospin dependence of the three-body force is the most uncertain term determining the density dependence of the $E_{sym}(\\rho)$ while effects of the tensor forces are normally not considered either. On the other hand, most of the more microscopic many-body theories using modern nucleon-nucleon interactions have incorporated all major ingredients affecting the $E_{sym}(\\rho)$ albeit at different levels. Because of the different many-body approaches and interactions used, although all these models are well established and transparent, it has been hard to identify the main causes for their different predictions for the $E_{sym}(\\rho)$. To our best knowledge, currently there is no community consensus regarding the underlying physics responsible for the uncertain density dependence of nuclear symmetry energy especially at supra-saturation densities. Why is the density dependence of nuclear symmetry energy so uncertain especially at supra-saturation densities? What are the effects of the tensor force? To help answer these questions, using simple and phenomenological approaches \\cite{Xu11,Li11}, we have recently investigated effects of the tensor force on the kinetic ($E_{sym}^{kin}(\\rho)$) and potential ($E_{sym}^{pot}(\\rho)$) parts of the symmetry energy, separately. The most striking finding is that, unlike the free Fermi gas model prediction $E_{sym}^{kin}(\\textrm{FG})(\\rho)\\equiv(2^{\\frac{2}{3}}-1)(\\frac{3}{5} \\frac{\\hbar^2 k_F^2}{2m})\\approx 12.5\\rho^{2/3}$ that has been widely used in both nuclear physics and astrophysics, the tensor force induced high momentum tail in the single-nucleon momentum distribution in SNM reduces significantly the $E_{sym}^{kin}(\\rho)$ to values much small than the $E_{sym}^{kin}(\\textrm{FG})(\\rho)$. In fact, the $E_{sym}^{kin}(\\rho)$ can become zero or even negative if more than about $15\\%$ nucleons populate the high-momentum tail above the Fermi surface as indicated by the recent experiments done at the Jefferson National Laboratory (J-Lab) by the CLAS Collaboration \\cite{CLAS}. It is very encouraging to note that not only this finding was very recently confirmed qualitatively by three independent studies using the state-of-the-art microscopic many-body theories \\cite{vid,carb,lov,vid12}, our calculation of the direct but second-order tensor contribution to the $E_{sym}^{pot}(\\rho)$ was also qualitatively verified by a more accurate calculation very recently \\cite{Wang12}. ", "conclusions": "Using phenomenological models, we explored effects of the tensor force on the density dependence of nuclear symmetry energy. The high momentum tail in symmetric nuclear matter induced by the tensor force acting between protons and neutron makes the kinetic part of the symmetry energy $E^{kin}_{sym}(\\rho)$ significantly smaller than the Fermi gas model prediction. With about $15\\%$ nucleons in the high momentum tail in SNM as indicated by the recent experiments at the J-Lab, the $E^{kin}_{sym}(\\rho)$ is negligibly small. It even becomes negative when more nucleons are in the high momentum tail in SNM. While at the mean-field level the tensor force has no contribution to the EOS, its second-order contribution to the potential part of the symmetry energy is large. To completely take into account effects of the tensor force, it is necessary to include not only its second-order potential contribution and effects on the kinetic part but also its effects on the central force contribution to the symmetry energy. Implications of these finding on extracting experimental constraints on the density dependence of nuclear symmetry energy are also discussed briefly. \\ack We would like to thank L.W. Chen, A. Lovato, A. Rios, I. Vidana and W. G. Newton for helpful discussions and communications. This work is supported in part by the US National Science Foundation grant PHY-0757839 and PHY-1068022, the National Aeronautics and Space Administration under grant NNX11AC41G issued through the Science Mission Directorate, the National Natural Science Foundation of China (Grants 10805026, 10905048 and 11175085) and the National Basic Research Program of China under Grant 2009CB824800." }, "1207/1207.5280_arXiv.txt": { "abstract": "We construct a parameterized model to explore the main properties of the star formation history of M33. We assume that the disk originates and grows by the primordial gas infall and adopt the simple form of gas accretion rate with one free parameter, the infall time-scale. We also include the contribution of gas outflow process. A major update of the model is that we adopt a molecular hydrogen correlated star formation law and calculate the evolution of the atomic and molecular gas separately. Comparisons between the model predictions and the observational data show that the model predictions are very sensitive to the adopted infall time-scale, while the gas outflow process mainly influences the metallicity profile. The model adopting a moderate outflow rate and an inside-out formation scenario can be in good agreement with most of observed constraints of M33 disk. We also compare the model predictions based on the molecular hydrogen correlated star formation law and that based on the Kennicutt star formation law. Our results imply that the molecular hydrogen correlated star formation law should be preferred to describe the evolution of the M33 disk, especially the radial distributions of both the cold gas and the stellar population. ", "introduction": "The NGC 598 (M33) galaxy is a low-luminosity, late-type disk galaxy in the Local Group. M33 is observed to be much smaller and less massive than the Milky Way galaxy, but has much larger gas fraction. It also shows no signs of recent mergers and no presence of prominent bulge and bar component (Regan \\& Vogel 1994; McLean \\& Liu 1996). In addition, due to its proximity, large angular size, and rather low inclination, M33 is an excellent target for detailed observations of its cold gas, metallicity, the star formation rate (SFR) and stellar population, and thus provides an excellent chance for testing the model of galactic chemical evolution. The star formation (SF) law is one of the important ingredients of the model. Based on the observed data of a sample of 97 nearby normal and star-burst galaxies, Kennicutt (1998) found a power-law correlation between the galaxy-averaged SFR surface density ($\\Psi(r,t)$) and the galaxy-averaged total gas surface density ($\\Sigma_{\\rm gas}(r,t)$), which was termed as the classical Kennicutt SF law. Later, observations of high spatial resolution (less than kpc-scale regions) showed that the SFR surface density correlated stronger with the surface density of the molecular hydrogen ($\\Sigma_{\\rm H_2}(r,t)$) than with that of the atomic hydrogen ($\\Sigma_{\\rm HI}(r,t)$) and the total gas (Wong \\& Blitz 2002; Kennicutt et al. 2007; Bigiel et al. 2008; Leroy et al. 2008). It was also shown that the SFR surface density is almost proportional to $\\Sigma_{\\rm H_2}(r,t)$: \\begin{equation} \\Psi(r,t)=\\Sigma_{\\rm{H_2}}(r,t)/t_{\\rm dep}, \\label{eq:h2sfr} \\end{equation} where $t_{\\rm dep}$ is the molecular hydrogen depletion time. Hereafter, Equ. \\ref{eq:h2sfr} is called as the $\\Sigma_{\\rm H_2}$-based SF law in this paper. Moreover, the gas outflow process may influence the evolution of M33. Garnett (2002) concluded that spiral galaxy with $V_{\\rm rot}\\leq125\\rm km\\,s^{-1}$ may lose some amount of gas in supernova-driven winds and Tremonti et al. (2004) also confirmed this conclusion. The results of Chang et al. (2010) indicated that the gas outflow process plays an important role in the chemical evolution of the disk galaxy since it can bring part of newly formed metal off the galactic disk. They show that the model assuming that the gas outflow efficiency increases as its stellar mass decreases can explain the observed mass-metallicity relation. The chemical evolution of M33 has been studied by several groups in previous studies (Moll\\'{a} et al. 1996; Magrini et al. 2007a; Marcon-Uchida et al. 2010). Magrini et al. (2007a) found that the model adopting an almost constant gas-infall rate can reproduce some of the observed properties of M33, especially the observed relatively high SFR and the shallow abundance gradients. Marcon-Uchida et al. (2010) compared the chemical evolution of the Milky Way, M31 and M33. They found that the model predictions of the Milky Way and that of M31 were in good agreement with the main features of observations, while the model of M33 failed to reproduce the present-day gas surface density in the inner disk. The oxygen abundance was also overestimated by 0.25 dex in the whole disk of M33. In this paper, we build a bridge between the observed data of M33 and its star formation history (SFH) by constructing a parameterized model of its formation and evolution. A major update of the model is that we adopt the $\\Sigma_{\\rm H_2}$-based SF law and this is maybe the first time for the model of M33 to calculate separately the evolution of the atomic and molecular gas. The paper is organized as follows. Section 2 describes the observed features of M33 disk, including the surface brightness, the SFR, the cold gas content and the metallicity etc.. The main assumptions and ingredients of our model are presented in Section 3. The comparisons between the model predictions and the observations are shown in Section 4, and Section 5 summarizes our main conclusions. ", "conclusions": "\\label{sect:summary} In this paper, we construct a parameterized model of the formation and evolution of M33 disk by assuming that the disk originates and grows by the primordial gas infall. The gas infall rate is described by a simple formula with one free parameter, the infall time-scale $\\tau$. We also include the contribution of gas outflow. A molecular hydrogen correlated SF law is adopted to describe how much the cold gas turns into the stellar mass. We numerically calculate the evolution of M33 and compare the model predictions with the observational data. The main results of our model can be summarized as follows: \\begin{enumerate} \\item Based on the observed $\\Sigma_{\\rm H_2}$ and the SFR, we estimate the depletion time of molecular hydrogen $t_{\\rm dep}$ along the disk of M33. It is shown $t_{\\rm dep}$ does not vary very much with the radius, which suggests that the SFE of M33 is almost constant along the disk. We also show that the SFE of M33 is higher than the average value derived by Leroy et al. (2008) on the basis of a large sample of galaxies. \\item Our results show that the model predictions are very sensitive to the adopted infall time-scale. A long infall time-scale will result in blue colors, low metallicity, high H$_{2}$ and H{\\sc i} mass surface densities, high SFR surface density. \\item We also find that the outflow has relatively little effects on the disk stellar population and cold gas content. But it has great influence in shaping the abundance profiles along the M33 disk since it takes a fraction of metals away from the M33 disk due to its low mass potential. \\item The model which adopts a moderate outflow rate and an inside-out formation scenario, that is, the infall time-scale increases with radius, can be in good agreement with the most of observed constraints of M33. Our results suggest that the formation of M33 is quiet and it may not form through violent accretion process. \\item It is shown that the model adopting the Kennicutt SF law predicts much flatter gradients of color, metallicity and mean stellar age and steeper gradients of cold gas than that adopting the $\\Sigma_{\\rm H_2}$-based SF law. Our results imply that, comparing to the Kennicutt SF law, the $\\Sigma_{\\rm H_2}$-based SF law would be more suitable to describe the evolution of the galactic disk, especially for the radial distributions of both the cold gas and the stellar population. \\end{enumerate}" }, "1207/1207.5249_arXiv.txt": { "abstract": "We completed the development of simulation code that is designed to study the behavior of a conjectured dark matter galactic halo that is in the form of a Bose-Einstein Condensate (BEC). The BEC is described by the Gross-Pitaevskii equation, which can be solved numerically using the Crank-Nicholson method. The gravitational potential, in turn, is described by Poisson's equation, that can be solved using the relaxation method. Our code combines these two methods to study the time evolution of a self-gravitating BEC. The inefficiency of the relaxation method is balanced by the fact that in subsequent time iterations, previously computed values of the gravitational field serve as very good initial estimates. The code is robust (as evidenced by its stability on coarse grids) and efficient enough to simulate the evolution of a system over the course of $10^{9}$ years using a finer ($100\\times 100\\times 100$) spatial grid, in less than a day of processor time on a contemporary desktop computer. ", "introduction": "\\label{} The rotation of spiral galaxies does not follow simple predictions based on Newton's laws. Instead, the rotational velocity curve of most spiral galaxies, plotted as a function of radial distance from the galaxy center, remains ``flat'' for a broad range of radii. The standard proposal to resolve this problem is to presume the existence of a ``dark matter halo'', which contains most of the mass of a spiral galaxy. To maintain consistency with the predictions of the most broadly accepted cosmological models, this halo must necessarily consist of ``exotic'' matter, i.e., matter predominantly composed of something other than baryons. The halo must also be collisionless and not interacting with baryonic matter \\cite{Weinberg2008}. The existence of such a halo with a suitable geometry can account for the observed rotation curves of visible matter. However, a difficult problem is to construct a dark matter halo that is gravitationally stable and does not predict excessive dark matter densities in the inner parts of the galaxy where most visible matter resides. This issue is known as the ``cuspy halo problem'' in the relevant literature \\cite{2010AdAst2010E...5D}. A recent proposal \\cite{1994PhRvD..50.3650S,1994PhRvD..50.3655J,2000PhRvL..85.1158H,Sahni2000,2007JCAP...06..025B,2011PhRvD..84d3532C} addresses the cusp problem by a dark matter halo that forms a Bose-Einstein condensate (BEC) \\cite{1924ZPhy...26..178B,Einstein1925}. A particularly intriguing argument is that the condensate dark matter is, in fact, axions \\cite{2012arXiv1210.0040S}. The dynamics of a BEC halo may be determined by the balance of the attractive force of gravity and a repulsive effective long-range interaction \\cite{Khlopov1985,Khlopov2000,Khlopov2002,Khlopov2005} (see also \\cite{Khlopov2004}). In particular, as the dark matter halo dominates the gravitational field of a spiral galaxy in its outer regions, a simulation that is restricted to just the halo should be sufficient to determine if a field can be obtained that yields the desired circular orbital velocities. In the present paper, we discuss a simulation tool that we constructed to explore the dynamics of a galactic BEC halo. The tool is not intended in its present form to study the core-cusp problem; however, we anticipate that it will be useful for investigating the rotational velocities of a galaxy surrounded by a BEC halo. Our work is based primarily on our previous simulation of BEC in laboratory conditions, described by the non-linear Schr\\\"odinger equation, also known in the literature as the Gross-Pitaevskii equation. Whereas in the laboratory, a BEC characterized by a repulsive interaction is held together by an artificially introduced trapping potential, in the case of a galaxy floating in empty space, the trapping potential must be replaced by self-gravity. A numerical solution must, therefore, simultaneously address the initial value problem of the Gross-Pitaevskii equation and the boundary condition problem of Poisson's equation \\cite{2012JCAP...02..011L}. In Sec.~\\ref{sec:GPE}, we introduce the dimensionless form of the Gross-Pitaevskii equation used in our computations, and the method used to solve this equation efficiently. In Sec.~\\ref{sec:PE} we discuss Poisson's equation for gravity and the relaxation method. In Sec.~\\ref{sec:UNITS} we elaborate on the use of physical units that are suitable for such a simulation in an astrophysical context. The problem of using suitable initial conditions to form a stable halo is briefly discussed in Sec. \\ref{sec:INIT}. In Sec.~\\ref{sec:CODE} we discuss the implementation of our method in FORTRAN and C++, and also comment on the possible use of GPUs for accelerated computation. Finally, our conclusions and outlook are presented in Sec.~\\ref{sec:END}. ", "conclusions": "\\label{sec:END} Development of the software code described in this paper is now complete, with the code yielding expected results for test cases. The next step is to find suitable initial conditions to model a BEC dark matter halo that may surround a real galaxy. The question is whether halo configurations can be found that are stable over time scales of $10^{10}$ years, and yield circular orbital velocities that remain approximately constant at different radii. Results of this on-going investigation will be reported when they become available." }, "1207/1207.2696_arXiv.txt": { "abstract": "{ UGC~4483 is a nearby Blue Compact Dwarf (BCD) galaxy. HST observations have resolved the galaxy into single stars and this has led to the derivation of its star formation history and to a direct estimate of its stellar mass. We have analysed archival VLA observations of the 21-cm line and found that UGC~4483 has a steeply-rising rotation curve which flattens in the outer parts at a velocity of $\\sim$20~km~s$^{-1}$. Radial motions of $\\sim$5 km~s$^{-1}$ may also be present. As far as we know, UGC~4483 is the lowest-mass galaxy with a differentially rotating \\hi disk. The steep rise of the rotation curve indicates that there is a strong central concentration of mass. We have built mass models using the HST information on the stellar mass to break the disk-halo degeneracy: old stars contribute $\\sim$50$\\%$ of the observed rotation velocity at 2.2 disk scale-lengths. Baryons (gas and stars) constitute an important fraction of the total dynamical mass. These are striking differences with respect to typical dwarf irregular galaxies (dIrrs), which usually have slowly-rising rotation curves and are thought to be entirely dominated by dark matter. BCDs appear to be different from non-starbursting dIrrs in terms of their \\hi and stellar distributions and their internal dynamics. To their high central surface brightnesses and high central \\hi densities correspond strong central rotation-velocity gradients. This implies that the starburst is closely related with the gravitational potential and the concentration of gas. We discuss the implications of our results on the properties of the progenitors/descendants of BCDs. } ", "introduction": "The mechanisms that trigger strong bursts of star formation in galaxies are poorly understood. In the Local Universe, starburst activity is mostly observed in low-mass galaxies, which are usually classified as blue compact dwarfs (BCDs) (e.g. \\citealt{GilDePaz2003}), amorphous dwarfs (e.g. \\citealt{Gallagher1987}), or \\hii galaxies (e.g. \\citealt{Taylor1995}). Hereafter, we will refer to any starbursting dwarf galaxy as a BCD. Several studies (e.g. \\citealt{GilDePaz2005}, \\citealt{Tosi2009} and references therein) have shown that BCDs are \\textit{not} young galaxies undergoing their first burst of star formation (as suggested by \\citealt{Searle1972}), as they also contain old stellar populations with ages $>$2-3~Gyr. In particular, HST has made it possible to resolve nearby BCDs into single stars and to derive colour-magnitude diagrams deep enough to provide the following information: i) accurate distances of the galaxies, ii) a direct estimate of their total stellar mass, and iii) their star formation history (SFH) (e.g. \\citealt{Tosi2009}). These SFHs show that the starburst is a short-lived phenomenon, typically sustained for a few 10$^{8}$~yr (\\citealt{McQuinn2010}). Thus, BCDs are \\textit{transition-type dwarfs} but the nature of their progenitors and descendants remains unclear. In particular, it is not known whether there are evolutionary connections with dwarf irregulars (dIrrs), spheroidals (dSphs), and/or ellipticals (dEs) (e.g. \\citealt{Papaderos1996, vanZee2001}). There are striking differences between BCDs and other types of dwarf galaxies: i) the old stellar component of BCDs generally has a smaller scale-length and higher central surface brightness than dIrrs and dEs/dSphs (e.g. \\citealt{Papaderos1996, GilDePaz2005}); ii) BCDs have strong concentrations of \\hi within the starburst region, where the column densities are typically 2-3 times higher than in dIrrs (e.g. \\citealt{vanZee1998b, vanZee2001}); iii) BCDs have steep central velocity gradients that are not observed in dIrrs (e.g. \\citealt{vanZee1998b, vanZee2001}). The steep velocity gradients may signify a steeply-rising rotation curve \\citep{vanZee2001, Lelli2012}, high velocity dispersion, or non-circular motions (e.g. \\citealt{Elson2011b}). Detailed studies of the gas kinematics are needed to determine the inner shape of the rotation curve. Recently, \\citet{Lelli2012} studied the BCD prototype I~Zw~18 and found that it has a flat rotation curve with a steep rise in the inner parts, indicating that there is a high central concentration of mass. Such a mass concentration is not observed in typical dIrrs. This points to a close connection between the starburst and the gravitational potential. It is also clear that a BCD like I~Zw~18 cannot evolve into a typical dIrrs at the end of the starburst, unless the central concentration of mass is removed. It is important, therefore, to investigate whether all BCDs have steeply-rising rotation curves, and to determine the relative contributions of gas, stars and dark matter to the gravitational potential. \\begin{table}[thp] \\caption{Properties of UGC~4483.} \\label{tab:prop} \\centering \\begin{tabular}{l c c} \\hline \\hline $\\alpha$ (J2000) & $08^{\\rm{h}} 37^{\\rm{m}} 3.^{\\rm{s}}1 \\pm 0.^{\\rm{s}}5$ \\\\ $\\delta$ (J2000) & $69^{\\circ} 46' 31'' \\pm 2''$ \\\\ Distance (Mpc) & $3.2 \\pm 0.2$ \\\\ $V_{\\rm{sys}}$ (km s$^{-1}$) & $158 \\pm 2$ \\\\ Position Angle ($^{\\circ}$) & $0 \\pm 5$ \\\\ Inclination Angle ($^{\\circ}$) & $58 \\pm 3$ \\\\ $V_{\\rm{rot}}$ (km s$^{-1}$) & $19 \\pm 2$ \\\\ $M_{\\rm{dyn}}$ (10$^{7}$ $M_{\\odot}$) & $16 \\pm 3$ \\\\ $M_{*}$ (10$^{7}$ $M_{\\odot}$) & $1.0 \\pm 0.3$ \\\\ $M_{\\hi}$ (10$^{7}$ $M_{\\odot}$) & $2.5 \\pm 0.3$ \\\\ $L_{\\rm{B}}$ (10$^{7}$ $L_{\\odot}$) & 1.4 \\\\ $L_{\\rm{R}}$ (10$^{7}$ $L_{\\odot}$) & 0.9 \\\\ \\hline \\end{tabular} \\tablefoot{Luminosities were calculated using the apparent magnitudes from \\citet{GilDePaz2003}, the distance from \\citet{Dolphin2001} and the solar absolute magnitudes from \\citet{BinneyMerrifield1998}. The stellar mass was calculated integrating the SFH from \\citet{McQuinn2010} and assuming a gas recycling efficiency of 30$\\%$. The dynamical mass was calculated taking into account the pressure-support.} \\end{table} We present a detailed study of the gas kinematics of UGC~4483, a starbursting dwarf galaxy located in the M81 group and resolved into individual stars by HST \\citep{Dolphin2001, Izotov2002}. UGC~4483 is extremely metal-poor (12+log(O/H)$\\simeq$7.5, see \\citealt{Skillman1994} and \\citealt{vanZee2006}) and may be classified as a ``cometary'' BCD, as its high-surface-brightness starburst region is located at the edge of an elongated low-surface-brightness stellar body (see Fig.~\\ref{fig:mosaic}, top-left). Previous \\hi studies (\\citealt{Lo1993, vanZee1998b}) showed that the galaxy has an extended \\hi disk, with a strong \\hi concentration near the starburst region and a steep central velocity gradient, but the \\hi kinematics was not studied in detail. We analysed archival VLA data and were able to derive a rotation curve which we used to investigate the distributions of luminous and dark matter in this galaxy. ", "conclusions": "We analysed archival \\hi observations of the blue compact dwarf galaxy UGC~4483 and built model datacubes to investigate its gas kinematics. Our main results can be summarized as follows: \\begin{itemize} \\item UGC~4483 has a steeply-rising rotation curve that flattens in the outer parts at a velocity of $\\sim$20 km s$^{-1}$. This is, to our knowledge, the lowest-mass galaxy with a differentially rotating \\hi disk. Radial motions of $\\sim$5 km~s$^{-1}$ may also be present. \\item The steep rise of the rotation curve indicates that there is a strong central concentration of mass. Mass models with a dark matter halo show that \\textit{old} stars contribute $\\sim$50$\\%$ of the observed rotation speed at 2.2 disk scale-lengths. Baryons (gas and stars) constitute an important fraction of the total dynamical mass. These conclusions are based on the stellar mass obtained from the color-magnitude diagram of the resolved stellar populations. \\item The maximum-disk solution requires a stellar mass 3 times higher than observed, that could be provided by molecules. A good solution is also found by scaling the \\hi contribution by a factor of $\\sim$5. These results suggest that the distribution of the dynamical mass is closely coupled to that of the baryons. \\end{itemize} UGC~4483, together with other BCDs like I~Zw~18 and NGC~1705, appears structurally different from typical dIrrs in terms of \\hi distribution, stellar distribution, and dynamics. In particular, a central concentration of mass (gas, stars, and dark matter) seems to be a characterizing property of BCDs. This implies that the starburst is closely related with the gravitational potential and the \\hi concentration. Our results also suggest that the progenitors/descendants of BCDs must be compact dwarf galaxies, unless a redistribution of mass (both luminous and dark) takes place before/after the starbursting phase." }, "1207/1207.7285_arXiv.txt": { "abstract": "We provide an analytic solution to the general wavelength integro-differential equation describing the damping of tensor modes of gravitational waves due to free streaming neutrinos in the early universe. Our result is expressed as a series of spherical Bessel functions whose coefficients are functions of the reduced wave number $Q$. ", "introduction": "Observations of the cosmic microwave background (CMB) have given increasing support to inflationary cosmological models. Density perturbations during this inflationary period are believed to have given rise to the large scale structures of the universe \\cite{GHS}. In addition to these scalar perturbations, a spectrum of gravitational waves is also produced \\cite{SP} (tensor perturbations) which could provide information about the early universe. In particular, the contribution of these tensor modes (measured in terms of a tensor-to-scalar ratio \\cite{K}) to the temperature anisotropy of the CMB and to B-mode polarization of CMB photons could be used to check the predictions of inflationary models. Here we shall consider the effect of anisotropic inertia on the gravitational radiation and confirm that it is non-negligible. Although it will be assumed here that the anisotropic inertia is dominated by neutrinos and antineutrinos \\cite{SW}, Ref. \\cite{WZ} has proposed a free streaming relativistic gas contribution and a method to constrain its fraction density through measurements of the CMB B-polarization spectra. In a 2004 paper \\cite{SW}, Weinberg derived an integro-differential equation for the propagation of cosmological gravitational waves. He writes a wave equation for the perturbation to the metric $h_{ij}({\\bf x},t)$ and defines $\\chi(u)$ as \\begin{equation} h_{ij}=h_{ij}(0) \\ts \\chi(u)\\,, \\end{equation} where $u$ is the conformal time multiplied by the wave number \\begin{equation} u=k \\int^{t} \\frac{dt'}{a(t')}\\,. \\end{equation} The wave equation for the perturbation leads to an integro-differential equation which the function $\\chi$ satisfies for general wavelengths. In the variable $y=a(t)/a_{EQ}$, where $a_{EQ}$ is the expansion parameter at matter-radiation equality, this is Eq.\\,(32) of \\cite{SW} \\begin{equation} (1+y)\\chi''(y) + \\left( \\frac{2(1+y)}{y} + \\frac{1}{2}\\right) \\chi'(y) + Q^{2}\\chi(y) = -\\frac{24 \\ts f_{\\nu}(0)}{y^{2}} \\int_{0}^{y} K(y,y') \\frac{d\\chi(y')}{dy'} dy'\\,, \\label{eq:yEQ} \\end{equation} with the initial conditions: \\begin{equation} \\chi(0)=1\\,, \\hspace{10mm} \\chi'(0)=0\\,. \\label{eq:yInitCond} \\end{equation} Here $f_{\\nu}(0)=0.40523$ is the fraction of the energy density in neutrinos and $Q=\\sqrt{2}k/k_{EQ}$. The kernel $K$ will be discussed below, and $y$ is related to $u$ in the following manner \\begin{equation} u=2Q\\left( \\sqrt{1+y}-1 \\right)\\equiv Q\\,s\\,. \\end{equation} We will provide an analytic solution of Eqs.\\,(\\ref{eq:yEQ}) and (\\ref{eq:yInitCond}) that is valid for general wavelengths. It will be shown that the effect of the anisotropic inertia is to damp the amplitude of the perturbation relative to the case where free-streaming neutrinos are absent, and that, in general, this damping depends on $Q$. Thus, it is important to have a solution capable of describing the damping for all wavelengths. For $Q^2\\gg 1$, it has been shown that the solution for $\\chi(u)$ can be written as a series of spherical Bessel functions whose coefficients are independent of $Q$ \\cite{DR}. In \\cite{SW}, Weinberg analyzed the $Q^2\\gg 1$ limit, provided results for the damping factor when $Q^2\\ll 1$ and made some observations about the damping for general value of $Q$. Our aim here is to show that it is possible to extend the spherical Bessel function expansion \\cite{DR,GS} for $\\chi(u)$ to the case of general $Q$. Having done this, we can recapture the $Q^2\\ll 1$ results of \\cite{SW}, the $Q^2\\gg 1$ results of \\cite{DR} and obtain precise results for intermediate values of $Q$. In the next section, we derive the equation that must be satisfied by a spherical Bessel function expansion for $\\chi(u)$, examine its form in the $Q^2\\gg 1$, obtain a recurrence relation for general $Q$, and use this relation to determine the damping factor for $Q\\ll 1$. Section 3 contains a discussion of general wavelength case and we end with some conclusions. Lists of the various expansion coefficients and the details related to the solution of the recurrence relation are contained in the Appendices. ", "conclusions": "We have shown that the treatment of gravitational wave damping by free streaming neutrinos can be framed in terms of a series of spherical Bessel functions for all values of the reduced wave number $Q$. The result for the coefficients of the spherical Bessel series when $Q \\gg 1$ \\cite{DR} emerges quite simply from the coefficient recurrence relation for a general $Q$. For the opposite limit, $Q \\ll 1$, the analysis can then be extended to arbitrarily small values of $Q$ by using the low $Q$ expansion of $j_n(Q\\,s_L)$. For intermediate values of $Q$, the coefficients obtained from the general recurrence relation can be used directly for any $Q$, provided that the factor $Q$ is retained in the argument of the spherical Bessel functions. There have been numerous successful numerical studies of the damping of the gravitational wave spectrum due to anisotropic inertia (Refs. \\cite{WK}, \\cite{Kuro}, \\cite{Latt}, \\cite{Kin}). The present approach represents a trade off between the use of reliable sophisticated numerical techniques to calculate $\\chi(s,Q)$ and the necessity of evaluating the exact series solution for $\\chi(s,Q)$ to sufficient accuracy. The main advantage of using the series expansion is that the convolution of the desired solution with the kernel can be evaluated exactly in terms of spherical Bessel functions. We used 20 terms in our expansions of $\\chi(s_L,Q)$ and $\\chi_0(s_L,Q)$. To illustrate the accuracy that this truncated series provides, Fig.\\,\\ref{diff}, compares the amplitude $\\chi(s_L,Q)$ computed with 20 terms and 100 terms. \\begin{figure}[h!] \\centering \\includegraphics[scale=0.75]{Fig2.eps} \\caption{The comparison between $\\chi(s_L,Q)$ with 20 terms (red dashed) and $\\chi(s_L,Q)$ with 100 terms (blue) is shown for $Q\\leq 20$.} \\label{diff} \\end{figure} As another test of the use of 20 terms rather than 100 terms in the series Eq.\\,(\\ref{eq:sBesselModel}), we compare the ratio of the damping factor $R=|\\chi'(s_L,Q)/\\chi'_0(s_L,Q)|^2$ for the two cases as a function of $Q$ in Table\\,\\ref{Delta}. \\begin{table}[h!] \\centering \\begin{tabular}{| c | c |} \\hline $Q$ & $R_{100}/R_{20}$ \\\\ \\hline 0.01 & $1.00003$ \\\\ \\hline 0.1 & $1.00003$ \\\\ \\hline 1 & $1.00002$ \\\\ \\hline 10 & $1.04926$ \\\\ \\hline $10^{2}$ & $1.37376$ \\\\ \\hline $10^{3}$ & $0.99772$ \\\\ \\hline $10^{4}$ & $1.00005$ \\\\ \\hline $10^{5}$ & $1.00002$ \\\\ \\hline $10^{6}$ & $0.99999$ \\\\ \\hline \\end{tabular} \\caption{The ratio of the damping factor $|\\chi'(s_L,Q)/\\chi'_0(s_L,Q)|^2$ is shown for 20 and 100 terms in the spherical Bessel function expansion.} \\label{Delta} \\end{table} The only region where the reduction in the number of terms significantly affects the answer is the transition region $10\\leq Q\\leq 100$. If this region is important, additional terms would be required, but are readily computable from the provided recurrence relations." }, "1207/1207.4579.txt": { "abstract": "The heliosphere represents a uniquely accessible domain of space, where fundamental physical processes common to solar, astrophysical and laboratory plasmas can be studied under conditions impossible to reproduce on Earth and unfeasible to observe from astronomical distances. Solar Orbiter, the first mission of ESA's Cosmic Vision 2015--2025 programme, will address the central question of heliophysics: How does the Sun create and control the heliosphere? In this paper, we present the scientific goals of the mission and provide an overview of the mission implementation. ", "introduction": "\\label{S-Introduction} We live in the extended atmosphere of the Sun, a region of space known as the heliosphere. Understanding the connections and the coupling between the Sun and the heliosphere is of fundamental importance to understanding how our solar system works. The results from current and past solar and heliospheric missions such as Helios \\cite{Porsche:1977aa,Schwenn:1990aa,Schwenn:1991ab}, Voyager \\cite{Stone:1977aa}, Ulysses \\cite{Wenzel:1992aa}, Yohkoh \\cite{Acton:1992aa}, SOHO \\cite{Domingo:1995aa}, TRACE \\cite{Handy:1999aa}, RHESSI \\cite{Lin:2002aa}, Hinode \\cite{Kosugi:2007aa}, STEREO \\cite{Kaiser:2008aa} and SDO \\cite{Pesnell:2012aa} have formed the foundation of our understanding of the solar corona, the solar wind, and the three-dimensional heliosphere. Each of these missions had a specific focus, being part of an overall strategy of coordinated solar and heliospheric research. However, none of these missions have been able to fully explore the interface region where the solar wind is born and heliospheric structures are formed with sufficient instrumentation to link solar wind structures back to their source regions at the Sun (Helios 1 and 2, e.g., carried no imaging instruments). This is the goal of Solar Orbiter, a mission of collaboration between ESA and NASA that was recently selected as the first medium (M)-class mission of ESA's Cosmic Vision 2015--2025 programme. With a combination of in-situ and remote-sensing instruments and its inner-heliospheric mission design, Solar Orbiter will address the central question of heliophysics: How does the Sun create and control the heliosphere? This primary, overarching scientific objective can be expanded into four interrelated top-level scientific questions that will be addressed by Solar Orbiter: \\begin{itemize} %\\item How and where do the solar wind plasma and magnetic field originate in the corona? \\item What drives the solar wind and where does the coronal magnetic field originate from? \\item How do solar transients drive heliospheric variability? \\item How do solar eruptions produce energetic particle radiation that fills the heliosphere? \\item How does the solar dynamo work and drive connections between the Sun and the heliosphere? \\end{itemize} These questions represent fundamental challenges in solar and heliospheric physics today. By addressing them, we expect to make major breakthroughs in our understanding of how the inner solar system works and is driven by solar activity. To answer these questions, it is essential to make in-situ measurements of the solar wind plasma, fields, waves, and energetic particles close enough to the Sun that they are still relatively pristine and have not had their properties modified by subsequent transport and propagation processes. This is one of the fundamental drivers for the Solar Orbiter mission, which will approach the Sun to as close as 0.28\\,AU. Relating these in-situ measurements back to their source regions and structures on the Sun requires simultaneous, high-resolution imaging and spectroscopic observations of the Sun in and out of the ecliptic plane. The resulting combination of in-situ and remote-sensing instruments on the same spacecraft, together with the new, inner-heliospheric perspective, distinguishes Solar Orbiter from all previous and current missions, enabling science which can be achieved in no other way. The following section introduces the science payload and mission design. Section~\\ref{S-Science} describes the mission's science objectives in detail: The present state of knowledge is presented for all major science questions of the mission, followed by descriptions of how Solar Orbiter will advance our understanding. Section~\\ref{S-Spacecraft} introduces the Solar Orbiter spacecraft, followed by an overview of the science operations in Section~\\ref{S-Operations}. Table \\ref{T-MissionSummary} gives a one-page mission summary. %%%%%%%%% \\begin{table} \\caption{Solar Orbiter Mission Summary.} \\label{T-MissionSummary} \\begin{tabular}{lp{7cm}} \\hline {\\bf Top-level Science Questions} & \\begin{itemize} %\\item How and where do the solar wind plasma and magnetic field originate in the corona? \\item What drives the solar wind and where does the coronal magnetic field originate from? \\item How do solar transients drive heliospheric variability? \\item How do solar eruptions produce energetic particle radiation that fills the heliosphere? \\item How does the solar dynamo work and drive connections between the Sun and the heliosphere? \\end{itemize}\\vspace{-6mm}\\\\ \\hline {\\bf Science Payload} & {\\bf In-Situ Instruments:} \\begin{itemize} \\item Energetic Particle Detector (EPD) \\item Magnetometer (MAG) \\item Radio and Plasma Wave analyser (RPW) \\item Solar Wind Analyser (SWA) \\end{itemize} {\\bf Remote-Sensing Instruments:} \\begin{itemize} \\item EUV full-Sun and high-resolution Imager (EUI) \\item Coronagraph (METIS) \\item Polarimetric and Helioseismic Imager (PHI) \\item Heliospheric Imager (SoloHI) \\item EUV spectral Imager (SPICE) \\item X-ray spectrometer/telescope (STIX) \\end{itemize}\\vspace{-6mm}\\\\ \\hline {\\bf Mission Profile} & \\vspace{-6mm} \\begin{itemize} \\item Launch on NASA-provided Evolved Expendable Launch Vehicle (Ariane 5 as back-up) \\item Interplanetary cruise with chemical propulsion and gravity assists at Earth and Venus \\item Venus resonance orbits with multiple gravity assists to increase inclination \\end{itemize} \\vspace{-6mm}\\\\ \\hline {\\bf Closest Perihelion} & 0.28\\,AU\\\\ \\hline {\\bf Max. Heliolatitude} & 25$^\\circ$ (nominal mission) / 34$^\\circ$--36$^\\circ$ (extended mission)\\\\ \\hline {\\bf Spacecraft} & 3-axis stabilized platform, heat shield, two adjustable, single-sided solar arrays, dimensions: 2.5 $\\times$ 3.0 $\\times$ 2.5 ${\\rm m}^3$ (launch configuration) \\\\ \\hline {\\bf Telemetry Band} & Dual X-band\\\\ \\hline {\\bf Data Downlink} & 150 kbit/s at 1\\,AU spacecraft--Earth distance\\\\ \\hline {\\bf Launch Date} & Jan-2017 (Mar-2017 and Sep-2018 back-ups)\\\\ \\hline {\\bf Nominal Mission Duration} & 7 years (incl.\\ cruise phase) \\\\ \\hline {\\bf Extended Mission Duration} & 3 years\\\\ \\hline \\end{tabular} \\end{table} %%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%% \\newpage ", "conclusions": "Understanding the connections and the coupling between the Sun and the heliosphere is of fundamental importance to understanding how our solar system works. To reach this goal, Solar Orbiter will make in-situ measurements of the solar wind plasma, fields, waves, and energetic particles as close as 0.28\\,AU from the Sun, simultaneously with high-resolution imaging and spectroscopic observations of the Sun in and out of the ecliptic plane. The combination of in-situ and remote sensing instruments on the same spacecraft, together with the new, inner-heliospheric perspective, distinguishes Solar Orbiter from all previous and current missions, enabling breakthrough science which can be achieved in no other way. In addition to delivering ground-breaking science in its own right, Solar Orbiter also has important synergies with NASA's Solar Probe Plus mission. Coordinated observations with this mission, combined with data from other missions operating in the inner heliosphere (or providing remote-sensing observations of the near-Sun environment), will contribute greatly to our understanding of the Sun and its environment. %%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%% \\begin{acks} Contributions to this paper were provided by the PIs and Co-PIs, the ESA Solar Orbiter Project Team, the NASA Solar Orbiter Collaboration Project Team, E.~Marsch (MPS Lindau), M.~Velli (JPL/U.\\ Firenze), C.~DeForest (SwRI Boulder), D.~Hassler (SwRI Boulder) and W.~Lewis (SwRI San Antonio). The authors would like to thank the referee, guest editor and journal editors for comments and suggestions, which helped to improve the quality of this work. \\end{acks}" }, "1207/1207.2469_arXiv.txt": { "abstract": "In this paper, we study small-$N$ gravitational dynamics involving up to six objects. We perform a large suite of numerical scattering experiments involving single, binary, and triple stars. This is done using the FEWBODY numerical scattering code, which we have upgraded to treat encounters involving triple stars. We focus on outcomes that result in direct physical collisions between stars, within the low angular momentum and high absolute orbital energy regime. The dependence of the collision probability on the number of objects involved in the interaction, $N$, is found for fixed total energy and angular momentum. Our results are consistent with a collision probability that increases approximately as $N^2$. Interestingly, this is also what is expected from the mean free path approximation in the limit of very large $N$. A more thorough exploration of parameter space will be required in future studies to fully explore this potentially intriguing connection. This study is meant as a first step in an on-going effort to extend our understanding of small-$N$ collisional dynamics beyond the three- and four-body problems and into the realm of larger-$N$. ", "introduction": "\\label{intro6} The gravitational three-body problem was first studied by Sir Isaac Newton in his Principia \\citep{newton1686}. After solving the two-body problem, Newton boldly added a third body into the mix and attempted to create a similar mathematical formalism to describe the motion of three celestial bodies under their mutual gravitational attraction. The solution evaded Newton until his end, leaving the problem at the mercy of his disciples. It was later taken up by Euler \\citep{euler1772}, Lagrange \\citep{lagrange1811}, Jacobi \\citep{jacobi1836}, Poincare \\citep{poincare1892}, Hill \\citep[e.g.][]{hill1878}, Henon \\citep[e.g.][]{henon69}, and a host of others, often with a focus on the Earth-Moon-Sun system \\citep[e.g.][]{valtonen06}. Despite the considerable efforts of numerous researchers, a simple analytic solution to the general three-body problem has never been found. \\citet{sundman1912} discovered a uniformly convergent infinite series involving familiar functions that technically solves the three-body problem. However, in order to attain a reasonable level of accuracy, the solution must contain on the order of $10^{8,000,000}$ terms \\citep{valtonen06}. More practical solutions require a number of simplifying assumptions to make the general three-body problem tractable \\citep[e.g.][]{henon69}. As a result, the most useful analytic solutions tend to be solely applicable to a very narrow subset of the total allowed parameter space. The introduction of computers within the last few decades revolutionalized our understanding of the three-body problem. This allowed for the direct integration of the equations of motion for each particle over small time steps, incrementally moving each body forward in an iterative fashion. Despite the fact that this approach completely transformed our tool-set for studying the three-body problem, it too has its short-comings. For example, the times required to run the simulations to completion are often quite long. The precise definition of a ``complete encounter'' can often be ambiguous as well. Quasi-stable configurations can form that remain bound for millions or even billions of years before eventually breaking apart \\citep[e.g.][]{mikkola86}. There is also the issue of errors in computing the trajectories of the particles in position-space which are introduced at each time-step. These arise as a result of moving one body forward at a time, instead of moving all bodies simultaneously. Such errors can be minimized by taking suitably short time-steps. However, this comes at the often considerable cost of simulation run-time. When only three bodies are involved, qualitative descriptions of the outcomes of dynamical interactions are often relatively straight-forward. They include ionizations, exchanges, fly-bys, and collisions \\citep[e.g.][]{hut83b, mikkola83, mikkola84b, hut84, mcmillan96, fregeau04}. However, the number of possible outcomes quickly increases with increasing $N$, where $N$ is the total number of objects involved in the interaction \\citep{leigh11}. This complicates descriptions of the interactions, and introduces a considerable challenge in developing a physical understanding of the evolution of the system. For example, nearly 100 generic outcomes are possible for encounters involving six objects. Additional bodies not only increase the parameter space to be explored, they also increase integration times for computer simulations. Consequently, most previous numerical scattering studies considered only single-binary and, to a lesser extent, binary-binary encounters \\citep[e.g.][]{heggie75, hut83a, mcmillan86, sigurdsson93, davies95, bacon96, fregeau04}. Many of the numerical scattering studies conducted to date considered equal-mass particles, and devoted their attention to studying the effects of varying the initial relative velocity or impact parameter \\citep[e.g.][]{hut83b}. For example, \\citet{hut83a} found analytic approximations for high-velocity encounters, and provided simple formulae for the corresponding cross-sections for exchanges and ionizations to occur. Similar analytic formulae were derived by \\citet{mikkola84a} for encounters involving two binaries having unequal orbital energies but equal masses. An extensive analysis that considered unequal mass particles was conducted by \\citet{sigurdsson93}, who studied the effects of dynamics on the stellar and remnant populations in a globular cluster. A number of other scattering experiments were designed purely to study the formation of different types of stellar exotica in globular clusters, including blue stragglers \\citep[e.g.][]{leonard89}, low-mass x-ray binaries \\citep[e.g.][]{sigurdsson95}, cataclysmic variables \\citep[e.g.][]{ivanova06}, and millisecond pulsars \\citep[e.g.][]{ivanova08}. Many of these also considered a range of particle masses. The effects of general relativity have also been studied in the context of the three- and four-body problems. For example, \\citet{mikkola90} and \\citet{valtonen94} considered interactions involving binary super-massive black holes during the mergers of galaxies, and identified several important trends arising from energy losses due to gravitational radiation. In this paper, we study gravitational interactions involving up to six objects. Our goal is to better understand how the outcome of an encounter depends on the number of interacting objects. This is a first step toward extending techniques developed for the three-body problem, where the vast majority of research efforts have thus far been concentrated, to treat larger-$N$ encounters. To this end, we perform $>10^7$ numerical scattering experiments involving single, binary and triple star systems. The number of possible encounter outcomes is a steeply increasing function of $N$. This presents a considerable challenge when trying to draw parallels between encounters involving different numbers of objects. To minimize this issue, we focus only on outcomes resulting in direct physical collisions, which we consider to have occurred when the radii of any two stars overlap. In Section~\\ref{method}, we describe the set-up for our numerical scattering experiments, including the range of initial conditions considered. Additionally, we develop an analytic formula for the collision probability as a function of $N$, that is based on a simple analytic model originally derived for 1+2 encounters. In Section~\\ref{results}, we present the results of our analysis of this large suite of single-binary (1+2), binary-binary (2+2), single-triple (1+3), binary-triple (2+3), and triple-triple (3+3) scattering experiments. Here, we fit the analytic formula to the results from these numerical scattering experiments, thereby deriving the $N$-dependence of the collision probability. The implications of our analysis for small-$N$ collisional dynamics are discussed in Section~\\ref{discussion}. Concluding remarks are given in Section~\\ref{summary}. Finally, in an appendix, we present a formalism for creating schematic diagrams that quantitatively depict the evolution of an interaction in energy-space, and briefly discuss some of their possible applications. ", "conclusions": "\\label{discussion} Within the angular momentum and energy regime studied here, we find that the probability of a direct physical collision occurring during an interaction increases roughly as $N^2$. One interpretation of this result can be understood as follows. Previous numerical scattering experiments of 1+2 interactions have shown that the collision probability is proportional to the average number of close approaches experienced by the system per crossing time, multiplied by the number of crossing times survived through \\citep{valtonen06}. Therefore, this $N^2$ dependence may arise physically from an $N^2$ dependence in the total number of close approaches experienced by the system. This can occur in one of (at least) two ways. Either the number of close approaches per crossing time scales as $N^2$ while the number of crossing times survived through is constant, or the number of crossing times the system survives through (until the time when the first collision occurs) scales as $N^2$ while the number of close approaches per crossing time is constant. We find that the average number of crossing times at the time of the first collision is the same to within roughly a factor of $\\approx$ 2 for all encounter types, and that there is no trend with $N$. Therefore, it is unlikely that the latter scenario described above is responsible for producing the $N^2$ dependence. This suggests that perhaps the $N^2$ dependence in the collision probability may arise from a similar dependence in the number of close approaches per crossing time. However, in practice, this hypothesis is much more difficult to test quantitatively. In particular it is not immediately clear how to define a ``close approach''. We attempt to quantify this effect by examining animations of a handful of encounters in position-space, and counting the number of close approaches per crossing time by eye. It is clear that the number of close approaches per crossing time increases with increasing $N$. However, this method is far from adequate to determine the precise $N$-dependence of this relation at any statistically significant level. In Appendix~\\ref{appendix}, we present a method that will improve upon this component of our analysis in future studies. Specifically, we present a prescription for generating schematic diagrams that depict the evolution of an interaction in energy-space. As is explicitly shown in Appendix~\\ref{appendix}, the advantage of this technique is to provide a straight-forward means of defining the criterion of ``close approach'' in terms of the fraction of the total energy exchanged between two individual stars. This directly relates the definition of close approach to the total energy and therefore the initial conditions of the encounter, which is not the case if this criterion is defined purely in terms of a distance. Also note that, in a sense, a collision could be interpreted as a very strict definition of a ``close approach''. Here multiple crossing times are required before this definition of close approach is satisfied. As we find that the number of crossing times until the time of first collision is roughly equivalent (to within a factor of $\\approx$ 2) for all encounter types, our result can potentially also be expressed as the number of ``close approaches'', defined in this manner, per crossing time scaling as $N^2$. We will explore a range of other criteria to define a close approach in a future study and investigate explicitly the dependence of the number of close approaches per crossing time on the total number of stars involved in the interaction. Intriguingly, the collision rate, as derived from the mean free path approximation \\citep[e.g.][]{leonard89}, also scales as $N^2$. In the limit of very large $N$, we would expect our relation for the collision probability to agree with what is predicted from the mean free path approximation. On the other hand, in the limit of small-$N$, it is not clear that the standard assumptions of the mean free path approximation should still hold. For the angular momentum and energy regime studied here, the collision probability appears consistent with that of the mean free path approximation, at least for $N = 4, 5, 6$. Further study is required to determine, e.g., how different energy and angular momentum regimes affect this relationship, what the minimum $N$ is for which the standard mean free path approximations are still valid, etc. The second most noticable similarity between all Runs is an exponential decline in the collision probability with increasing angular momentum that is steeper for larger $N$. This effect is small, however, relative to the total drop in the collision probability as $N$ decreases within a given Run (see Figures~\\ref{fig:prob-merge-run1}, \\ref{fig:prob-merge-run2}, and~\\ref{fig:prob-merge-run3}). One possible explanation that is consistent with our results is that the number of close approaches per crossing time may decrease with increasing angular momentum more steeply for larger $N$ encounters. Our results are applicable to a regime of low angular momentum and high absolute orbital energy. As the integration time increases with increasing angular momentum, we focussed our attention on the low-angular momentum regime in this paper to maximize the number of simulations performed for each Run, and thereby increase the statistical significance of our results. We cannot probe lower angular momentum encounters without violating our assumptions. There are two reasons for this. First, the lower boundary for the long-term stability of triple systems corresponds to a ratio between the inner and outer orbital separations $\\gtrsim 7$ \\citep{mardling01}. Therefore, we cannot lower the total angular momentum by decreasing the semi-major axis of the outer orbit of the triple without simultaneously reducing the semi-major axis of its inner orbit. This brings us to our second requirement, namely that $q \\ll a_0$ (where $a_0$ is the semi-major axis of the shortest-period orbit initially involved in the interaction and $q$ is the physical sizes of the objects involved in the encounter). We performed two additional Runs with $a_0 = 0.1$ AU to test the limit of the assumption $q \\ll a_0$. In both cases, the resulting chi-squared values found using the best-fits for our free parameters in Equation~\\ref{eqn:prob-coll-N} were significantly larger than we found for our other Runs. We interpret this as being due to a breakdown of the requirement $q \\ll a_0$. Our results suggest that the assumption $q \\ll a_0$ holds provided $a_0 \\gtrsim 100q$. This is only a rough estimate for the lower limit for the ratio $q/a_0$, and more work will be needed to better constrain its precise value. In future work, we intend to address the $N$-dependence of the collision probability for higher angular momentum encounters, as well as different mass-ratios and eccentricities. A non-circular orbit provides an additional free parameter that can be changed to affect the total angular momentum, but not the total energy. Therefore, the addition of a non-zero eccentricty should in principle allow us to include 1+2 encounters in our estimates for the collision probability using the same or similar normalization method (i.e. fixing the total angular momentum and energy) as adopted in this paper to compare between the different encounter types." }, "1207/1207.0831_arXiv.txt": { "abstract": "{We report on a detailed study of the Fe K emission/absorption complex in the nearby, bright Seyfert 1 galaxy Mrk~509. The study is part of an extensive \\xmm\\ monitoring consisting of 10 pointings ($\\sim60$ ks each) about once every four days, and includes also a reanalysis of previous \\xmm\\ and \\chandra\\ observations.} {We aim at understanding the origin and location of the Fe K emission and absorption regions.} {We combine the results of time-resolved spectral analysis on both short and long time-scales including model independent rms spectra.} {Mrk~509 shows a clear (EW$=58\\pm4$ eV) neutral \\FeKa\\ emission line that can be decomposed into a narrow ($\\sigma=0.027$ keV) component (found in the \\chandra\\ HETG data) plus a resolved ($\\sigma=0.22$ keV) component. We find the first successful measurement of a linear correlation between the intensity of the resolved line component and the 3-10 keV flux variations on time-scales of years down to a few days. The Fe K$\\alpha$ reverberates the hard X-ray continuum without any measurable lag, suggesting that the region producing the resolved \\FeKa\\ component is located within a few light days-week (r $\\lsimeq~10^3$ r$_{\\rm g}$) from the Black Hole (BH). The lack of a redshifted wing in the line poses a lower limit of $\\geq$40 r$_{\\rm g}$ for its distance from the BH. The Fe K$\\alpha$ could thus be emitted from the inner regions of the BLR, i.e. within the $\\sim$80 light days indicated by the H$\\beta$ line measurements. In addition to these two neutral \\FeKa\\ components, we confirm the detection of weak (EW$\\sim8-20$ eV) ionised Fe K emission. This ionised line can be modeled with either a blend of two narrow \\Fevc\\ and \\Fevs\\ emission lines (possibly produced by scattering from distant material) or with a single relativistic line produced, in an ionised disc, down to a few r$_{\\rm g}$ from the BH. In the latter interpretation, the presence of an ionised standard $\\alpha$-disc, down to a few r$_{\\rm g}$, is consistent with the source high Eddington ratio. Finally, we observe a weakening/disappearing of the medium and high velocity high ionisation Fe K wind features found in previous \\xmm\\ observations.} {This campaign has made possible the first reverberation measurement of the resolved component of the Fe K$\\alpha$ line, from which we can infer a location for the bulk of its emission at a distance of r$\\sim40-1000$ r$_{\\rm g}$ from the BH.} ", "introduction": "X-ray observations of AGN have shown the almost ubiquitous presence of the \\FeKa\\ line at 6.4 keV (Yaqoob et al. 2004; Nandra et al. 1997; 2007; Bianchi et al. 2007; de la Calle et al. 2010). Unlike the optical-UV lines that are emitted by distant material only, the \\FeKa\\ line traces reflection not only from distant material (such as the inner wall of the molecular torus, the broad line region and/or the outer disc) but also from regions as close as a few r$_{\\rm g}$ (where r$_{\\rm g}$=GM/c$^2$) from the BH (Fabian et al. 2000). The powerful reverberation mapping technique, routinely exploited on optical-UV lines (Clavel et al. 1991; Peterson 1993; Kaspi et al. 2000; Peterson et al. 2004), can also be applied to X-ray lines such as the \\FeKa\\ line. This kind of analysis has a tremendous potential, allowing us to map the geometry of matter surrounding the BH, starting from distances of a few gravitational radii up to light years. However, each \\FeKa\\ component is expected to respond on a different characteristic time (years-decades for the torus, several days-months for the BLR-outer disc, and tens of seconds to a few hours for the inner accretion disc) and current X-ray instruments cannot easily disentangle the different components. Indeed, reverberation mapping of all \\FeKa\\ emission components represents an enormous observational challenge, and specially tailored monitoring campaigns (to sample the proper time scales) have to be designed. Since the detection of the first clear example of a broad and skewed Fe line profile in the spectrum of an AGN (indicating that most of the line emission is produced within few tens of r$_{\\rm g}$; e.g. MCG-6-30-15, Tanaka et al. 1995) the quest to understand how the broad \\FeKa\\ line varies with the continuum is ongoing. Indeed, close to the BH the simple one-to-one correlation between continuum and reflection line is distorted by General and Special relativistic effects. Several papers present extensive theoretical computations to describe the inner disc reverberation to the continuum taking into account all relativistic effects (Reynolds et al. 1999; Fabian et al. 2000; Reynolds \\& Nowak 2003). Several techniques have been employed to measure the variability-reverberation of the relativistic \\FeKa\\ line. However, for the best cases such as MCG-6-30-15, the relativistic Fe line showed a complex behaviour, having a variable intensity at low fluxes (Ponti et al. 2004; Reynolds et al. 2004) while showing a constant intensity at higher fluxes (Vaughan et al. 2003; 2004; see also the case of NG4051: Ponti et al. 2006). This puzzling and unexpected behaviour has been interpreted as due to strong light bending effects by some authors (Miniutti et al. 2003; 2004) or, alternatively, as the evidence that the broad wing of the \\FeKa\\ line is produced by strong and complex absorption effects (Miller et al. 2008). Thanks to the application of \\FeKa\\ excess emission maps (Iwasawa et al. 2005; Dovciak et al. 2004; De Marco et al. 2009), it has been possible to track weaker coherent patterns of \\FeKa\\ variations. In a few sources \\FeKa\\ variations are consistent with being produced by orbiting spots at a few r$_g$ from the BH (Iwasawa et al. 2004; Turner et al. 2006; Petrucci et al. 2007; Tombesi et al. 2007). Future larger area telescopes are needed to finally assess if these features are present only sporadically during peculiar periods or if, instead, although weak, are always present and can be used to map the inner disc (see e.g. Vaughan et al. 2008; De Marco et al. 2009). A leap forward in X-ray reverberation studies occurred thanks to the application of pure timing techniques to the long \\xmm\\ observation of 1H0707-495 that allowed the discovery of a \"reverberation lag\" between the direct X-ray continuum and the soft excess, probably dominated by FeL line emission (Fabian et al. 2009; Zoghbi et al. 2010). Soon after similar delays were seen in a few other objects (Ponti et al. 2010; De Marco 2011; Emmanouloulos et al. 2011; Zoghbi \\& Fabian 2011; Turner et al. 2011). Recently, De Marco et al. (2012) showed that these lags are ubiquitous in AGN, that they scale with M$_{\\rm BH}$ and have amplitudes of the order of the light crossing time of a few r$_g$, thus suggesting a reverberation origin of the delay (but see also Miller et al. 2010). Another fundamental step forward will be to combine these timing techniques to detect reverberation lags in the Fe K band (see Zoghbi et al. 2012). Reverberation from distant material has the advantage that the intensity of the \\FeKa\\ line and the continuum are expected to follow a simple one-to-one correlation, however, the expected delays between the reflection component and the direct emission are usually too large for a typical X-ray exposure. In fact, reflection from the inner walls of a molecular torus is expected to be delayed by a few years up to several decades and thus requires a very long monitoring campaign. Reflection from the BLR and/or outer disc is more accessible, the delay between continuum and reflection is expected to be between a few days up to few months. Thus a properly tailored monitoring campaign on a bright AGN with \\xmm, \\chandra\\ or \\suzaku\\ could achieve this goal. Several attempts have been made (Markowitz et al. 2003; Yaqoob et al. 2005; Liu et al. 2010). However, the 15-20 \\% or larger error on the flux of the \\FeKa\\ line and the low-sampling frequency of the X-ray observations have made the application of reverberation of the \\FeKa\\ line on weeks-months timescales basically impossible, until now. Mrk~509 (z=0.034397) is one of the brightest Seyfert 1 galaxies of the (2--100 keV) X-ray sky (Malizia et al. 1999; Revnivtsev et al. 2004; Sazonov et al. 2007), thus it has been observed by all major X-ray/Gamma-ray satellites. The \\chandra\\ HETG spectrum shows a narrow component of the Fe K line with an Equivalent Width (EW) of 50 eV (Yaqoob et al. 2004). \\xmm\\ and \\suzaku\\ data provide evidence for a second broader ($\\sigma=0.12$ keV) neutral Fe K line (Ponti et al. 2009) as well as a weak ionized emission feature between 6.7--6.9 keV (Pounds et al. 2001; Page et al. 2003; Ponti et al. 2009). The ionised emission can be fit either using a relativistically broadened ionised line or an outflowing photo-ionised gas component. Imprinted on the Fe K band emission of Mrk 509 are the fingerprints of two kinds of ionised absorption components, one marginally consistent with a medium velocity outflow (v$\\sim14000$ km s$^{-1}$; Ponti et al. 2009) and the others out(in)flowing with relativistic velocities (Cappi et al. 2009; Dadina et al. 2005; Tombesi et al. 2010). Here, we present the spectral and variability analysis of the Fe K complex energy band of Mrk~509 using the set of 10 \\xmm\\ observations (60 ks each), about one every fours days and spanning more than 1 month, which we obtained in 2009 (see the 3-10 keV light curve in Fig. \\ref{lc}). We also re-analyse the previous 5 \\xmm\\ observations. Thanks to this extensive monitoring campaign we can measure correlated variations between the Fe K line intensity and X-ray continuum flux, allowing us, for the first time, to perform a reverberation mapping study on this X-ray emission line. In addition we can study the presence of highly ionised matter from the innermost regions around the BH. The paper is organised as follows. \\S 2 is devoted to the description of the observations and data reduction. In \\S 3 a first parametrisation (with a single Gaussian profile for the Fe K$\\alpha$ line) of the total summed spectrum of the 2009 campaign is presented. Section 4 is dedicated to the detailed study of the \\FeKa\\ emission. We first present the study of the \\FeKa\\ line variability, assuming a single Gaussian profile (\\S 4.1) we then use the \\chandra\\ HETG data (\\S 4.2) to decompose the Fe K$\\alpha$ line in two Gaussian (narrow and resolved) components. \\S 4.3 presents the correlation between \\FeKa\\ intensity and the 3-10 keV continuum (once the Fe K$\\alpha$ line is fitted with 2 Gaussian lines) which is confirmed, in a model independent way, by the rms spectrum (\\S 4.4). In \\S 4.5 we discuss the possible origin of the \\FeKa\\ line. \\S 5 presents the study and the discussion of the origin of the ionised Fe K emission/absorption. Conclusions are in \\S 6. ", "conclusions": "We investigated the spectral variability of the Fe K band in the nearby, bright Seyfert 1 galaxy Mrk 509, using the 10 observations of the 2009 \\xmm\\ monitoring campaign as well as all the previous \\xmm\\ observations, totalling an exposure of more than 900 ks in about 10 years, resulting in one of the best quality Fe K spectra ever taken of a Seyfert 1 galaxy. This allows us, for the first time, to perform reverberation mapping of the resolved Fe K$\\alpha$ line. Figure \\ref{sketch} summarise in a sketch a possible scenario for the production of the Fe K emission in Mrk 509. The width of the narrow core of the Fe K$\\alpha$ line suggests an origin from distant material, possibly the inner wall of the molecular torus located at 0.2-few pc. The correlated variations (on a few days time-scales) between the 3-10 keV continuum and the intensity of the resolved component of the Fe K$\\alpha$ suggest an origin between several tens and few thousands of r$_{\\rm g}$ from the BH. The resolved Fe K$\\alpha$ emission can be produced in the disc, but we favour an origin at the base of a stratified broad line region. We note that none of the X-ray or UV absorption components with measured location is co-spatial with the resolved Fe K$\\alpha$ emitting region. Moreover, the properties of the X-ray and UV absorbers appear to differ from the ones required to produce the resolved Fe K$\\alpha$ line, suggesting that this emitting material is outside the line of sight, possibly in the form of an equatorial disc flattened wind such as observed in stellar mass black holes in the soft state (Ponti et. al. 2012) and neutron stars (Diaz-Trigo et al. 2006). The ionised Fe K emission might be produced either by photo-ionisation from distant material, such as the narrow line region and/or the ionised skin of the torus, or in the ionised inner accretion disc. \\begin{figure} \\begin{center} \\includegraphics[width=0.42\\textwidth,height=0.52\\textwidth,angle=-90]{SketchFeKalpha.ps} \\end{center} \\caption{Sketch of a possible locations for the different regions producing Fe K emission (diagram not to scale). The star represents the primary X-ray source, located close to the BH.} \\label{sketch} \\end{figure} The results of this study show that: \\begin{itemize} \\item{} The \\xmm\\ spectrum of Mrk 509 shows an evident \\FeKa\\ line with total EW$=58\\pm4$ eV. Fitted with a single Gaussian line the width is $\\sigma=0.092\\pm0.012$ keV. The line intensity increases with the 3-10 keV flux, but not as strongly as expected in a constant EW scenario, suggesting the presence of a constant and a variable \\FeKa\\ line component. \\item{} The \\chandra\\ HETG spectrum has enough energy resolution to resolve the narrow component of the \\FeKa\\ line ($\\sigma=0.027^{+0.018}_{-0.010}$ keV; line intensity ($1.5\\pm0.2$)$\\times10^{-5}$ ph cm$^{-2}$ s$^{-1}$). The width of the narrow component of the line suggests an origin at around 0.2-0.5 pc ($\\sim30000$ r$_{\\rm g}$) from the BH. This value is of the same order of magnitude as the molecular sublimation radius, suggesting that the narrow component of the \\FeKa\\ line might be produced as reflection from the inner walls of the molecular torus. If so, because of light travelling effects, the intensity of this component has to be constant on years time-scales. We assume the presence of a constant narrow \\FeKa\\ line (as observed by \\chandra\\ HETG) and add a second, resolved (now observed to be broader $\\sigma=0.22\\pm0.04$ keV), \\FeKa\\ component with EW$=42^{+9}_{-4}$ eV. There is excess emission at 7.06 keV, consistent with being produced (at least in part) by the associated Fe K$\\beta$ emission. \\item{} For the first time reverberation mapping of the resolved component of the \\FeKa\\ line on timescales of several days-years was successfully performed. The intensity of the resolved \\FeKa\\ component shows a significant ($\\sim4 \\sigma$) 1-to-1 correlation with the 3-10 keV flux variability, however the EW stays constant during the 9 years \\xmm\\ observed the source. The robustness of this result is confirmed by the results of the rms spectra which, in a model independent way, show an excess of variability at E$=6.45\\pm0.08$ keV. This excess of variability is consistent with being the resolved component of the \\FeKa\\ line ($\\sigma ~ \\lsimeq ~ 0.2$ keV) varying in such a way as to keep a constant EW$=71\\pm36$ eV. No measurable lag of the reflected component is observed. \\item{} The width of the resolved component of the \\FeKa\\ line suggests an origin between 300 and 1000 r$_{\\rm g}$ from the BH. This location is consistent with the observed \\FeKa\\ variability on days to a week timescale and the lack of measurable lag. The lack of a relativistic red wing of the \\FeKa\\ line suggests an inner radius for the line production larger than several tens of r$_{\\rm g}$ ($\\sim40$ r$_{\\rm g}$). \\item{} The EW$=42^{+9}_{-4}$ eV of the resolved \\FeKa\\ line suggests a larger covering factor of the primary X-ray sources (assumed to have altitudes of few r$_{\\rm g}$ above the BH) compared to the one expected from a flat disc annulus, indicating a possible azimuthal distribution above the disc of the reflecting material. A possibility is that the material producing the resolved \\FeKa\\ emission might be in the form of clouds, perhaps associated to the inner BLR (see Costantini et al. 2012). This geometry is further reinforced by the consistency between the \\FeKa\\ line width ($\\sigma=0.22$ keV) and those from the broadest components of the UV broad emission lines (Kriss et al. 2011). We also observe that the location of the reverberating \\FeKa\\ emission does not correspond to that of any X-ray or UV absorption components (Detmers et al. 2011; Kaastra et al. 2012; Kriss et al. 2011; 2012; Ebrero et al. 2011). \\item{} Significant, but weak (15-20 eV) ionised Fe K emission is observed. The ionised emission can be fit equally well with two narrow emission lines (from both \\Fevc\\ and \\Fevs), possibly from a photo-ionised or collisionally ionised gas, or by a single broad relativistic emission line (either \\Fevc\\ or \\Fevs). We note that the source high Eddington ratio suggests the presence of a standard thin $\\alpha$-disc down to a few r$_{\\rm g}$ from the BH. However, the neutral \\FeKa\\ line has no redshifted wing with no neutral emission closer than $\\sim40$ r$_{\\rm g}$ from the BH. This suggests that the surface of the inner accretion disc in Mrk 509 might be highly ionised. For these reasons, although the two interpretations for the origin of the ionised Fe K emission are equivalent on a statistical ground, we slightly prefer the latter interpretation on physical grounds. The picture of an higher ionised disc in the inner few tens of r$_{\\rm g}$ from the BH and less ionised outside is in line with the presence of a compact hard X-ray corona, providing there a high flux of hard X-ray photons, and a soft more extended one, as proposed by Petrucci et al. (2012). \\item{} A highly ionised, medium outflow velocity ($v\\sim0.048\\pm0.013$ c) Fe K absorption component detected in previous observations (EW$=-13^{+5.9}_{-2.9}$ eV at $\\sim~4~\\sigma$ significance) appears much weaker (EW$=-3.2^{+2.6}_{-2.8}$ eV, $\\sim4$ times weaker), if not absent, during the 2009 campaign. \\item{} Previous \\xmm\\ observations showed evidence for highly ionised high outflow velocity ($v\\sim0.05-0.2$ c) absorbers based on a total exposure of $\\sim300$ ks. We find no convincing ($>3 \\sigma$) evidence for these features during the 2009 \\xmm\\ long (600 ks) monitoring campaign. \\end{itemize}" }, "1207/1207.2005_arXiv.txt": { "abstract": "We present an estimation of the lower limits of local magnetic fields in quiescent, activated, and active (surges) promineces, based on reconstructed 3-dimensional (3D) trajectories of individual prominence knots. The 3D trajectories, velocities, tangential and centripetal accelerations of the knots were reconstructed using observational data collected with a single ground-based telescope equipped with a \\textit{Multi-channel Subtractive Double Pass} imaging spectrograph. Lower limits of magnetic fields channeling observed plasma flows were estimated under assumption of the equipartition principle. Assuming approximate electron densities of the plasma $n_e = 5\\times 10^{11}$ cm$^{-3}$ in surges and $n_e = 5\\times 10^{10}$ cm$^{-3}$ in quiescent/activated prominences, we found that the magnetic fields channeling two observed surges range from 16 to 40 Gauss, while in quiescent and activated prominences they were less than 10 Gauss. Our results are consistent with previous detections of weak local magnetic fields in the solar prominences. ", "introduction": "Solar quiescent prominences observed in the hydrogen H$\\alpha$ line ($\\lambda$=6562.8~\\AA{}) are long living and globally stable structures, having a typical density of the order of $10^{-13}$ g cm$^{-3}$ ($n_{e}$=10$^{10}$-10$^{11}$ cm$^{-3}$) and a typical temperature of the order of 10$^{4}$ K \\cite{W1993,TH1995,H2010}. Prominence plasma is coupled to magnetic fields, which channel plasma flows, carry plasma away and/or support it against solar gravity \\cite{AD2003,BC2010}. Magnetic fields insulate also relatively cold and dense prominence plasma from the surrounding hot and much more dilute coronal plasma. Although the overall structure of the quiescent prominence is relatively stable, numerous small-scale structures observed inside the prominences, like well outlined knots of plasma or extended flows, are locally highly dynamic and move along complicated, curved trajectories. Precise measurements of magnetic fields in solar prominences are crucial for understanding their structure, magnetic support, and evolution. Direct investigations of the magnetic fields in prominences started in the early '60s, using at first the Zeeman effect \\cite{ZS1961,R1967,TH2011}, subsequently replaced by spectropolarimetry of prominences applying Zeeman and Hanle effect observations \\cite{L1977,BSB1978}. The majority of the absolute field strength measurements in quiescent prominences revealed fields of less than 10 Gauss, but some stronger local fields up to 70 Gauss were also reported \\cite{Q1985,TH1995,P2001,AD2003,Cae2003,WB2003,M2006,Ch2010,H2010}. It is worth to stress, that Kuckein and co-workers reported recently 600-700 Gauss magnetic fields in a filament (\\citeauthor{K2009}, \\citeyear{K2009}; see also \\citeauthor{J2010}, \\citeyear{J2010}), ascribing low values of the earlier measurements to the lack of full Stokes polarimetry. Surges are active prominences that look like long straight or curved columns of plasma shot out from the solar surface with velocities up to a few hundred km~s$^{-1}$ \\cite{G1982}, their typical density could be estimated as roughly $5 \\times 10^{-12}$ g cm$^{-3}$ ($n_{e}$=10$^{11}$-10$^{12}$ cm$^{-3}$) \\cite{Bru1977,TH1995}. The magnetic fields in active prominences (like surges or sprays) are less known than magnetic fields in quiescent ones \\cite{TH2011}. However, in the case of surges there exist some estimations of magnetic field strengths, ranging from a few up to a few tens of Gauss \\cite {L2004,Bea2007,SAN2008}. Direct observations and modeling of the magnetic fields in solar prominences are still difficult while the resolving power of the existing relevant observing facilities is substantially lower than the diameters of prominence subtle structures (\\textit{i.e.} threads). Thus, each opportunity for independent check of the obtained results should be explored. In our previous paper \\cite{Zapior2010} we presented in detail a method of restoration of the true 3D trajectories of prominence knots using ground-based observations taken with a single telescope equipped with a \\textit{Multi-channel Subtractive Double Pass} (MSDP) imaging spectrograph \\cite{Mein77}. We have shown that the method allows an evaluation of the true 3D trajectories of the prominence knots without any assumptions concerning the shape of trajectories or dynamics of the motion. In the present work we exploit the method for the determination of accelerations acting on observed plasma knots and for the estimation of lower limits of the magnetic fields channeling the observed plasma flows under the assumption that magnetic field configurations fully control plasma trajectories. ", "conclusions": "We have presented 3D trajectories, velocities, tangential and centripetal accelerations of the plasma knots observed in solar prominences of various kinds (surges, activated, and quiescent prominences), restored using single ground-based telescope and MSDP spectrograph observations. The established parameters of the 3D (spatial) motions of the individual knots were applied to estimate lower limits of the magnetic fields channeling observed plasma flows under the assumption of the equipartition principle. The applied assumption is very strong, but it is consistent with contemporary prominence models assuming that local magnetic fields guide plasma flows and determine their geometry, being revealed by trajectories of the visible individual plasma knots or streams \\cite{AD2003}. We found also that the lower limits of the magnetic fields varied along the trajectories (\\textit{i.e.} along the magnetic ropes). However, while the diameter variations of the magnetic ropes channeling investigated knots are unknown, no unique values of the magnetic field were calculated for the whole trajectories, but instant values along the trajectories only. The tangential accelerations of the knots observed in active and quiescent prominences ranged roughly from -100 m s$^{-2}$ to 80 m s$^{-2}$, while in the case of two surges they range from -250 m s$^{-2}$ to -20 m s$^{-2}$. The centripetal accelerations of the knots observed in active and quiescent prominences ranged roughly from 0~m~s$^{-2}$ to 100 m s$^{-2}$, while for the two surges they ranged from 30 m s$^{-2}$ to $\\sim250$ m s$^{-2}$. The measured accelerations are roughly similar to previously presented results \\cite{KP1985,R1990,V1990,A2012}. Assuming approximate typical electron densities of the plasma $n_e = 5\\times 10^{11}$~cm$^{-3}$ in surges and $n_e = 5\\times 10^{10}$ cm$^{-3}$ in quiescent/activated prominences \\cite{Bru1977,Lab2010} we found that lower limits of the local magnetic fields in surges were equal to 16-40 Gauss, while in activated and quiescent prominences were less than 10 Gauss. In case if actual densities of the investigated prominences were smaller than assumed typical values applied in calculations, the obtained lower limits of the magnetic fields controlling prominences are slightly overestimated. Most of the direct measurements of the solar prominence magnetic fields, made using various methods also revealed rather weak local magnetic fields, ranging from a few up to a few tens of Gauss (see the discussion in the Introduction and references therein). However, as it was mentioned before, there were also some recent observations and numerical models reporting magnetic field strengths in the filaments of the order of 600-700 Gauss \\cite{K2009,J2010}. Our result is consistent with previous detections of weak local magnetic fields in solar prominences, but, establishing the lower limits of the magnetic fields for a limited number of prominence knots only, we cannot discard a possibility of strong fields in prominences. \\begin{acks} MZ is supported by Human Capital - European Social Fund. \\end{acks}" }, "1207/1207.3003_arXiv.txt": { "abstract": "{The {\\it Swift} era has posed a challenge to the standard blast-wave model of Gamma Ray Burst (GRB) afterglows. The key observational features expected within the model are rarely observed, such as the achromatic steepening (`jet-break') of the light curves. The observed afterglow light curves showcase additional complex features requiring modifications within the standard model . Here we present optical/{\\it NIR} observations, millimeter upper limits and comprehensive broadband modelling of the afterglow of the bright GRB 0505025A, detected by {\\it Swift}. This afterglow cannot be explained by the simplistic form of the standard blast-wave model. We attempt modelling the multi-wavelength light curves using (i) a forward-reverse shock model, (ii) a two-component outflow model and (iii) blast-wave model with a wind termination shock. The forward-reverse shock model cannot explain the evolution of the afterglow. The two component model is able to explain the average behaviour of the afterglow very well but cannot reproduce the fluctuations in the early X-ray light curve. The wind termination shock model reproduces the early light curves well but deviates from the global behaviour of the late-time afterglow. } ", "introduction": "\\label{intro} Gamma Ray Bursts (GRBs) are extremely energetic cosmic explosions which outshine the entire $\\gamma$-ray sky for a few seconds. The launch of {\\it Swift} (Gehrels et al. 2004), a dedicated satellite to detect GRBs and rapidly follow-up their afterglow emission, has revolutionized the study of the most energetic cosmic explosions in the Universe. In the standard blast-wave model for GRB afterglows (\\citealt{1992MNRAS.258P..41R,1993ApJ...418L...5P}; also see \\citealt{1999PhR...314..575P} for a review), a relativistic shock decelerates through uniform circumburst medium, heats up the matter, accelerates particles and enhances the magnetic field downstream. Synchrotron radiation from the shocked particles is observed as the afterglow. The snap-shot synchrotron spectrum can be characterized by four spectral parameters (apart from the electron index $p$): the injection frequency $\\nu_m$, the cooling frequency $\\nu_c$, the self-absorption frequency $\\nu_a$, and the peak flux $f_m$. The spectral parameters can be mapped to four physical parameters: the isotropic equivalent energy $E_{\\rm{iso}}$, the ambient medium density (parameterized as number density $n_0$ for a constant density medium and as $A_\\star$ for a wind driven medium for which $\\rho(r) = 5.5 \\times 10^{11} \\frac{A_{\\star}}{\\rm{g \\, cm}^{-1} } \\Lb \\frac{r}{\\rm {cm} } \\Rb^{-2}$ as in \\citealt{1999ApJ...520L..29C}), and the fractional energy content in non-thermal electrons and magnetic field ($\\epsilon_e$ and $\\epsilon_B$ respectively). Jet-break, a simultaneous steepening seen in the multi-frequency light curves considered as a signature of the collimated outflow from the burst \\citep{1999ApJ...525..737R}, if observed (at time $t_j$ since burst), gives a handle on the initial collimation angle ($\\theta_j$) of the explosion, and hence to the total kinetic energy involved ($E_{\\rm{tot}}$). The model was largely successful in explaining the pre-{\\it Swift} observations of GRB afterglows. However, {\\it{Swift}} with its ability to locate the afterglow within minutes of the burst and follow it up in {\\it UV}, optical and X-ray bands has revealed complexity in the early afterglow emission that is not predicted by the model. The X-ray light curves in the {\\it Swift} era have been rather dramatic, with steep decays, plateaus and flares, not witnessed earlier \\citep{2006ApJ...642..389N,2007ApJ...671.1903C}. Moreover, in several afterglows, the flux evolution did not follow the predicted spectral-temporal relations \\citep{2008ApJ...675..528L}. This has led to the conclusion that afterglow light curves differ drastically from burst to burst , owing to various physical processes shaping the flux evolution in various bands \\citep{2006ApJ...642..354Z}. Another open issue in the {\\it Swift} era is the absence of a jet-break. These complications often make it a demanding task to extract the physics of the burst and its surroundings from afterglow data. The bright low redshift ($z = 0.606$, \\citealt{2005GCN..3483....1F}) Gamma Ray Burst 050525A was detected by the {\\it Swift}-BAT on 2005 May 25 at 00:02:53 UT \\citep{2005GCN..3466....1B}. We refer to the burst trigger time as $t_0$. An isotropic equivalent $\\gamma$-ray energy of $2.3 \\times 10^{52}$~erg is inferred for the observed BAT fluence at a distance of $3.57$~Gpc (assuming $\\Omega_m = 0.3, \\Omega_\\Lambda = 0.7$ and $H_0 = 70$~km s$^{-1}$ Mpc$^{-1}$). The UVOT {\\it V}-band observations started at $T=t-t_0\\sim$65~s and XRT observations began $T\\sim 75$~s leading to well sampled early afterglow light curves. The proximity of the burst and the brightness of the afterglow made it a very good target for multi-wavelength observations. Ground based optical observations including robotic telescopes \\citep{2005A&A...439L..35K,2006ApJ...642L.103D}, radio observations in multiple frequencies by the Very Large Array \\citep{2005GCN..3495....1C} and Spitzer observations at $\\sim 2$~days in multiple {\\it IR}-bands \\citep{2008ApJ...681.1116H} have been reported in the literature. However, a detailed modelling involving the full evolution of the relativistic shock to infer the physical parameters has not been attempted for this burst. In this paper we present a new set of {\\it VRIJH} observations using eight different optical telescopes and millimeter upper limits from IRAM Plateau de Bure Interferometer (PdBI). We supplement our data with {\\it Swift}-UVOT and XRT data reported by \\cite{2006ApJ...637..901B}, the optical data reported by \\cite{2005A&A...439L..35K} and \\cite{2006ApJ...642L.103D} and the radio data from the VLA afterglow repository\\footnote{http://www.aoc.nrao.edu/$\\sim$dfrail/allgrb\\_table.shtml} to study the broad band evolution of the afterglow. We model the afterglow using three different extensions of the standard blast-wave model e.g. the forward-reverse shock model \\citep{1999MNRAS.306L..39M}, the two-component model \\citep{2003Natur.426..154B} and a model including a wind termination shock \\citep{2006ApJ...643.1036P}. Section \\ref{observations} gives a description of the data acquired from different telescopes and the analysis techniques. Multi-wavelength modelling under various premises are described in Section \\ref{modelling}. Section \\ref{conclusion} provides a summary of the multi-wavelength modelling results. ", "conclusions": "\\label{conclusion} We have presented the optical afterglow of GRB~050525A in {\\it VRIJH} photometric bands. Our data fill some gaps in the optical multi-wavelength light curves beyond $0.01$~days and provide a better constraint on the optical decay index. Our {\\it IR} observations, though confined to a narrow time bin, provide additional spectral constraints. The millimeter upper limits contributes to a better picture of the low frequency behaviour of the afterglow. We have undertaken a comprehensive multi-wavelength modelling of the afterglow, and tested various models against the data. The afterglow behaviour is too complex for a simple blast-wave model. We find that including emission from a possible reverse shock component is not sufficient to explain the afterglow evolution. Either the outflow is structured as a two-component jet or the ambient medium has a complex structure involving a variation in the density profile. Our two-component jet model is able to reproduce the overall behaviour of the afterglow, except the finer fluctuation in the early X-ray light curve. The wind termination shock model succeeds in explaining the early phase including the short time-scale features but deviates from the late-time data. The radio light curves are moderately well explained by both models. Afterglow modelling is necessary to unravel the nature of the outflow and the structure of the medium around the burst. The complexity of the observed light curves demands that the most realistic models are used. Dense sampling and long monitoring campaigns are also required in conjunction with afterglow modelling." }, "1207/1207.4848_arXiv.txt": { "abstract": "Space-based gravitational wave interferometers are sensitive to the galactic population of ultra-compact binaries. An important subset of the ultra-compact binary population are those stars that can be individually resolved by both gravitational wave interferometers and electromagnetic telescopes. The aim of this paper is to quantify the multi-messenger potential of space-based interferometers with arm-lengths between 1 and 5 Gm. The Fisher Information Matrix is used to estimate the number of binaries from a model of the Milky Way which are localized on the sky by the gravitational wave detector to within 1 and 10 deg$^2$ and bright enough to be detected by a magnitude limited survey. We find, depending on the choice of GW detector characteristics, limiting magnitude, and observing strategy, that up to several hundred gravitational wave sources could be detected in electromagnetic follow-up observations. ", "introduction": "A variety of detector concepts for space-based gravitational wave interferometers have been proposed, the most well studied concept being LISA\\citep{LISA}. It was understood early on that the most numerous source class radiating in the band covered by LISA-like detectors will be the galactic population of ultra-compact binaries (UCBs) comprised of pairs of stellar remnants: white dwarfs, neutron stars or black holes. The gravitational radiation from these UCBs will be the dominant signal in the frequency band covered by LISA-like detectors. Early estimates of the composite signal from the UCBs \\citep{Evans1987, HBW,HB1997} demonstrated that the signals of the vast majority of the galactic binaries will overlap and be unresolvable from one another, forming a limiting foreground (or ``confusion noise'') for space-based gravitational wave detectors. Later studies based on population synthesis \\citep{Nelemans, Benacquista, Edlund2005, TRC, RBBL} have borne this expectation out. Detailed data analysis studies have shown that $\\sim 10^{4}$ individual binaries could be resolved out of the foreground by a gravitational wave observatory like LISA \\citep{TRC, Crowder2007, Littenberg2011, Nissanke2012}. A subset of the resolvable binaries will be detectable electromagnetically. The purpose of this work is to assess the multi-messenger potential for different space-based detectors spanning the trade-space of future mission designs. This builds off previous work \\citep{Cooray2003, Nelemans2006, Nelemans2009} demonstrating the feasibility of follow-up observations for high-frequency UCB sources. We estimate the total number of multi-messenger sources by beginning with a population synthesis model of the galaxy \\citep{Nelemans2004}, complete with optical magnitudes. From this we produce a magnitude limited source catalog, then estimate how well each system will be localized on the sky by different gravitational wave detector configurations. Using hundreds of Monte Carlo realizations over the spatial distribution of the galaxy and the UCB orientations, we find tens to hundreds of sources that can be observed both electromagnetically and gravitationally. The information encoded about the UCBs in each of the two spectrums is highly complementary, enabling tests of general relativity, full measurement of the physical parameters enabling constraints on binary synthesis channels, and new methods of probing the close interaction dynamics of the compact stars \\citep{Cutler2003, Stroeer2005}. ", "conclusions": "\\label{sec.Discussion} We conclude that space-based gravitational wave detectors will be useful observatories for discovering new UCBs in the galaxy that could be observed electromagnetically, though deep, wide field, optical surveys may be required to produce large catalogs. We reach this verdict by considering a range of plausible near-future space-based gravitational wave detector concepts, and assess their measurement capabilities for magnitude limited catalogs of UCBs. Magnitudes for the constituents of each binary were derived from the population synthesis simulations, and the gravitational wave measurement capabilities were estimated using the Fisher Information Matrix. Any UCBs that were brighter than our chosen magnitude limits (18-24) and located on the sky by the gravitational wave detector to within angular resolution $d\\Omega$ were considered multi-messenger candidates. We estimated the multi-messenger catalog sizes for both $d\\Omega \\leq 1$ and 10 deg$^2$. At the pessimistic end, we consider magnitude 18 limited catalogs, and single-vertex interferometers with 1 Gm arm-lengths. The best scenario considered the classic LISA design and an optical telescope limited at $24^{\\rm th}$ magnitude. The number of multi-messenger candidates was anywhere from a few to several hundreds over that range of detector capabilities. If we put on the additional constraint that the sources must be eclipsing to allow for electromagnetic observation the counts were reduced by a factor of $\\sim 3$. While most of the known verification binaries are AM CVn type stars, our study only considered detached white dwarf binaries, thus providing a very complimentary catalog of UCB multi-messenger systems. This work considered a conservative approach to finding multi-messenger UCBs, with competing criteria that strongly affect the expected population of systems detectable in both spectrums. Electromagnetic detections are most strongly affected by the magnitude limit of the detection survey, a function of telescope aperture. By contrast, the gravitational wave detection catalogs of UCBs are expected to have thousands of systems in them; most will be too faint to be detectable by any electromagnetic survey. However the gravitational wave localization criterion is a strong constraint on the multi-messenger catalog. We find that wide-field surveys ($d\\Omega\\leq10$ deg$^2$) yield more candidates than more narrow fields of view ($d\\Omega\\leq1$ deg$^{2}$) by 50-100\\% for the full catalogs, and by a factor of 2-4 for the eclipsing binaries. We have estimated the number of UCB multi-messenger candidates without considering what could be done with joint GW and EM observations. Our follow-on effort will consider the science yield from joint observations of both the known verification binaries -- mostly mass-transferring systems -- and the close, detached binaries that will be discovered by space-borne gravitational wave detectors." }, "1207/1207.1695_arXiv.txt": { "abstract": "We present preliminary reconstructions of the EUV from 26 to 34\\,nm from February 1997 to May 2005, covering most of solar cycle 23. The reconstruction is based on synthetic EUV spectra calculated with the spectral synthesis code Solar Modeling in 3D (SolMod3D). These spectra are weighted by the relative area coverage of the coronal features as identified from EIT images. The calculations are based on one-dimensional atmospheric structures that represent a temporal and spatial mean of the chromosphere, transition region, and corona. The employed segmentation analysis considers coronal holes, the quiet corona, and active regions identified on the solar disk. The reconstructed EUV irradiance shows a good agreement with observations taken with the CELIAS/SEM instrument onboard SOHO. Further improvement of the reconstruction including more solar features as well as the off-limb detection of activity features will be addressed in the near future. ", "introduction": "The solar EUV irradiance is the main energy input for the upper Earth's atmosphere with important effects on the ionosphere and thermosphere. The solar energy output changes on short time-scales of minutes to hours as well as longer times-scales such as the 27-day solar rotation cycle or the 11-year solar cycle. There is also indication that the EUV irradiance might shows a secular trend (\\cite[{Didkovsky} {et~al.} 2010]{Didkovsky2010ASPC}). The EUV radiation incident on the upper Earth's atmosphere leads to a change in its temperature and density (see e.g. \\cite[{Solomon} {et~al.} 2010]{Solomon2010}). In order to understand the effects of the changing EUV radiation on the Earth's atmosphere a continuous data set covering the short-term and long-term variations is essential. However, as space instruments are limited with regard to their temporal and spectral coverage, reliable models are needed to fill the gaps of the observational data sets. {Several reconstruction approaches involve the use of proxies to describe the EUV variability. \\cite[{Lean} {et~al.} (2011)]{Lean2011JGR} employ the two and three component NRLSSI model based on the Mg\\,II and F$_{10.7}$ index to characterize and forecast the EUV variations. A further example for an empirical model is SOLAR2000 (\\cite[{Tobiska} {et~al.} 2000]{Tobiska2000}), a model based on an extensive number of irradiance proxies. There is also ongoing work to determine which proxies, or spectral lines, are the best representatives for the variations of the entire EUV spectrum (see e.g. \\cite[{Kretzschmar} {et~al.} 2009]{Kretzschmar2009AcGeo}, \\cite[{Dudok de Wit} {et~al.} 2009]{DudokdeWit2009GRL}). Proxy models have been quite successful in describing the EUV variations, however, in order to understand the complete physical processes driving the irradiance variations, it is important to model the variations of the full solar spectra. \\cite[Warren (2006)]{Warren2006} utilizes differential emission measure distributions derived from spatially and spectrally resolved solar observations and full-disk solar images. The reconstruction presented here follows the same principle as used by \\cite[{Haberreiter} {et~al.} (2005)]{Haberreiter2005ASpR} and \\cite[{Shapiro} {et~al.} (2011)]{Shapiro2011AA}. It includes the calculation of synthetic spectra with a radiative transfer code for various activity features on the solar disk. Weighting the spectra by filling factors derived from the relative area coverage of these activity features or from proxy data then yields the time dependent irradiance spectrum. In the following section the spectral synthesis code Solar Modeling in 3D (SolMod3D) is introduced. Then, in Section\\,\\ref{sec:rec} the reconstruction approach is described briefly. Finally, in Section\\,\\ref{sec:results} our results are compared with the EUV irradiance observations carried out with the CELIAS/SEM instrument (\\cite[{Hovestadt} {et~al.} 1995]{SEM}) onboard the SOHO mission. ", "conclusions": "We have presented a preliminary version of the reconstruction of the EUV variations from Feb 14, 1997 to May 1, 2005 for the wavelength range between 26 and 34\\,nm and compared it with observations taken with the CELIAS/SEM instrument. This work is based on the calculation of spectra in the EUV with the SolMod3D code and the filling factors for three coronal components. The good agreement of the reconstruction with the observed irradiance variations shows that our approach is suitable for the study of the EUV variability. Nevertheless, further analysis is required, in particular the study of the off-limb contribution to the EUV irradiance and the consideration of additional coronal activity features." }, "1207/1207.2249_arXiv.txt": { "abstract": "We present the results of dust scattering simulations carried out for the Orion Eridanus Superbubble region by comparing them with observations made in the far-ultraviolet. The albedo and the phase function asymmetry factor (\\textit{g}-factor) of interstellar grains were estimated, as were the distance and thickness of the dust layers. The results are as follows: 0.43$^{+0.02}_{-0.04}$ for the albedo and 0.45$^{+0.2}_{-0.2}$ for the \\textit{g}-factor, in good agreement with previous determinations and theoretical predictions. The distance of the assumed single dust layer, modeled for the Orion Molecular Cloud Complex, was estimated to be $\\sim$110 pc and the thickness ranged from $\\sim$130 at the core to $\\sim$50 pc at the boundary for the region of the present interest, implying that the dust cloud is located in front of the superbubble. The simulation result also indicates that a thin ($\\sim$10 pc) dust shell surrounds the inner X-ray cavities of hot gas at a distance of $\\sim$70-90 pc. ", "introduction": "As stellar ultraviolet (UV) radiations scattered by interstellar dust grains are regarded as the most dominant source of the diffuse Galactic light \\citep{seo11}, the optical properties of the dust grains, generally characterized by the albedo and the phase function asymmetry factor \\textit{g} $\\equiv$ $<\\cos\\theta>$, can be inferred from the observations of the diffuse Galactic light. The observations can also be compared with model calculations. For example, \\citet{dra03} estimated the albedo and \\textit{g}-factor for the carbonaceous-silicate grains, and obtained $\\sim$0.40 and $\\sim$0.65 for the albedo and the \\textit{g}-factor at $\\lambda$ $\\sim$ 1600 {\\AA}, respectively. While the values of the albedo and g-factor are clearly dependent on the models of dust grains, accurate photometry and modeling of dust clouds are also required for reliable estimations of these optical properties from observations. Until now, there have been numerous studies in this line of research from the early 70's, especially at far-ultraviolet (FUV) and near-ultraviolet (NUV) wavelengths \\citep{wit73,wit78,lil76,mor78,mor82,mor80,ona84,wit92,wit97,gor94,cal95}. For example, the observations made from the \\textit{Voyager 2} spacecraft have been used for the wavelengths shorter than 1200 {\\AA} \\citep{wit93,mur93,bur02,sha04,sha06,suj05,suj07}. A comparative review of past work can be found in \\citet{gor04}. Recently, \\textit{GALEX} observations provided diffuse UV background images with high spatial resolution, which have been used for the determination of the albedo and the \\textit{g}-factor in both the FUV (1350-1750 {\\AA}) and NUV (1750-2850 {\\AA}) bands \\citep{suj09,suj10}. Most of these studies showed that the interstellar grain causes a strong forward scattering in the FUV, with a moderate albedo, which is in good agreement with theoretical predictions. The Orion-Eridanus Superbubble (henceforth OES) is a large hot bubble with a diameter of $\\sim$30$^\\circ$ across the sky located in the Orion and Eridanus constellations. Its distance is known to be $\\sim$155 pc to the near-side and $\\sim$586 pc to the far-side of the bubble, and its near edge is considered to interact with the expanding shell of the Local Bubble \\citep{bur96}. The OES has been observed extensively in various wavelengths, ranging from X-ray to radio. Enhanced soft X-ray emission seen at its inner cavities is the evidence of the hot gas (T $\\sim$ 10$^6$ K) occupying its inner parts \\citep{hei99}. Two prominent filamentary H$\\alpha$ shells, known as arcs A and B, have been observed at the boundaries of these hot cavities \\citep{rey98,bou01}. The presence of 21 cm radiation of neutral hydrogen (T $\\sim$ 10$^{2-3}$ K) was also noted outside the hot bubble (Brown et al. 1995). The dust extinction (\\textit{E}(\\textit{B-V})) map, made from the Schlegel, Finkbeiner and Davis (henceforth SFD) Dust Survey \\citep{sch98}, shows a strong anti-correlation with the soft X-ray emission image \\citep{jo11}. Recent spectroscopic observations made in the FUV revealed the existence of \\ion{C}{4} and \\ion{Si}{2}* emission lines as well as molecular fluorescence H$_2$ lines, implying that the OES is truly a multi-phase object, consisting of hot gases and cold dust that interact with each other \\citep{kre06,ryu06,ryu08,jo11}. The dust properties of the OES region have not been explored well because of the lack of sufficient diffuse FUV observations. Based on the data available from the Voyager measurement, \\citet{mur93} suggested a lower limit of $\\sim$0.3 for the albedo assuming isotropic scattering, and an upper limit of $\\sim$0.8 for the phase function asymmetry factor with perfectly reflecting grains. However, it should be noted that the Voyager observations provided limited information, and only for two regions with a field of view of 0$^\\circ$.10 $\\times$ 0$^\\circ$.87 and a spectral band of 38 {\\AA}. In this paper, we report the results of our study on the dust scattering properties of the OES region. We performed Monte Carlo simulations and compared the results with the diffuse emission map made from the recent FUV imaging spectrograph mission. We obtained the average albedo and phase function asymmetry factor, and estimated the distance and thickness of the dust cloud. \\begin{figure} \\begin{center} \\includegraphics[width=7.3cm]{f1a.eps}\\\\ \\includegraphics[width=7cm]{f1b.eps}\\\\ \\end{center} \\caption{(a) FUV continuum map overplotted with dust contours, in equatorial coordinates. The continuum intensity is given in CU (photons s$^{-1}$ cm$^{-2}$ sr$^{-1}$ {\\AA}$^{-1}$) and the levels of dust extinction contours are in \\textit{E}(\\textit{B-V}). (b) A scatter plot for the FUV intensity against dust extinction, obtained from pixel-by-pixel comparison of Figure \\ref{fig:fims}(a). \\label{fig:fims}} \\end{figure} ", "conclusions": "We have performed Monte Carlo simulations for the dust scattering of FUV emissions in the OES region and compared the results with the diffuse emission map made from the recent FUV imaging spectrograph mission, FIMS. We were able to obtain the optical parameters of interstellar dust grains for the OES region: the average albedo is 0.43$^{+0.02}_{-0.04}$ and the phase function asymmetry factor is 0.45$^{+0.2}_{-0.2}$, in agreement with previous observational and theoretical estimations. Furthermore, the simulation results indicate that the dust clouds are located in front of the OES, with its distance of $\\sim$110 pc and thickness of $\\sim$50-130 pc, while the hot X-ray cavity is bounded by a thin ($\\sim$10 pc) dust shell toward the Sun." }, "1207/1207.2555_arXiv.txt": { "abstract": "We study the radio--FIR correlation between the nonthermal (synchrotron) radio continuum emission at $\\lambda90$ cm (333 MHz) and the far infrared emission due to cool ($\\sim20$ K) dust at $\\lambda70~\\mu$m in spatially resolved normal galaxies at scales of $\\sim$1 kpc. The slope of the radio--FIR correlation significantly differs between the arm and interarm regions. However, this change is not evident at a lower wavelength of $\\lambda20$ cm (1.4 GHz). We find the slope of the correlation in the arm to be $0.8\\pm0.12$ and we use this to determine the coupling between equipartition magnetic field ($B_{\\rm eq}$) and gas density ($\\rho_{\\rm gas}$) as $B_{\\rm eq} \\propto \\rho_{\\rm gas}^{0.51 \\pm 0.12}$. This is close to what is predicted by MHD simulations of turbulent ISM, provided the same region produces both the radio and far infrared emission. We argue that at 1 kpc scales this condition is satisfied for radio emission at 1.4 GHz and may not be satisfied at 333 MHz. Change of slope observed in the interarm region could be caused by propagation of low energy ($\\sim$ 1.5 GeV) and long lived ($\\sim10^8$ yr) cosmic ray electrons at 333 MHz. ", "introduction": "\\label{intro} The radio--far infrared (FIR) correlation in normal galaxies was first observed by \\cite{kruit71, kruit73} and later extended by the IRAS mission. Subsequently it was established that the correlation holds good (within a factor of 2) over five orders of magnitude in radio and FIR luminosity \\citep{condo92,yun01} for a wide morphological class of galaxies like, spirals, irregulars and dwarfs \\citep{wunde87,dress88,price92} on global scales. Based on spatially resolved studies of normal and irregular galaxies it is seen that the correlation holds even at scales of few tens to hundreds of parsecs \\citep[see e.g][]{beck88, xu92, hoern98, murgi05, tabat07a, hughe06, murph06a, palad06, palad09, dumas11}. The basic model that connects these two regimes of emission is via star formation \\citep{harwi75}. The radio continuum emission arises due to synchrotron emission (henceforth nonthermal emission) from relativistic electrons, produced in supernova remnants. A good fraction of them originate from massive ($\\gtrsim 10~ \\rm M_\\odot$), short lived ($\\lesssim10^6$ yr) stars. The FIR emission arises from re-radiation by dust heated due to ultra violet (UV) photons emitted by the above population of stars. Though the cause of the correlation is well understood, the tightness over several orders of magnitude still remains puzzling. Many models explaining the correlation require close coupling between the magnetic field ($B$) and the gas density ($\\rho_{\\rm gas}$) of the form, $B\\propto\\rho_{\\rm gas}^{\\kappa}$ \\citep[see e.g.,] []{helou93, nikla97, thomp06}. Such a coupling can be established by magnetohydrodynamic (MHD) turbulence of the interstellar medium (ISM) \\citep[see][]{chand53, cho00, cho03, grove03}. Numerical simulations by \\citet{cho00} revealed that $\\kappa = 0.5$ is a manifestation of the equipartition condition, i.e, in steady MHD turbulence the magnetic field energy density and the energy density of the gas are similar. Similar values of $\\kappa$ have been found through observations of magnetic field by Zeeman splitting observations in molecular clouds by \\cite{crutc99}, also by using equipartition magnetic field and molecular gas observations in external galaxies by \\citet{nikla97} and in Milky Way and M31 by \\citet{berkh97}. Alternatively, the slope of the radio--FIR correlation has been used to find $\\kappa$, where $\\kappa\\sim$ 0.4--0.6 \\citep{nikla97, hoern98, dumas11}. \\begin{table*} \\begin{centering} \\caption{The sample galaxies.} \\begin{tabular}{@{}lcccccccc@{}} \\hline Name & Morphological &Angular & $i$ & Distance & FIR & &Radio~~~~~~~~~~~~~~~~~~~~~~~~~~~~~ \\\\ &type & size (D$_{25}$)($^\\prime$)& ($^\\circ$) & (Mpc) & $\\lambda70\\mu$m &$\\lambda90$cm&$\\lambda20$cm\\\\ (1) &(2) & (3) & (4) & (5) & (6) & (7) &(8)\\\\ \\hline NGC 4736 & SAab & 11.2$\\times$9.1 & 41 & 4.66$^1$ &SINGS&GMRT& Westerbork\\footnote{The Westerbork Synthesis Radio Telescope (WSRT) is operated by the Netherlands Foundation for Research in Astronomy (NFRA) with financial support from the Netherlands Organization for scientific research (NWO).} SINGS (1374.5 MHz)$^4$ \\\\ NGC 5055 & SAbc & 12.6$\\times$7.2 & 59 & 9.2$^\\dagger$&SINGS&GMRT& Westerbork SINGS (1696 MHz)$^4$ \\\\ NGC 5236 & SABc & 11.2$\\times$11 & 24 & 4.51$^2$ &SINGS&GMRT& VLA\\footnote{The Very Large Array (VLA) is operated by the NRAO. The NRAO is a facility of the National Science Foundation operated under cooperative agreement by Associated Universities, Inc.} CD array (1452 MHz)$^5$ \\\\ NGC 6946 & SABcd & 11.5$\\times$9.8 & 33 & 6.8$^3$ &SINGS&GMRT& VLA C$+$D array (1465 MHz)$^6$ \\\\ \\hline \\end{tabular}\\\\ In column (3) D$_{25}$ refers to the optical diameter measured at the 25 magnitude arcsec$^{-2}$ contour from \\cite{vauco91}. Column (4) gives the inclination angle ($i$) defined such that $0^\\circ$ is face-on. Distances in column (5) are taken from: $^1$ \\cite{karac03}, $^2$ \\cite{karac02}, $^3$ \\cite{karac00} and the NED $^\\dagger$. Column (6) and (7) are the sources of data for the FIR maps and 333 MHz ($\\lambda90$cm) maps respectively. Column 8 are the data available at a higher frequency near 1 GHz ($\\lambda20$cm): $^4$ \\cite{braun07}, $^5$ VLA archival data using the CD array configuration (project code : AS325), $^6$ VLA archival map by combining interferometric data from C and D array, \\cite{beck07}. \\end{centering} \\label{tablesample} \\end{table*} So far, spatially resolved and global study of the correlation has been done primarily using radio emission at 1.4 GHz and higher frequencies. The only low frequency study done at 150 MHz \\citep{cox88}, confirms that on global scales the radio--FIR correlation holds good and is similar to what is seen at 1.4 GHz. To our knowledge, no low frequency ($<$ 1.4 GHz, such as 333 MHz) spatially resolved study of the radio--FIR correlation exists in the literature. The motivation to do such a study arises from the fact that at lower frequencies the emission is largely nonthermal, hence better exhibiting the relation between magnetic field and star formation. Secondly, since the cosmic ray electrons (CRe) propagate larger distances in the galaxies at lower frequencies, it is important to assess how that affects the form of the radio--FIR correlation. In this paper, we present spatially resolved study of the radio--FIR correlation for four normal galaxies, NGC 4736, NGC 5055, NGC 5236 and NGC 6946 at spatial resolution of $\\sim$1--1.5 kpc with radio observations made at 333 MHz ($\\lambda90$ cm) and 1.4 GHz ($\\lambda20$ cm). We also estimate the value of $\\kappa$ and verify the equipartition assumptions. In Section~\\ref{data} we discuss the various sources of maps used in this work and also define the parameter `$q$' which is used to quantify the correlation. In Section~\\ref{results} we present our results on spatially resolved radio--FIR correlation using far infrared emission at $\\lambda70~\\mu$m and radio emission at $\\lambda20$ cm and $\\lambda90$ cm. We discuss our results in Section~\\ref{discussion}. ", "conclusions": "\\label{discussion} We have studied the radio--FIR correlation at $\\sim$1 kpc scales for four normal galaxies using nonthermal radio maps at $\\lambda90$ cm and $\\lambda20$ cm and the far infrared maps at $\\lambda70~\\mu$m. From the basic synchrotron theory (e.g., \\citealt{moffe75}) and considering the radio emission from CRe emitting at critical frequencies, the energy of CRe at $\\lambda90$ cm is $\\sim$1.5 GeV and at $\\lambda20$ cm is $\\sim$3 GeV when they are gyrating in a typical magnetic field of $\\sim10 ~\\mu$G. The far infrared emission at $\\lambda70~\\mu$m originates from cool dust at $\\sim$20 K heated by the interstellar radiation field (ISRF) due to $\\rm \\sim5-20~M_\\odot$ stars \\citep{dever89, xu90, xu96, dumas11}. We separately examine these correlations for the arm and the interarm regions, that is, regions of high and low thermal fractions respectively. The results of the various parameters as discussed in Section~\\ref{results} are given in Table 3 for individual galaxies, and here we discuss the average properties. The dispersion on the parameter $q_\\lambda$ is a measure of the tightness of the radio--FIR correlation, which for the arm region is found to be less than 10 percent around the mean $q_\\lambda$ for both $\\lambda20$ cm and $\\lambda90$ cm. For the interarm region the dispersion is seen to increase to around 20 percent for both the frequencies. Further we find the slope of the radio--FIR correlation for the arm regions (also the high thermal fraction regions) remains similar at both the radio frequencies (see Table 3). It should be noted that a large number of global scale radio--FIR correlation studies exist, where the observed slope is steeper and closer to unity \\citep[see e.g.,][and the references therein]{price92, yun01}. However, the spatially resolved studies relating FIR cool dust emission to $\\lambda$20 cm radio emission, yields a value of the slope $\\sim0.6-0.9$ for LMC \\citep[]{hughe06} and $0.80\\pm0.09$ for M31 \\citep[]{hoern98}. It is difficult to compare the slopes obtained in global studies with the spatially resolved case. The flux in global studies are averaged over both arm and interarm regions and we are uncertain about the contribution from each component. Multifrequency spatially resolved studies can provide an understanding of the relation between global scale and spatially resolved studies. For the present case, in the interarm regions (regions of low thermal fraction) for $\\lambda20$ cm the slope is slightly flatter as compared to the arms (see Eq. 1 and 2). However, at $\\lambda90$ cm, the slopes become distinctly flatter than the arm regions (see Fig.~\\ref{radio-fir} and Eq. 3 and 4). Our results can be used to determine the coupling between magnetic field ($B$) and the gas density ($\\rho_{\\rm gas}$) as discussed in the introduction and thereby validating the `equipartition' assumptions in these galaxies at 1 kpc scales. \\citet{dumas11} showed that the slope of the radio--FIR correlation relates to $\\kappa$ as, \\begin{numcases}{\\kappa =} \\frac{n~b}{3-\\ant}, & optically thick dust \\\\ \\frac{(n+1)~b}{3-\\ant}, & optically thin dust \\end{numcases} where, $n = 1.4$$\\pm0.15$ is the Kennicutt-Schmidt law index \\citep[see e.g.,][]{kenni98}, $b$ is the slope of the radio--FIR correlation and $\\ant$ is the nonthermal spectral index. For these face-on galaxies we use the assumption of optically thin dust to UV photons to estimate $\\kappa$. We find that $\\kappa = 0.51\\pm0.1$ at $\\lambda20$ cm and $\\kappa = 0.4\\pm0.1$ at $\\lambda90$ cm. Similarly, for interarm regions due to a large range of $\\ant$ we find $\\kappa$ in the range 0.41 -- 0.5 at $\\lambda20$ cm and between 0.18 -- 0.22 at $\\lambda90$ cm. Our estimated values of $\\kappa$, using the correlation between $\\lambda20$ cm and $\\lambda70~\\mu$m, are consistent with the predictions of numerical MHD simulations of different ISM tubulence models, where $\\kappa \\sim 0.4-0.6$ \\citep[see e.g.,][]{fiedl93,kim01, thomp06, grove03}. In the arm regions, the slope and thus $\\kappa$ remains similar for both $\\lambda20$ cm and $\\lambda90$ cm. Note that the above prescription to determine $\\kappa$ is valid provided the radio and the FIR emission arises from the same emitting volume, with a diameter of about 1 kpc for most of the observations reported here. In the arm regions the UV photon has a mean free path of $\\sim$100 pc within which most of the FIR emission arises. On the other hand, the CRe which gives rise to the radio emission diffuse farther away to $\\sim1$ kpc at 1400 MHz and $\\sim2$ kpc at 333 MHz in a galactic magnetic field of $\\sim10~\\mu$G. Hence in order to have a similar slope with frequency, the energy spectrum of the CRe giving rise to the radio emission should be independent of the volume element. This can only happen if the timescale for CRe diffusion/propagation ($\\tau_{\\rm diff}$) is significantly larger than their generation timescale ($\\tau_{\\rm gen}$). It turns out that the $\\tau_{\\rm diff}$ is about $8\\times10^7$ yr at 333 MHz and $4\\times10^7$ years at 1400 MHz which is significantly larger than the $\\tau_{\\rm gen}$ as evident from the supernova rates, which is one every $10^4 - 10^5$ yr kpc$^{-2}$ in Milky Way. We assume the same rate for these galaxies. \\begin{figure} \\begin{center} \\begin{tabular}{c} \\includegraphics[width=7cm]{ngc6946_diffuse.ps} \\end{tabular} \\caption{The H$\\alpha$ image of the galaxy NGC 6946 (KPNO 2-m telescope, filter: KP1563) obtained from the ancillary data at SINGS website. The solid circles represents the diffusion scale of $\\sim$1 kpc for $\\sim3$ GeV CRe at $\\lambda20$ cm, while the dashed circles represents the diffusion scales of $\\sim$2 kpc for $\\sim$1.5 GeV CRe at $\\lambda90$ cm. See text for the details.} \\label{spiral} \\end{center} \\end{figure} The slope of the radio--FIR correlation in the interarm (low thermal fraction) region is similar to that of the arm at $\\lambda20$ cm, however it becomes distinctly flatter at $\\lambda90$ cm. The flattening primarily happens due to relative increase in radio flux at $\\lambda90$ cm as compared to $\\lambda20$ cm, which has the effect that $\\ant$ gradually becomes steeper in the interarms. This relative increase in $\\lambda90$ cm flux can be explained by continuous generation of CRe in the arm, which subsequently propagates into the interarm (e.g., from A to B or from farther regions in arms like C to B in Fig.~\\ref{spiral}). The propagation timescale for these CRe are few times $10^7$ years assuming Alfv$\\acute{\\rm e}$n velocity of 100 km s$^{-1}$ and typical arm to interarm distance of 1--2 kpc. In such a scenario, using Equation 6 of \\citep{karda62}, in a typical galactic magnetic field of $\\sim10~\\mu$G, there would be a break in the energy spectrum for electrons above $\\sim$2 GeV. This break frequency lies below $\\lambda90$ cm or above 333 MHz. Such breaks have been seen at $\\sim900$ MHz and $\\sim$1 GHz for similar normal galaxies, NGC 3627 and NGC 7331 respectively \\citep{palad09}. Thus the CRe emitting at $\\lambda90$ cm, which lie above the break, do not loose significant amount of energy as compared to their higher energy counterparts. Hence, this results in increasing the relative flux at $\\lambda90$ cm. For the slope to remain similar between arms and interarm regions at $\\lambda20$ cm (below the break), the ratio of the radio to FIR flux densities should remain similar. Observed radio flux between arm and interarm changes by a factor of $\\sim$2--2.5. Similar ratio of flux density between arm and interarm regions at $\\lambda20$ cm can be caused due to steeping of the spectral index to $\\lesssim -1.1$ as compared $\\sim-0.6 ~\\rm to -0.8$ in the arms. This implies the FIR flux should change by a factor of $\\sim$2.5--3 between arm and interarm regions for radio--FIR slope of $\\sim$0.8. The FIR flux density ($F_\\lambda$) depends on the dust temperature ($T_{\\rm dust}$) and its density ($\\rho_{\\rm dust}$) as, $F_\\lambda \\propto \\rho_{\\rm dust} Q_{\\rm abs}(a, \\lambda) B_\\lambda(T_{\\rm dust})$, where $Q_{\\rm abs}(a, \\lambda)$ is the FIR wavelength ($\\lambda$) dependent absorption coefficient for gain radius, $a$ \\citep{drain84, alton04}. The temperature do not change significantly between arm and interarm for these galaxies \\citep{basu12}. For a constant gas-to-dust ratio, i.e, $\\rho_{\\rm dust} \\propto \\rho_{\\rm gas}$, a factor of 2--4 drop in average gas density between arm and interarm regions \\citep[found using the CO$_{\\rm J:2\\to1}$ maps from Heracles;][]{leroy09} would therefore cause the factor of 2--3 drop in FIR emission. The slope of 0.8$\\pm$0.1 of the radio--FIR correlation indicates that the energy equipartition assumption between cosmic ray particles and magnetic field may be valid in the gas rich arms of the galaxies at our spatial resolution of $\\sim$1 kpc. For the interarm regions at $\\lambda20$ cm the slope is similar to what is seen in arms, and thereby satisfying the equipartition conditions. The flattening of the slope at $\\lambda90$ cm does not indicate any break down of equipartition condition, but results due to overlapping emissions from adjacent regions." }, "1207/1207.2139_arXiv.txt": { "abstract": "We explore the weak lensing effect by line-of-sight halos and sub-halos with a mass of $M \\lesssim 10^7\\,\\ms$ in Quasi-Stellar Object(QSO)-galaxy strong lens systems with quadruple images in a concordant $\\Lambda$ cold dark matter universe. Using a polynomially fitted non-linear power spectrum $P(k)$ obtained from $N$-body simulations that can resolve halos with a mass of $M \\sim 10^5 \\ms$, or structures with a comoving wavenumber of $k\\sim 3\\times 10^2\\,h \\textrm{Mpc}^{-1}$, we find that the ratio of magnification perturbation due to intervening halos to that of a primary lens is typically $\\sim 10$ per cent and the predicted values agree well with the estimated values for 6 observed QSO-galaxy lens systems with quadruple images in the mid-infrared band without considering the effects of substructures inside a primary lens. We also find that the estimated amplitudes of convergence perturbation for the 6 lenses increase with the source redshift as predicted by theoretical models. Using an extrapolated matter power spectrum, we demonstrate that small halos or sub-halos in the line-of-sight with a mass of $M=10^3-10^7 \\ms$, or structures with a comoving wavenumber of $k=3\\times 10^2-10^4\\, h \\textrm{Mpc}^{-1}$ can significantly affect the magnification ratios of the lensed images. Flux ratio anomalies in QSO-galaxy strong lens systems offer us a unique probe into clustering property of minihalos with a mass of $M < 10^6 \\, \\ms$. ", "introduction": "Gravitational lensing is one of the most powerful tools for directly probing the structure and the distribution of dark matter. The remarkable agreement between the predicted and the observed weak lensing effects by large-scale structures or clusters provides independent and consistent estimates of clustering property of dark matter on cosmic scales $\\gtrsim 10\\, h^{-1} \\textrm{Mpc}$. However, we do not fully understand the clustering property on scales below $\\sim 1\\,h^{-1} \\textrm{Mpc}$, which correspond to individual galaxy halos. Although the cold dark matter (CDM) model predicts a large population of mini-halos ($\\lesssim 10^7\\, \\ms$), the observed number of dwarf galaxies in our galaxy seems too low in comparison with the predicted value. The discrepancy may be alleviated by some baryonic process, such as suppression of star formation by background UV radiation in the reionization epoch (e.g., \\citet{bullock2000}, \\citet{busha2010}), or tidal disruption due to a galactic disk \\citep{donghia2010}. Alternatively, the suppression of the number count might be associated with super-weakly interacting massive particles (super-WIMPs) or warm dark matter which has a larger free-streaming length than CDM \\citep{hisano2006}. In order to probe the clustering property of dark matter at mass scales of $\\lesssim 1\\,h^{-1} \\textrm{Mpc}$, strong QSO-galaxy lensing systems with quadruple images have been used in literature \\citep{metcalf2001,chiba2002}. In fact, the flux ratios in some quadruply lensed QSOs disagree with the prediction of best-fit lens models with a potential whose fluctuation scale is larger than the separation between the lensed images. Such a discrepancy called the ``anomalous flux ratio'' has been considered as an imprint of substructure inside a lensing galaxy \\citep{mao1998,metcalf2001,metcalf2004,chiba2005,sugai2007,mckean2007, more2009,minezaki2009,macleod2009}. However, recent studies based on high resolution simulations suggested that the predicted substructure population is too low to explain the observed anomalous flux ratios \\citep{maccio2006,amara2006, xu2009,xu2010,chen2009,chen2011}. More detailed modeling of gravitational potential of the lens on scales comparable to or larger than the distance between the lensed images might also improve the fit \\citep{wong2011}. However, the origin of the anomalous flux ratios in some quadruple image systems such as B1422+231 and MG0414+0534 has been veiled in mystery \\citep{chiba2005, minezaki2009}. In addition to substructures in lensing galaxy, any intergalactic halos along the entire line-of-sight from the source to the observer can perturb the lensing potential. Therefore, they may change the flux ratios of the lensed images. \\citet{chen2003} have found that the contribution from intergalactic halos modeled as singular isothermal spheres would be $\\lesssim10\\,\\%$ of that from substructures within the lensing halo. \\citet{metcalf2005a} performed ray-tracing simulations for intergalactic halos with a mass of $10^6\\, \\ms \\le M \\le 10^9\\, \\ms$. Assuming that the halos have Navarro, Frenk \\& White (NFW) \\citep{navarro1997} profiles and the number density is given by the Press-Schechter mass function \\citep{press1974}, he found that four radio lensed QSOs that shows a strong cusp-caustic violation are consistent with the predicted values without any substructures in the lensing galaxy. Assuming that halo profiles are modeled as singular isothermal spheres and the number density is given by the Sheth-Tormen mass function \\citep{sheth2002}, \\citet{miranda2007} obtained a similar conclusion for three radio and two optical/IR lensed QSOs. Using a $N$-body simulation that can resolve halos with a mass of $> 10^8\\, h^{-1} \\ms$, and halos with a mass ($10^6\\,\\ms \\le M \\le \\, 10^8\\, \\ms$) whose number density obeys the Sheth-Tormen mass function, \\citet{xu2012} obtained a result that violation of the cusp-caustic relation caused by line-of-sight halos are comparable to (even larger than) those caused by intrinsic substructures though it depends sensitively on the halo profile. In order to estimate the magnification perturbation due to intervening halos more precisely, it is important to take into account various effects that have been overlooked in literature. Firstly, if the shifts in relative positions of images and lens due to line-of-sight halos are too large, fitting a model with a smooth potential to the observed data becomes difficult since such a change is a consequence of a local effect. Moreover, even if the individual perturbing halo is not so massive, clustering halos could produce larger image shifts. Therefore, we need to incorporate the effects of clustering as well as the shifts of position of images and lens. Accuracy in observed positions of lensed images and lens would give an upper limit on the mass scale of perturbing halos. Secondly, in some lens systems, violation of the cusp-caustic relation might be caused by relatively massive faint satellite galaxies in the neighborhood of the lensing galaxy \\citep{mckean2007,shin2008,macleod2009}. Therefore, application of the cusp-caustic relation to generic lensed QSO systems may not be appropriate. Instead, we need to use other statistics to fit the model. Thirdly, the effects of massive line-of-sight halos should be subtracted off since they can contribute to low-order components in magnification tensor such as a constant convergence and an external shear in the lens model. Otherwise, we would estimate anomalies in the flux ratios systematically large because of double counting. In this paper, we explore the weak lensing effect due to line-of-sight halos in QSO-galaxy lensing systems taking these three effects into account and study how it will affect the flux ratios of lensed QSOs with quadruple images. To take into account of halo clustering, we use $N$-body simulations to calculate the non-linear power spectrum of matter fluctuations down to mass scales of $\\sim 10^5\\, h^{-1 } \\ms$. For simplicity, however, we do not put baryons in our $N$-body simulations. Then we estimate the magnification perturbation using the obtained non-linear power spectrum and study wheather observed lensed QSO systems with quadruple images are consistent with our model prediction. In section 2, we describe magnification perturbation due to line-of-sight halos. In section 3, we derive analytic formulae for the power spectrum of convergence due to line-of-sight halos constrained from perturbations in image shifts. In section 4, we describe our $N$-body simulations for obtaining the non-linear power spectrum. In section 5, image shifts and magnification perturbation are investigated using a semi-analytic method developed in section 3. In Section 6, we describe 6 samples of QSO-galaxy lensing systems with quadruple images observed in the mid infrared (MIR) band. In section 7, we present our results on the flux ratio anomalies using these lens sysetems. In section 8, we conclude and discuss some relevant issues. In what follows, we assume a cosmology with a matter density $\\Omega_m=0.272$, a baryon density $\\Omega_b=0.046$, a cosmological constant $\\Omega_\\Lambda=0.728$, the Hubble constant $H_0=70, \\textrm{km}/\\textrm{s}/\\textrm{Mpc}$, the spectrum index $n_s=0.97$, and the root-mean-square (rms) amplitude of matter fluctuations at $8 h^{-1}\\, \\textrm{Mpc}$, $\\sigma_8=0.81$, which are obtained from the observed CMB (WMAP 7yr result, \\citep{jarosik2011}), the baryon acoustic oscillations (Percival et~al. 2010), and $H_0$ \\citep{riess2009}. ", "conclusions": "We have studied the weak lensing effect by line-of-sight halos and sub-halos with a mass of $M \\lesssim 10^7\\,\\ms$ in QSO-galaxy strong lens systems with quadruple images in a concordant $\\Lambda$CDM universe. Using a polynomially fitted non-linear power spectrum $P(k)$ obtained from $N$-body simulations that can resolve halos with a mass of $M \\sim 10^5 \\ms$, or structures with a comoving wavenumber $k = 3.2 \\times 10^2\\, h{\\rm Mpc}^{-1}$, we find that the ratio of magnification perturbation due to intervening halos to that of a primary lens is typically $\\eta \\sim 0.1 $ and the predicted values agree with the estimated values for 6 QSO-galaxy lens systems (continuum emission for 5 lenses, line emission from NLR for 1 lens) with quadruple images in the mid-infrared band without considering the effects of substructures inside the primary lens. The estimated amplitudes of convergence perturbation for the 6 lenses increase with the source redshift as predicted by our semi-analytical model. This feature strongly supports a hypothesis that the observed flux ratio anomalies are caused by intervening halos rather than substructures associated with the primary lens. However, we do not exclude minor effects from substructures especially for systems with low lens redshift $z_L$ in which the weak lensing effect is small. Using an extrapolated matter power spectrum, we have demonstrated that small halos with a mass of $M=10^3-10^7 \\ms$ can significantly affect the magnification ratios of lensed images. Instead of mass $M$, we have used comoving wavenumber $k$ for parametrizing cut off scale of matter fluctuations due to intervening halos. We have considered two types of cut off, $k_{cut}$ and $k_{lens}$. $k_{cut}$ is determined from accuracy in positions of lensed images and the primary lens since intervening halos would induce shifts in relative positions of images. $k_{lens}$ is given by the (effective) Einstein radius of the primary lens. Because large scale fluctuations are taken into account as a constant convergence and a constant shear in lens models, fluctuations that are larger than the Einstein radius should be neglected. Neglecting the shift of the center of a primary lens, we find that $k_{lens}\\lesssim k_{cut}=O[10^2]\\, h{\\rm Mpc}^{-1}$ for 5 MIR lenses and $k_{cut}\\ll k_{lens}=O[10^3]\\, h{\\rm Mpc}^{-1}$ for 1 MIR lens in our sample. We have not used the cusp-caustic relation $R_{cusp}$ in order to measure the strength of flux ratio anomalies since most of our lens systems have either a complex structure (a luminous satellite) or a broad opening angle $\\theta > 30^\\circ$. Instead, we have devised a new statistic $\\eta$, to quantify the magnification perturbation. As we need a detailed lens model that fits the observed positions of images and lens, it may sounds less generic than using $R_{cusp}$. In fact, the mass-sheet degeneracy yields ambiguity in estimating the magnification perturbation. Different models with different radial profiles would certainly give different predictions. However, this is not a problem. As we have discussed, perturbations in convergence and shear $\\delta \\kappa, \\delta \\gamma$ can be measured from extended sources surrounding the MIR continuum emitting region \\citep{inoue2005b}. From observed $\\eta$, $\\delta \\kappa $, and $\\delta \\gamma$, we would be able to break the mass-sheet degeneracy. This means such an ambiguity can be removed by estimating shifts of lensed images with spatial structures with respect to unperturbed ones. Because we have used a new statistic $\\eta$ instead of $R_{cusp}$ it is difficult to directly compare our result with previous studies \\citep {metcalf2005a,xu2012} in which the effect of clustering halos is considered to be minor. However, as our numerically obtained non-linear power spectrum incorporates all the effects of clustering halos and that of their substructures, our result indicates that clustering effect on mass scales of $M\\lesssim 10^7 \\ms$ is much important than considered in previous studies. In fact, we observed that our new statistic $\\eta$ is systematically reduced by $20\\sim 30$ per cent for $k_{max} \\le 1000\\,h{\\rm Mpc}^{-1}$ and $z_S>2.6$ if no correlation between lensed images is not taken into account. Moreover, Xu et al. 2012 considered only the case $z_S=2.0$ theoretically though $z_S$ in our lens systems varies from $0.658$ to $3.62$. We think that the restriction on the source redshift is one of the weak point in their analysis as the source redshift dependence is the most important factor to probe the contribution from the line- of-sight halos. We have first shown that observed MIR lenses indeed show lens systems with high redshift sources tend to exhibit more anomalous flux ratios than those with low redshift sources. Omitting effects of source redshift dependence, clustering of halos tend to reduce the signal of anomalous flux ratios, on the other hand, neglecting constraints from astrometric shifts or contribution from a constant convergence and shear due to line-of-sight halos (yielding upper limit of mass) tends to increase the signal. Thus, it is difficult to compare our result with the previous works in literature though the conclusion may look similar. In order to estimate the magnification perturbation constrained from shifts of positions of images and lens, we have considered a ``sharp k-space filter'' for cutting off the fluctuations on large scales. If we use ``Gaussian filters'' that are sufficiently smooth, variance in convergence can be systematically decreased than using the ``sharp k-space filter''. However, if we also consider a cut off due to modeling of a primary lens, such an effect may be negligible as large scale modes are taken into account as a constant convergence or shear. It should be noted that we have neglected effects of 3 point or 4 point correlation of matter fluctuations, which may enhance the flux ratio anomalies. In order to incorporate these effects and check validity of our approximation, we need to implement ray-tracing simulation based on $N$-body simulations. If we include effects of baryons, we naively expect further enhancement in magnification perturbation as baryon cooling would steepen the gravitational potential of halos at small scales \\citep{rudd2008,semboloni2011,vandaalen2011}. Then our result would give a lower limit of the amplitude of perturbation in magnification ratios. However, feedback from supernovae or super massive black holes could suppress such a steepening near the center of halo due to outflows \\citep{booth2009}. This might eventually suppress the magnification perturbation due to line-of-sight halos. Thus in order to improve our $N$-body simulations using only collisionless dark matter particles, it is very important to incorporate baryonic physics down to mass scale of $\\sim 10^3 \\ms$ or less. In other words, small scale baryonic physics which is relevant to galaxy formation might be gravitationally probed by the weak lensing effect in QSO-galaxy strong lensing system in the near IR or MIR band. Next generation telescopes such as the European Extreme Large Telescope (E-ELT) \\citep{gilmozzi2007} or Thirty Meter Telescope (TMT) \\citep{crampton2006} can be used to probe hundreds of such strong lens systems that are too faint for currently available largest telescopes to observe. They will provide us a unique probe into clustering property of mini-halos with a mass of $M<10^6 \\,\\ms$." }, "1207/1207.3415_arXiv.txt": { "abstract": "We have obtained \\emph{Spitzer Space Telescope} Multiband Imaging Photometer for \\emph{Spitzer} (MIPS) 24 $\\mu$m and 70 $\\mu$m observations of 215 nearby, \\emph{Hipparcos} B- and A-type common proper motion single and binary systems in the nearest OB association, Scorpius-Centaurus. Combining our MIPS observations with those of other ScoCen stars in the literature, we estimate 24 $\\mu$m B+A-type disk fractions of 17/67 (25$^{+6}_{-5}$\\%), 36/131 (27$^{+4}_{-4}$\\%), and 23/95 (24$^{+5}_{-4}$\\%) for Upper Scorpius ($\\sim$11 Myr), Upper Centaurus Lupus ($\\sim$15 Myr), and Lower Centaurus Crux ($\\sim$17 Myr), respectively, somewhat smaller disk fractions than previously obtained for F- and G-type members. We confirm previous \\emph{IRAS} excess detections and present new discoveries of 51 protoplanetary and debris disk systems, with fractional infrared luminosities ranging from $L_{IR}/L_{*}$ = 10$^{-6}$ to 10$^{-2}$ and grain temperatures ranging from $T_{gr}$ = 40 - 300 K. In addition, we confirm that the 24 $\\mu$m and 70 $\\mu$m excesses (or fractional infrared luminosities) around B+A type stars are smaller than those measured toward F+G type stars and hypothesize that the observed disk property dependence on stellar mass may be the result of a higher stellar companion fraction around B- and A-type stars at 10 - 200 AU and/or the presence of Jupiter-mass companions in the disks around F- and G-type stars. Finally, we note that the majority of the ScoCen 24 $\\mu$m excess sources also possess 12 $\\mu$m excess, indicating that Earth-like planets may be forming via collisions in the terrestrial planet zone at $\\sim$10 - 100 Myr. ", "introduction": "Radial velocity studies of solar-like stars and \"retired\" intermediate-mass stars have discovered hundreds of Jovian-like planets and revealed correlations between host star properties and planet frequency. Measurements of [Fe/H] suggest that stars with super-solar metal abundance ($>$0.3 dex) are almost 10$\\times$ more likely to possess gas giant companions than those with subsolar abundance (-0.5$<$[Fe/H]$<$0.0) \\citep{fischer05}. Studies of host star mass suggest that the fraction of stars with Jovian planet companions increases as a function of stellar mass with intermediate mass stars (2 $M_{\\sun}$) possessing twice as many companions on average compared to solar-like stars \\citep{johnson10}. Since planets are believed to form in circumstellar disks around young stars, planetary demographics in conjunction with disk observations, are expected to place constraints on the processes by which planets form. We have conducted a search for dusty disks around young stars in ScoCen to determine whether disk properties (e.g. frequency, fractional infrared luminosity, grain temperature) are consistent with oligarchic growth models and dependent on stellar host properties. The Scorpius-Centaurus OB association (ScoCen), with typical stellar distances of $\\sim$100 - 200 pc, is the closest OB association to the Sun and contains three subgroups: Upper Scorpius (US), Upper Centaurus Lupus (UCL), and Lower Centaurus Crux (LCC), with estimated ages of $\\sim$11 Myr \\citep{pecaut12}, $\\sim$15 Myr, and $\\sim$17 Myr \\citep{mamajek02}, respectively. The close proximity of ScoCen and the age of its constituent stars make this association an excellent laboratory for studying the formation and evolution of planetary systems. Several hundred candidate members have been identified to date; although, the association probably contains thousands of low-mass members. Member stars with spectral type F and earlier have been identified using moving group analysis of \\emph{Hipparcos} positions, parallaxes, and proper motions \\citep{dezeeuw99}, while later type members have been identified using youth indicators (i.e., high coronal X-ray activity and large lithium abundance; \\cite{preibisch08,slesnick06}). Infrared surveys of $\\sim$10 - 20 Myr old moving groups and clusters indicate that young stars possess disks with a wide variety of properties. A \\emph{Spitzer} IRAC and IRS peak-up survey of 204 B- through M-type (0.1 - 20 M$_{\\sun}$) stars in Upper Sco found that the disks around late-type K- and M-type members are optically thick and accreting while those around B- and A-type members are optically thin and gas-poor, consistent with the expectation that disks evolve faster around higher mass stars \\citep{carpenter06}. High-resolution imaging studies of TW Hya and HR 4796A in the $\\sim$10 Myr old TW Hya Association suggest that 10 - 20 Myr old disks may be sculpted by planets regardless of the mass of the central star or the evolutionary state of the disk. VLA thermal emission mapping has revealed the presence of a $\\sim$4 AU inner hole in the TW Hya disk that may be dynamically cleared by a forming giant planet \\citep{hughes07}. \\emph{HST} STIS scattered light imaging has revealed the presence of a steep inner truncation at $\\sim$65 AU in the ring-like debris disk around HR 4796A that is consistent with the presence of one or more planetary-mass companions \\citep{schneider09}. \\cite{kenyon04,kenyon08} have developed \"self-stirred\" disk models to describe the formation of oligarchs and their impact on disk evolution around solar to intermediate mass stars with ages $>$few million years. They assume that disks are initially gas-rich and possess planetesimals with radii 1 m - 1 km at 30 - 150 AU from the central star. Since the disks are initially gas rich, the relative velocities between planetesimals are small, leading to constructive collisions between planetesimals that continue to grow until they have formed oligarchs (1000 km-sized bodies) that have depleted their feeding zone. The oligarchs are then massive enough to gravitationally perturb remnant planetesimals into eccentric orbits where they collide at high relative velocities, generating micron-size sized grains that may be detected via thermal emission. \\cite{kenyon04,kenyon08} predict that the number of micron-sized debris particles and therefore the infrared excess emission rises rapidly at 5 Myr and peaks at 10-15 Myr. Their models assume that the total planetesimal mass is proportional to the stellar mass, suggesting that the disks around intermediate-mass stars are dustier and brighter than those around solar-like stars. Their models also indicate that evolutionary timescale for disks around intermediate-mass stars is faster than that around solar-like stars because the dynamical timescale around higher mass stars is shorter. Several groups have used \\emph{Spitzer} MIPS at 24 $\\mu$m to search for infrared excess trends expected from self-stirred disks around intermediate mass stars at 5 - 100 Myr. The first such study was carried out by \\citet{hernandez06} who compared measurements of 26 early-type (F5 or earlier) stars in the $\\sim$5 Myr Orion OB1a and 34 early-type stars in the $\\sim$10 Myr Orion OB1b at $\\sim$400 pc with observations of intermediate-mass stars in young clusters. Based on a modest sample of stars, they concluded that 24 $\\mu$m excess peaks at 10 Myr, consistent with the onset of oligarchic growth at 30-150 AU. \\citet{currie08} carried out a larger survey of $\\sim$600 intermediate and high-mass stars in h and $\\chi$ Persei at $\\sim$2.5 kpc. Their statistical analysis also indicated a peak in 24 $\\mu$m excess emission at 10-15 Myr, followed by a decline up to 1 Gyr. More recently, \\citet{carpenter09} observed 62 B- and A-type members of the $\\sim$11 Myr Upper Sco. They found that an increase in 24 $\\mu$m excess emission for stars with ages between 5 and 17 Myr was statistically insignificant ($<$2$\\sigma$). We report the results of a \\textit{Spitzer} MIPS 24 $\\mu$m and 70 $\\mu$m survey of 215 B- and A-type \\emph{Hipparcos} common proper motion members of ScoCen stars, expanding on our survey of 182 F- and G-type members \\citep{chen11}. We list the targets for the full sample, along with their spectral types, distances, and subgroup memberships in Table \\ref{tab:starprops}. We compare our data with (1) \\cite{kenyon04,kenyon08} self-stirred disk models to determine whether models of 10-20 Myr disks are consistent with our data and (2) MIPS 24 $\\mu$m and 70 $\\mu$m photometry of F- and G-type members \\citep{chen11} to search for evidence of stellar mass dependent disk evolution, expected based on results from radial velocity surveys. A recent cursory analysis of the preliminary WISE ScoCen data indicates that the disks around intermediate type members are less massive than those around solar-like members \\citep{rizzuto12}. ", "conclusions": "We have obtained \\emph{Spitzer} MIPS 24 and 70 $\\mu$m photometry of 215 candidate B- and A-type members of Scorpius-Centaurus. We conclude the following: 1. The ScoCen subgroups US, UCL, and LCC possess statistically indistinguishable B+A 24 $\\mu$m excess fractions of 25$^{+6}_{-5}$\\%, 27$^{+4}_{-4}$\\%, and 24$^{+5}_{-4}$\\%, somewhat lower than the disk fractions observed by \\cite{chen11} for F+G stars. 2. The $F_{\\nu}$(24 $\\mu$m)/$F_{*}$(24 $\\mu$m) and $F_{\\nu}$(70 $\\mu$m)/$F_{*}$(70 $\\mu$m) excesses of B+A stars is systematically smaller than that measured toward F+G stars, consistent with recent WISE observations \\citep{rizzuto12}. Estimates of fractional infrared luminosity and grain temperature suggest that $L_{IR}/L_{*}$ decreases with increasing stellar mass and $T_{gr}$ is not dependent on stellar mass. 3. For 2.5 $M_{\\sun}$ stars, the debris disk fraction does not appear to change statistically between 10 Myr and 15-20 Myr; however the 1st quartile of the MIPS $F_{\\nu}$(24 $\\mu$m)/$F_{*}$(24 $\\mu$m) increases as expected from collisions between oligarchs at 30-150 AU in self-stirred disks. For 2.0 $M_{\\sun}$ stars, the disk fraction does not appear to change statistically between 10 Myr and 15-20 Myr and the 1st quartile of the MIPS $F_{\\nu}$(24 $\\mu$m)/$F_{*}$(24 $\\mu$m) increases only if primordial disks are retained in the sample. 4. The known fraction of stellar companions at 10 - 200 AU around B+A stars is approximately a factor of two higher than that around F+G stars; therefore, the lower disk fraction and the lower fractional infrared excesses associated with detected disks may be the result of disk truncation. 5. ScoCen B- through G-type members with MIPS 24 $\\mu$m excess also possess weak WISE 12 $\\mu$m excess, indicating the presence of additional warm dust in the terrestrial planet forming zone, suggesting that terrestrial planets may still be forming around ScoCen stars at 10 - 100 Myr. We would like to thank D. Hines, G. Kennedy, S. Kenyon, S. Lubow, and M. Wyatt for their helpful comments and suggestions. This work is based on observations made with the \\emph{Spitzer Space Telescope}, which is operated by JPL/Caltech under a contract with NASA. Support for this work was provided by NASA through an award issued by JPL/Caltech. This research made use of the SIMBAD database, operated at CDS, Strasbourg, France, and data products from the 2MASS, which is a joint project of the U. Massachusetts and the Infrared Processing and Analysis Center/Caltech, funded by NASA and the NSF." }, "1207/1207.6006_arXiv.txt": { "abstract": "We study the temporal evolution of umbral dots (\\uds) using measurements from the CRISP imaging spectropolarimeter at the Swedish 1-m Solar Telescope. Scans of the magnetically sensitive 630~nm iron lines were performed under stable atmospheric conditions for 71~min with a cadence of 63~s. These observations allow us to investigate the magnetic field and velocity in and around \\uds\\ at a resolution approaching 0\\farcs13. From the analysis of 339 UDs, we", "introduction": "\\label{sec:introduction} Umbral dots (\\uds) are transient brightenings observed in sunspot umbrae and pores, with typical sizes of 300~km and lifetimes of 10~min \\citep[e.g.,][]{1997A&A...328..682S, 1997A&A...328..689S}. They cover only 3--10\\%\\ of the umbral area, but contribute 10--20\\%\\ of its brightness. For this reason, \\uds\\ are believed to play a vital role in the energy balance of sunspots \\citep{1965ApJ...141..548D, 1993ApJ...415..832S, 2010SoPh..267....1M}. \\uds\\ exhibit systematic proper motions in mature sunspots: those appearing in the central umbral region are static, while \\uds\\ in peripheral regions move inward with an average velocity of 1.0\\,\\kms\\/ \\citep{2007PASJ...59S.585K, 2008A&A...492..233R}. Some peripheral \\uds\\ are the continuation of penumbral grains---bright elongated structures at the head of penumbral filaments that move toward the center of the sunspot with speeds of about 0.4\\,\\kms\\ \\citep{1999A&A...348..621S, 2006ApJ...646..593R}. When the migration front detaches into a circular bright point, the tip of the penumbral grain becomes an \\ud. It is believed that the mechanism behind \\uds\\ is convection interacting with the strong vertical field of the umbra, and many observational results support this idea \\citep{2008ApJ...678L.157R, 2010A&A...510A..12B, 2010SoPh..266....5W}. In the formation phase of sunspots, \\uds\\ are akin to granules but their apparent motion is more stochastic because of the surrounding magnetic field \\citep{1999ApJ...511..436S}. In developed sunspots, \\uds\\ are small and quiescent due to the stronger suppression of convection. UD research is entering a new phase in which computer simulations guide observational efforts. The innovative simulations by \\citet{2006ApJ...641L..73S} predicted UDs with central dark lanes and small localized downflow patches at their ends. A clear detection of those features would immediately validate the numerical models, so they have been the target of recent observations. The dark lanes are the result of enhanced density in the upper central part of UDs, caused by the piling up of hot gas that rises from deeper down. Once the gas reaches the surface, it cools by radiative losses and descends in narrow downflow channels at the end of the dark lanes. \\citet{2007ApJ...669L..57B} observed a dark lane in a big \\ud. However, the very large size of this UD ($>$1000\\,km) suggests that it could actually have been a cluster of several \\uds. The first detection of downflows surrounding an UD was presented by \\citet{2007ApJ...665L..79B}. Subsequently, \\citet{2010ApJ...713.1282O} reported solid evidence of dark lanes and localized downflows based on spectropolarimetric observations taken at the Swedish 1-m Solar Telescope. The sizes of the dark lanes and downflowing patches found by \\citet{2010ApJ...713.1282O} are near the diffraction limit of the telescope, with the substructures keeping their identity for periods of only a few minutes. These authors also reported enhanced net circular polarization at the site of the downflows. \\begin{figure*}[bhtp] \\centerline{\\includegraphics[width=1\\textwidth, bb=0 0 1020 651]{f1_final.pdf}} \\caption{Filtergrams from the best scan of the data set, taken at 08:30\\,UT. The full FOV is shown. From left to right: Stokes $I$ in the blue wing of 6301.5\\,\\AA\\ (line position 0), Stokes $I$ at line center (line position 7), and Stokes $V$ in the blue wing (line position 5). The direction to disk center (DC) is displayed with an arrow. The white rectangle indicates the FOV of Figure~\\ref{fig:cut_image}. The bottom row shows the Stokes profiles emerging from the \\ud\\ marked with triangles in the upper panels. The $+$-symbols indicate the measured signals. } \\label{fig:fig1} \\end{figure*} The evolution of \\uds\\ and their magnetic fields is difficult to study---and hence poorly known---because one needs full vector spectropolarimetric measurements at very high temporal and spatial resolution. To the best of our knowledge, the magnetic properties of \\Uds\\ have never been investigated at the required cadence and spatial resolution \\citep[but see][]{2009ApJ...694.1080S}. \\citet{2010ApJ...713.1282O} performed a preliminary analysis of the temporal evolution of six \\uds, and this work should be considered a substantial extension of their study. Both the cadence and the polarimetric sensitivity of our measurements are improved with respect to those of \\citet{2010ApJ...713.1282O}, as is the total duration of the observations, 71 minutes, during which the seeing conditions were excellent and stable. We use this unique data set to investigate the evolution of the magnetic and velocity fields in and around \\uds. The paper is organized as follows. The observations are described in Section~\\ref{sec:observation}, followed by an account of the methods used for the detection of \\uds\\ and derivation of the velocity and magnetic information (Section~\\ref{sec:reduction}). In Section~\\ref{sec:convection} we quantify how convection is modified in the umbra. In Section~\\ref{sec:UDs} we describe the evolution of some typical \\uds, and present the results of our statistical analysis. Finally, based on these results, we discuss the physical properties of \\uds\\ in Section~\\ref{sec:discussion}. ", "conclusions": "\\label{sec:discussion} In this paper we have performed a detailed analysis of \\uds\\ in a mature sunspot using data from the CRISP spectropolarimeter. The excellent spatial resolution, temporal cadence, and polarimetric sensitivity of the measurements are ideal for \\ud\\ studies. The perturbations caused by \\uds\\ are usually very small, and thus an statistical approach has to be followed to reveal their common properties. Our work addresses for the first time the temporal evolution of velocity and magnetic fields in and around \\uds, using a statistically significant sample of \\uds. \\paragraph{Convection in the umbra} \\Uds\\ are considered to be the manifestation of convection in the presence of the strong umbral field, while more vigorous convection occurs in the quiet Sun in the form of granules. The morphological differences between granules and \\uds\\ can be seen in Figure~\\ref{fig:convection_morphology}. Driven by overshooting cellular convection, granules are characterized by sharp edges and irregular polygonal shapes \\citep{1990ARA&A..28..263S}. \\Uds, on the other hand, show Gaussian brightness profiles. Linear theory reveals that the preferred horizontal scale of convection decreases with increasing field strength \\citep{1990MNRAS.245..434W}, which explains why \\uds\\ are smaller than granules. The convective origin of \\uds\\ seems well established, as many papers including ours found a good correlation between upflows and brightness. However, in the umbra there are other diffuse areas with enhanced brightness whose origin is still unknown. Our speculation that the diffuse bright areas of the umbra are also caused by convection could not be unambiguously confirmed on the basis of a unique relation between brightness and velocities (Figure~\\ref{fig:bisector_scatter}). \\paragraph{Photometric properties} \\citet{2009ApJ...702.1048W} found constant lifetimes regardless of the \\ud\\ type and the structure of magnetic field at the position of the UD. This is in agreement with our results (Section\\,\\ref{sec:scatter}). However other studies report longer lifetimes for brighter \\uds\\ \\citep{2002A&A...388.1048T, 2010A&A...510A..12B}, and we speculate this is partly due to the fact that bright \\uds\\ tend to reoccur at the same position. A correlation between shorter lifetime and faster proper motion is reported for the first time in this paper (Figure~\\ref{fig:scatter_9parameter}). If the energy dissipation rate is proportional to the UD speed, the lifetime can be expected to be reduced for fast-moving UDs, as observed. We also find that the travelled distance depends linearly on lifetime, i.e., long-lived \\uds\\ travel longer distances even though they move at lower speeds. The same conclusion can be obtained from a similar analysis of the data set of \\citet{2009ApJ...702.1048W}. We found \\ud\\ diameters consistent with the values reported in the literature, i.e., about 400\\,km on average. The Gaussian shape of the histogram (Figure~\\ref{fig:histogram_4parameter}) and the lack of dependence of the diameter on brightness ratio (Figure~\\ref{fig:scatter_9parameter}) suggest that \\uds\\ indeed have a ``typical'' size, regardless of their type. This common \\ud\\ size is probably determined by a universal near-surface stratification in mature sunspots. However, the scatter plot analysis performed in Section~\\ref{sec:scatter} did not reveal any physical parameter having a strong correlation with the \\ud\\ size. The intensity oscillations in UDs reported by, e.g., \\citet{1997ApJ...490..458R} have not been studied in this paper because of the uncertainties that residual seeing fluctuations may introduce. However it is true that many \\uds\\ show recurrence, as observed also by \\citet{2012ApJ...752..109L}. For example, the peripheral UD\\#C displayed in Figure~\\ref{fig:PUD_036_plot} reappeared twice within a time interval of 13~minutes. This timescale is comparable to the oscillatory period of the \\uds\\ shown in \\citet{1997ApJ...490..458R} and \\citet{2009PASJ...61..193W}. \\paragraph{Categorization of \\uds} We classified the observed \\uds\\ in central, peripheral, and grain-origin UDs according to their place of birth. Do these categories represent physically different structures or different manifestations of the same phenomenon? Grain-origin \\uds\\ have larger brightness ratios, larger sizes, and faster proper motions than the other UDs. The temporal evolution of grain-origin \\uds\\ is smooth and shows monotonic changes, while central and peripheral \\uds\\ show mound-shaped evolutionary curves (Figure~\\ref{fig:typical_lightcurve}). Despite this, the scatter plots presented in Figure~\\ref{fig:scatter_9parameter} suggest that the properties of grain-origin \\uds\\ lie on the extension lines of those of central and peripheral \\uds. The differences between them may arise from stronger convection in grain-origin \\uds\\ rendered possible by the weaker background field. \\citet{2006A&A...447..343S} speculated that field-free convection can explain both \\uds\\ and penumbral grains. The computer simulations of \\citet{2007ApJ...669.1390H} and \\citet{2009ApJ...691..640R} succeeded in reproducing basic properties of penumbral filaments and \\uds\\ as weakly-magnetized convective structures. Our results are in general agreement with the predictions of this scenario, although they also indicate that UDs are far from being completely field-free (at least in the photospheric layers accessible to the observations). \\paragraph{Substructures} Localized downflow patches at the periphery of UDs are considered to be a signature of overturning convection in \\uds\\ \\citep{2006ApJ...641L..73S}. In Section\\,\\ref{sec:verydeep} we presented some examples of downflow patches without performing a full statistical analysis. The patches of Figure~\\ref{fig:verydeep} have sizes of 0\\farcs2 and redshifts of up to 0.75\\,\\kms. Both the size and the velocity are in good agreement with those reported by \\citet{2010ApJ...713.1282O} in a pore. The fact that the downflow patches are observable only at very high bisector levels (i.e., deep photospheric layers) is also consistent with the results of those authors. However, many \\uds\\ do not show downflow patches in our data. This lack of detection could be due to: \\begin{enumerate} \\item Insufficient spatial resolution. \\item The existence of downflows only in deep layers that cannot be probed by the \\ion{Fe}{1} 6301 and 6302 \\AA\\/ lines. \\item The transient nature of the downflows, which could appear only in a particular phase of the UD's evolution. \\item The possibility that the convective energy escapes to the upper layers instead of returning to deep layers \\citep[see the narrow jet-like upflows above the cusp in][]{2006ApJ...641L..73S}. \\end{enumerate} The two \\uds\\ featured in Figure~\\ref{fig:verydeep} are in the peak brightness phase when they show downflows. Possibly the speed of the downflows reaches a maximum when the brightness is also maximum. \\paragraph{Velocities in \\uds} Upflows and brightness follow similar evolutionary patterns in UDs (Figure~\\ref{fig:typical_lightcurve}). This correlation supports the convective nature of \\uds\\ \\citep{2009ApJ...694.1080S, 2010SoPh..266....5W}. The scatter plots of velocity vs brightness shown in Figure~\\ref{fig:scatter_9parameter} also point to a convective origin of UDs. For central UDs, however, we observe stronger blueshifts in darker structures, although the tendency is not very pronounced. We suspect this is an artifact caused by systematic blueshifts in very cold umbral areas. The velocities observed within \\uds\\ are likely to represent field-aligned flows, because upflows are readily found on the disk-center side where the field is closer to the line-of-sight direction, compared to the limb side (Section\\,\\ref{sec:grain-origin}). The same effect can be observed in the Evershed flow (Figure~\\ref{fig:data_reduction}), which is also a field-aligned flow. \\paragraph{Magnetic field in and around \\uds} The field-free convection model of \\uds\\ \\citep{2006ApJ...641L..73S} predicts weaker and more inclined fields in \\uds. Our scatter analysis (Figure~\\ref{fig:scatter_9parameter}) confirms these properties, but only for \\uds\\ with fields below 2000\\,G. In strongly magnetized \\uds\\ ($>2000$~G), the magnetic field is enhanced and more vertical compared to the surroundings. For grain-origin UDs, the physical conditions also depend on the phase of evolution: in the first half of their lifetime, weaker and more inclined fields appear, while stronger and more vertical fields are observed in the latter half as the UDs intrude into the umbra. To the best of our knowledge, the enhanced and more vertical fields of strongly magnetized \\uds\\ cannot be explained by currently available \\ud\\ models. We observe strong field regions at the migration front of grain-origin UDs for the first time (see the evolution of UD\\#D and UD\\#E in Section~\\ref{sec:grain-origin}). These strong field regions seem to impede the migration of the UDs. The situation is reminiscent of that modeled by \\citet{1998A&A...337..897S}, where a weakly magnetized penumbral flux tube pushes and compresses the pre-existing vertical field at the leading edge. However, also a weakly magnetized convective structure would produce a compression of the adjacent magnetic field, which has to wrap around the field-free gas. The MHD simulations performed by \\citet{2007ApJ...669.1390H} predict the existence of enhanced field regions surrounding grain-origin \\ud\\ only in layers deeper than the continuum forming region (see Figure 3 in that paper), but on both the leading and the tail sides. Our observation did not find field enhancements on the tail side. \\paragraph{Final remarks} This work extends our knowledge of the temporal evolution of velocities and magnetic fields in UDs. We found some new and unanswered results that may provide constraints to future modeling efforts. A pioneering comparison of observational and computer-simulated \\uds\\ has been performed by \\citet{2010A&A...510A..12B}, and this kind of studies should be extended. At the same time, more spectropolarimetric observations of UDs at high cadence should be performed. The temporal resolution of our data, 63\\,s, seems appropriate to track the evolution of \\uds, but is insufficient for resolving the evolution of UD substructures. \\Uds\\ will remain one of the most challenging targets for solar observations in the coming years." }, "1207/1207.4235_arXiv.txt": { "abstract": "{ In this paper, we study an extension of the standard model with a vector-like generation of leptons. This model provides a viable dark matter candidate and a possibility to enhance the Higgs decay rate into a pair of photons. We evaluate constraints from electroweak precision tests and from vacuum stability, and find that the latter provide an upper limit on the lepton Yukawa couplings. A large enhancement of the Higgs di-photon rate can therefore only be obtained when the mass of the lightest charged lepton is close to the LEP limit. The relic density constraint suggests a co-annihilation scenario with a small mass difference between the lightest charged and neutral leptons, which also weakens the LEP limit on the lightest charged lepton mass and allows for larger Higgs di-photon decay rates. Cross sections for direct detection of the dark matter candidate are calculated, and prospects for detecting the new particles at the LHC are discussed briefly. } ", "introduction": "\\label{sec:into} The discovery of a scalar resonance consistent with the SM Higgs boson~\\cite{Gianotti:gia12,Incandela:inc12} marks a great success in the early run of the CERN Large Hadron Collider (LHC) and its associated experiments. It is now conceivable that electroweak symmetry is indeed broken by the Higgs mechanism, namely through the vacuum expectation value of a fundamental Higgs doublet. With a Higgs mass of about 125~GeV, this model is fully renormalizable and potentially stable up to very high scales. However some mysteries remain that suggest additional particles being present at the electroweak scale. First, the existence of dark matter has been established through astrophysical observations beyond reasonable doubt, however its nature and properties remain largely unknown. With all known matter content in the SM being related to the weak scale, it is at least a reasonable assumption that also the dark matter particle is connected to the electroweak scale. The second mystery is provided by the Higgs boson itself. While the existence of a new resonance has very recently been established, it remains to be seen whether its properties, in particular the production cross sections and decay branching ratios, agree with the very precise predictions from the SM. The decay channels that are most sensitive to new physics effects are the loop induced decays of the Higgs to pairs of photons or to a photon and a Z boson. And indeed the rate of di-photon events shows the largest and most consistent deviation from SM expectations, with both 2011~\\cite{ATLAS:2012ad,Chatrchyan:2012tw} and 2012~\\cite{Atlasnote:-2012-091,CMSnote:12-015} data from both the Atlas and CMS experiments suggesting an enhancement of the branching ratio of about 50\\% (but still within SM expectations at the two sigma level\\footnote{Very recent analyses of the combined ATLAS and CMS data sets suggest that the deviation of the di-photon signal from SM expectations is at the $2\\sigma$-$2.5\\sigma$ level~\\cite{Giardino:2012dp,Espinosa:2012im,Low:2012rj,Ellis:2012hz,Carmi:2012in}.}). Motivated by the above, here we study a simple extension of the SM that provides a dark matter candidate and allows for an enhancement of the Higgs branching ratio to photon pairs. The model introduces a vector-like fourth generation of leptons, namely SU(2) doublets $\\ell = (\\ell_{\\rm L}', \\ell_{\\rm R}'')$, charged SU(2) singlets $e = (e_{\\rm R}', e_{\\rm L}'')$ and neutral singlets $\\nu = (\\nu'_{\\rm R}, \\nu''_{\\rm L})$. In the limit where the vector-like masses vanish the particle content is that of a fourth generation of leptons (indicated by a single prime) and an exact copy with opposite chirality (double primed). The remainder of the paper is organized as follows. In the next section, the particle content of the model is introduced and the mixing in the charged and neutral sector is discussed. Electroweak precision tests and other constraints on the model are presented in Sec.~\\ref{sec:EWP} while the modified Higgs boson properties are analyzed in detail in Sec.~\\ref{sec:higgs}. In Sec.~\\ref{sec:rge} constraints on the magnitude of the Yukawa couplings from vacuum stability considerations are derived. Dark matter properties are explored in Sec.~\\ref{sec:dm} before we conclude in Sec.~\\ref{sec:conclusions}. Formulas for the electroweak precision observables are presented in detail in the appendix. ", "conclusions": "\\label{sec:conclusions} In this work, we have presented a compact extension of the SM by a vector-like family of leptons, and shown that it can provide a viable dark matter candidate and a source for the enhancement in the Higgs branching ratio to photon pairs, thus explaining two of the remaining mysteries of the SM. We have performed a detailed analysis of electroweak precision constraints, outlining the decoupling effects of the vector-like masses that suppress contributions to the S and T parameters even in the presence of sizable custodial symmetry breaking from Yukawa couplings. It is shown that the Higgs branching ratio can be enhanced if there is mixing in the charged lepton sector, while the case without mixing is excluded at the $3\\sigma$ level by the observation of the Higgs boson in the di-photon channel. The running of the Higgs quartic coupling places an upper bound on the Yukawa couplings. If one demands that the model is stable up to energies accessible at the LHC, then $Y_c', Y_c'' \\lesssim 1$, assuming that the effects of the neutral Yukawas are negligible. This leads to a strong correlation between the maximal attainable value for $R_{\\gamma \\gamma}$ and the mass of the lightest charged lepton, which has to be ${\\cal O}(100~$GeV) in order to obtain a 50\\% enhancement. Such a light charged particle will eventually be in reach of the LHC experiments, thus making the model testable in the near future. In the absence of mixing between the new lepton sector and the SM leptons, the lightest neutrino is stable. To satisfy the relic density constraint while keeping the neutral Yukawa couplings small, either resonant annihilation or co-annihilation with the lightest charged lepton is needed. The latter possibility is preferred since it also avoids invisible Higgs decays to pairs of dark matter. The preferred mass difference $\\Delta M = M_{E_1} - M_{N_1}$ for the co-annihilation scenario is between 5~GeV and 10~GeV. In this regime the strongest LEP limits on $M_{E_1}$ do not apply, and charged lepton masses below $100$~GeV and values of $R_{\\gamma \\gamma} \\sim 2$ become possible. At the LHC, the most promising signal will come from Drell-Yan production of $E_1^+ E_1^-$ pairs that subsequently decay through a virtual $W$ boson, $E_1 \\to W^{*} N_1$, giving rise to final states with leptons, jets and missing energy. In the co-annihilation regime the leptons and jets from the $W^{*}$ decay become increasingly soft, so that only events where sufficiently hard jets are radiated from the initial state partons will be detectable. In this case better sensitivity can be obtained from $E_1 N_2$ and $N_1 N_2$ production, followed by $N_2 \\to W E_1$ or $N_2 \\to Z N_1$ decays where now the $Z$ and $W$ bosons can be on-shell. Furthermore monojet and mono-photon + missing energy searches can be employed to place a bound on direct $E_1^+ E_1^-$ and $N_1 N_1$ production rates. These searches might be able to probe the $m_{E_1}\\lesssim 120$~GeV regime of our model with a few tens of fb$^{-1}$ in the 14~TeV LHC run~\\cite{ArkaniHamed:2012kq}, provided that systematic uncertainties are under control. In addition the dark matter candidate can be searched for in direct detection experiments. The SI and SD cross sections are in reach of the next generation of experiments, but the search will be challenging in the co-annihilation regime. It remains to say that the model is mostly phenomenologically motivated, and should be viewed as an effective theory at the weak scale that will eventually be embedded in a more complete UV theory, which should also ensure vacuum stability beyond the 10~TeV scale. One option would certainly be a supersymmetric completion of the model with a somewhat higher supersymmetry breaking scale and superpartners above the TeV scale. Another intriguing possibility is the completion with a larger number of vector-like families of quarks and leptons that can give successful gauge coupling unification when three complete families are added to the SM~\\cite{Dermisek:2012as}." }, "1207/1207.6230_arXiv.txt": { "abstract": "We calculate the particle production rate in an expanding universe with a three-torus topology. We discuss also the complete evolution of the size of such a universe. The energy density of particles created through the nonzero modes is computed for selected masses. The unique contribution of the zero mode and its properties are also analyzed. \\vspace{8mm} ", "introduction": "Although current astrophysical observations provide precise information on the geometry of the universe \\cite{Komatsu:2010fb}, its topology remains a mystery. We don't even know whether the universe is compact or infinite. Nevertheless, lower bounds can be put on its size for each compact topology (see, \\cite{Glen} and references therein). Amongst the possible topologies for the universe those with some or all spatial dimensions compactified are especially interesting, since then Casimir energies provide an additional contribution to the energy density. This may lead to an interesting vacuum structure for the standard model coupled to gravity which is insensitive to quantum gravity effects \\cite{ArkaniHamed:2007gg, AFW, FW}. The simplest flat topology with all dimensions compactified is a three-torus, and that is the topology we will concentrate on throughout this paper. The scenario of our universe having a three-torus topology was investigated by many authors (see, \\cite{Linde} and references therein). Probably one of the most appealing features of such a model is that the creation of a three-torus universe is much more likely to occur than that of an infinite flat or closed universe \\cite{Zeldov,Linde}. In addition, it has also been shown that a three-torus topology can provide convenient initial conditions for inflation \\cite{Zeldov}. Here we consider gravitational particle creation in an expanding toroidal universe. The particle production formalism was developed in \\cite{Parker1,Parker2,Parker3} and investigated in great detail in later works (see, \\cite{Zeldovich,Starob,Grib} and references therein). Since then, it has been thoroughly studied in the case of a FRW cosmology, including its implications for dark matter creation around the inflationary epoch \\cite{Chung}. However, particle production in a toroidal universe hasn't been extensively studied (see, \\cite{Berger} for some work on the subject). In this paper we provide a detailed numerical calculation of the particle production in a three-torus universe. We start with introducing the relevant formalism. We then discuss the evolution of the size and energy density of the universe from the Planck time to the present time. Next, we derive analytical formulae for the particle number and energy density at early and late times. We then find full numerical solutions and confirm that they agree with the analytical approximations in the appropriate regions. The particle production through the nonzero modes is somewhat similar as in the closed universe case discussed in \\cite{Starob}. However, the three-torus particle creation includes an additional contribution from the zero mode, which strongly depends on the choice of initial conditions. ", "conclusions": "In the present work we calculated the rate for gravitational particle creation in a three-torus universe. We performed the calculation assuming the metric on the three-torus for which the shape moduli are stable. We adopted the current size of the universe of ten Hubble radii and estimated its full evolution in time. We argued that the Casimir energies might dominate the total energy density before inflation, resulting in the expansion of the universe like in the radiation era. We then showed that there are two types of contributions to the particle production rate -- one from the nonzero modes on a torus, and the other coming from the zero mode. The nonzero mode contribution is interesting because at early times it is characterized by a ``quasi-vacuum-like'' equation of state for the produced matter. The particle production itself continues until $t \\approx 1/m$, after which it essentially stops. It turns out that the energy density of particles created through the nonzero modes is tiny compared to the total energy density in the universe at all stages of its evolution. The production through the zero mode is more unique. The corresponding energy density of the created particles is very sensitive to initial conditions. If we set those at the Planck time, the energy density of the particles is again small with respect to the total energy density of the universe. We showed that although the particles produced through the zero mode have a ``quasi-accelerating'' equation of state before inflation, they are described by a ``quasi-radiative'' equation afterwards. \\subsection*{Acknowledgment} The author is extremely grateful to Mark Wise for stimulating discussions and helpful comments. The work was supported in part by the U.S. Department of Energy under contract No. DE-FG02-92ER40701." }, "1207/1207.6140_arXiv.txt": { "abstract": "We have identified a previously unrecognized population of very compact, embedded low-mass Galactic stellar clusters. These tight (r$~\\approx~$0.14 pc) groupings appear as bright singular objects at the few arcsec resolution of the {\\it Spitzer Space Telescope} at 8 and 24 $\\mu$m but become resolved in the sub-arcsecond UKIDSS images. They average six stars per cluster surrounded by diffuse infrared emission and coincide with 100 -- 300 M$_{\\sun}$ clumps of molecular material within a larger molecular cloud. The magnitudes of the brightest stars are consistent with mid- to early-B stars anchoring $\\sim$80 M$_{\\sun}$ star clusters. Their evolutionary descendants are likely to be Herbig Ae/Be pre-main sequence clusters. These ultra-compact embedded clusters (UCECs) may fill part of the low-mass void in the embedded cluster mass function. We provide an initial catalog of 18 UCECs drawn from infrared Galactic Plane surveys. ", "introduction": "Most, if not all, stars are born in stellar clusters. It has been estimated that 96\\% of massive OB ($>$8 M$_{\\sun}$) stars are associated with clusters \\citep{dW05}. For nearby ($\\lesssim$2 kpc) embedded star clusters, \\citet{la03} found a flat mass distribution function, implying a power law ($\\alpha=-2$) distribution by number for young clusters. Their relation exhibits a sharp turnover for clusters with total stellar masses less than $\\sim$50 M$_{\\sun}$ and suggests that $>$90\\%\\ of all stars form in clusters more massive than this lower limit. \\citet{gu09} surveyed 36 star clusters in young star-forming regions (SFRs), mostly within 1 kpc of the sun, and found an average of 26 members per cluster with mean radii of 0.39 pc. These clusters are quite young, as evidenced by their high incidence of young stellar objects (YSOs). Despite recent advances and observations, the census of the smallest embedded clusters is still incomplete owing to the limited depth and angular resolution of large scale infrared (IR) surveys. During a study of SFRs in the Galactic Plane \\citep{al12}, we serendipitously identified two compact stellar clusters. These objects have pointlike or marginally resolved morphologies in the few arcsecond resolution mid-IR {\\it Spitzer Space Telescope} images at [8.0] and [24] \\micron, but are resolved in the sub-arcsecond $JHK$ images from the United Kingdom Infrared Deep Sky Survey (UKIDSS; \\citealt[][]{lu08}). The mid-IR images are typically dominated by a single bright object that exhibits a steeply rising spectral energy distribution (SED) through the mid-IR. The putative clusters are often found within infrared dark clouds (IRDCs) surrounded by nebulosities tracing the hot dust, reflection nebulosity, and PAH emission characteristic of young embedded clusters. Using brightness and color criteria derived from the prototype sources, we searched for similar objects in the Galactic Legacy Infrared MidPlane Survey Extraordinaire (GLIMPSE; \\citealt[][]{be03}) Point Source Catalogs (PSCs) as well as a subsample of massive YSO (MYSO) candidates from the {\\it MidCourse Space Experiment} ({\\it MSX}) Red MSX Source (RMS) catalog \\citep{ur11}. The search yielded additional candidates, and we present an initial (incomplete) sample of 18 ultra-compact embedded clusters (UCECs) and their properties. ", "conclusions": "Our initial search shows that young, embedded compact clusters can be selected by an IR color and magnitude cut with follow-up visual inspection. Undoubtedly, there are additional compact clusters that remain unidentified because of very high extinction, large field star densities, extreme compactness, or incompleteness stemming from high diffuse background levels. Our analysis revealed that all 12 of the UCECs from the MYSO sample were missing either a 2MASS $K_{S}$ or [8.0] catalog entry. The missing detection at $K_{S}$ is probably from source blending, while at [8.0] saturation likely kept the object out of the point source catalogs. Therefore, this color selection technique is limited to sources faint enough to be unsaturated in GLIMPSE. If we were to relax the selection criteria to include fainter sources with $[8.0]~=~$5--6 mag, the number of selected objects increases to more than a thousand. These could include even lower-mass clusters and those too faint to be included in the MSX catalog, but visual classification of such a large sample is beyond the scope of this work. Sub-arcsecond mid-IR imaging was performed on 14 MYSO candidates at 24.5 \\micron\\ \\citep{dw09} and on 346 MYSOs at 10.3 \\micron\\ \\citep{mo07}. These studies found that approximately 20 -- 25\\%\\ of MYSO candidates have multiple detections and/or extended diffuse emission. These wavelengths primarily detect stars with strong IR excesses rather than stellar photospheres and may miss sources that lack or have only a weak excess. However, their results are consistent with a portion of MSYO candidates being compact stellar clusters. Only four of 18 UCECs have possible radio detections. UCEC \\#11 was detected at 3.6 and 1.3 cm by \\citet{sa08} and they estimate a spectral type of B2 -- B3 for the exciting source. Three others (\\#4, \\#5, and \\#6) appear to be associated with faint 20 cm emission from the MAGPIS survey \\citep{he06}, however they do not appear in the MAGPIS point source catalog \\citep{wh05}. The estimated completeness limit at 20~cm is 14~mJy, which is sensitive enough to detect an O9.5V at 10~kpc based on the Lyman continuum flux \\citep{ma05}. This indicates that the most massive star(s) within the majority the UCECs is an early-B star producing relatively few Lyman continuum photons, consistent with the inferences drawn from the CMDs in Figure~\\ref{cmd}. Another possibility is that the stars earlier than B0 are so young that they have not yet formed a detectable H{\\scriptsize II} region \\citep{ur11}. In some cases a single, bright source appears to dominate the UCECs and may in fact be a MYSO, but in others the IR flux is more evenly distributed among cluster members, in which case the most massive star may be an early- to mid-B star. \\citet{te97} identified small clusters of PMS stars around Herbig Ae/Be stars. These clusters have radii of about 0.2 pc, typically contain 4 -- 12 stars, and $<$few Myr old. In these clusters the maximum stellar mass is correlated with the K band source counts. Field-star-subtracted source counts (6--8 stars) from Figure~\\ref{cmd} are roughly consistent with those found in Herbig Be PMS clusters (2--16) \\citep{te99}. It is likely that UCECs suffer a higher level of extinction owing to their highly embedded nature and, as a result, source counts within UCECs may be underestimated compared to more exposed and evolved Herbig Ae/Be clusters. This evidence suggests that UCECs represent a younger, more heavily embedded phase destined to evolve into Herbig Ae/Be clusters after a few Myr. \\citet{we10} found a correlation between the most-massive cluster member (m$_{max}$) and the total cluster mass (M$_{ecl}$) that cannot be explained by a random sampling of the IMF. The m$_{max}$-M$_{ecl}$ relation is incomplete below 100 M$_{\\sun}$ \\citep{we10}, but is supported by later investigations \\citep{ki11}. An accurate determination of M$_{ecl}$ depends on cluster age because of the increasing probability of losing cluster members over time \\citep{bo03}. UCECs are ideal objects for further probing the m$_{max}$-M$_{ecl}$ relation because they are likely to have M$_{ecl}~<~100$ M$_{\\sun}$ and young enough to have not lost a significant number of cluster stars. Figure~\\ref{cmd} indicates that m$_{max}$ is near an $\\sim$8 M$_{\\sun}$ B2V star, which implies M$_{ecl}~\\sim~80$ M$_{\\sun}$, while an O9V star ($\\sim$20 M$_{\\sun}$) and a B8V ($\\sim$4 M$_{\\sun}$) would have M$_{ecl}$ of 251 and 26 M$_{\\sun}$, respectively \\citep{we10}. If UCECs are 80 M$_{\\sun}$ clusters the SF efficiency, SFE = M$_{*}$/(M$_{*}$~+~M$_{gas}$) would be 0.52, 0.27, and 0.39 for \\#8, \\#13, and \\#16, respectively, while the median gas mass of the entire sample (141 M$_{\\sun}$) produces a SFE of 0.36. These SFE values are slightly higher, though still consistent, with studies of other Galactic clusters and SFRs \\citep{la03,al07}. The CMDs in Figure~\\ref{cmd} show that the clusters may contain more than one early-B star, which would increase the implied number of unseen low-mass stars for a standard IMF. These sources are absent either because they fall below the UKIDSS detection limit or the IMF is truncated in these types of objects. After stars form within a cluster, they immediately begin to expel the surrounding ISM. The rapid explusion of gas alters the cluster's potential well and may cause clusters to dissolve \\citep{bo03}. \\citet{la03} estimate up to 95\\%\\ of embedded clusters will disperse in under 5 -- 10 Myr, and those that do survive longer typically have masses over 500 M$_{\\sun}$. This puts an upper limit on the lifetime of UCECs ($<$ a few Myr) and suggests that they will quickly disperse. Such clusters, in any case, would be difficult to identify after a few Myr once the large IR luminosity arising in circumstellar and intracluster dust diminishes. \\citet{ja10} suggest that IRDCs are the precursors to massive stars and star clusters. The presence of UCECs embedded within IRDCs supports this hypothesis. After several Myr, it is likely that the IRDC will have dissipated and SF ceased. Small clusters, including UCECs, will be distrupted and may appear as a single loose cluster or stellar association of a few hundred solar masses. Thus, large stellar associations may be comprised of the distributed remnants of many smaller clusters born out of the same IRDC. The gas-free merger of small clusters may explain why the m$_{max}$-M$_{ecl}$ relation differs from random IMF sampling (for clusters $>$100 M$_{\\sun}$) by limiting accretion and growth of the most massive members, except in the most massive molecular clouds \\citep{we10}. UCECs may represent an unrecognized but significant population of low-mass stellar clusters destined to quickly disperse into the Galactic stellar field. In large enough numbers, these types of objects may be numerous enough to steepen the low-mass end of the embedded cluster mass function." }, "1207/1207.3079_arXiv.txt": { "abstract": "I present three methods to determine the distance to the Galactic centre $\\Rsun$, the solar azimuthal velocity in the Galactic rest frame $\\Vtot$ and hence the local circular speed $V_c$ at $\\Rsun$. These simple, model-independent strategies reduce the set of assumptions to near axisymmetry of the disc and are designed for kinematically hot stars, which are less affected by spiral arms and other effects. The first two methods use the position-dependent rotational streaming in the heliocentric radial velocities ($U$). The resulting rotation estimate $\\theta$ from $U$ velocities does not depend on $\\Vtot$. The first approach compares this with rotation from the galactic azimuthal velocities to constrain $\\Vtot$ at an assumed $\\Rsun$. Both $\\Vtot$ and $\\Rsun$ can be determined using the proper motion of Sgr $A^{*}$ as a second constraint. The second strategy makes use of $\\theta$ being roughly proportional to $\\Rsun$. Therefore a wrong $\\Rsun$ can be detected by an unphysical trend of $\\Vtot$ with the intrinsic rotation of different populations. From these two strategies I estimate $\\Rsun = (8.27 \\pm 0.29) \\kpc$ and $\\Vtot = (250 \\pm 9) \\kms$ for a stellar sample from SEGUE, or respectively $V_c = (238 \\pm 9) \\kms$. The result is consistent with the third estimator, where I use the angle of the mean motion of stars, which should follow the geometry of the Galactic disc. This method also gives the Solar radial motion with high accuracy. The rotation effect on $U$ velocities must not be neglected when measuring the Solar radial velocity $\\Usun$. It biases $\\Usun$ in any extended sample that is lop-sided in position angle $\\alpha$ by of order $10 \\kms$. Combining different methods I find $\\Usun \\sim 14 \\kms$, moderately higher than previous results from the Geneva-Copenhagen Survey. ", "introduction": "Among the central questions in Galactic structure and parameters are the Solar motion, Solar Galactocentric radius $R_0$ and the local circular velocity $V_c$ of our Galaxy. Galactic rotation curves are found to be generally quite flat over a vast range of Galactocentric radii \\citep[][]{Krumm79}. As common for Galaxies with exponential discs \\citep[][]{Freeman70} there is some evidence for a radial trend of the Galactic circular velocity near the Sun, but it is very moderate \\citep[][]{Feast97, McMillanB09}, so that the circular velocity at solar Galactocentric radius $\\Rsun$ characterises well the entire potential. Initially local kinematic data from stellar samples were the main source to extract Galactic parameters including $V_c$, e.g. by the use of the Oort constants \\citep[][]{Oort27}. The kinematic heat of stellar populations requires large sample sizes, so that studies determining the Local Standard of Rest (LSR) are still primarily based on stars, but some classical strategies like the position determination of the Galactic centre pioneered by \\cite{Shapley18} reached their limits by the constraints from geometric parallaxes, by the number of available luminous standard candles, and by the magnitude requirements of stellar spectroscopy. Apart from some more recent attempts to use luminous stars \\citep[e.g.][]{Burton74}, most current evidence on $\\Rsun$ and $V_c$ derives from modelling streams in the Galactic halo \\citep[see e.g.][]{Ibata01, Majewski06} and from radio observations of the $HI$ terminal velocity \\citep[see][]{McMillan11}, the Galactic centre, molecular clouds and MASERs \\citep[e.g.][]{Reid04}. While the first branch relies strongly on assumptions such as distance scale and shape of the Galactic potential \\citep[cf. the discussion in ][]{Majewski06}, there are further complications like the failure of tidal streams to delineate stellar orbits \\citep[][]{EyreB09}. Radio observations \\citep[][]{Reid04} have accurately determined the proper motion of the radio source Sgr $A^*$, which is identified with the central black hole of the Milky Way \\citep[for a discussion of possible uncertainties see also][]{Broderick11}. It tightly constrains the ratio of the solar speed in the azimuthal direction $\\Vtot$ to $\\Rsun$, but further information is required to obtain both quantities, letting aside the need for an independent measurement. Recently parallaxes to objects in the central Galactic regions have become available \\citep[see the discussion in][]{Reid09}, and values for the Galactic circular speed have been derived from $HI$ motions and from MASERs \\citep[][]{Rygl10}. Despite the decent errors in the determined kinematics of MASERs, the small sample sizes impose considerable systematic uncertainties: they are not on pure circular orbits and more importantly they are intimately connected to the intense star formation in spiral arms, where the kinematic distortions are largest. As the youngest strategy, studies of orbits in the Galactic centre have gained high precision, but weakly constrain $R_0$ due to a strong degeneracy with the black hole mass \\citep[][]{Ghez09}. Hence the values for $\\Rsun$ and $\\Vtot$ remain under debate. An independent determination of Galactic parameters is facilitated by the new large spectroscopic surveys like RAVE \\citep[][]{RAVEI} and SEGUE \\citep[][]{Yanny09}. I will show that already now the stellar samples, which so far have been primarily used for the exploration of substructure \\citep[][]{Belokurov07, Hahn11} give results for $V_c$ competitive with radio observations. Common ways to determine Galactic parameters from the motion of stars require a significant bundle of assumptions and modelling efforts to evaluate the asymmetric drift in a subpopulation compared to the velocity dispersion and other measurements and assumptions, e.g. detailed angular momentum and energy distributions, the validity of the theoretical approximations, etc. On the contrary a good measurement strategy for Galactic parameters should have the following properties: \\begin{itemize} \\item Do not rely on dispersions and other quantities that require accurate knowledge of measurement errors. \\item Do not rely on kinematically cold objects in the disc plane, which are particularly prone to perturbations from Galactic structure. \\item Do not rely on any specific models with hidden assumptions and parameters and require as few assumptions as possible. \\end{itemize} In an attempt to approximate these conditions this note concentrates on the use of kinematically hot (thick disc, halo) populations. The use of mean motions and the assumption of approximate axisymmetry in our Galaxy will be sufficient to fix the Galactocentric radius of the Sun as well as the Solar motion and Galactic circular speed. In Section \\ref{sec:p7general} I lay out the general method before describing the sample selection and treatment in Section \\ref{sec:p7sample}. The latter comprises a discussion of proper motion systematics, a discussion of distances in subsection \\ref{sec:p7dist} and a discussion of line-of-sight velocities and the vertical motion of the Sun. In Section \\ref{sec:p7globrot} I lay out the radial velocity based rotation measurement using SEGUE, discuss the radial velocity of the Sun in Subsection \\ref{sec:p7LSR} and shortly discuss possible trends. From \\ref{sec:p7Galparam} on I present three methods to to extract Galactic parameters from stellar samples, the first two methods relying on the radial velocity based rotation measurement and the third on the mean direction of stellar motions. In Section \\ref{sec:p7conclude} I summarise the results. ", "conclusions": "\\label{sec:p7conclude} The most important outcomes of this work are three modelling-free and simple estimators for the Solar azimuthal velocity and hence the local circular speed $V_c$ and for the Solar Galactocentric radius $\\Rsun$. On this course I developed the idea that in a spatially extended sample the absolute rotation of stellar components can be measured from systematic streaming in the heliocentric radial direction. The stars on one side of the Galactic centre show an opposite heliocentric $U$ velocity to those on the other side. This value has a lower formal precision than the classically used azimuthal velocity, but can boost accuracy by its relative independence from assumptions about the velocity of the Sun. The rotation in any extended sample severely affects determinations of the Local Standard of Rest: $\\Usun$ and $\\Wsun$ are frequently determined via simple sample averages in each component. The rotation bias in $\\Usun$ gets more important with increasing distance and impacts all presently available big surveys, since they are lopsided, i.e. asymmetric in Galactic longitude. For SEGUE dwarfs it amounts to $\\sim 10 \\kms$. Accounting for rotation the otherwise observed difference between disc and halo stars disappears and combining this classic method with the estimate from the mean direction of motion described below, I find $\\Usun = (14.0 \\pm 0.3)\\kms$ with an additional systematic uncertainty of about $1.5 \\kms$. The value is $\\sim 3 \\kms$ larger than from the Geneva-Copenhagen Survey, but still within the error margin. While the GCS is clearly affected by stellar streams, the presented value may be distorted by the problematic Sloan proper motions and possibly residual distance errors. While there could of course be some interesting physics involved, a systematic difference of $\\sim 4 \\kms$ in the average $W$ motion between cones towards the Galactic North and South Poles points to a systematic error in the line-of-sight velocities by $\\sim 2 \\kms$. This is not implausible in light of the adhoc shift of $7.3 \\kms$ applied in \\citep[][]{SloanDR6,SloanDR8}. The correction reconciles $\\Wsun$ to reasonable agreement with Hipparcos and the Geneva-Copenhagen Survey \\citep[][]{Holmberg09,AB09} albeit at a lower value of about $6 \\kms$. Comparison of the absolute rotation measure $\\theta$ based on heliocentric $U$ velocities to the mean azimuthal velocities $V_g$ in a sample delivers the solar azimuthal velocity $\\Vtot$. This measurement is correlated with the assumed Galactocentric radius $\\Rsun$. Combining this relation with another datum like the proper motion of Sgr $A^{*}$ one can determine both $\\Rsun$ and $\\Vtot$. For DR8 I obtain $\\Rsun = 8.26^{+0.37}_{-0.33} \\kpc$ and $\\Vtot = 250^{+11}_{-10} \\kms$. By dissecting the sample via metallicity into slow and fast rotating subgroups I can independently infer the Galactocentric radius from their comparison: A larger $\\Rsun$ reduces $\\alpha$ and hence nearly proportionally increases $\\theta$. Thus fast rotators experience a larger absolute change in the rotation speed $\\theta$, putting their value of $\\Vtot$ at odds with that from the slow rotators when $\\Rsun$ is wrong. Enforcing consistency provides $\\Rsun = 8.29^{+0.63}_{-0.54} \\kpc$, and in combination with the simple rotation measure and the proper motion of Sgr $A^{*}$ I get $\\Rsun = (8.27 \\pm 0.29) \\kpc$ and $\\Vtot = (250 \\pm 9) \\kms$ in excellent agreement with the values from \\cite{McMillan11} or \\cite{Gillessen09}. The circular speed $V_c = \\Vtot - \\Vsun$ using the LSR value of $\\Vsun = (12.24 \\pm 0.47 \\pm 2) \\kms$ from \\cite{SBD} is then $V_c = (238 \\pm 9) \\kms$. The third approach uses just the direction of motion in the Galaxy to estimate $\\Rsun$. Assuming near axisymmetry again the direction of motion in a sample should always point in the direction of the azimuth. Fitting the Galactic angle throughout the plane to the angle of motion provides $\\Rsun$. This measurement displays a strong dependence on the solar radial motion apart from the dependence on $\\Vtot$. Fortunately $\\Usun$ and $\\Rsun$ are not degenerate and I obtain again a relatively large $\\Usun = (13.84 \\pm 0.27) \\kms$, confirming the previous conventional analysis, but with an even smaller formal error. As the resulting relationship between $\\Vtot$ and $\\Rsun$ is only weakly inclined against the result from Sgr $A^{*}$, this approach will be of relevance for larger samples and here just provides some reassurance. The laid out strategies are extremely simple. They do not rely on any modelling with possible hidden assumptions, so that possible biases are readily understood. They rely on the sole assumption of near axisymmetry of the Galactic disc. They are designed for thick disc stars that reside high above the Galactic plane and are by their position and high random energy far less affected by perturbations of the Galactic potential, like the spiral pattern. Their modelling independence is supported by the fact that by construction none of the estimators depends on the detailed changes of sample composition with position in the Galaxy or on the radial dependence of $V_c$. Still some uncertainties must be pointed out: The methods are vulnerable to large-scale distortions from axisymmetry in the Galactic disc. The bar region should be avoided, and there may be a residual signal from spiral arms. I could not detect significant structures connected to this, so the consequences should be rather small, particularly as I stay out of the region dominated by the bar and as the sample is sufficiently extended not to cover just one side of a spiral arm. The currently available data add to the possible biases by their uncertain distances, radial velocities and especially astrometry. The Segue proper motions display a distinctive systematic pattern on quasar samples. While I use the quasar catalogues for corrections, the solution remains unsatisfactory: I could not quantify separately the physical reasons, i.e. chromatic aberrations, astrometric ``frame-dragging'' by stars mistaken for Galaxies or even focal plane distortions of the telescopes, and also there might be some undetected dependence on the stellar colours. Further some minor uncertainty in parameters is caused by possible biases in line-of-sight velocities. I made extensive use of the distance corrections developed by \\cite{SBA}. Besides that this project would have been futile without the accuracy achieved by the statistical corrections, the outcome is prone to all the weaknesses of that method. Especially streams and wrong assumptions about the velocity ellipsoid can induce systematic distance errors of order $5 \\%$. The contamination in the sample will vary to some extent with Galactic coordinates, an effect that I did not model. This can lead to a bias, because the distribution of weights in the rotation estimators and the distance estimator is different. The accuracy could be far better could I use more stars and had I better control of the systematic errors by another dataset. Part of this will be resolved by use of additional samples like RAVE. Fortunately the estimators show very different behaviour on the different biases: The proper motion problem affects almost exclusively the first estimator for rotation with a correction to $\\Vtot$ of $\\sim 6 \\kms$, while the difference between fast and slow rotators is nearly untouched. Vice versa, the difference estimator reacts strongly to distance changes, while the first estimator does not - a $10 \\%$ change shifts it by $\\sim 2 \\kms$. Despite their benign distribution the systematic uncertainties on $\\Vtot$ and $\\Rsun$ may come close to the formal errors. If the reader should take only one point from this note then let it be this: With the advent of the large surveys stellar samples are regaining their place as a primary source to obtain not only the Local Standard of Rest, but global Galactic parameters. Already a sample of $\\sim 50000$ stars from SEGUE provides formal accuracies for the galactocentric radius and the solar azimuthal velocity that are competitive with any other known approach and without any need for modelling." }, "1207/1207.3799.txt": { "abstract": "% % Text of abstract % % Recent dynamical studies have identified pairs of asteroids that reside in nearly identical heliocentric orbits. Possible formation scenarios for these systems include dissociation of binary asteroids, collisional disruption of a single parent body, or spin-up and rotational fission of a rubble-pile. Aside from detailed dynamical analyses and measurement of rotational light curves, little work has been done to investigate the colors or spectra of these unusual objects. A photometric and spectroscopic survey was conducted to determine the reflectance properties of asteroid pairs. New observations were obtained for a total of 34 individual asteroids. Additional photometric measurements were retrieved from the Sloan Digital Sky Survey Moving Object Catalog. Colors or spectra for a total of 42 pair components are presented here. The main findings of this work are: (1) the components in the observed pair systems have the same colors within the uncertainties of this survey, and (2) the color distribution of asteroid pairs appears indistinguishable from that of all Main Belt asteroids. These findings support a scenario of pair formation from a common progenitor and suggest that pair formation is likely a compositionally independent process. In agreement with previous studies, this is most consistent with an origin via binary disruption and/or rotational fission. ", "introduction": "} Analyses of osculating orbital elements of Main Belt asteroids have revealed over 80 pairs of asteroids that reside in nearly identical heliocentric orbits \\citep{Vok08,Pravec09,Rozek11}. These objects are distinct from binary asteroids as they are not on bound orbits around a common center of mass, and it is unlikely that their proximity is due to random fluctuations of asteroid densities in orbital element space. Backwards integration of these pairs' heliocentric orbits suggests they may have separated recently into an unbound state, in some cases much less than a Myr ago \\citep{Vok09,Vok09b,Pravec10,Vok11,Duddy12}. As such these are interesting objects for studying phenomena, such as space weathering and radiation pressure forces, that are relevant to the ongoing dynamical, physical and chemical evolution of Main Belt asteroids. The components of known pairs are typically a few km in size and consist of a primary and a secondary (respectively defined as the larger and smaller components based on measured absolute magnitudes). One formation scenario for these systems \\citep{Scheeres07,Pravec10} involves parent asteroids that were spun up to a critical frequency by the YORP effect, i.e. a change in angular momentum due to anisotropic emission of thermal photons \\citep{Rubincam00,Bottke06}. At this critical frequency the parent would fission into a proto-binary system and eventually disrupt under its own internal dynamics to form an unbound asteroid pair \\citep{Jacobson11}. The estimated size ratios and observed rotational properties of known pair systems are consistent with a formation scenario via rotational fission \\citep{Pravec10}. Progressive mass shedding due to YORP spin-up and accretion of a dynamically unstable proto-satellite offers a similar pathway to pair formation \\citep{Walsh08}. The close spectral similarity between components in one pair system \\citep{Duddy12} and the photometric similarity between components in another \\citep{Willman10} support these scenarios. Though a fission origin is consistent with the size ratios and rotation properties for a large number of pairs \\citep{Pravec10}, collisions provide another possible formation mechanism that may explain a subset of systems. In this scenario a catastrophic collision would produce a distribution of fragments, of which only the largest two are observed as an associated pair. Collisions between small bodies can result in compositionally complex outcomes \\citep[e.g.][]{Leinhardt09}, but it is unclear how collisional formation would affect the relative colors and/or spectra of the km-scale objects found in pair systems. Aside from any spectroscopic or photometric implications, hydrodynamic simulations of impact events make predictions about the resulting orbital properties and size ratios of collisional fragments \\citep{Nesvorny06,Durda07}. Unfortunately, due to incompleteness for sub-km bodies in the Main Belt, it is currently not possible to fully test these predictions \\citep{Vok08}. A final formation mechanism involves the dynamical dissociation of bound binary systems. Perhaps the best evidence for ongoing pair formation via binary disruption comes from the system of asteroids associated with 3749 Balam. Adaptive optics and light curve observations \\citep{Merline02,Marchis08a,Marchis08b,Polishook11} have shown that Balam has two bound satellites, making it a rare triple system. One of its two satellites is on a highly eccentric orbit ($e\\sim0.9$), while the other orbits at a distance of only 20 km (Balam itself is about 5-10 km in size). In addition, this triple system has a dynamically associated pair, asteroid (312497) 2009 BR60. Backwards numerical integrations show that the orbits of Balam and 312497 converge within the past 0.5 Myr \\citep{Vok09b}. Numerical models suggest that a cascade of fragments, like that seen in the Balam system, can result from repeated rotational fission events \\citep{Scheeres11}. It seems in this case that YORP fission and binary dissociation may be closely related processes. However, it is unclear how the compositions of asteroid pairs would reflect an origin due to the dissociation of binaries. Component-resolved spectra have been obtained for a very small number of bound binary systems \\citep[e.g.][]{Polishook09,Marchis11,DeMeo11}. This is in part due to the need for either adaptive optics or space based observations to resolve the individual components. Available spectra suggest that the components of binaries are compositionally similar, however data is scarce and more information is required before generalized statements can be made. In light of these various formation mechanisms, there remain several unaddressed questions regarding the formation and evolution of asteroid pairs. For instance, the relationship between pairs and bound multi-component systems is unclear. In addition, little is known about their compositions/taxonomic types. Different formation mechanisms could affect the relative compositions of pair components in different ways. However, for each of these mechanisms the physics of separation should predominantly depend on the internal (rubble pile) structure and density of the parent asteroid. Objects in the size range of asteroid pairs are expected to be rubble piles \\citep{Pravec00}. It is possible that a variety of formation mechanisms are responsible for the formation of the ensemble population of asteroid pairs. Reflectance spectroscopy or photometric colors may help to provide arguments regarding the mechanism of formation on a case-by-case basis. Here we present a survey of asteroid pairs to constrain their spectro-photometric properties. The observations, data reduction and error analysis are presented in \\S2. This data set provides the means to investigate the relative reflectance properties of pair primaries and secondaries, and facilitates a direct comparison to the color distribution of ordinary Main Belt asteroids (\\S3). The results and implications of this survey are discussed in \\S4. ", "conclusions": "} We have presented results and analysis of a spectro-photometric survey of dynamically associated asteroid pairs. A combination of new observations and archival data from the SDSS MOC have provided insight on the reflectance properties of 44 individual asteroids in 30 pair systems and one spurious pair. Data were obtained for both components in 12 pair systems. These data suggest a correlation between the colors of primary and secondary components (Fig. \\ref{fig.pca}) at greater than 98\\% significance. We suggest this argues in favor of a common origin for these pairs. The components in one of the observed systems, 34380-216177, have significantly different colors (Fig. \\ref{fig.34380}). However, updated dynamical integrations have revealed that this is a spurious pair (P. Pravec, private communication, April 2012) and thus should not have been included in our survey. This highlights the need for future follow-up observations of other dynamically identified pairs. Asteroids 69142-127502 are a second pair whose $a^*$ colors are not the same within the estimated $\\pm0.1$ magnitude systematic uncertainties of this survey. However, within the rather large photometric errors for these objects, the spectral profiles and $a^*$ colors are the same (Fig. \\ref{fig.69142}). The difference in filter sets used to observe this system may be the primary cause of its discrepant $a^*$ colors. Follow-up spectroscopy could confirm or refute any taxonomic or compositional link between these objects. The results from several other studies can be used to compare pair reflectance properties. \\citet{Duddy12} showed that the components in the pair 7343-154634 have very similar reflectance spectra. Data from \\citet{Masiero11} show that the albedos of the pair 38395-141513, as determined by observations from the Wide-field Infrared Survey Explorer mission, are nearly indistinguishable with values of 0.0638 and 0.0623 respectively. This pair was included in our sample and has $a^*$ colors that differ by less than 0.03 magnitudes. These additional results support a compositional link between components and thus a common origin for pair systems. Further comparison can be made to the survey of \\citet{Ye11}. As part of a larger sample they observed 12 asteroids in 10 pair systems with data collected for two complete pairs (1979-13732 and 11842-228747). Unfortunately the data for one component in each of the completed systems were unreliable due to instrumental problems in one case and proximity to a bright field star in the second. As such it is difficult to draw conclusions regarding the relative colors of the components in these two systems. Four of the asteroids discussed here (2110, 4765, 15107, and 54041) were also part of the \\citet{Ye11} survey. With two exceptions the data agree within the error bars. The $V-I$ colors for 15107 are significantly different: we measured $V-I=0.77 \\pm 0.06$ whereas \\citet{Ye11} measured $V-I=1.016 \\pm 0.021$. The cause of this offset is not clear, but we note that our measured $a^*$ colors for 15107 and its companion 291188 are identical. The second discrepancy is for asteroid 4765: the data from \\citet{Ye11} suggest $a^*=0.06$ while SDSS MOC data suggest $a^*=-0.07$. Follow-up observations would help to clarify this inconsistency. We have also shown that the $a^*$ distribution of pairs is similar to that of all Main Belt asteroids (Fig. \\ref{fig.histcomp}). There appears to be no bias towards a single taxonomic complex. This strongly suggests that formation of pairs is independent of composition, and instead depends solely on the mechanical properties of the parent bodies. This is consistent with the findings of \\citet{Pravec10}. Taken as a whole our results are most consistent with pair formation via rotational fissioning and/or binary disruption. It is expected that a collisional formation between compositionally distinct bodies would produce at least some primaries and secondaries with disparate colors, though this presumption should be numerically investigated in detail. It is unclear how a collisional formation scenario would influence the color distribution of pairs in Figure \\ref{fig.histcomp}. Density differences between C- and S-complex asteroids \\citep{Britt02} might be a reason for expecting different pair formation efficiencies from disruptive collisions. Several avenues for future work would help to further constrain the origin of these objects. Spectra or photometric colors of the components in binary systems could determine whether binary disruption can produce a population of pairs whose primaries and secondaries have similar reflectance properties. New models that address the compositional implications of pair formation via rotational fission and via collisions would be useful. Additional spectroscopic observations (particularly at near-infrared wavelengths) could provide further insight into the composition, extent of weathering and surface properties of these interesting systems. %% Using an acknowledgements command is not in the Elsevier template, %% but it can be used. \\ack I would like to thank Scott Sheppard and Mark Willman for their assistance with observing several of the objects presented here and in their helpful comments on early drafts of this manuscript. Thoughtful comments on the manuscript were also provided by David Polishook. Insightful reviews were kindly provided by David Vokrouhlick\\'y and an anonymous referee. This work includes data obtained at the Magellan 6.5m and DuPont 2.5m telescopes located at Las Campanas Observatory in Chile, and at the University of Hawaii 2.2m telescope located on Mauna Kea in Hawaii. Support for this project was provided by the Carnegie Institution of Washington and by the National Aeronautics and Space Administration through the NASA Astrobiology Institute (NAI) under Cooperative Agreement No. NNA04CC09A. \\label{lastpage} % Bibliographic references with the natbib package: % Parenthetical: \\citep{Bai92} produces (Bailyn 1992). % Textual: \\citet{Bai95} produces Bailyn et al. (1995). % An affix and part of a reference: % \\citep[e.g.][Ch. 2]{Bar76} % produces (e.g. Barnes et al. 1976, Ch. 2).- %" }, "1207/1207.6648_arXiv.txt": { "abstract": "Black hole accretion disks can form through the collapse of rotating massive stars. These disks produce large numbers of neutrinos and antineutrinos of electron flavor that can influence energetics and nucleosynthesis. Neutrinos are produced in sufficient numbers that, after they are emitted, they can undergo flavor transformation facilitated by the neutrino self interaction. We show that some of the neutrino flavor transformation phenomenology for accretion disks is similar to that of the supernova case, but also, we find the disk geometry lends itself to different transformation behaviors. These transformations strongly influence the nucleosynthetic outcome of disk winds. ", "introduction": "\\label{Section: Introduction} Accretion disks are a compelling astrophysical setting, the properties of which are still being understood. Sophisticated numerical models have examined the origins of the disks and their structure. The disks may arise out of mergers between neutron stars \\cite{Eichler:1989ve, Ruffert:2001gf}, between a neutron star and a black hole \\cite{Narayan:1992iy}, or as a result of some stellar collapses \\cite{Paczynski:1997yg, MacFadyen:1998vz, MacFadyen:1999mk}. It has been suggested that accretion disks could play an important role in the production of gamma ray bursts \\cite{Ruffert:1998qg, Meszaros:2001vi} and they have also been studied as possible sites for nucleosynthesis in the absence of neutrino oscillations \\cite{Pruet:2002ky, Pruet:2003yn, Surman:2003qt, Surman:2008qf, Kizivat:2010ea, Metzger:2010sy,Caballero:2011dw, Wanajo:2011vy}. Disks that originate from stellar collapse with sufficiently high accretion rates are understood to have a region of trapped neutrinos with mean energies of tens of MeV. These disks emit primarily electron and anti-electron type neutrinos and relatively little mu and tau type. The emission surface of the neutrinos generally exceeds that of the antineutrinos, but the antineutrinos have higher temperatures \\cite{Matteo:2002ck, Surman:2003qt, Chen:2006rra}. We focus on these features of the stellar collapse case. Beginning in the early 1990s, it was realized \\cite{Pantaleone:1992eq} that coherent forward scattering of neutrinos could impact neutrino oscillation, resulting in coherence between neutrinos of different energies and parametric resonances \\cite{Samuel:1993uw}. Therefore, we expect that neutrino-neutrino interactions will play an important role in neutrino oscillations above disks. Neutrino oscillations involving neutrino-neutrino interactions have been extensively studied in general and in the contexts of the early universe and supernovae \\cite{Balantekin:2006tg, Barbieri:1990vx, Qian:1995ua, % Fogli:2007bk, Gava:2009pj, Galais:2009wi, Duan:2007fw, Duan:2010bf, Duan:2010bg,Duan:2009cd, EstebanPretel:2007bz, Duan:2006an, Fuller:2005ae,Banerjee:2011fj, Sawyer:2008zs}. When restricted to two flavors, the numerical results in regions of high neutrino density have been analyzed in terms of an analogy to precession and nutation of spins in a magnetic field \\cite{Pastor:2001iu,Duan:2006jv,Hannestad:2006nj}. This analogy is called the Neutrino Flavor Isospin (NFIS) picture. Another approach was to use the techniques of BCS theory \\cite{Pehlivan:2011hp}. The early calculations assumed that interacting neutrinos shared a common history, in what is known as the single angle approximation \\cite{Qian:1994wh}. Calculations that compute the histories of neutrinos across many trajectories resulted in effects not seen in the single angle approximation, including the decoherence of different energy modes \\cite{Duan:2005cp,EstebanPretel:2007ec}. Numerical calculations \\cite{Dasgupta:2009mg, Dasgupta:2010cd, Fogli:2009rd, Duan:2008za} showed that oscillations split the neutrino spectra among flavors at discrete energies. These so-called spectral splits depend on the adiabaticity of the oscillations \\cite{Raffelt:2007xt,Pehlivan:2011hp}. They can be understood in terms of the NFIS picture as a magnetic resonance phenomenon \\cite{Galais:2011gh}. Neutrino-neutrino interactions above disks were studied by Dasgupta et al. \\cite{Dasgupta:2008cu}. They generalized self-maintained coherence behavior --where neutrinos of different energies oscillate as a single mode-- to a disk geometry. For their disk geometry, Dasgupta et al. studied two identical, circular disks, without black holes at the center, emitting both neutrinos and antineutrinos, with number densities fixed to a constant ratio, so that neutrinos dominated the self-interaction term. Coherent forward scattering of neutrinos on electrons (matter) was included by assuming a small mixing angle. With this model, a single ``nutation'' region, similar to that found in supernovae, was found when the neutrino densities became sufficiently low. In this paper, we create disk models that reflect the qualitative understanding we have of black hole accretion disks that originate from stellar collapse and include an explicit matter interaction term. We choose our neutrino and antineutrino disks to have different temperatures and different radii, so that the ratio between the neutrino flux densities will not be fixed, but will vary with position. Typically, the antineutrino disk will be hotter than the neutrino disk, but smaller. Therefore, there are two basic relationships the densities may have above the disks. Neutrinos may always dominate or antineutrinos dominate near the disk while neutrinos dominate further away. Our disks have ``holes'' in the center, from which no neutrinos or antineutrinos are emitted. This region corresponds to the space close to the black hole within its last stable orbit. To these disks, we add a nontrivial electron density. We also consider the role that neutrino flavor transformation plays in nucleosynthesis. The impact of matter-enhanced neutrino flavor transformation core collapse supernovae nucleosynthesis has been considered in e.g. \\cite{Yoshida:2006qz, Yoshida:2006sk, Qian:1993dg, Fuller:1998kb} and the impact of self interaction on supernova nucleosynthesis has been considered in \\cite{Duan:2010af, Chakraborty:2009ej}. However, the impact of flavor transformation in disk nucleosynthesis has not been previously evaluated. In section \\ref{sec:disk-model} we outline our two general disk models and in section \\ref{neutrinos} we outline our method. In section \\ref{calculation}, we perform a single angle numerical computation of three flavor neutrino oscillations in these disks. In section \\ref{analysis}, we analyze the results, and find a new type of transition region in accretion disks. In section \\ref{nucleosynthesis} we discuss the ramifications for nucleosynthesis and in section \\ref{conclusions} we conclude. ", "conclusions": "We have studied models of neutrino and antineutrino emission for accretion disks that encapsulate the qualitative behavior of the neutrino fluxes leaving the disk. Two models were examined in detail: those dominated by neutrinos (model A) and those that begin dominated by antineutrinos and end up dominated by neutrinos (model B). The neutrino dominated disks in the normal hierarchy result in little flavor transition until the neutrino interaction strengths become close to the vacuum strengths. In both hierarchies they exhibit oscillations in type (\\ref{nutations}) (nutation/bipolar) regions. On the other hand, disks that begin antineutrino dominated and end up neutrino dominated produce large flavor transition when the neutrinos flux is about the same as the antineutrino flux. These transitions are associated with the cancellation of the neutrino and antineutrino terms with the neutrino-matter interaction strength, i.e. a type (\\ref{cancellationOscillations}) (matter-neutrino enhanced) region. The calculations described here can be expanded to more complex scenarios. We considered disks of a single temperature, but one should consider also disks with a temperature distribution such that hotter neutrinos are emitted at the center and cooler neutrinos are emitted at the edges. The expected effect would be to shift the interesting transformation behavior nearer to the disk. We performed our calculations in the single angle approximation, but it would be worthwhile to expand this to multi-angle scenario. Based on the arguments in \\cite{Duan:2010bg} we expect that single angle calculations will work well as a description of type (\\ref{nutations}) transitions just as in the supernova case, although there will be some situations akin to those studied in \\cite{Duan:2010bf} when multiangle calculations are necessary. Type (\\ref{cancellationOscillations}) transitions have not been studied from this perspective before. We compared the position of the type (\\ref{cancellationOscillations}) region for the radial neutrino with those coming from various positions on the disk and find that this region is at a similar position for all neutrinos. However, about 30\\% of these neutrinos have multiple type (\\ref{nutations}) regions. Thus multi-angle calculations are warranted. Finally, halo effects have been suggested as a mechanism to alter the simple picture of type \\ref{nutations} transformations \\cite{Cherry:2012zw}, although no complete calculations exist yet. Such effects may influence accretion disk neutrinos as well. Transitions close to the disk, like those we see in model A, are particularly important because they occur at a time when the neutrinos are influencing nucleosynthesis. Using a disk which approximates the type of disk found in a ``collapsar'' scenario \\cite{MacFadyen:1998vz}, i.e. one that has trapped electron neutrinos and antineutrinos only, with different sized trapping regions, we find that the addition of neutrino oscillations enables the formation of $r$-process elements. These transitions will typically occur close to the disk, where the neutron to proton ratio is being set. The removal of the electron neutrinos as a consequence of this transition, allows the neutron to proton ratio to remain sufficiently high to allow the production of the r-process elements. This effect should be typical of disks that have type (\\ref{cancellationOscillations}) transitions." }, "1207/1207.1774_arXiv.txt": { "abstract": "{Using an \\emph{empirical} description of a prompt GRB pulse, we analyze the individual pulses of all Fermi/GBM GRBs with known redshifts, till July 2009. This description is simultaneous in time and energy and allows one to determine the peak energy of Band spectrum at zero fluence ($E_{peak,0}$). We demonstrate, for the first time, that the $E_{peak,0}$ bears a very strong correlation with the isotropic energy of the individual pulses, and hence, each pulse can be used as a luminosity indicator. As a physical description is needed in order to use GRB pulses for cosmological purposes, we explore other physical spectral models. As pulses are the building blocks of a GRB, we choose another sample of Fermi/GBM GRBs having bright, long and single/ separable pulse(s) and fit the time-resolved spectra of the individual pulses with the Band model and a model consisting of a blackbody and a power-law. Both these models give acceptable fits. We find that the peak energy/ temperature always decreases exponentially with fluence in the later part of a pulse. We investigate multiple spectral components in the initial rising part and provide a comprehensive empirical description of the spectral and timing behaviour of prompt GRB pulses. This work strongly extends the possibility of using GRB pulses as standard candles and the spectral parameters as proxy for redshift.} \\FullConference{Gamma-Ray Bursts 2012 Conference -GRB2012,\\\\ May 07-11, 2012\\\\ Munich, Germany} \\begin{document} ", "introduction": "The spectral model suggested by Band et al. (1993) gives very good fits to the GRB spectra obtained from a variety of instruments, but this model is difficult to reconcile with any physical scenario, like synchrotron emission. In recent times, researchers have also reported multiple spectral components while fitting the prompt emission spectra in a wider band (Zhang et al. 2011) --- the Band model is not sufficient to capture the whole electromagnetic spectrum. Ryde (2004) have shown that instantaneous spectra of GRB pulses sometimes disagree with the Band model, but can be described by a combination of a blackbody and a power-law (BBPL). The temperature of the BBPL model evolves with time. Ryde et al. (2010) have also found that the time resolved spectra of GRB 090902B do not agree even with the BBPL model, but a multi-colour BB with a power-law fits the data. Hence, no unified picture of GRB prompt spectrum has emerged yet. It is very important to obtain a correct description of the GRB spectrum, not only to understand the emission mechanism, but, a comprehensive physical description is essential to use GRBs for cosmological purposes. Basak \\& Rao (2012a) have shown that the individual \\emph{pulses} of a GRB can be parametrized by a set of variables, using an empirical law of hard-to-soft spectral evolution. Basak \\& Rao (2012b) have used this model to find peak energy at the very beginning of the pulses ($E_{peak,0}$) of a set of Fermi/GBM GRBs with measured redshifts. They have shown that $E_{peak,0}$ shows a very strong correlation with the isotropic energy ($E_{\\gamma,iso}$) of the pulses. This shows that the individual pulses of a GRB can be separately treated, and each of them can be used as a standard candle, instead of the full GRB. The pulse description in Basak \\& Rao (2012b), however, is an empirical one. It is worthwhile to find whether other physical model(s) is (are) consistent with the data. The realization that the pulses are the building blocks of a GRB also motivates one to try various models and spectral evolution in a single pulse. The strategy is to establish a comprehensive description for a single pulse, then use them for more complicated pulses of GRBs (having known redshifts), and use these individual pulses as standard candles in cosmology. We collect such a set of 11 GRBs from the catalogue of Nava et al. (2011). In this paper we present the main conclusions of the early works (Basak \\& Rao 2012a; b) and then focus on the new findings. The detailed results will be published elsewhere (Basak \\& Rao, in preparation; Rao et al. ApJ submitted). The final aim is to use these comprehensive descriptions to parametrize prompt GRB pulses by physical models. ", "conclusions": "We summarize our work as follows: (a) A simultaneous timing and spectral description of GRB pulses has been developed, which correctly predicts derived parameters, e.g., width, spectral lag. (b) One of the major outcome of our model is finding a better correlation of $E_{peak,0}$, compared to $E_{peak}$, with $E_{\\gamma,iso}$. Hence we conclude that $E_{peak,0}$ is the correct parameter to use for Amati-type correlation study. (c) Alternative spectral model, e.g., BBPL, is consistent with the data. The temperature of the BB falls exponentially with the running fluence similar to the $E_{peak}$ of Band model. (d) If we force double BBs in the initial part of a GRB pulse, then the temperature variation is either smooth throughout or constant till the break. The source of these two blackbodies is speculative. For example, they might be coming from different regions of photosphere boosted by different multiples of the bulk Lorentz factor ($\\Gamma$). Alternately, if they are thermal supernova photons boosted by a cannon ball (CB; Dado et al. 2007), ejected by the central engine, then one identifies one of the components as the thermal inverse Compton and the other might be the bremsstrahlung photons radiated by the electrons. In future, a comprehensive physical model of prompt GRB pulses can be constructed from a detailed calculation of the origin of the observed spectral behaviours. The current model is an empirical one. Once a pulse model is determined, the same can be used for the constituent pulses of a complicated GRB with known redshift. This raises an enormous hope of using GRB pulses as standard candles. The fact that each pulse can be used for this purpose, gives an extra constraint on the derived redshift and other related parameters." }, "1207/1207.3547_arXiv.txt": { "abstract": "We study the compact stars internal structure and observable characteristics alterations due to the quark deconfinement phase transition. To proceed with, we investigate the properties of isospin- asymmetric nuclear matter in the improved relativistic mean-field (RMF) theory, including a scalar-isovector $\\delta$-meson effective field. In order to describe the quark phase, we use the improved version of the MIT bag model, in which the interactions between $u$, $d $ and $s$ quarks inside the bag are taken into account in the one-gluon exchange approximation. We compute the amount of energy released by the corequake for both cases of deconfinement phase transition scenarios, corresponding to the Maxwellian type ordinary first-order phase transition and the phase transition with formation of a mixed quark-hadron phase (Glendenning scenario). ", "introduction": "Neutron stars are objects with a very complicated, many-component and many-layered structure. The surface layer of such objects consists of an ordinary matter with atomic and molecular structure. In the inner layers of different depths, the conditions both for the rearrangement of the structural formations and for the creation of new constituents of matter are satisfied. In the central region of neutron star the density of matter reaches such high values that makes possible the appearance of various exotic particle species and phases such as hyperons, deconfined $u$, $d$, $s$ quarks and $\\pi$, $K$ meson condensates. The existence of compact stars consisting of matter in a deconfined quark phase was predicted long ago\\cite{Iv-Kurd}\\cdash\\cite{Bod}. Over the past few decades many researchers had been intensively studied various aspects related to the formation of exotic degrees of freedom in neutron stars and clarification of the observationally testing of dynamic processes, confirming the existence in the interiors of stars such constituents (for review see, e.g., Refs.~\\refcite{Gl_bk,Haen_bk} and references therein). Phase transitions accompanied by discontinuities of the thermodynamic potentials are the most interesting because it leads to a dynamical rearrangement of neutron stars. Depending on the value of surface tension $\\sigma_{s}$, the phase transition of nuclear matter into quark matter can occur in two scenarios\\cite{Heis,Hj}: either ordinary first order phase transition at constant pressure with a density jump (Maxwell construction), or formation of mixed hadron-quark matter with a continuous variation of pressure and density (Glendenning construction)\\cite{Gl1992}. The question of whether the formation of a mixed phase is energetically favorable, given the finite dimensions of the quark structures inside nuclear matter, the Coulomb interaction, as well as the surface energy, has been examined in Refs.~\\refcite{Hj,Vos,Maruy}. It was shown that the mixed phase is energetically favorable for small values of the surface tension between the quark matter and the nuclear matter. Uncertainty of the surface tension values makes it impossible to determine the phase transition scenario which actually takes place. An important manifestation of the hadron-quark phase transition in the compact stars is a dynamical process of accretion of matter onto the surface of neutron stars with the hadron structure, which leads to fulfillment of conditions for the formation in the center of the star a new phase containing deconfined quarks. This type of transition can also occur in case of a rotating neutron star that is slowing down when the pressure in the center rises and exceeds the threshold value. The process of catastrophic rearrangement with formation of a quark core of finite radius at the star's center will be accompanied by release of a colossal amount of energy comparable to the energy release during a supernova explosion. These features of accreting neutron star near the critical configuration make it a potential candidate both for the gamma-ray bursts (GRBs)\\cite{Bom,Drag} and for the gravitational wave (GW) emission sources\\cite{Marang,Lin}. Note that a similar process of both restructuring and energy release takes place also in the case of pion condensation in the cores of neutron stars\\cite{Mig}\\cdash\\cite{Ber}. Recent series of our articles\\cite{Alav1}\\cdash\\cite{Alav5} were devoted to a detailed investigation of quark deconfinement phase transition of neutron star matter, when the nuclear matter is described in the relativistic mean-field (RMF) theory with the scalar-isovector $\\delta$-meson effective field. The calculation results of the mixed phase structure (Glendenning construction) are compared with the results of a usual first-order phase transition (Maxwell construction). This article is a continuation of this series. Here we aim to investigate the energy release and the change of integral parameters of compact stars due to the quark phase transition in the two alternative scenarios and to identify possible differences in the manifestations of this phenomenon. ", "conclusions": "In this article we have studied the changes of the internal structure and observable characteristics of the near-critical configurations of compact stars and associated energy release due to the quark deconfinement phase transition. For description of isospin-asymmetric nuclear matter we use the RMF theory, including a scalar-isovector $\\delta$-meson effective field. The quark phase is described in frame of the improved version of the MIT bag model. Using the obtained characteristics of hadronic and strange quark matter phases, we calculate the neutron star matter EOS with quark-deconfinement phase transitions, corresponding to the Maxwell and Glendenning scenarios. We find the dependence of conversion energy on the baryonic mass of neutron stars and analyze the changes in stellar radiuses due to the deconfinement phase transition. We show that for a fixed value of the baryonic mass $M_0$ of star the conversion energy in the case of Glendenning construction more than in the case of Maxwell construction. The minimum required baryonic mass for the catastrophic rearrangement of the neutron star and the formation of a quark core in the center of the star greater in the case of the Maxwell construction compared to the Glendenning construction case. In the case considered here, the quark deconfinement phase transition in the neutron star interior leads to the energy release of the order $10^{50}\\div 10^{52}$~erg. Our obtained results will give the opportunity to clarify the observational differences between the two scenarios of quark-deconfinement phase transition and to formulate a specific test for determining the phase transition scenario taking place in reality." }, "1207/1207.5227_arXiv.txt": { "abstract": "In the present work we propose an innovative estimation method for the minimum Doppler factor and energy content of the $\\gamma$-ray emitting region of quasar 3C 279, using a standard proton synchrotron blazar model and the principles of automatic photon quenching. The latter becomes relevant for high enough magnetic fields and results in spontaneous annihilation of $\\gamma$-rays. The absorbed energy is then redistributed into electron-positron pairs and soft radiation. We show that as quenching sets an upper value for the source rest-frame $\\gamma$-ray luminosity, one has, by neccessity, to resort to Doppler factors that lie above a certain value in order to explain the TeV observations. The existence of this lower limit for the Doppler factor has also implications on the energetics of the emitting region. In this aspect, the proposed method can be regarded as an extension of the widely used one for estimating the equipartition magnetic field using radio observations. In our case, the leptonic synchrotron component is replaced by the proton synchrotron emission and the radio by the VHE $\\gamma$-ray observations. We show specifically that one can model the TeV observations by using parameter values that minimize both the energy density and the jet power at the cost of high-values of the Doppler factor. On the other hand, the modelling can also be done by using the minimum possible Doppler factor; this, however, leads to a particle dominated region and high jet power for a wide range of magnetic field values. Despite the fact that we have focused on the case of 3C 279, our analysis can be of relevance to all TeV blazars favoring hadronic modelling that have, moreover, simultaneous X-ray observations. ", "introduction": "Blazars, a subclass of Active Galactic Nuclei, emit non-thermal, highly variable radiation across the whole electromagnetic spectrum. According to the standard scenario, particles accelerate to relativistic energies in the jets of these objects which point, within a small angle, towards our direction and the resulting photon emission is boosted due to relativistic beaming. Detailed modelling of observations, especially in the $\\gamma$- and X-ray regimes, makes possible the estimation of the physical parameters of the emitting region. Thus quantities like the source size, the magnetic field strength, the bulk Lorentz factor and the particle energy density can nowadays be routinely calculated. Furthermore, from the values of these quantities one could obtain meaningful estimates for the particle and Poynting fluxes and use them to make comparisons, for example, to the Eddington luminosity of the source, connecting thus the black hole energetics with the jet power. One major uncertainty of the modelling is the nature of the radiating particles. While there seems to be a consensus that the emission from radio to X-rays comes from the synchrotron radiation of a population of relativistic electrons, there are still open questions regarding the $\\gamma$-ray emission of these objects. Broadly speaking, the models fall into two categories, the leptonic ones (e.g. \\citealt*{dermeretal92, maraschietal92, dermerschlick93}) which assume that the same electrons which radiate at lower frequencies via synchrotron produce the $\\gamma$-rays by inverse Compton scattering and the hadronic ones (\\citealt*{mannheimbiermann93, mueckeprotheroe01, boettcher09}; \\citealt{muecke03}) which postulate that an extra population of relativistic protons produce the high energy radiation as a result of hadronically induced and electromagnetic processes. Due to the very different radiating mechanisms involved, the two classes of models can result in very different parameters for the source. Thus, for example, typical leptonic synchrotron self-compton models require low magnetic field strengths, ranging from $B \\simeq 0.01 - 1$ G for high synchrotron-peaked BL Lacs \\citep{tavecchio11, murase12} up to $B\\simeq 1-10$ G for Flat Spectrum Radio Quasars, and low jet power ($\\simeq 10^{47}$ erg/sec) (e.g. \\citealt{celotti08}). On the other hand, the corresponding values for the hadronic models are higher by at least one order of magnitude (e.g. \\citealt{protheroe01, boettcher09}). While both can fit reasonably well the observations, they both face some problems. For instance, it has been argued that one problem with leptonic models is the high ratio between the required relativistic electron energy density to the magnetic one, implying large departures from equipartition. On the other hand, the hadronic models imply jet powers which can be uncomfortably high, especially when compared to the accretion luminosity. In previous work \\citep{petropoulou12} we have argued that automatic $\\gamma$-ray quenching, i.e. a loop of processes which result in spontaneous $\\gamma-$ray absorption accompanied by production of electron-positron pairs and soft radiation, can become instrumental in the modelling of high energy sources. Its application to any $\\gamma$-ray emitting region, makes it relevant to both leptonic and hadronic models; the fact, however, that quenching requires rather high magnetic fields makes it more relevant for the latter. In the present paper we revisit the hadronic model taking into account the effects of non-linear photon quenching. This, as we shall show, has as a result to exclude many sets of parameters, which otherwise would give good fits to observations, as the non-linear cascade growth modifies drastically the produced multiwavelength spectrum. As a typical example we focus on the quasar 3C 279. This has been detected for the first time in very high-energy (VHE) $\\gamma$-rays at $> 100$ GeV by the MAGIC telescope \\citep{albert08} and ever since it remains the most distant VHE $\\gamma$-ray source with a well measured redshift. Both leptonic and (lepto)hadronic models have been applied \\citep{boettcher09} and it has been pointed out that the former require the system to be well out of equipartition. In the present work we fit only the VHE part of the spectrum using a proton-synchrotron blazar model and use the contemporary X-ray data (see e.g. \\cite{chatterjee08}) as an upper limit; the modelling of the complete multiwavelength spectrum requires a primary leptonic population radiating in the IR/X-ray energy bands which, we assume, that can always be determined. We show that one can derive a lower limit for the Doppler factor of the radiating blob by just combining (i) the values of the proton luminosity and the Doppler factor for a given magnetic field that provide a good fit to the MAGIC data and (ii) the fact that for high values of the proton injection compactness the absorption of $\\gamma$-rays becomes non-linear leading to an overproduction of softer photons and consequently violating the X-ray observations {\\sl even in the absence of a leptonic component}. We also show that the minimum possible value of the Doppler factor is largely independent of the magnetic field and this has some interesting implications for the energetics of 3C 279. The paper is structured as follows. In \\S2 we give some simple, qualitative estimates on the effects of quenching on fiting the multiwavelength data of blazars, in \\S3 we apply these, using a numerical code, to the case of 3C~279 and we conclude in \\S4 with a discussion of our results. For numerical results we adopt working in the $\\Lambda$CDM cosmology with $H_0=70$ km s$^{-1}$ Mpc$^{-1}$, $\\Omega_{\\rm m}=0.3$ and $\\Omega_{\\Lambda}=0.7$. The redshift of 3C 279 $z=0.536$ corresponds to a luminosity distance $D_{\\rm L}=3.08$ Gpc. ", "conclusions": "Hadronic models have been used extensively for fitting the $\\gamma$-ray emission from Active Galactic Nuclei. Usually, detailed fitting to the multiwavelength spectrum of these objects requires, in addition to a population of relativistic protons, the presence of an extra leptonic component which is responsible for the emission at lower energies (radio to UV or X-rays). As it was shown recently \\citep{SK07, PM11,petropoulou12} compact $\\gamma$-ray sources can be subject to photon quenching which can result, if certain conditions are met, to automatic $\\gamma$-ray absorption and redistribution of the absorbed $\\gamma$-ray luminosity to electron-positron pairs and lower energy radiation. This is a non-linear loop which can operate even in the absence of an initial soft photon population and is expected to have direct consequences on the aforementioned models. For the present application other feedback loops, such as the Pair-Production-Synchrotron one \\citep{mastkirk92}, are less relevant as they usually operate at lower proton energies and at higher energy densities \\citep{dimitrakoudisetal12}. Aim of the present paper is to set a general framework for investigating the effects of photon quenching on the parameter space available for modelling $\\gamma$-rays in the context of a hadronic model. As an example we concentrated on the February 2006 observations of the blazar 3C 279 which were performed simultaneously in the TeV and X-ray regimes. To this end we have focused only on fitting the high energy spectrum using a synchrotron proton emission and treating conservatively the X-ray observations as an upper limit. We found that for a wide range of parameters, automatic photon quenching plays a crucial role as its onset produces X-rays which violate the observations. To be able to assess its impact on the parameter space, we have relaxed the usual goodness-of-fit method by accepting fits with $\\chi_{\\rm red}^2<1.5$. Our results indicate, in agreement with other researchers in the field, that the hadronic model requires in general high magnetic fields. Additionally we find that the presence of automatic quenching limits the proton luminosity (or, equivalently the proton energy density in the case of non substantial proton cooling) which, in turn, results in a minimum value of the Doppler factor $\\dmin$. Interestingly enough, these ideas do not apply to leptonic models because they favour much lower values for the magnetic field. For these values photon quenching does not operate, since its feedback criterion is not satisfied. The latter is a necessary condition for the emergence of automatic quenching and requires for a given $\\gamma$-ray energy a certain value of the magnetic field $\\Bq$ , for the absoprtion loop to operate. We have shown both analytically and numerically, (see figures \\ref{dmin-ana} and \\ref{dmin-num} respectively), that the minimum Doppler depends on the magnetic field strength in a different way, depending on the relative relation of $B$ and $\\Bq$. Specifically, if the magnetic field strength is above $\\Bq$, then $\\dmin \\propto B^{-1/7}$. For values of the magnetic field with $ B < \\Bq$ quenching is not relevant. However, we have shown using arguments based on the particles gyroradii, that also in this case a lower limit to the Doppler factor exists, which depends strongly on $B$, i.e. $\\dmin \\propto B^{-3}$. Therefore if one wants to adopt a low magnetic field for the radiating region, fits to the TeV $\\gamma$-rays are still possible but at the expense of a very high value of the Doppler factor. The fact that quenching does not allow the Doppler factor to become smaller than some value is intriguing and leads naturally to the investigation of the proton energy density inside a blob which moves with this characteristic value. In this case, we showed that there are two values of the magnetic field that minimize the energy content, one corresponding to the $\\delta$ branch with $B<\\Bq$ and the other to the one with $B> \\Bq$. For values of the magnetic field between the two equipartition magnetic fields, the emission region is particle dominated and the ratio $\\up / \\ub$ can be as large as $10^3$ (see figure \\ref{densities}). Note however, that if a Doppler factor twice the minimum one is adopted, then the calculated proton energy density for the same magnetic field will be lower by almost an order of magnitude -- see equation (\\ref{Lg}). Furthermore, we have calculated the jet power in the case where $\\delta=\\dmin$ and shown both analytically and numerically that it is a function of $B$. The jet power is rather high ($10^{47} - 10^{49}$ erg/sec) for the whole range of $B$ values, as expected in the context of hadronic modelling. We have repeated the above calculations by relaxing the condition $\\delta=\\dmin$ while requiring the parameters to be such as to minimize the power of the jet. In this case we have shown that the energy content of the blob is also minimized (see figures \\ref{minpjet} and \\ref{upub2}). For adopted $B$ values lying between the two equipartition values, the jet power can be minimized, albeit at the cost of a high $\\delta$ value. However, we have shown that for magnetic field values outside of this range, jet power minimization is not possible as this can be achieved only if $\\delta<\\dmin$. Thus, the existence of a minimum value for $\\delta$ has indirect implications on the energetics of the system. An interesting question is whether a detailed fitting to the X-ray observations of 3C~279 with the addition of an extra leptonic component would bring any change to the basic ideas presented here. In this case apart from automatic quenching, the linear absorption of $\\gamma$-rays on the X-ray photons emitted by the leptonic component, is also at work. However, by including this component and repeating our numerical calculations of \\S3, we found that our results do not change. This is due to the fact that the X-ray luminosity of 3C~279 is rather low and therefore the effects of linear $\\gamma\\gamma$ absorption are minimal, at least up to the compactnesses above which the automatic photon quenching sets in. Apart from the absorption of $\\gamma$-rays on the synchrotron photons emitted by the `extra' leptonic component discussed above, inverse compton scattering of these photons to higher energies by the same leptonic component would be an additional mechanism at work. The upscattered photons would lie in the hard X-ray and $\\gamma$-ray energy range and would affect our calculations only if their luminosity $L_{\\rm ssc}$ would be comparable to that carried by the synchrotron component $L_{\\rm syn}$. While estimating the ratio $u_{\\rm syn}/ \\ub$ for a wide range of $B$ and $\\dmin$ values used in the present work (see figure \\ref{dmin-num}) we find that this is always much smaller than unity, and therefore the inverse Compton scattering of an extra leptonic component which would fit the radio to X-ray observations does not interfere to the quenching mechanism. The estimation method proposed in the present work can be regarded as an extension of the widely used method for estimating the equipartition magnetic field using radio observations. In our case, the leptonic synchrotron component is replaced by the proton synchrotron emission and the radio by the VHE $\\gamma$-ray observations. The innovative feature of our method is the estimation of a minimum Doppler factor, which is the result of automatic photon quenching. This can be more of relevance to TeV observations rather than the Fermi ones as quenching at GeV energies requires very high values of the magnetic field -- see figure \\ref{bfield}. The fact that our numerical results are in very good agreement with the analytical calculations offers a fast yet robust way for estimating many physical quantities of the $\\gamma$-ray emitting region of TeV blazar jets. We caution however the reader, that the effects of automatic quenching, i.e. when no soft photons are initially present in the emitting region, can be seen only in a self-consistent treatment of the radiative transfer problem." }, "1207/1207.5966_arXiv.txt": { "abstract": "Analyzing available photometry from Hipparcos, ASAS, Pi of the sky and Super WASP, we found that the system SY Phe is a detached eclipsing binary with similar components and orbital period about 5.27089~day. It has a slightly eccentric orbit, however the apsidal motion is probably very slow. The system undergoes an additional photometric variation on longer time scales superimposed on the eclipsing light curve. It also contains one distant component, hence the third light was also considered. ", "introduction": "The system SY Phe (= HD 9283 = HIP 7024) was discovered as a variable by \\cite{Hoff1949}. Later, \\cite{1958VeSon...3..333H} noted that the Algol-type curve is of BO~Cep-type, and the additional variability is discussed. The Hipparcos satellite \\citep{HIP} reveals two clearly-shaped eclipses and found the period about 5.27140~day. The photometric amplitude is about 0.5~mag. It is therefore remarkable that SY~Phe is still classified as a \"Variable Star with rapid variations\", according to Simbad, or GCVS \\citep{GCVS}. Spectral type of the system was firstly derived as F8 by \\cite{Spencer}, while \\cite{1958VeSon...3..333H} gave the type F4. Later \\cite{1978mcts.book.....H} noted the type F3/F5V. According to the Tycho data \\citep{1997A&A...323L..57H} the photometric index is $B_T-V_T = 0.512$~mag, and \\cite{2007A&A...474..653V} derived the parallax of the system $\\pi = 4.69 \\pm 1.49$~mas. However, the system consists of two visual components separated about 4$^{\\prime\\prime}$ on the sky. According to the Washington Double Star Catalog (hereafter WDS\\footnote{\\href{http://ad.usno.navy.mil/wds/}{http://ad.usno.navy.mil/wds/}}, \\citealt{WDS}), the astrometric observations of this double do not show any significant change of the position angle. Therefore, the pair is only weakly gravitationally bounded and its semi-major axis is rather large. The Hipparcos observations indicate that the eclipsing variable is the A component (the brighter one). \\begin{figure} \\includegraphics[width=14cm]{SYPhe1.eps} \\caption{Available light curves of SY Phe.} \\label{FigLCs} \\end{figure} ", "conclusions": "The first light curve and period analyses of the system SY Phe were performed. Dealing with no spectroscopy of the target, one has to consider some assumptions and many of the physical parameters are still affected by relatively large errors. However, we can roughly estimate the internal structure constants for both apsidal motion solutions. Assuming the two eclipsing components have masses of about 1.36~M$_\\odot$ (spectral type F4), then the semimajor axis of the orbit is about 17.8~R$_\\odot$, which yields the values of radii of both components (see Table \\ref{TableLC}). Using these values, one can calculate the internal structure constants and compare these values with the theoretical ones, e.g. by \\cite{2004A&A...424..919C}. The Solution I gives $\\log$ k$_2$ value of about 0.686, while Solution II gives $\\log$ k$_2$ = -2.198. The latter value agrees well with the theoretical values by \\cite{2004A&A...424..919C}, assuming the F4 star is on main sequence and is about 5 $\\cdot 10^8$~yr old. The light curve solution also provides the first rough estimation about the third light from the distant component. Our result of the third light value (15\\% of the total light) is in excellent agreement with the $\\Delta$M value provided by the WDS catalogue. The nature of the long-term photometric variations with period about 249 days still remains an open question. Therefore, new more detailed analysis is still needed, especially based on new spectroscopic data together with the photometric observations obtained in various photometric filters." }, "1207/1207.2294_arXiv.txt": { "abstract": "We obtain the non-linear relation between cosmological density and velocity perturbations by examining their joint dynamics in a two dimensional density-velocity divergence phase space. We restrict to spatially flat cosmologies consisting of pressureless matter and non-clustering dark energy characterised by a constant equation of state $w$. Using the spherical top-hat model, we derive the coupled equations that govern the joint evolution of density and velocity perturbations and examine the flow generated by this system. In general, the initial density and velocity are independent, but requiring that the perturbations vanish at the big bang time sets a relation between the two. This traces out a curve in the instantaneous phase space, which we call the `Zel'dovich curve'. We show that this curve acts like an attracting solution for the phase space dynamics and is the desired non-linear extension of the density-velocity divergence relation. We obtain a fitting formula which is a combination of the formulae by Bernardeau and Bilicki \\& Chodorowski, generalised to include the dependence on $w$. We find that as in the linear regime, the explicit dependence on the dark energy equation of state stays weak even in the non-linear regime.Although the result itself is somewhat expected, the new feature of this work is the interpretation of the relation in the phase space picture and the generality of the method. Finally, as an observational implication, we examine the evolution of galaxy cluster profiles using the spherical infall model for different values of $w$. We demonstrate that using only the density or the velocity information to constrain $w$ is subject to degeneracies in other parameters such as $\\sigma_8$ but plotting observations onto the joint density-velocity phase space can help lift this degeneracy. ", "introduction": "The distribution of matter on very large scales ($\\sim 100$ Mpc) is fairly homogenous (e.g., \\citealt{hogg_cosmic_2005}, \\citealt{sarkar_scale_2009}, \\citealt{scrimgeour_wigglez_2012}), but on smaller scales it is far from it. The fractional overdensity ($\\delta$) and the peculiar velocity (${\\bf v}$) are the two variables that characterise this inhomogeneity and are related via the continuity equation, the Euler equation and the law of gravitation. When the inhomogeneities are small, the equations can be linearised and ignoring the decaying mode results in the local relation $\\nabla \\cdot {\\bf v} = -f H\\delta$, where $H$ is the Hubble parameter and $f$ is the linear growth factor. In the theory of gravitational instability, except in the orbit-crossing regions, the peculiar velocity is curl free and the above equation completely characterises the relation between the two fields in the linear regime. $f$ mainly depends on the matter density parameter $\\Omega_m$, and is usually expressed as $\\Omega_m^\\gamma$, where $\\gamma$ is the growth index. Given $\\gamma$, a comparison of the data from redshift and peculiar velocity surveys constrains $\\Omega_m$, or some combination of $\\Omega_m$ and the galaxy bias parameter. Early studies in this field mostly assumed a pure matter universe and used $\\gamma \\approx 0.6$ (\\citealt{peebles_peculiar_1976}), to either get bias independent measures of mass from velocity fields or to constrain $\\Omega_m$ (see review articles by \\citealt{dekel_dynamics_1994} and \\citealt{strauss_density_1995}). The estimate by Peebles, given more than thirty five years ago, turned out to be a good approximation for a large range of models. The dependence of the linear growth rate on the cosmological constant (for e.g., \\citealt{martel_linear_1991}; \\citealt{lahav_dynamical_1991}) or on the dark energy equation of state $w$ (\\citealt{wang_cluster_1998}; \\citealt{linder_cosmic_2005}) was shown to be relatively weak. However, more recently, it has been pointed out that models with similar expansion histories but with different gravitational dynamics can be distinguished by their growth indices; $\\gamma$ is 0.55 for the $\\Lambda$CDM model, but 0.67 for certain modified theories of gravity (\\citealt*{lue_probing_2004}; \\citealt*{bueno_belloso_parametrization_2011}). Accordingly, modern observational efforts are focussed on measuring the linear growth rate with the added aim of constraining $\\gamma$ (for e.g., \\citealt{guzzo_test_2008}; \\citealt{majerotto_probing_2012}). The method of velocity reconstruction from redshift surveys using linear theory breaks down when the perturbations are of order unity. Quasi-linear effects need to be included even to get an accurate determination of the linear growth rate (\\citealt*{nusser_new_2012}). Various analytic quasi-linear and non-linear extensions using the Zel'dovich approximation (\\citealt{nusser_cosmological_1991}; \\citealt{gramann_improved_1993}), second order Eulerian perturbation theory (\\citealt{bernardeau_quasi-gaussian_1992}, hereafter B92), and higher order Lagrangian perturbation theory (\\citealt{gramann_second-order_1993}; \\citealt{chodorowski_weakly_1997}; \\citealt{chodorowski_recovery_1998}; \\citealt{kitaura_estimating_2012}) have been proposed. Many of these analytic methods were also tested with numerical simulations (\\citealt{mancinelli_nonlinear_1993}; \\citealt{mancinelli_local_1995}). N-body codes, although accurate, give mass-weighted instead of volume-weighted estimates (\\citealt*{dekel_potential_1990}) making it difficult to compare them with analytical answers. Refined volume averaged estimates were given using uniform-grid codes (\\citealt{kudlicki_reconstructing_2000}; \\citealt{ciecielg_gaussianity_2003}) and the method of tessellations (\\citealt{bernardeau_new_1996}; \\citealt{bernardeau_omega_1997}, \\citealt{bernardeau_non-linearity_1999}). In general, these simulations are slow and are applicable to a limited range of models. It is usually assumed that the weak dark energy dependence of the linear relation extends to the non-linear regime and often the results are quoted in terms of the scaled velocity divergence $\\nabla \\cdot {\\bf v}/f(\\Omega_m)$, absorbing the cosmology dependence into the linear growth rate. While this assumption has been tested for different values of the cosmological constant $\\Lambda$ (for e.g., \\citealt{lahav_dynamical_1991}; \\citealt{bouchet_perturbative_1995}; \\citealt{nusser_omega_1998}), the explicit dependence on the dark energy equation of state $w$ has not yet been derived. The spherical top-hat system is an alternate way to model evolution in the non-linear regime. The model is simple to solve; equations of motion reduce to ordinary second order differential equations for the evolution of the scale factor (\\citealt{Peebles80}). The main drawback is that it does not take into account interaction between scales; the price paid for computational ease. Nevertheless, it has been successful in predicting non-linear growth until virialization (\\citealt{engineer_non-linear_2000}; \\citealt{shaw_improved_2008}). Recently \\cite{bilicki_velocity-density_2008}, hereafter BC08, obtained the non-linear density-velocity relation from the spherical top-hat. They utilised the top-hat's known exact analytic solutions and hence were restricted to pure matter cosmologies. We adopt a different approach, one that can be generalised to a range of background cosmologies. We numerically investigate how generic density and velocity perturbations evolve in a two dimensional density-velocity divergence phase space. We identify a special curve which is obtained by imposing the condition that the perturbations vanish at the big bang time and show that it acts like an attracting solution for the dynamics of the system. We refer to this as the `Zel'dovich curve' and demonstrate that it is the desired density-velocity relation in the non-linear regime. This approach was first put forward in a recent paper (\\citealt{nadkarni-ghosh_extending_2011}, hereafter NC), but the analysis there was restricted to a EdS (with $\\Omega_m =1$) cosmology. Here we extend the same idea to flat cosmological models with pressureless matter and non-clustering dark energy described by a constant equation of state $w$. The paper is organised as follows. \\S \\ref{sec:dynamics} gives the equations governing the spherical top hat, introduces the `Zel'dovich curve' and demonstrates its importance in the evolution of perturbations in the joint density-velocity phase space. \\S \\ref{sec:fits} gives a fit to the curve by generalising the formulae proposed by B92 and BC08 to include the dark energy term and this results in a new parametrization for the linear growth index $\\gamma$ as a function of $w$. \\S \\ref{sec:cluster} discusses the evolution of galaxy cluster profiles using the spherical infall model and examines this curve in the context of observationally relevant quantities. \\S \\ref{sec:conc} gives the summary and conclusion. ", "conclusions": "\\label{sec:conc} We have obtained the non-linear density-velocity divergence relation by examining the dynamics of perturbations in the joint $\\delta-\\delta_v$ phase space. Although we were restricted to spherical top-hat models, unlike standard practice, we did not use exact solutions of the top-hat in our derivation of the result. Instead, at each epoch we obtained pairs of initial $(\\delta_i, \\delta_{v,i})$ which satisfied the condition `no perturbations at the big bang'. These pairs traced out a curve in the $\\delta -\\delta_v$ phase space which we refer to as the `Zel'dovich curve'. Because of the non-autonomous nature of the $\\delta-\\delta_v$ evolution equations, this curve is a time dependent entity. We demonstrated that the curve acts like an `attractor' for the dynamics of the perturbations. Small amplitude perturbations, as they evolve through the phase space, asymptotically approach this curve and perturbations that start along the curve stay along it with a very high accuracy. Thus, the Zel'dovich curve gives the long term behaviour of density and velocity perturbations and is exact non-linear extension of the density-velocity curve for the spherical top-hat. We obtained a 3\\% fit to the Zel'dovich curve in the range $-1\\leq \\delta \\leq10$ by generalising existing formulae of B92 and BC08 to account for the inclusion of dark energy. From this we obtained a new parametrization for the linear growth index $\\gamma = 0.56(-w)^{-0.08} -0.01 (-w)^{-1.18}$, which agrees with the well known result of \\citet{linder_cosmic_2005} for $\\Lambda$CDM models and deviates by at most $1\\%$ for values of $w$ between $-3/2\\leq w \\leq -1/2$. As expected, the explicit dependence of the $\\delta-\\delta_v$ relation on $w$ is weak, both in the linear regime and in the non-linear regime. Nevertheless there is a implicit dependence on $w$ through the evolution of $\\Omega_m$. As a practical application, we considered the evolution of density and velocity profiles of galaxies in the quasi-linear regime ($\\delta \\sim 5$), before effects of virialization become important, and investigated the dependence on $w$. We found that the deviations in cluster profiles due to a 50\\% change in $w$ are about $\\sim100$ km/s which are barely within the current observational errors in the local universe (\\citealt{courtois_cosmic_2012}, \\citealt{davis_local_2011}). We also demonstrated that the degeneracy in the density and velocity profiles arising due to parameters such as $\\sigma_8$ can be broken if the same information is plotted on the $\\delta-\\delta_v$ plot. The attracting behaviour of the Zel'dovich curve makes it insensitive to small changes in the initial parameters and the points get classified according to $w$. However, this is works best only at redshifts $z\\sim 1$, where the curves for different $w$ have the maximum spread. Given the observational difficulties in isolating peculiar velocities at these redshifts, this result probably is only of theoretical value at the moment. The main new feature of this work is the interpretation of the non-linear density-velocity divergence relation in the phase space picture and the generality of the method. In this paper we focussed only on the simplest possible phenomenological model for dark energy: constant equation of state, not necessarily $-1$. It would be interesting to consider more generalised models such as varying equation of state, coupled dark energy models or alternate models of gravity. Spherical collapse models have already been worked out for various dark energy or quintessence models (for e.g., \\citealt{mota_spherical_2004}; \\citealt{maor_spherical_2007}; \\citealt{pace_spherical_2010}) and modified gravity scenarios (for e.g., \\citealt*{dai_consequences_2008}; \\citealt{schafer_spherical_2008}). The general method remains the same: obtain the correct equation for the evolution of the scale factor, impose the `equal age' condition to get constrains on the initial density and velocity parameters and examine the behaviour of the resulting curve in the phase space dynamics. However several potential caveats and questions may arise. For example in modified gravity scenarios the evolution of a spherical shell enclosing a fixed mass is not independent of the internal distribution of the mass or the environment; Birkoff's theorem does not hold true and spherical top-hats do not remain top-hats as the evolution proceeds (\\citealt{dai_consequences_2008}; \\citealt*{borisov_spherical_2012}). Clearly the $\\delta-\\delta_v$ relation will have to be generalised to treat such situations. It remains to be investigated whether there exist variables related to density and velocity that satisfy a unique relation in phase-space. Recent simulations (\\citealt{schmidt_nonlinear_2009}; \\citealt{lombriser_cluster_2012}) have shown that halo density profiles calculated in $f(R)$ models of modified gravity show an enhancement at a few virial radii when compared to those evaluated with `standard' $\\Lambda$CDM. It would be interesting to investigate if the these features translate to a signature in some appropriate Zel'dovich relation. The Zel'dovich curves were obtained using a spherical top-hat model. In truly inhomogenous systems, the $\\delta-\\delta_v$ relation is no longer one to one and the joint probability distribution of the two variables provides a more complete description. Whether or not the Zel'dovich relation is obeyed even in a spherically averaged sense is not yet examined. Numerical simulations will be required to give more refined theoretical answers. From an observational point of view too, there are many sources of deviation from spherical symmetry; tidal interactions give rise to transverse forces, infall of smaller systems superimposes a random velocity on the radial infall (e.g., \\citealt{diaferio_infall_1997}) etc. In addition to the difficulties associated with isolating peculiar velocities, these pose further constraints on the application of the Zel'dovich relation to real systems. Nevertheless, the Zel'dovich curve can provide a pivot point to analyse more complicated systems and we hope that its significance in the density-velocity phase space dynamics can be exploited to constrain cosmological parameters in the future." }, "1207/1207.2777_arXiv.txt": { "abstract": "We examine {\\it Herschel Space Observatory} images of one nearby prototypical outer ring galaxy, NGC\\,1291, and show that the ring becomes more prominent at wavelengths longer than 160\\,$\\micron$. The mass of cool dust in the ring dominates the total dust mass of the galaxy, accounting for at least 70\\% of it. The temperature of the emitting dust in the ring (T$=19.5\\pm0.3$\\,K) is cooler than that of the inner galaxy (T$=25.7\\pm0.7$\\,K). We discuss several explanations for the difference in dust temperature, including age and density differences in the stellar populations of the ring versus the bulge. ", "introduction": "Outer or external rings are typically large, low surface brightness features of barred and weakly barred galaxies, prominent at optical wavelengths. They are most frequently observed in early-type spirals (S0/a) and, as with gaseous rings and pseudorings, are most often believed to be associated with outer Lindblad resonances (Buta \\& Combes 1996). These resonances are assumed to arise with bars or other pertubations (see simulations by, e.g., Bagley et al.\\ 2009). For outer rings in galaxies without bars, explanations vary from ring creation through tidal forcing via interactions with companions or bars that have since dissipated, to spiral density waves in the disk (Rautiainen \\& Salo 2000). There are also outer rings that presumably formed in collisions (e.g., Appleton \\& Struck-Marcell 1996) or through prograde major mergers of gas-rich disk galaxies (Brook et al.\\ 2007). Outer rings can have major axes that are twice the size of the bar component (Schwarz 1981) and, therefore, they dominate the outer areas of the disks of the galaxies that contain them. The study of outer disks as a path to understanding galaxy evolution has had recent renewed interest due to the discovery of very extended disks at ultraviolet (Gil de Paz 2005; Thilker et al.\\ 2005, 2007; Munoz-Mateos et al.\\ 2007), optical (de Jong et al.\\ 2007; McConnachie et al.\\ 2009), infrared (Engelbracht et al.\\ 2004; Hinz et al.\\ 2004, 2006) and submillimeter (Planck Collaboration et al.\\ 2011) wavelengths. In particular, implications for the production source of dust and the mechanisms for heating and transporting the dust in galaxy outskirts may reveal much about the growth and chemical enrichment of disks. The proximity of NGC\\,1291 and the wealth of available ancillary and space-based data make it an ideal test case for understanding dust emission in these extended disk structures. The first large outer ring structure to be discovered was, in fact, the one in NGC\\,1291 (Perrine 1922). The galaxy is classified as an (R)SB(s)0/a (de Vaucouleurs et al.\\ 1991) with an inclination of $i=35\\arcdeg\\pm7\\arcdeg$ (Prescott et al.\\ 2007) and is at a distance estimated to be between $7-10.4$\\,Mpc (Masters 2005; Kennicutt et al.\\ 2008); for the remainder of the paper we use the 10.4\\,Mpc value, following Kennicutt et al.\\ (2011). In addition to its outer ring, at a radius of $\\sim9$\\,kpc, at optical wavelengths it is characterized by a bright inner lens, a primary bar, and a small secondary bar misaligned by $\\sim30\\arcdeg$ (de Vaucouleurs 1975; P{\\'e}rez \\& Freeman 2006; for more general information on the classification of rings and lenses in S0's, see, e.g., Michard \\& Marchal 1993). Its star formation has been studied via H\\,$\\alpha$ emission (Caldwell et al.\\ 1991; Crocker et al.\\ 1996; Meurer et al.\\ 2006). It is part of the {\\it Spitzer} Survey of Stellar Structure in Galaxies (S$^4$G; Sheth et al.\\ 2010) sample and, as such, has been re-classified in the Infrared Array Camera (IRAC; Fazio et al.\\ 2004) bands as (R)SAB(l,ub)0$^+$ (Buta et al.\\ 2010). NGC\\,1291 is easily detected in the ultraviolet, but it is not classified as exhibiting an extended UV disk because optical emission is found to be coincident with the UV knots (Thilker et al.\\ 2007). H\\,{\\sc i} measurements of the galaxy (van Driel et al.\\ 1988) show that the atomic gas is concentrated in the outer ring with a pronounced central hole. The H\\,{\\sc i} gas mass is $0.81\\times10^9$\\,M$_{\\odot}$, relatively gas-rich for an S0/a galaxy (e.g., Li et al.\\ 2011). Using images of the galaxy taken as part of the {\\it Spitzer} Infrared Nearby Galaxies Survey (SINGS; Kennicutt et al.\\ 2003) Legacy program, Bendo (2006) noted that the nucleus of NGC\\,1291 is the dominant source of 8\\,$\\micron$ emission, as well as 24 and 70\\,$\\micron$ warm dust emission, and that the 8 and 24\\,$\\micron$ emission are well correlated. However, at 160\\,$\\micron$, assumed to be associated with cool (T\\,$\\sim20$\\,K) dust emission, the outer ring is a stronger source than the central portion of the galaxy. This was confirmed by Balloon-borne Large Aperture Submillimeter Telescope ({\\it BLAST}; Griffin et al.\\ 2007) observations of NGC\\,1291, presented as part of a paper on resolved galaxies that also have {\\it Spitzer} images (Wiebe et al.\\ 2009). Both the central core and outer ring were detected by {\\it BLAST} in a total of four observations of a $\\sim0.4$\\,deg$^2$ area. In a continuing effort to understand the role dust plays in the evolution of galaxies, we present {\\it Herschel Space Observatory} images of NGC\\,1291, complementary to the {\\it Spitzer} and {\\it BLAST} images described above. Section 2 describes the observations and data reduction, and Section 3 the analysis and comparison with previous data. Section 4 contains a discussion of these results, while Section 5 contains a summary. ", "conclusions": "Several explanations for the existence of an outer ring with cooler dust may be plausible. It is possible that all spiral galaxies have a dust temperature gradient between their inner and outer disks. This might be due to the concentration of old stars or star formation occurring at the center of the galaxy in comparison to the more sparse stellar population of the comparatively large outer disk, i.e., that the density of the radiation field in the center of a galaxy is higher than that in the outer parts. Another explanation could be that the nucleus of the galaxy contains an AGN which heats the dust to higher temperatures in the inner region in comparison to the outskirts. Perhaps the inner bar of the galaxy fuels star formation in that area by funneling gas continuously, such that rapid star formation heats the inner dust to warm temperatures. Or it could be that the stellar population of the outer ring is older than that of the inner features. In this case, the older stars maintain the dust at cooler temperatures than areas with active star formation. We now discuss each of these possible explanations in turn. \\subsection{Temperature Gradients} First, we explore the idea that all galaxies have similar dust temperature gradients to NGC\\,1291. Engelbracht et al.\\ (2010) showed that the central areas of early-type spiral galaxies generally have enhanced dust heating of their cool dust components compared to their disks. They find that, on average, the cool dust temperature of the central component is $15\\pm3\\%$ hotter than the disk. Similarly, Pohlen et al.\\ (2010) showed that SPIRE surface brightness ratios seem to decrease with radius in spirals, implying that the dust in the outer regions is colder than dust in the centers of galaxies. Galametz et al.\\ (2012; in prep.) study 11 galaxies in the KINGFISH sample, including NGC\\,1291, fitting {\\it Spitzer} and {\\it Herschel} data SEDs with two modified blackbodies and creating spatially-resolved maps of their dust properties. They also see systematic drops in dust temperature with radius for disk galaxies, also on the order of 10-15\\,K from inner to outer galaxy. However, it is not obvious that NGC\\,1291 should follow this pattern and that outer rings should be dominated by cool dust emission. UV rings in S0 and early-type spiral galaxies are often associated with sites of recent star formation, sometimes marked by clumpy H\\,{\\sc ii} regions, and can even be the only sites of star formation in some galaxies (e.g., Donovan et al.\\ 2009; Thilker et al.\\ 2010). A recent study of five nearby barred S0 galaxies with rings by Marino et al.\\ (2011) showed that these outer rings account for a majority (up to 70\\%) of the flux at UV wavelengths, indicating their young stellar population. Buta et al.\\ (2010) specifically discuss the morphology of NGC\\,1291 based on $B$-band and 3.6\\,$\\micron$ IRAC images (Sheth et al.\\ 2010), saying that the outer ring is ``where most of the recent star formation is taking place'' and is prominent in the $B$-band (Figure 1). At 3.6\\,$\\micron$, however, the ring does not stand out and is instead a broad ellipse in a rounder diffuse background (Figure 1 and their Figure 7). Therefore, we might expect that, if the 24\\,$\\micron$ emission is associated with the warm (T$\\sim50$\\,K) dust heated by the young stellar population (seen at UV and $B$-band to be prominent), then the warm dust at 24\\,$\\micron$ should be brighter in the ring as well, yet the wavelengths associated with cool (T$\\sim20$\\,K) dust are brighter relative to the flux from the rest of the galaxy. To study the dust temperature gradient in NGC\\,1291 in more detail, Figure 6 displays circles of radii $2\\arcmin$, $3\\arcmin$, $4\\arcmin$, $5\\arcmin$, and $6\\arcmin$ overlaid on the SPIRE 250\\,$\\micron$ image, along with the estimated temperature of the dust within each annular aperture, using the same blackbody fitting technique as before. (These annuli match the dust mass annuli calculations found in Table 1. Unlike Table 1, the data for Figure 6 are not convolved to the 160\\,$\\micron$ resolution.) We remind the reader that such temperatures are representative, characterizing large grains that dominate the emission at wavelengths greater than 70\\,$\\micron$, and that, in fact, such large annuli contain a wide range of temperatures, including small grains undergoing temperature fluctuations. Keeping this in mind, if all disks, with or without rings, have similar dust temperature gradients due to the gradually increasing distance between star formation regions with radius, then we should be able to identify the same drop in temperature for spirals other than NGC\\,1291. We choose, from the KINGFISH sample, NGC\\,628 (SAc; $d=7.3$\\,Mpc), which does not contain any rings but is face-on, is at a similar distance to NGC\\,1291 and has a similar angular size and total infrared luminosity to NGC\\,1291. We show in Figure 6 the dust temperatures for this galaxy for the same annular apertures as NGC\\,1291. For this galaxy, we see a gradual decline in temperature with some scatter. We see neither the dramatic dip in temperature at $3\\arcmin$ that is seen for NGC\\,1291, nor do we see such a large spread in temperatures, with NGC\\,628 having a temperature difference between inner and outer areas of 4.8\\,K versus 8.2\\,K for NGC\\,1291. Thus, it appears that the dust temperature from the ring is cooler than would be expected from a simple temperature decline with radius for a normal (non-ringed) spiral galaxy. Comparisons of spatially resolved dust temperatures with a larger number of KINGFISH galaxies are shown in Galametz et al.\\ (2012). \\begin{figure} \\plottwo{figure6a.eps}{figure6b.eps} \\caption{{\\it Left:} The 250\\,$\\micron$ image of NGC\\,1291 (grayscale) overlaid with apertures at $2\\arcmin$, $3\\arcmin$, $4\\arcmin$, $5\\arcmin$ and $6\\arcmin$ radii (green circles). {\\it Right:} The dust temperature deduced from the infrared flux densities within each aperture. Solid dots represent NGC\\,1291 while open squares represent NGC\\,628, a comparison face-on spiral galaxy. The shaded region indicates the approximate area occupied by the ring of NGC\\,1291 at SPIRE wavelengths.} \\end{figure} \\subsection{AGN Heating} The next possibility is that non-stellar sources such as an AGN could heat the dust in the inner portion of the galaxy, where we find a warmer dust temperature compared to the outer ring. This is suggested by the fact that outer rings are found to be much more frequent in galaxies with Seyfert nuclei (Hunt \\& Malkan 1999). NGC\\,1291 is an X-ray source, with emission detected within $2\\farcm5$ (Bregman et al.\\ 1995), which appears to be consistent with two energy components with their origin in a stellar component and an extended hot gas component (Hogg et al.\\ 2001). Irwin et al.\\ (2002) found an excess of soft emission similar to that seen in several other low-luminosity AGNs. Moustakas et al.\\ (2010) classify NGC\\,1291 as an ``AGN'' type galaxy based on optical spectra of the nuclear region, although they note that one or more emission line(s) failed their S/N$>2$ requirement for such a classification. Recent analysis of X-ray spectra of the nuclear source also indicates a low-luminosity AGN with moderate obscuration (Luo et al.\\ 2012) which could be a source of dust heating. \\subsection{Transport via the Bar} The third possibility is that the bar in NGC\\,1291, seen at the various wavelengths in Figure 1, funnels gas into the inner region of the galaxy, replenishing the supply necessary for star formation. This compact region of star formation could then generate warmer dust temperatures than in the outer ring (see also Engelbracht et al.\\ 2010). If gas is being transported inwards in this manner, the observational evidence would likely be in the form of ``hotspots'' at the ends of the bar - features that are brighter than the rest of the bar formed by shock heating of the gas as it is drawn to the center of the galaxy, causing a ``pile-up'' near the inner Lindblad resonance (Combes \\& Gerin 1985; Shlosman et al.\\ 1989). This is not seen in the H\\,{\\sc i} data, which has a hole at the center (van Driel et al.\\ 1988), or in the other wavelengths that are available. However, we cannot rule out this possibility, as such hotspots are thought to be small, very faint features. Hunt \\& Malkan (1999) argue that rings alone could be signs of inward transport of material, essentially a second ``pile-up'' of dust and gas in the outer resonance of the galaxy; the star formation in the ring prominent in the UV, presumably indicating stellar ages of $\\sim100$\\,Myr, could be a residual signature of such a build-up. We note that transport of material via the bar may be related to the AGN hypothesis in the previous section, in that AGNs as well as star formation can be fueled via inflow. \\subsection{Stellar Age Gradient} The last explanation for the change in cool dust temperature between the inner galaxy and outer ring is that the ring is composed mainly of an old population of stars with some current star formation. In this case, the dust in the ring was produced by generations of stars in the past and is heated now mainly by the old ($\\sim10$\\,Gyr) stellar population, shown best by the 3.6\\,$\\micron$ image (Buta et al.\\ 2010 and Figure 1) where the ring is a large, diffuse entity spanning much more physical space than seen at far-infrared wavelengths. Such ``old population rings'' have been suggested by comparisons between $B$ and $I$-band images of nearby galaxies. Buta (1995) showed a typical galaxy of this type (IC\\,1438) which has a nuclear ring, an oval inner ring, and an outer ring that is much brighter in the $I$ passband than in the $B$. The explanation given for this difference in brightness was that the outer ring component formed first and left behind a stellar remnant as the other two inner features formed. An old outer ring would be unusual in that the dynamical timescales for forming rings or arms are much shorter at smaller disk radii (e.g., Freeman et al.\\ 2010). A remnant of this type could be the source of dust heating for the emission we see at far-infrared wavelengths. In addition, there would also be a contribution to the dust heating from ionizing ultraviolet photons which have escaped from the existing (but weaker, in terms of contribution) star forming regions. This combination of dust heating sources has been well modeled by, e.g., Misiriotis et al.\\ (2001) and Popescu et al.\\ (2002; 2011), and proposed to explain observations of other nearby galaxies (Hinz et al.\\ 2004, Tabatabaei et al.\\ 2007, Bendo et al.\\ 2010). (Though see Groves et al. 2012 for a case where the old stellar population heats dust to warm temperatures.) Inspection of the $B$-$I$ image (R.\\ Buta, private communication) shows such a color difference and will be investigated further in Buta et al.\\ (2012). Simulations by Freeman et al.\\ (2010) addressing Hoag's object, an outer ring galaxy with no inner morphological features, show that an outer ring can be formed via perturbations from a bar which then dissipates slowly over time. In the environment of these instabilities, the majority of the simulated gas particles fall to the center, while the remaining particles create an outer ring, with the region between the two essentially empty. If this is the manner in which NGC\\,1291 is forming its ring, then the implication is that the ring is unlikely to be older than the inner regions. Noll et al.\\ (2009) use SED fitting techniques to study the star-formation history of NGC\\,1291 and posit two exponential starburst events for the galaxy, one of which occurred 10\\,Gyr ago and the second of which occurred 200\\,Myr ago. This analysis was performed for the entire galaxy. A similar study for the inner portion and the ring separately would help disentangle their stellar histories." }, "1207/1207.1950.txt": { "abstract": "We present new, accurate positions, spectral classifications, radial and rotational velocities, H$\\alpha$ fluxes, equivalent widths and B,V,I,R magnitudes for 579 hot emission-line stars (classes B0 - F9) in the Large Magellanic Cloud (LMC) which include 469 new discoveries. Candidate emission line stars were discovered using a deep, high resolution H$\\alpha$ map of the central 25 deg$^{2}$ of the LMC obtained by median stacking a dozen 2 hour H$\\alpha$ exposures taken with the UK Schmidt Telescope (UKST). Spectroscopic follow-up observations on the Anglo-Australian Telescope (AAT), the UKST, the Very Large Telescope (VLT), the South African Astronomical Observatory (SAAO) 1.9m and the 2.3m telescope at Siding Spring Observatory have established the identity of these faint sources down to magnitude $R_{\\textrm{equiv}}\\sim$23 for H$\\alpha$ ($4.5\\times10^{-17}\\textrm{ergs~cm}^{-2}~\\textrm{s}^{-1}~$\\AA$^{-1}$). Confirmed emission-line stars have been assigned an underlying spectral classification through cross-correlation against 131 absorption line template spectra covering the range O1 to F8. We confirm 111 previously identified emission line stars and 64 previously known variable stars with spectral types hotter than F8. The majority of hot stars identified (518 stars or 89\\%) are class B. Of all the hot emission-line stars in classes B-F, 130 or 22\\% are type B[e], characterised by the presence of forbidden emission lines such as \\SII, \\NII~and \\OII. We report on the physical location of these stars with reference to possible contamination from ambient \\HII~emission. Only 13 of the emission-line stars require additional high resolution spectroscopic observations in order to assign a spectroscopic classification. They have nonetheless been added to the catalogue. Along with flux calibration of the H$\\alpha$ emission we provide the first H$\\alpha$ luminosity function for selected sub-samples after correction for any possible nebula or ambient contamination. We find a moderate correlation between the intensity of H$\\alpha$ emission and the V magnitude of the central star based on SuperCOSMOS magnitudes and the Optical Gravitational Lensing Experiment (OGLE-II) photometry where possible. Cool stars from classes G-S, with and without strong H$\\alpha$ emission, will be the focus of part 2 in this series. ", "introduction": "%%%%%%%%%% LMC The Large Magellanic Cloud (LMC) is a unique laboratory in which to study the peculiar characteristics of massive and luminous emission-line stars. At a known distance of $\\sim$50 kpc (see Reid \\& Parker 2010 and references therein) to all LMC members, modest inclination angle to the line of sight ($\\sim$21deg) and with relatively low interstellar extinction (Rv = 3.41 $\\pm$ 0.06; Gordon et al. 2003), apparent brightness is a good indicator of absolute luminosity to within a few tenths of a magnitude. We take advantage of these benefits as we identify and begin basic analysis of emission-line stars in the LMC. The most prominent observational feature of the emission-line stellar group is the presence of the H$\\alpha$ line. The presence of this emission feature has been widely used as an identifier in the many previous searches for emission-line stars in the LMC (eg. Feast et al. 1960; Henize 1956; Lindsay 1963,1974; Bohannan \\& Epps 1974; Grebel 1997; Keller et al. 1999; Grebel \\& Chu 2000; Keller et al 2000; Olsen et al. 2001). None of these surveys went particularly deep. More recently, the OGLE II database has prevailed as the main tool used to uncover emission-line star candidates (Sabogal et al. 2005). %%%%%%%%%% H alpha survey The UKST H$\\alpha$ survey of the central 25deg$^{2}$ of the LMC has changed this situation. It was adjunct to the successful Southern Galactic Plane H$\\alpha$ survey (Parker et al. 2005) and has revealed large numbers of various emission objects. In addition to revealing 460 new planetary nebulae within the survey region which were confirmed spectroscopically (Reid \\& Parker, 2006a,b), spectroscopic followup and careful analysis has revealed 579 hot emission-line stars with spectral classes B-F out of a total sample of 1,062 emission-line stars of all spectral types uncovered. Only 111 of these were previously known or identified while 469 are newly discovered. The majority are Be, B[e], Bpe and HAeBe stars but two are Luminous Blue Variable (LBV) candidates. %The H$\\alpha$ luminosity and continuum %emission from these stars provides estimates of mass loss rates, indicating %those with eruptive mass loss. Identifying these objects will assist our understanding of the main sequence evolution of massive stars. We have also identified 6 new and 33 previously known Wolf-Rayet stars, which are not included in this number but will be the special focus of a follow-up paper. Be stars are known to be variables which undergo active and quiescent stages (Telting 2000; Bjorkman et al. 2002). A single epoch survey could miss many of these stars if they were undergoing a quiescent stage. This problem has already been demonstrated by several follow-up investigations (Hummel et al. 1999; Keller et al. 1999; Wisniewski and Bjorkman 2006) which were unable to identify all of the previously identified Be stars in the Magellanic Clouds and in the Galaxy. In addition, these same follow-up studies revealed previously unidentified Be stars. Our H$\\alpha$ survey, utilising 12 H$\\alpha$ exposures taken over a three year period has largely alleviated such problems and revealed a large number of emission-line stars in the survey region to a magnitude of $R_{\\textrm{equiv}}\\sim$22 for H$\\alpha$. %%%%%%%% Accurate positions of ELS and improvements on previous published positions In order to study the Balmer emission we have measured the Equivalent Width (EW) and Full Width Half Maximum (FWHM) of the H$\\alpha$ emission-lines. In addition, we include H$\\alpha$ fluxes from medium resolution spectroscopy of 575 (99.3\\%) of the detected emission-line stars within the survey area. Our follow-up spectroscopy was conducted from November 2004 to February 2005 on a variety of telescopes, allowing us to re-observe several known variable stars and detect minor changes in spectral characteristics. All but 2 candidate emission-line stars found in the H$\\alpha$ survey had some degree of H$\\alpha$ emission detectable in their spectrum at the time of observation. After describing flux calibration (section~\\ref{section3}), we explain the method used to assign a spectral classification and luminosity class to each star using cross-correlation against well-established templates (section~\\ref{section4}). Section~\\ref{section6} describes our method for deriving the rotational velocities and section~\\ref{section7} outlines a simple method for correcting or at least estimating nebula contribution in the spectrum. Section~\\ref{section8} details our routine for assigning accurate positions to each star. In section~\\ref{section9} we describe the method used for measuring the radial velocity of each star. Velocities accurate to $\\sim$4 km s$^{-1}$ have been found for 572 emission-line stars using both the weighted emission-line and cross-correlation techniques on our higher dispersion spectroscopic data. These velocities can be used to search for kinematical substructures in the LMC disk, create a 3D kinematic map of the LMC for comparison with the \\HI~disk, assist studies of age-metallicity dispersion and distribution, potentially find stellar associations and streams, and compare medium to old age populations such as planetary nebulae within the LMC (Reid \\& Parker 2006b). In section~\\ref{section10} we show the projected distribution of emission-line stars and late-type stars across the survey field of the LMC. In section~\\ref{section11} we measure the intensity of the H$\\alpha$ emission considering ambient sky and any nebula contamination in order to create the first luminosity function for these stars in the LMC. Then, in section 13 we assess the emission by comparing BVI photometry from SuperCOSMOS and OGLE-II data where available. We discuss the stellar photometry, its reliability and problems associated with variability. In section~\\ref{section12} we briefly discuss the variability already found in many of the candidate emission-line stars. The full catalogue of emission-line stars is described in section~\\ref{section13} and presented in the appendix. Individual spectra and H$\\alpha$ images will be available (2nd half 2012) on a dedicated web page hosted by the Astronomy Department at Macquarie University. %%%%%%%% Velocities are given %%%%%%%% How Be stars form ", "conclusions": "" }, "1207/1207.1681_arXiv.txt": { "abstract": "The number of plausible associations of extended VHE (TeV) sources with pulsars has been steadily growing, suggesting that many of these sources are pulsar wind nebulae (PWNe). Here we overview the recent progress in X-ray and TeV observations of PWNe and summarize their properties. ", "introduction": " ", "conclusions": "" }, "1207/1207.3960_arXiv.txt": { "abstract": "In this paper a spatially resolved, fully self-consistent SSC model is presented. The observable spectral energy distribution (SED) evolves entirely from a low energetic delta distribution of injected electrons by means of the implemented microphysics of the jet. These are in particular the properties of the shock and the ambient plasma, which can be varied along the jet axis. Hence a large variety of scenarios can be computed, e.g. the acceleration of particles via multiple shocks. Two acceleration processes, shock acceleration and stochastic acceleration, are taken into account. From the resulting electron distribution the SED is calculated taking into account synchrotron radiation, inverse Compton scattering (full cross section) and synchrotron self absorption. The model can explain SEDs where cooling processes are crucial. It can verify high variability results from acausal simulations and produce variability not only via injection of particles, but due to the presence of multiple shocks. Furthermore a fit of the data, obtained in the 2010 multi-frequency campaign of Mrk501, is presented. ", "introduction": "Synchrotron Self Compton (SSC) models have been quite successful in explaining the emissions of blazars. However, recent observational results have led to the conclusion that the approach of modeling blobs in blazar jets as homogeneous regions employing SSC codes is not sufficient. Especially the observation of intra day variability introduces strong constraints for both the size of the emission region and the Doppler factor that usually lack observational evidence or even contradict those (e.g. Ref.~\\refcite{mrk501_vlbi} and Ref.~\\refcite{rapid_tev_var}). Furthermore the observed time delays between light curves of different frequencies, especially soft lags\\cite{soft_xray_lag}, are hardly explainable by such models. Two-zone-models like Ref.~\\refcite{weidinger2} can produce scenarios that can reproduce observational data by merely varying physical parameters. Although this is an enormous improvement in comparison to one-zone-models, into which the electron distribution enters as a couple of unphysical parameters, there are still problems: The stochastic character of the dominant acceleration process (Fermi-I) is not taken into account. It is modeled as a convection in momentum space and enters as a term in the Fokker-Planck equation. Furthermore the radiation zone is modeled homogeneously. Therefore variations, for example introduced via particle injection from the acceleration zone, will affect the entire space immediately. This behavior will lead to shortened timescales. Finally the cooling of the particles happens everywhere quasi locally and not depending on distance to the acceleration site. Additional to addressing the aforementioned issues spatially resolved models can predict the morphology of the emission region. Using VLBI this is in principle observable. ", "conclusions": "With the here presented model it is possible to confirm results regarding the particle acceleration from one- and two-zone-models and therefore trace them back to the jets microphysics. The number of parameters is not higher than that of homogenous SSC models and could be further reduced by connecting momentum diffusion and the scattering parameter, since both processes occur due to the presence of Alfv\\'en waves. In the fit of the Mrk501 data the difference to homogeneous models is most prominent in the radio range. This effect is presumably due to the emission of electrons far away from the shock that are already cooled. Furthermore our model is able to explain very short time variability since the lower limit for the timescale, due to the light crossing time of the emission region, no longer holds. Fitting of actual lightcurves will be subject to future publications." }, "1207/1207.1224_arXiv.txt": { "abstract": "We present the results of an analysis of {\\Kepler} data covering 1.5 years of the dwarf nova V447 Lyr. We detect eclipses of the accretion disk by the mass donating secondary star every 3.74 hrs which is the binary orbital period. V447 Lyr is therefore the first dwarf nova in the {\\Kepler} field to show eclipses. We also detect five long outbursts and six short outbursts showing V447 Lyr is a U Gem type dwarf nova. We show that the orbital phase of the mid-eclipse occurs earlier during outbursts compared to quiescence and that the width of the eclipse is greater during outburst. This suggests that the bright spot is more prominent during quiescence and that the disk is larger during outburst than quiescence. This is consistent with an expansion of the outer disk radius due to the presence of high viscosity material associated with the outburst, followed by a contraction in quiescence due to the accretion of low angular momentum material. We note that the long outbursts appear to be triggered by a short outburst, which is also observed in the super-outbursts of SU UMa dwarf novae as observed using {\\Kepler}. ", "introduction": "Cataclysmic Variable (CV) binary systems contain a white dwarf primary star that accretes mass from a Roche lobe-filling late-type main sequence secondary star. Mass loss from the secondary through the inner Lagrange point (L1) forms an accretion disk about the primary, and viscosity within the accretion disk acts to transfer angular momentum outward in radius allowing mass to migrate inward to the surface of the white dwarf. The disk and bright spot associated with the accretion stream impact point are typically the brightest components in a CV system (Warner 1995, Hellier 2001, Frank, King \\& Raine 2002). The mean disk luminosity is ultimately provided by the release of gravitational potential energy as the material migrates inward through the disk given by $L_{\\rm disk} \\sim GM_1\\dot M_1/R_1$, where $\\dot M_1$ is the mass transfer rate onto the primary of mass $M_1$ and radius $R_1$. The mass flow rate through L1 is governed by the long-term evolution of the binary separation and the secondary star itself, but the mass flow rate through the disk and onto the primary is a function of the viscosity of the disk -- the higher the viscosity, the higher the inward mass flow rate. V447 Lyrae (KIC 8415928, $r$=18.4, Brown et al. 2011) is a little-studied CV in the NASA {\\Kepler} field of view. Its sky co-ordinates are $\\alpha=19^\\textrm{h}00^{\\rm m}19\\fs92$ $\\delta = +44\\degr 27'44\\farcs9$ (2000.0). It was discovered and announced as GR 247 by Romano (1972) who noted a maximum photographic magnitude of 17.2 and a minimum fainter than 18.5. These observations were included in Downes, Webbink \\& Shara (1997) but the system was noted as undetected in the 2MASS survey (Hoard et al. 2002) nor is it a known X-ray source. \\begin{table*} \\begin{center} \\begin{tabular}{lllll} \\hline Quarter & \\multicolumn{2}{c}{Start} & \\multicolumn{2}{c}{End} \\\\ & MJD & UT & MJD & UT\\\\ \\hline Q6 (LC) & 55371.947 & 2010 Jun 24 22:46 & 55461.794 & 2010 Sep 22 19:04 \\\\ Q7 (LC) & 55462.673 & 2010 Sep 23 16:10 & 55552.049 & 2010 Dec 22 01:09 \\\\ Q8 (SC) & 55567.855 & 2011 Jan 06 20:42 & 55634.856 & 2011 Mar 14 20:15 \\\\ Q9 (SC) & 55641.007 & 2011 Mar 21 00:23 & 55738.434 & 2011 Jun 26 10:13 \\\\ Q10 (LC) & 55739.343 & 2011 Jun 27 08:16 & 55832.766 & 2011 Sep 28 18:24 \\\\ Q11 (SC) & 55833.696 & 2011 Sep 29 16:58 & 55930.837 & 2012 Jan 04 20:48 \\\\ \\hline \\end{tabular} \\end{center} \\caption{Journal of Observations. The start and end MJD and UT dates are the mid-point of the first and final cadence of the LC time series for each quarter respectively.} \\label{log} \\end{table*} This work is the fifth in a series of publications focussing on the CVs in the {\\Kepler} mission field of view. In Still et al. (2010), we presented preliminary results for the periods observed in the Q2 data for V344 Lyr, a previously little-studied CV in the {\\Kepler} field. In Cannizzo et al. (2010), we presented results of the thermal viscous disk instability model for CVs applied to the Q2--Q4 V344 Lyr outburst time series data. In Wood et al. (2011) we presented a detailed analysis of the orbital and superhump periods present in the V344 Lyr Q2--Q4 time series data. In Cannizzo et al. (2012) we discussed the outburst properties of V1504 Cyg and V344 Lyr over the first two years of {\\Kepler} observations and most recently (Barclay et al. 2012) we report the serendipitous discovery of an SU UMa dwarf nova within 7 arcsec of a G-type star. Here we report {\\Kepler} observations of V447 Lyr covering 1.5 years. ", "conclusions": "We report observations of V447 Lyr which show that it is the first dwarf nova in the {\\Kepler} field to show eclipses. It has an orbital period of 3.74 hrs and shows almost equal numbers of long and short outbursts, which makes it a near twin of the well studied dwarf nova U Gem. By fitting the mean eclipse profile of three successive eclipses we find that the phase of the mid-eclipse occurs earlier during outbursts compared to quiescence and that the width of the eclipse is greater during an outburst. This suggests that the accretion disk has a larger radius during outburst compared to during quiescence and is consistent with an expansion of the outer disk radius due to the presence of high viscosity material associated with the outburst, followed by a contraction in quiescence due to the accretion of low angular momentum material. {\\Kepler} observations of dwarf novae outbursts have found that super-outbursts in the shorter orbital period dwarf novae appear to be triggered by a normal outburst. We find that long outbursts also appear to be triggered by short outbursts in V447 Lyr. This indicates that this is a general phenomena found in CVs which any outburst model will have to reproduce." }, "1207/1207.1823_arXiv.txt": { "abstract": "We discuss weak lensing characteristics in the gravitational field of a compact object in the low-energy approximation of fourth order $f\\left( R\\right) $\\ gravity theory. The particular solution is characterized by a gravitational strength parameter $\\sigma $\\ and a distance scale $r_{c}$\\ much larger than the Schwarzschild radius. Above $r_{c}$ gravity is strengthened and as a consequence weak lensing features are modified compared to the Schwarzschild case. We find a critical impact parameter (depending upon $r_{c}$) for which the behavior of the deflection angle changes. Using the Virbhadra-Ellis lens equation we improve the computation of the image positions, Einstein ring radii, magnification factors and the magnification ratio. We demonstrate that the magnification ratio as function of image separation obeys a power-law depending on the parameter $\\sigma $, with a double degeneracy. No $\\sigma \\neq 0$ value gives the same power as the one characterizing Schwarzschild black holes. As the magnification ratio and the image separation are the lensing quantities most conveniently determined by direct measurements, future lensing surveys will be able to constrain the parameter $\\sigma $ based on this prediction. ", "introduction": "The recent advent of the so-called \\textquotedblleft Precision Cosmology\\textquotedblright\\ along with galactic observations indicate that General Relativity (GR) with standard matter sources disagrees with an increasing number of observational data, e. g. those coming from IA-type Supernovae, used as standard candles, large scale structure ranging from galaxies up to superclusters \\cite{SNeIa,CMBR,WMAP}, and galactic rotation curves. In addition, from a theoretical point of view, being not renormalizable, GR fails to be quantized in any \\textit{standard} way (see \\cite{uti}). Therefore at the extreme ultra-violet and infrared scales GR is not and cannot be the definitive theory of Gravitation despite the fact that it successfully addresses a wide range of phenomena and the Newtonian weak field limit is correctly recovered. In order to interpret the recent observational data in the framework of GR, the introduction of unknown \\textit{dark matter} (DM) (to address dynamical phenomena as the formation of self-gravitating astrophysical structures) and \\textit{dark energy} (DE) (to address the problem of cosmic acceleration) seems to be necessary: however, the price of preserving the simplicity of the Hilbert-Einstein Lagrangian has been the introduction of rather odd-behaving physical entities that, up to now, have not been revealed by any experiment at fundamental scales. This situation has led to several attempts devoted either to recover the validity of GR at any scale, or to construct alternative gravity theories that suitably generalize the Einsteinian one. The philosophy of these two schemes is that in the former case one has to modify the matter sector introducing DM and DE, in the latter approach the dynamics of the geometry (i.e. the left-hand-side of the Einstein equations) is modified but with the constraint to recover GR at local scales. Higher-order theories of gravity (both in metric \\cite% {review,book,capozcurv,sante,MetricRn} and Palatini \\cite% {PalRn,lnR,Allemandi} formulations) represent an interesting approach able to fruitfully cope with both dark matter and dark energy problems. A further approach is based on scalar\\thinspace -\\thinspace tensor theories of gravity but it can be shown that higher-order theories and scalar tensor ones can be related by conformal transformations (see, e.g., \\cite{book,faraoni} and references therein). It has been widely demonstrated that such theories can agree with the cosmological observations of the Hubble flow \\cite{noiijmpd,noifranca} and the large scale structure evolution \\cite{frlss}. In addition, in the weak field limit, the gravitational potential turns out to be modified \\cite% {stelle,schmidt,mannheim,noipla} in such a way that interesting consequences to galactic dynamics may be achieved without violating, at the same time, the constraints on the parametrized post-Newtonian (PPN) parameters coming from Solar System tests \\cite{grgOdi}. If alternative theories of gravitation are able to explain both cosmological and local observations without the introduction of exotic energy-momentum sources, then one might ask how would the differences between these alternative theories be compared to GR with the unusual sources. It is proposed that gravitational lensing might be able to act a means for determining which theory governs the gravitational interaction, through measurements of image separations and the brightnesses of those multiple images. It is well known that the deflection of light observed during the Solar eclipse of 1919 was one of the first experimental confirmations of Einsteinian GR. \\textit{Gravitational lensing}, i.e. the deflection of light rays crossing the gravitational field of a compact object referred to as the \\textit{lens}, has become one of the most astonishing successes of GR and it represents nowadays a powerful tool capable of putting constraints on the dynamics of gravitational structures at different scales, from stars to galaxies and clusters of galaxies, and from the large scale structure to cosmological parameters \\cite{SEF,petters}. If we modify the Lagrangian of the gravitational field, it is obvious that gravitational lensing should be affected. It is therefore mandatory to investigate how gravitational lensing behaves in the framework of alternative theories of gravity to develop a further check on the existence of DM and DE. In particular, one has to verify that the phenomenology of standard gravitational lensing is recovered in the limit as the modified theory of gravity reduces to GR, since several observations point to the validity of GR. However, it is worth stressing that the presence of DM has to be invoked in such cases, in particular for large-scale structure (see e.g. the case of Bullet Cluster \\cite{bullet}). On the other hand, it is worth exploring whether deviations from the classical results for the main lensing quantities could be detected and act as clear signatures for modified theories of gravity. A preliminary study in this direction is in \\cite{stab1} where the gravitational lensing, in the Newtonian limit and in the Jordan frame, for a generic $f(R,R_{\\mu \\nu }^{\\mu \\nu },R_{\\mu \\nu \\alpha \\beta }R^{\\mu \\nu \\alpha \\beta })$ is considered. In this paper, the modifications of the Hilbert-Einstein action are induced by corrections to the Newtonian potential due to the Riemann tensor. As a first step towards such an ambitious task we consider here power\\thinspace -\\thinspace law fourth order theories, i.e. we replace the Ricci scalar $R$ in the gravity Lagrangian with the function $f(R)\\propto R^{n}$ (with $n$ a positive integer) and investigate how this affects the gravitational lensing in the theory. In Ref. \\cite{LTNJC} the gravitational lensing in $f\\left( R\\right) $ theories with a Yukawa-type correction in the potential was investigated and the conclusion reached was that weak lensing could not discern between these theories and general relativity. We revisit this issue by using better lens equations together with a post-Newtonian approximation, both of which can account for higher order deviations away from general relativity. We discuss weak lensing due to compact objects of the $R^{n}$ theory, which have been already used to analyze rotation curves of galaxies in \\cite{FRdark} and in \\cite{stab2}. The paper is organized as follows. The next section provides a short summary of $f(R)$ theories, focusing on the special $R^{n}$ case and we review the post-Newtonian solution representing the compact object. The modifications from GR can be interpreted as an effective energy-momentum tensor. We discuss the referring energy conditions in the Appendix. We discuss in Section \\ref{f} how the predictions of weak gravitational lensing are different in the fourth order theory and in general relativity. For this we determine the image locations, Einstein ring radii, magnification factors and the flux ratio, for various model parameters. We then demonstrate that the flux ratio as function of image separation has a different power-law dependence for each model parameter. We summarize our findings in the Concluding Remarks. ", "conclusions": "In this paper we have analyzed the weak lensing signatures of a fourth order [$f\\left( R\\right) =R^{n}$] gravity compact object, with gravitational potential given in the post-Newtonian regime by Eqs. (\\ref{potential}) and (% \\ref{sigma}). This introduces a new parameter $\\sigma \\in \\lbrack 0,1)$, which governs the deviation from the Newtonian gravitational potential for different values of $n$. General relativity is contained as the special case $n=1$ (corresponding to the model parameter $\\sigma =0$). For any other value of the parameter $n$ (or $\\sigma $) the gravitational attraction increases at distances larger than $r_{c}$ as compared to the prediction of Newtonian gravity. The lensing properties of such compact objects were analyzed before in Ref. \\cite{FRdefl}, based on the small angle lens equation \\cite{Bozza}. In this paper we have improved upon this approach, by employing the first order accurate Virbhadra-Ellis lens equation (\\ref{lens_Ellis}), or with equivalent results the D\\'{a}browski-Schunck lens equation (\\ref{ds}). We analyzed the dependences upon $\\sigma $ and upon the impact parameter $b$ of the deflection angle (\\ref{deflection.angle}). The deflection angle decreases with increasing impact parameter for all $\\sigma $. There is a transition at a critical value $\\left( b/r_{c}\\right) _{crit}=2$, below which the deflection angle monotonically decreases with increasing $\\sigma $% , and above which there is a single maximum. This maximum value increases with the impact parameter. The image positions as a function of the lensing mass and source position, also the image magnifications and their ratio as function of source position show features similar to those in the Schwarzschild case. Nevertheless, in contrast with previous claims in the literature \\cite{stab1}, \\cite{LTNJC} these lensing quantities depend upon $\\sigma $. We have computed the image positions for two values of $\\sigma $. For the larger value of $\\sigma $, the image separation grows faster with an increase in the mass and grows more slowly as the source moves away from the optical axis. For the same source position the magnification factors of the images increase with $\\sigma $, especially the one for the primary image. The increases in their ratio $\\mu _{1}/\\mu _{2}$ is even more significant. Using the most easily measurable lensing observables, the ratio of the magnifications is shown to have a power-law dependence on the image separations, with the power strongly depending on $\\sigma $. The power is the smallest for Schwarzschild black holes ($\\sigma =0$), then it increases with $\\sigma $ to a critical value, after which it decreases again. This behavior provides a means for future gravitational lensing observations to either establish the value of $\\sigma $ up to a double degeneracy or falsify the power-law gravitational potential discussed in this paper if $\\sigma =0$ is confirmed. Given that the next generation of radio telescopes will easily be able to resolve images to less than milliarcsecond (in fact tens of microarcseconds\\textbf{)} accuracy, the different rates at which the ratio of the magnifications changes should be able to provide a significant observational signature constraining the validity of $f\\left( R\\right) $ gravitational theories." }, "1207/1207.1012_arXiv.txt": { "abstract": "We announce the discovery of a low-mass planet orbiting the super metal-rich K0V star HD77338 as part of our on-going Calan-Hertfordshire Extrasolar Planet Search. The best fit planet solution has an orbital period of 5.7361$\\pm$0.0015~days and with a radial velocity semi-amplitude of only 5.96$\\pm$1.74~\\ms, we find a minimum mass of 15.9$^{+4.7}_{-5.3}$~\\me. The best fit eccentricity from this solution is 0.09$^{+0.25}_{-0.09}$, and we find agreement for this data set using a Bayesian analysis and a periodogram analysis. We measure a metallicity for the star of +0.35$\\pm$0.06~dex, whereas another recent work (\\citealp{trevisan11}) finds +0.47$\\pm$0.05~dex. Thus HD77338$b$ is one of the most metal-rich planet host stars known and the most metal-rich star hosting a sub-Neptune mass planet. We searched for a transit signature of HD77338$b$ but none was detected. We also highlight an emerging trend where metallicity and mass seem to correlate at very low masses, a discovery that would be in agreement with the core accretion model of planet formation. The trend appears to show that for Neptune-mass planets and below, higher masses are preferred when the host star is more metal-rich. Also a lower boundary is apparent in the super metal-rich regime where there are no very low-mass planets yet discovered in comparison to the sub-solar metallicity regime. A Monte Carlo analysis shows that this \\emph{low-mass planet desert} is statistically significant with the current sample of 36 planets at $\\sim$4.5$\\sigma$ level. In addition, results from Kepler strengthen the claim for this paucity of the lowest-mass planets in super metal-rich systems. Finally, this discovery adds to the growing population of low-mass planets around low-mass and metal-rich stars and shows that very low-mass planets can now be discovered with a relatively small number of data points using stable instrumentation. ", "introduction": "One of the first correlations announced between two parameters that dealt with exoplanets and their host stars was the abundance of heavy elements in the stellar atmospheres. \\citet{gonzalez97} first noted that all three exoplanet host stars known at that time were over abundant in iron. The paper indicated that the metallicity of these stars were all in the super solar regime. This feature has since been studied by a number of authors, most notably by \\citet{fischer05} who defined a relationship between the host star metallicity ([Fe/H]) and the probability of a star hosting a gas giant planet. This metallicity bias was one of the major features that helped to confirm that the core accretion scenario of planet formation (e.g. \\citealp{ida04}; \\citealp{mordasini09}), coupled with planetary migration (\\citealp{lin86}), appears to be the dominant mechanism to build these planetary systems. A further validation of the core accretion mechanism again comes from stellar metallicity since it appears that low-mass Neptunes and super-Earths are not predominantly found around metal-rich stars (\\citealp{udry06}). The Calan-Hertfordshire Extrasolar Planet Search (CHEPS) is a program to mainly monitor metal-rich stars to primarily hunt for short period planets across a wide range of masses and better constrain the statistics of planets around these stars, whilst also searching for planetary transits of bright and nearby stars in the southern hemisphere. In \\citet{jenkins08} we discussed the sample selection for the CHEPS, using previous observations made with the ESO-FEROS instrument to measure the chromospheric activity and metallicity of a sample of a few hundred stars. This sample has recently been increased using data discussed in \\citet{jenkins11} and current metallicity analyses of these stars is still ongoing using our new methods for measuring accurate atomic abundances in stars like the Sun (\\citealp{pavlenko12}; \\citealp{jenkins12}). In \\citet{jenkins09} we published the discovery of an eccentric brown dwarf, or extreme Jovian planet, orbiting the metal-rich star HD191760, along with new orbits for three other recently discovered southern hemisphere metal-rich planets. The CHEPS data has mostly been obtained using the HARPS instrument but significant observing time on Coralie has also recently been acquired to pre-select interesting new targets for HARPS follow-up (see \\citealp{jenkins10b}). ", "conclusions": "\\begin{figure} \\vspace{5.0cm} \\hspace{-4.0cm} \\special{psfile=fig12.eps hscale=30 vscale=30 angle=90 hoffset=340 voffset=-25} \\vspace{0.5cm} \\caption{Plot of radial velocity detected planets host star metallicity against minimum mass of the planet, taken from the Exoplanet Database (\\citealp{wright11}). Only planets around dwarf stars were included.} \\label{rv_met_mass} \\end{figure} HD77338 is one of the most metal-rich stars currently known to host a planet. If we take the recent \\citet{trevisan11} [Fe/H] value of +0.47$\\pm$0.05~dex for HD77338 then only one planet host star has a metallicity higher than this, HD126614~A with a metallicity of +0.56$\\pm$0.04~dex (\\citealp{howard10}). However, HD126614~A has an M dwarf companion at only 33~AU separation from the host star, meaning HD77338$b$ would be the most metal-rich single star known to host an exoplanet. Also, the low-mass nature of HD77338$b$ (\\msini~$<$~0.05M$_{\\rm{J}}$) helps to populate the low-mass metal-rich bin and shows that metal-rich, very low-mass planets may be plentiful. \\begin{table*} \\tiny \\center \\caption{Orbital and stellar parameters for all radial velocity detected host stars with sub-Neptune mass planets.} \\label{tab:planets} \\begin{tabular}{ccccccccccc} \\hline \\multicolumn{1}{c}{Planet} & \\multicolumn{1}{c}{\\msini} & \\multicolumn{1}{c}{Semimajor Axis} & \\multicolumn{1}{c}{Period} & \\multicolumn{1}{c}{eccentricity} & \\multicolumn{1}{c}{K} & \\multicolumn{1}{c}{$M_{\\rm{\\star}}$} & \\multicolumn{1}{c}{[Fe/H]} & \\multicolumn{1}{c}{$V$} & \\multicolumn{1}{c}{SpT} & \\multicolumn{1}{c}{U,V,W} \\\\ \\multicolumn{1}{c}{} & \\multicolumn{1}{c}{M$_{\\rm{J}}$} & \\multicolumn{1}{c}{AU} & \\multicolumn{1}{c}{days} & \\multicolumn{1}{c}{} & \\multicolumn{1}{c}{\\ms} & \\multicolumn{1}{c}{$M_{\\rm{\\odot}}$} & \\multicolumn{1}{c}{dex} & \\multicolumn{1}{c}{mags} & \\multicolumn{1}{c}{} & \\multicolumn{1}{c}{\\kms} \\\\ \\hline GJ3634$b^1$ & 0.022 & 0.029 & 2.65 & 0.08 & 5.57 & 0.45 & -- & 11.90 & M2.5 & -- \\\\ GJ667C$b^2$ & 0.018 & 0.049 & 7.20 & 0.17 & 3.90 & 0.31 & -0.59 & 10.22 & M1.5V & 20,29,-27$^a$ \\\\ GJ667C$c^2$ & 0.014 & 0.123 & 28.16 & $<$0.27 & 2.02 & 0.31 & -0.59 & 10.22 & M1.5V & 20,29,-27$^a$ \\\\ HD20794$b^3$ & 0.008 & 0.121 & 18.32 & 0.00 & 0.83 & 0.70 &-0.40 & 4.26 & G8V & -79,-93,-29$^b$ \\\\ HD20794$c^3$ & 0.007 & 0.204 & 40.11 & 0.00 & 0.56 & 0.70 & -0.40 & 4.26 & G8V &-79,-93,-29$^b$ \\\\ HD20794$d^3$ & 0.015 & 0.350 & 90.31 & 0.00 & 0.85 & 0.70 & -0.40 & 4.26 & G8V & -79,-93,-29$^b$\\\\ HD85512$b^3$ & 0.011 & 0.260 & 58.43 & 0.11 & 0.77 & 0.69 &-0.33 & 7.67 & K5V & 34,11,-5$^b$ \\\\ HD4308$b^4$ & 0.048 & 0.119 & 15.56 & 0.00 & 4.07 & 0.93 & -0.31 & 6.55 & G3V & 50,-110,-27$^b$ \\\\ HD40307$b^5$ & 0.013 & 0.047 & 4.31 & 0.00 & 1.97 & 0.74 &-0.31 & 7.17 &K2.5V & 3,-25,-18$^b$ \\\\ HD40307$c^5$ & 0.021 & 0.080 & 9.62 & 0.00 & 2.47 & 0.74 & -0.31 & 7.17 & K2.5V & 3,-25,-18$^b$ \\\\ HD40307$d^5$ & 0.028 & 0.132 & 20.46 & 0.00 & 2.55 & 0.74 & -0.31 & 7.17 & K2.5V & 3,-25,-18$^b$ \\\\ HD97658$b^6$ & 0.020 & 0.080 & 9.50 & 0.13 & 2.36 & 0.75 & -0.30 & 7.78 & K1V & -11,-1,-2$^c$ \\\\ GJ674$b^7$ & 0.035 & 0.039 & 4.69 & 0.20 & 8.70 & 0.35 &-0.28 & 9.36 & M2.5V & -15,-5,-19$^d$ \\\\ HD7924$b^8$ & 0.029 & 0.057 & 5.40 & 0.17 & 3.87 & 0.83 & -0.15 & 7.18 & K0V & 13,-17,-9$^e$ \\\\ GJ176$b^9$ & 0.026 & 0.066 & 8.78 & 0.00 & 4.12 & 0.49 & -0.10 & 9.97 & M2V & -26,-62,-13$^e$ \\\\ GJ581$b^{10}$ & 0.050 & 0.041 & 5.37 & 0.03 &12.65 & 0.31 &-0.10 & 10.60 & M5 & -25,-26,12$^d$ \\\\ GJ581$c^{10}$ & 0.017 &0.073 &12.92 & 0.07 &3.18 & 0.31 & -0.10 & 10.60 & M5 & -25,-26,12$^d$ \\\\ GJ581$d^{10}$ & 0.019 & 0.218 & 66.64 & 0.25 & 2.16 & 0.31 & -0.10 & 10.60 & M5 & -25,-26,12$^d$ \\\\ GJ581$e^{10}$ & 0.006 & 0.028 & 3.15 & 0.32 & 1.96 & 0.31 &-0.10 & 10.60 & M5 & -25,-26,12$^d$ \\\\ BD-082823$b^{11}$ & 0.046 &0.056 & 5.60 & 0.15 &6.50 & 0.74 & -0.07 & 9.99 & K3V & --\\\\ HD69830$b^{4}$ & 0.032 & 0.078 & 8.67 & 0.10 & 3.51 & 0.85 &-0.06 & 5.95 & K0V & 29,-61,-10$^b$ \\\\ HD69830$c^{4}$ & 0.037 &0.185 &31.56 & 0.13 &2.66 & 0.85 & -0.06 & 5.95 & K0V & 29,-61,-10$^b$ \\\\ HD125595$b^{12}$ & 0.042 &0.081 & 9.67 & 0.00 & 4.79 & 0.76 &0.02 & 9.03 & K3/K4V & -34,-42,9$^e$ \\\\ 61Vir$b^{13}$ & 0.016 &0.050 & 4.21 & 0.12 &2.12 & 0.94 &0.05 & 4.87 & G5V & -24,-47,-32$^b$ \\\\ 61Vir$c^{13}$ & 0.033 & 0.217 & 38.02 & 0.14 & 2.12 & 0.94 & 0.05 & 4.87 & G5V & -24,-47,-32$^b$ \\\\ HD156668$b^{14}$ & 0.013 & 0.050 & 4.65 & 0.00 & 1.89 & 0.77 & 0.05 & 8.42 & K3V & -- \\\\ HD10180$c^{15}$ & 0.042 & 0.064 & 5.76 & 0.08 & 4.54 & 1.06 & 0.08 & 7.33 & G1V & 9,-16,-30$^b$ \\\\ HD10180$d^{15}$ & 0.038 & 0.129 & 16.36 & 0.14 & 2.93 & 1.06 & 0.08 & 7.33 & G1V & 9,-16,-30$^b$ \\\\ HIP57274$b^{16}$ & 0.037 & 0.071 & 8.14 & 0.19 & 4.64 & 0.73 & 0.09 & 8.96 & K8 & 2,-34,27$^e$ \\\\ HD1461$b^{17}$ & 0.024 & 0.064 & 5.77 & 0.14 & 2.70 & 1.03 & 0.18 & 6.60 & G0V & -31,-39,-1$^b$ \\\\ HD215497$b^{18}$ & 0.021 & 0.047 & 3.93 & 0.16 & 2.98 & 0.87 & 0.23 & 9.10 & K3V & -- \\\\ $\\mu$Ara$d^{4}$ & 0.035 & 0.093 & 9.64 & 0.17 & 3.06 & 1.15 & 0.29 & 5.12 & G3V & -15,-7,-3$^e$ \\\\ 55Cnc$e^{10}$ & 0.025 &0.015 & 0.74 & 0.06 &5.92 & 0.90 &0.31 & 5.96 & G8V & -37,-18,-8$^d$\\\\ GJ876$d^{19}$ & 0.018 &0.021 & 1.94 & 0.21 &6.56 & 0.32 &0.37 & 10.20 & M5 & -13,-20,-12$^d$ \\\\ GJ876$e^{19}$ & 0.039 &0.333 &124.26 & 0.05 &3.42 & 0.32 &0.37 & 10.20 & M5 & -13,-20,-12$^d$ \\\\ \\hline \\end{tabular} \\medskip $^1$\\citet{bonfils11}; $^2$\\citet{anglada-escude12}; $^3$\\citet{pepe11}; $^4$\\citet{valenti05}; $^5$\\citet{mayor09}; $^6$\\citet{henry11}; $^7$\\citet{bonfils07} $^8$\\citet{howard09}; $^9$\\citet{forveille09}; $^{10}$\\citet{vonbraun11}; $^{11}$\\citet{hebrard10}; $^{12}$\\citet{segransan11}; $^{13}$\\citet{vogt10}; $^{14}$\\citet{howard11} $^{15}$\\citet{lovis11}; $^{16}$\\citet{fischer12}; $^{17}$\\citet{rivera10}; $^{18}$\\citet{locurto10}; $^{19}$\\citet{johnson09} \\newline $^a$\\citet{anglada-escude12}; $^b$\\citet{holmberg09}; $^c$\\citet{marsakov88}; $^d$\\citet{montes01}; $^e$\\citet{woolley70} \\end{table*} The high abundance of all the elements we have analysed also indicates that the proto-planetary disk left over from the formation of HD77338 was rich in planet building material. The established preponderance for gas giant planets to favour metal-rich stars (\\citealp{gonzalez97}; \\citealp{fischer05}; \\citealp{sousa11}) can be seen in Fig.~\\ref{rv_met_mass}. We also see that metal-rich stars tend to cover the entire phase space of planetary masses. Above a host star metallicity of $\\sim$0.2~dex there are planets covering the whole regime from low-masses to high-masses, in dense clusters. However, for lower metallicity systems, particularly at sub-solar metallicities, there are regions free from any planet detections, or at least less densely packed. This shows why metal-rich stars are so highly prized in the hunt for exoplanets and for better understanding the nature of planet formation and migration. \\begin{figure} \\vspace{7.0cm} \\hspace{-4.0cm} \\special{psfile=fig13_1.eps hscale=40 vscale=40 angle=90 hoffset=385 voffset=-25} \\vspace{0cm} \\caption{Metallicity against minimum mass for all radial velocity detected planets with a minimum mass less than that of Neptune (\\msini~$\\le$~0.05~M$_{\\rm{J}}$). The horizontal dotted line highlights the canonical boundary where runaway gas accretion onto the growing core is expected to occur. The dashed line marks the lower boundary in mass and the data points have been scaled in size by their orbital period. The filled circles connected by a straight line show the position of HD77338$b$ using our metallicity and the metallicity of \\citet{trevisan11}, respectively.} \\label{rocky_met_mass} \\end{figure} \\subsection{Low-mass Systems} Lower mass systems may show a somewhat different metallicity distribution in comparison to gas giants (\\citealp{udry06}; \\citealp{neves09}), with no clear metallicity bias. Models indicate that gas giant planets form outside the ice line boundary, taken to be $\\sim$5~AU, in metal-poor disks (\\citealp{mordasini12}) meaning smaller rocky/icy planets can form interior to this boundary, giving rise to the observed population of short period rocky planets around stars with a lower metal content. Fig.~\\ref{rocky_met_mass} shows the distribution of sub-Neptune mass planets (\\msini$\\le$0.05M$_{\\rm{J}}$) as a function of metallicity using data taken from the Exoplanet Database\\footnote{http://www.exoplanets.org} as of February 2012. All values we have used are listed in Table~\\ref{tab:planets}. One of the most striking features we see is the lack of planets with high metallicities and very low masses, in the lower right corner of the plot. Even though the numbers are relatively small at present, there appears a lower boundary that increases with metallicity, which would be contrary to the presumed biases in current radial velocity surveys given the number of metal-rich programs currently running. HD77338 also appears to follow this trend given that it has a very high metallicity and a minimum mass commensurate with that of Uranus in our solar system. A strong bias that could be manifest here is from the distribution of the orbital periods of the sample. Given the high fraction of short period gas giants in metal-rich systems, it is possible that lower mass planets reside at longer orbital periods, meaning they will be biased against due to the well known radial velocity bias against longer period and low amplitude objects (see \\citealp{cumming04}). To test this we scale the data points in Fig.~\\ref{rocky_met_mass} by the ratio between each planet's period and that of the longest period planet in the set, GJ876$e$. This shows us that this is not the case and that as expected, the planets are mostly short period planets at the perimeter of this boundary. It is also noticeable that beyond the cluster of planets near the perimeter of the boundary we mark by the dashed line, there is a possible valley in the metallicity plane, then the rest of the sample is found. A metallicity-mass correlation like this, at low-masses, would fit well with the core accretion model for planet formation, where the low-mass planets around super metal-rich stars attain a higher mass due to the abundance of planet forming material in the disk (\\citealp{ida04b}; \\citealp{mordasini12}). What is not obvious is why there would be a lower boundary. Possibly the planetesimals quickly accrete material and grow much faster than at lower metallicities, meaning the boundary traces the planet formation timescales for a given disk metallicity. In any case, if there is a lower boundary that holds to long period orbits for super metal-rich stars, this has consequences for the fraction of habitable Earth-mass planets in the galaxy, as super metal-rich stars will not be abundant in such planets. \\subsubsection{Monte Carlo Test} To test if the lack of planets in the lower right corner of the metallicity-mass plane is statistically significant, and hence a metallicity-mass correlation for low-mass rocky planets is statistically significant, we perform a Monte Carlo analysis. We setup the model in two parts where we first generate a random realisation of the sample of 36 low-mass planets we have tested in Fig.~\\ref{rocky_met_mass}. To generate the masses in a robust fashion we take the observationally driven mass distribution for lower-mass planets (\\citealp{butler06}; \\citealp{cumming08}). Lopez \\& Jenkins (2012, in preparation) show that a power law model such as this explains the observed turnover of the mass distribution at the lowest masses and hence the model is a robust representation of the current distribution of exoplanets. Eq$^{\\rm{n}}$~\\ref{eq:massfn} shows the form of the mass function we consider. \\begin{equation} \\label{eq:massfn} dN/dM = k M^{-1.2} \\end{equation} Here dN/dM represents the frequency of planets as a function of mass, M is the mass bin considered, and $k$ is a scaling constant to fit the observed population. We can then integrate this mass function equation (Eq$^{\\rm{n}}$~\\ref{eq:massfnint1}) to get the result shown in Eq$^{\\rm{n}}$~\\ref{eq:massfnint}, and by taking the inverse we can generate the cumulative distribution function (CDF; Eq$^{\\rm{n}}$~\\ref{eq:massfncdf}) that can be used to draw masses randomly from the observed population of exoplanets. \\begin{equation} \\label{eq:massfnint1} dN = k \\int_{-\\infty}^{M_{o}} M^{-1.2} dM \\end{equation} \\begin{equation} \\label{eq:massfnint} N = \\frac{-5k}{M^{0.2}} \\end{equation} Now if we set the constant $k$ equal to -0.014 this will normalize the function such that we have a probability density function and the corresponding CDF is given by: \\begin{equation} \\label{eq:massfncdf} CDF(M) = \\left(\\frac{0.069}{N} \\right)^5 \\end{equation} The CDF allows us to draw random masses for planets and restrict the values to be within the range between 0.006 to 0.050~M$_{\\rm{J}}$ to match our region of interest. We can then randomly generate metallicities, and in this test we simply draw from a uniform distribution of [Fe/H] values in the range from -0.5 to +0.5~dex, which covers all the values of metallicities in our test. Once we have built a random sample of planets we can then add the boundary and test how many times there are planets that reside under the boundary region we identify in Fig.~\\ref{rocky_met_mass}. We run the code 1'000'000 times to ensure a high level of robustness in the final probability measurement. Our test reveals that for a sample of 36 stars, assuming a uniform metallicity distribution, 99.9993\\% of the time there are planets to be found below the boundary region. Therefore, this \\emph{low-mass planet desert} is statistically significant at almost the $\\sim$4.5$\\sigma$ level, under these test conditions. For a sample of only 10 planets we still find that 96\\% of the time there are planets below the boundary region, and for a sample of 5 planets we find a percentage of 81\\%. These tests indicate that there is some significant correlation between metallicity and mass for low-mass planets, at least in the metal-rich regime, that would be important to quantify in the future. We do note that our results can vary due to the assumed metallicity model that we use to draw the sample of random metallicities from. For instance, the result will become more significant for a distribution that follows the current observed distribution for more massive exoplanets (see \\citealp{sousa11}), and vice-versa for a distribution shaped in the opposite fashion. As for the possible valley, the sample is as yet too small to draw statistically significant conclusions and therefore difficult to test without detailed modeling of the core accretion method of forming planets, therefore we are only pointing out the possibility that there are two classes of low-mass planets that have a metallicity dependency. In particular, magnitude or distance limited samples will be biased towards sub-solar metallicities as they will be governed by the metallicity distribution of the local galactic neighbourhood (\\citealp{holmberg09}). However, to fully probe these samples it is necessary to simulate the radial velocity data for each star individually (\\citealp{otoole09}) to better understand the completeness of the bins we have discussed. \\subsubsection{What does Kepler Say?} \\begin{figure} \\vspace{7.0cm} \\hspace{-4.0cm} \\special{psfile=fig13_2.eps hscale=40 vscale=40 angle=90 hoffset=385 voffset=-25} \\vspace{0cm} \\caption{Kepler results for 156 planet-hosting stars detected by Kepler in the metallicity - radius plane. The dashed box region in the lower right corner highlights the lack of low-mass Kepler transits in comparison to the more metal-poor region.} \\label{kepler_met_mass} \\end{figure} The Kepler Space Telescope has recently released a plethora of transiting planet candidates, that range in radius from below Earth radii up into the gas giant planet arena (\\citealp{borucki11}). Follow-up of Kepler targets are difficult due to the relative faintness of the candidate stars, however, \\citet{buchhave12} have published the first large set of spectroscopic metallicity for 156 Kepler planet-host stars that can be used to draw statistically significant conclusions on Kepler host star abundance patterns. In Fig.~\\ref{kepler_met_mass} we show the Kepler results published in Buchhave et al. in the metallicity - radius plane. We only show the stars that host planets with radii of 4$R_{\\oplus}$ or less, and we do not distinguish between multiplanet systems and single planet host stars. The 4$R_{\\oplus}$ limit was chosen simply because typical mass-to-radius relationships, e.g. $M_{\\rm{p}}$~=~$R_{\\rm{p}}^{2.06}$ (\\citealp{lissauer11}), would give rise to around the Neptune-mass limit highlighted in Fig.~\\ref{rocky_met_mass}. The paucity of planets we find in the metallicity - mass plane from the radial velocity data could therefore be manifest in the Kepler data too. However, we note that a single mass-to-radius relationship is surely unrealistic due to the possible diversity of rocky planets (see \\citealp{seager07}). Studying the lower right corner of Fig.~\\ref{kepler_met_mass}, the high metallicity and low mass regime, we do find a paucity of planets when we compared to the lower left corner, or the low metallicity and low mass region. To better highlight this relative $planet desert$ we bound the region by a dashed box and also a similar rising boundary line (solid curve) as we show in the metallicity - mass plane. The Kepler data is in good agreement with the radial velocity sample in respect to this paucity of planets, although not likely in absolute value, with only two planets significantly below the boundary region we highlight. In fact, within the boxed region, parameterised by the host star having a metallicity of solar or above and the transiting planet having an Earth radius or lower, there are only 8 planets, and only one of these planets is currently in a single system. Therefore, it could be that really low-mass planets can exist in the super metal-rich regime if they are part of multi-planet systems, since under core accretion the planetesimals will compete for material to form and this will mean less material for each of the planetesimals to reach higher core masses. If these features however are real, then for a given stellar/disk metallicity it appears that planetesimals quickly grow to their boundary region and stop there, or they then quickly make the transition to around 2 to 3 times their boundary mass and either continue growing to higher mass planets before the disk dissipates, or stay where they are due to depletion of the disk material." }, "1207/1207.2157_arXiv.txt": { "abstract": "We investigate the relation between total X-ray emission from star-forming galaxies and their star formation activity. Using nearby late-type galaxies and ULIRGs from Paper~I and star-forming galaxies from {\\it Chandra} Deep Fields, we construct a sample of 66 galaxies spanning the redshift range $z\\approx 0-1.3$ and the star-formation rate (SFR) range $\\sim 0.1-10^3\\,M_{\\odot}\\,\\rmn{yr}^{-1}$. In agreement with previous results, we find that the $L_{\\rmn{X}}-\\rmn{SFR}$ relation is consistent with a linear law both at $z=0$ and for the $z=0.1-1.3$ CDF galaxies, within the statistical accuracy of $\\sim 0.1$ in the slope of the $L_{\\rmn{X}}-\\rmn{SFR}$ relation. For the total sample, we find a linear scaling relation $L_{\\rmn{X}}/\\rmn{SFR}\\approx (4.0\\pm 0.4) \\times 10^{39}(\\rmn{erg}\\,\\rmn{s}^{-1})/(M_{\\odot}\\,\\rmn{yr}^{-1})$, with a scatter of $\\approx 0.4$ dex. About $\\sim 2/3$ of the 0.5--8 keV luminosity generated per unit SFR is expected to be due to HMXBs. We find no statistically significant trends in the mean $L_{\\rmn{X}}/\\rmn{SFR}$ ratio with the redshift or star formation rate and constrain the amplitude of its variations by $\\lesssim 0.1-0.2$ dex. These properties make X-ray observations a powerful tool to measure the star formation rate in normal star-forming galaxies that dominate the source counts at faint fluxes. ", "introduction": "High-mass X-ray binaries (HMXBs) and the hot ionized inter-stellar medium (ISM) are the main contributors to the total X-ray output of normal (i.e. not containing a luminous active galactic nucleus (AGN)) star-forming galaxies. It is well established that the collective X-ray luminosity of HMXBs well correlates with the star formation activity of the host galaxy \\citep[][hereafter Paper~I]{2003MNRAS.339..793G, 2003A&A...399...39R, 2010ApJ...724..559L, 2011AN....332..349M, 2012MNRAS.419.2095M}. The hot ionized ISM contributes about $\\sim 1/4$ to the observed X-ray emission from late-type galaxies in the standard X-ray band (0.5--8 keV), its luminosity has been also shown to scale linearly with star formation rate (SFR) \\citep[][hereafter Paper~II]{2005ApJ...628..187G, 2012arXiv1210.2997L, 2012MNRAS.426.1870M}. Thus, it has been proposed that the total, integrated X-ray luminosity from star-forming galaxies can be used as a proxy of the SFR \\citep{2003MNRAS.339..793G, 2003A&A...399...39R}. Although not entirely free from its own systematic uncertainties and contaminations, the X-ray based SFR proxy is less affected by the interstellar extinction and cosmological passband redshift, than conventional SFR indicators. Furthermore, the $L_{\\rmn{X}}-\\rmn{SFR}$ scaling relation does not experience significant cosmological evolution up to redshifts of $z\\sim 1-2$. This has been initially suggested based on direct measurements of the $L_{\\rmn{X}}/\\rmn{SFR}$ ratios for several galaxies in Chandra Deep Fields \\citep{2003MNRAS.339..793G, 2008ApJ...681.1163L, 2010ApJ...724..559L, 2012MNRAS.419.2095M} and was further supported by calculations of the maximal contribution of X-ray faint star-forming galaxies to the unresolved part of the Cosmic X-ray background \\citep{2012MNRAS.421..213D} and stacking analysis results \\citep{2012ApJ...748...50C}. These properties make the X-ray based SFR proxy a powerful tool to measure the star formation rate in distant galaxies. The most significant systematic effect which can compromise the X-ray-based SFR measurements is the contamination by the emission of the central supermassive black hole, the AGN. Indeed, even low luminosity AGN, with $\\log(L_{\\rmn{X}})\\sim 42$ can outshine a $\\sim 100 M_\\odot$/yr starburst. As populations of bright galaxies are mainly composed of AGN \\citep[for recent results, see e.g.][]{2011ApJS..195...10X, 2012ApJ...752...46L}, SFR measurements using X-ray luminosity can be applied only to a relatively small fraction of bright galaxies. In this case, a careful investigation of the nature of each galaxy is required in order to separate late-type from early-type galaxy populations. On the contrary, among faint sources, $F_{\\rmn{X}}\\la 10^{-17}$ erg~cm$^{-2}$~s$^{-1}$, the majority are star-forming galaxies located at moderate and large redshifts, $z\\sim 0.5-3$ \\citep{2012ApJ...748...50C,2012ApJ...752...46L}. This makes the X-ray based SFR proxy a powerful tool to measure the star formation rate in faint galaxies, where it can be used {\\em en masse}, to infer the cosmic star formation history \\citep[e.g.][]{2012ApJ...748...50C}. The aim of this paper is to obtain the $L_{\\rmn{X}}-\\rmn{SFR}$ scaling relation for the {\\em total} X-ray luminosity and to investigate its behavior in a broad range of redshifts. For the redshift $z=0$, we use the sample of star-forming galaxies and ULIRGs (ultra-luminous infrared galaxies) from Papers I and II. We then select normal star-forming galaxies from the {\\it Chandra} Deep Fields (CDFs) expanding the local sample towards cosmologically interesting redshifts and high star formation rates. We combine these data in order to calibrate the $L_{\\rmn{X}}-\\rmn{SFR}$ scaling relation over a broad range of redshifts and star formation rates. The structure of the paper is as follows. In Section 2 we briefly summarize the selection criteria and properties of the local sample. In Section 3 we describe selection of late-type galaxies from the CDF data and the procedures used to calculate their X-ray luminosities and star formation rates. The $L_{\\rmn{X}}-\\rmn{SFR}$ relation is derived in Section 4 and its redshift and SFR dependences are investigated in Section 5. In Section 6 we summarize our findings. Throughout this paper we assume a flat $\\Lambda$CDM cosmology with $H_0 = 70$ km/s/Mpc, $\\Omega_M = 0.3$ and $\\Omega_\\Lambda = 0.7$. ", "conclusions": "Based on the sample of nearby resolved galaxies more distant ULIRGs at intermediate distances and star-forming galaxies from the CDFs, we construct a sample of 54 star-forming galaxies spanning the range of redshifts from $z=0$ up to $z=1.3$ and the range of star formation rates $\\rmn{SFR}\\sim 0.1-10^3$ M$_\\odot$/yr (Fig.\\ref{fig:m-sfr}). Using this sample, we calibrate the $L_{\\rmn{X}}-\\rmn{SFR}$ relation for the 0.5--8 keV band luminosity (Fig.\\ref{fig:ltot_sfr}). We find that $L_{\\rmn{X}}-\\rmn{SFR}$ dependences for the local and CDF samples are consistent with linear relations with the typical accuracy of $\\sim 0.1$ in the slope. The linear $L_{\\rmn{X}}-\\rmn{SFR}$ relation obtained for the entire sample is given by eq.(\\ref{eq:ltot_sfr_all}). We did not find any statistically significant trends in the scaling relation with the redshift and star formation rate with the upper limit on the possible variations in the $L_{\\rmn{X}}/\\rmn{SFR}$ ratio of $\\sim 0.1-0.2$ dex (a factor of $\\sim 1.3-2.6$) (Fig.\\ref{fig:z_dep}). This property makes the X-ray emission a powerful tool to measure star formation rate in a broad range of redshifts and star formation regimes, which can be applied {\\em en masse} to faint distant galaxies." }, "1207/1207.2682_arXiv.txt": { "abstract": "Molecular line emission from protoplanetary disks is a powerful tool to constrain their physical and chemical structure. Nevertheless, only a few molecules have been detected in disks so far. We take advantage of the enhanced capabilities of the IRAM 30m telescope by using the new broad band correlator (FTS) to search for so far undetected molecules in the protoplanetary disks surrounding the TTauri stars DM Tau, GO Tau, LkCa15 and the Herbig Ae star MWC\\,480. We report the first detection of HC$_3$N at 5$\\sigma$ in the GO Tau and MWC 480 disks with the IRAM 30-m, and in the LkCa 15 disk (5 $\\sigma$), using the IRAM array, with derived column densities of the order of $10^{12}$cm$^{-2}$. We also obtain stringent upper limits on CCS (N $< 1.5 \\times\\ 10^{12} \\textrm{cm}^{-3}$). We discuss the observational results by comparing them to column densities derived from existing chemical disk models (computed using the chemical code Nautilus) and based on previous nitrogen and sulfur-bearing molecule observations. The observed column densities of HC$_3$N are typically two orders of magnitude lower than the existing predictions and appear to be lower in the presence of strong UV flux, suggesting that the molecular chemistry is sensitive to the UV penetration through the disk. The CCS upper limits reinforce our model with low elemental abundance of sulfur derived from other sulfur-bearing molecules (CS, H$_2$S and SO). ", "introduction": "Understanding the formation of planetary systems requires an in-depth study of the initial conditions, i.e. the structures of the protoplanetary disks. Several theoretical works have been done \\citep[e.g.][and references therein]{Bergin_etal2007}, leading to the current picture of a flared disk consisting of three layers \\citep[e.g.][]{vanZadelhoff_etal2001,Bergin_etal2007}. However, key model parameters such as gas density and temperature remain so far poorly constrained by observations due to the limited sensitivity and spatial resolution of current instruments. Observations of molecular lines have proven to be an excellent tool to study the physical and chemical structure and the dynamics of protoplanetary disks \\citep[e.g.][]{Dutrey_etal1997,Kastner_etal1997,Guilloteau_Dutrey_1998,Pietu_etal2007,Qi+etal2008,Semenov_Wiebe2011,Oberg_etal2012}. Depending on the molecule and transition, the observations trace different physical conditions and, therefore, sample chemically different regions in the disks. In gas-rich protoplanetary disks the \\cp\\ emission likely comes from the ionized upper part of disks \\citep{Semenov_etal2004, jonkheid_etal2007, chapillon_etal2010, panic_etal2010, kamp_etal2011, bruderer_etal2012}. Atomic species such as O{\\small I} have also been detected by the Herschel satellite but its origin (warm atmosphere, disk wind or jet?) is still debated \\citep{Mathews+2010}. CO and its main isotopomers are characteristics of the molecular layers which are located up to a few scale heights only above the mid-plane \\citep{Dartois_etal2003}. Near the disk plane, the gas is very cold, so that molecules are expected to be depleted on grains. Such regions may be only traced by \\hhd\\ \\citep{chapillon_etal2012, oberg_etal2011-h2d}. To retrieve useful information and to reconstruct the 3D (physical and chemical) structure of disks, molecular observations need to be compared with chemical models dedicated to protoplanetary disks \\citep{Semenov_etal2010, Vasyunin_etal2011}. Such studies help characterise the dominant processes in the disk leading to planet formation (e.g., grain growth, molecule formation and destruction, gas-to-dust ratio evolution). So far, in the mm/submm range which characterizes the bulk of the gas disk, only a few molecules have been firmly detected in protoplanetary disks: CO and its main isotopomers (\\tco~and \\cdo), HCO$^+$ and H$^{13}$CO$^+$, DCO$^+$, H$_2$CO, H$_2$O, CS, C$_2$H, N$_2$H$^+$, HCN, HNC, CN and DCN \\citep{Dutrey_etal1997,vanDishoeck+etal2003,dutrey_etal2007,Qi+etal2008,Henning_etal2010, oberg_etal2011,Hogerheijde_etal2011, Oberg_etal2012}. These are simple, light (mass number $< 44$) molecules. Heavier or more complex molecules remain undetected due to their low abundances and the lack of sensitivity of current instruments. Within the framework of the CID (\"Chemistry In Disks\") collaboration, we take here advantage of the improved performance of the IRAM-30m telescope to carry out a molecular survey of four well known large (R$_{out}>500$\\,AU) protoplanetary disks. The main driver of our observational project was to search for heavier molecules. We report in this observational paper the detection of the new molecule \\hcccn , in the three disks surrounding the T\\,Tauri stars GO Tau and LkCa15 and the Herbig Ae star MWC\\,480. We also present the upper limits obtained on CCS after a deep integration. A more complete chemical analysis will be presented in a forthcoming paper (Wakelam et al. in prep.). In Section 2, we present the observations and the methods used for the data reduction. Results (derivation of the best-fit model and column densities) are described in Section 3. In Section 4 the implications of our new results are discussed. Summary and conclusion are presented in Section 5. ", "conclusions": "\\label{sec:dis} \\subsection{Chemical models of protoplanetary disks} Chemical modeling is essential to retrieve the vertical structure of the disk from the observable column densities. Generic chemical models of circumstellar disks have been published by \\citet{aikawa_herbst1999}, \\citet{Willacy+Langer_2000}, \\citet{Aikawa_etal2002}, \\citet{vanZadelhoff_etal2003}, \\citet{Aikawa+Nomura_2006}, \\citet{Willacy+etal_2006} and \\citet{Fogel+etal_2011, Semenov_Wiebe2011, Vasyunin_etal2011}. Each of these works make different assumption and hypothesis. \\citet{aikawa_herbst1999} consider gas phase chemistry, along with sticking and desorption, and assume a vertically isothermal disk model following the ``minimum mass solar nebula'' extended to 800 AU. The study from \\citet{Willacy+Langer_2000} use the thermal structure derived from the \\citet{Chiang+Goldreich_1997} two-layer approximation, as well as enhanced photodesorption yields following \\citet{Westley+etal_1995}. \\citet{Aikawa_etal2002} use a D'Alessio disk model including vertical temperature gradients and self-consistently variable flaring, although dust and gas are assumed to be fully thermally coupled. All three models use a 1+1D approximation for radiative transfer, in which the stellar UV is attenuated along the line of sight to the star, while the ISRF impacts isotropically on the disk surface. \\citet{vanZadelhoff_etal2003} improved the UV treatment by using a 2-D radiative transfer code to solve for the UV field inside the disk, as did \\citet{Fogel+etal_2011}. Shielding by H$_2$ is treated in an approximate way, however. \\citet{Aikawa+Nomura_2006} expanded the models further by considering the effect of grain growth as \\citet{Woitke_etal2009}. \\citet{Woitke_etal2009} and \\citet{Fogel+etal_2011} investigated the effects of dust settling, \\citet{Fogel+etal_2011}, \\citet{Willacy+etal_2006} and \\citet{Walsh_etal2010,Walsh_etal2012} also include grain-surface chemistry, which may be important in such environments. 1D turbulent diffusion is included in the model from \\citet{Willacy+etal_2006}, and 2D in the models from \\citet{Semenov_etal2006} and \\citet{Semenov_Wiebe2011}. \\citet{Vasyunin_etal2011} have investigated dust coagulation, fragmentation, sedimentation, and turbulent stirring.. Making comparisons with our observations is difficult for several reasons. All these chemical models are not necessarily tailored to the physical conditions relevant to the specific sources we have observed. The chemical networks used are different from one code to another and the large uncertainties on reaction rates may explain some intrinsic differences \\citep[e.g.][]{Daranlot_etal2012}. In many cases the initial physical and chemical conditions are also different. \\subsection{The Nautilus models} Following our previous studies, we performed a chemical modeling using Nautilus, a gas-grain chemistry model adapted for the disk physics \\citep{Hersant_etal2009}. Nautilus computes the abundances of 460 gas-phase and 195 surface species as a function of time using the rate equation method \\citep{1992ApJS...82..167H} and the chemical network contains 4406 gas-phase reactions and 1733 reactions involving grains, including adsorption and desorption processes and grain-surface reactions (we assume 4580 K for the energy of desorption of CCS and \\hcccn). The gas-phase network is regularly updated (both by optimizing the reaction rates and by adding new reactions) according to the recommendations from the KIDA\\footnote{http://kida.obs.u-bordeaux1.fr} experts. In CIDV, we used Nautilus to perform a full chemistry modeling and only discussed the sulfur chemistry. The initial abundances were obtained computing the chemical composition of the parent cloud as discussed in CIDV. We compare here our new observational results with our best model from CIDV (model C). For the four disks, the physical conditions and elemental abundances used are summarized in Fig.2 and Table 4 of CIDV, respectively. Model C was obtained assuming a C/O ratio of 1.2 \\citep[following][]{Hincelin+etal_2011}, a sulfur and nitrogen abundance with respect to H of $8 \\times 10^{-9}$ and $6.2 \\times 10^{-5}$ \\citep{Jenkins_2009}, respectively and a grain size of 0.1 $\\mu$m with an initial cloud density of $2 \\times 10^{5}$ H.cm$^{-3}$ and a cloud age of 1 Myr. Our results are shown in Fig.\\ref{fig:chemistry}. At 1~Myr, we have an abundance of $10^{-10}$ for \\hcccn\\ and $4 \\times 10^{-13}$ for CCS in the gas-phase and of $2.5 \\times 10^{-12}$ for \\hcccn\\ and $4 \\times 10^{-11}$ for CCS, on grains. The main conclusion from CIDV was that although our chemical model reproduces the SO and CS column densities reasonably well, it fails to reproduce the upper limits obtained on H$_2$S by at least one order of magnitude. This suggests that a fraction of Sulfur may be depleted in mantles or refractory grains. At the high densities and low temperatures encountered around disk mid-planes, H$_2$S may remain locked onto the grain surfaces and react to form other species preventing desorption of H$_2$S.\\\\ \\subsection{Comparison with observations} In protoplanetary disks, only four N bearing molecules have been previously detected. NH$_3$ is not detected but N$_2$H$^+$ is observed with an abundance relative to H$_2$ of $\\sim 10^{-12} $ \\citep{dutrey_etal2007, oberg_etal2011}. CN is easy to detect \\citep{Dutrey_etal1997} and likely results from a chemistry which is more sophisticated than the simple photo-dissociation of HCN under the stellar and ambient UV fields \\citep{chapillon_etal2012}. As an example, in the DM Tau disk, the observed surface densities of CN and HCN at 300 AU are $3.5 \\cdot 10^{13}$ cm${^{-2}}$ and $6.5 \\cdot 10^{12}$ cm${^{-2}}$, respectively. HNC has been reported in DM Tau only by \\citet{Dutrey_etal1997}. \\hcccn\\ is the fifth nitrogen-bearing molecule detected in protoplanetary disks. We observe this molecule in MWC\\,480, GO Tau and LkCa 15. Fig.\\ref{fig:chemistry} presents the modeled and observed surface densities of \\hcccn\\ at 300 AU and Fig.\\ref{fig:abun} presents the modeled abundances. DM Tau and GO Tau are assumed to share the same disk physical structure and therefore the chemical predictions are identical. Fig.\\ref{fig:chemistry} (left panel) shows that our non detection of CCS remains, at the 3$\\sigma$ level, in very good agreement with the model C from CIDV. Furthermore, our CCS upper limits are incompatible with the other models we explored in CIDV. By increasing the elemental abundance of sulfur by a factor 10 (models A and B), we would produce ten times more CCS with respect to H$_2$. Fig.\\ref{fig:chemistry} (right panel) also shows that CCS is expected to be formed around 1.5 scale-heights in the outer disk (R$\\sim 300$~AU) of both T\\,Tauri and Herbig Ae stars. \\citet{Semenov_Wiebe2011} have calculated typical column densities of gaseous CCS at 300~AU to be about $10^{11}$~cm$^{-2}$ and $7\\,10^{11}$~cm$^{-2}$ in the laminar and 2D-mixing models of the DM Tau disk. The CCS layer in these models is also located around 1 pressure scale height from the midplane. Contrary to the observed behavior of sulfur-bearing molecules (CS and H$_2$S), whose surface densities are of the same order of magnitude in both sources, \\hcccn\\ is only detected in the disk surrounding GO Tau. In DM Tau, the 3$\\sigma$ upper limit is a factor 4 lower than the detection in GO Tau. Some intrinsic differences in the physical/chemical properties have to be explored to understand the origin of such a difference. \\subsection{Results} For the T\\,Tauri disks, our best chemical model predicts \\hcccn\\ surface densities which are typically two orders of magnitude larger than the observed one. Our chemical model show that \\hcccn\\ is found in a layer around 1-2 scale height in the T\\,Tauri disks (see Fig.\\ref{fig:chemistry}). In this layer, the density is of the order of $\\sim 3\\times 10^{6}$cm$^{-3}$, i.e. comparable or slightly above the expected critical densities of the transitions \\citep{Wernli_etal2007}, especially for J=16-15 transition. As we assumed LTE, this may affect our determination of the surface densities by a factor of at most a few, but not as large as $\\sim$ 100. The chemical model however appears in good agreement for the disk surrounding the Herbig Ae star MWC\\,480 in which the predicted abundance is much smaller than in the other sources. We checked that the difference is due to the UV flux, significantly stronger in the case of MWC\\,480, which photo-dissociates \\hcccn. The \\hcccn\\ abundance appears better reproduced in the presence of a strong UV field. Enhancing the UV penetration by changing vertically the dust grain properties (grain growth and vertical settling) may decrease the predicted \\hcccn\\ column densities even in sources with low or medium UV fluxes. The stronger UV flux also explains that most of the \\hcccn\\ is formed at an altitude which is lower in the Herbig Ae disk ($Z/H \\simeq 1$) than in T\\,Tauri disks ($Z/H \\simeq 1.5$, see Fig.\\ref{fig:chemistry}, right panel). As are result, the observed \\hcccn\\ transitions should be more easily thermalized in the Herbig Ae disk.\\\\ Investigating the vertical abundance variation is also very intersting. Figure \\ref{fig:abun} shows the predicted vertical abundances of \\hcccn\\ and CCS, in the gas-phase and at the surface of the grains at 300 AU in the four disks. These two molecules are present in a layer between 1 and 3 scale heights in DM Tau and at lower height in the two other disks where the stellar UV flux is stronger (410, 2550, 8500 $\\chi_0$ for DM Tau, LkCa 15 and MWC 480 respectively, see CIDV). In all models, CCS is formed in the gas-phase only and can only be destroyed on grain surfaces, after being accreted. In DM Tau, the surface abundance of CCS is smaller than in LkCa 15 because more complex S-bearing species such as C$_3$S are formed. In LkCa 15, the formation of such molecules is prevented by the stronger UV field. In MWC 480, the CCS abundance in the ices is even smaller than $10^{-12}$. The gas phase abundance of \\hcccn\\ for DM Tau is similar to the one in LkCa 15, whereas the abundance in the grain surface is much larger. The smaller surface abundance of \\hcccn\\ in LkCa 15 is due to direct photo-dissociations by UV photons. The situation in MWC 480 is very different, with a very large abundance of \\hcccn\\ in the ices and a very small one in the gas. Although the UV field is stronger than in the two other sources, the dust temperature is larger and allows the diffusion of larger radicals on the surfaces. This enhanced diffusion produces more \\hcccn\\ on grains.\\\\ In Table \\ref{tab:results-ratio}, we compare our results with those from \\citet{chapillon_etal2012}. If our prediction of the \\hcccn\\ column density in MWC\\,480 is in good agreement with the observation, it fails to reproduce the \\hcccn/CN and \\hcccn/HCN ratios. For the T Tauri disks, the ratios are in roughly good agreement with the observations while the predicted absolute values are off by two orders of magnitude. This is not surprising since this model was aimed at reproducing sulfur-bearing molecule chemistry. Using a lower elemental abundance of nitrogen would produce less N bearing molecules, \\citep[see][]{Semenov_Wiebe2011}. Moreover, some neutral-neutral reactions involving atomic nitrogen are currently being experimentally studied \\citep{Bergeat_etal2009, Daranlot_etal2011, Daranlot_etal2012}. As soon as these new experimental rates are known, a more detailed study of nitrogen-chemistry in disks will be performed \\citep{Wakelam_etal2012}. We also note that the observed CN/HCN ratio is of the order of 5 in the T Tauri stars disks, and of 10 in the disks of Herbig stars. The observed \\hcccn/HCN ratio in the three disks is very similar to that observed in the cold molecular core L134N \\citep[HC$_3$N/HCN $\\sim$ 0.07,][]{Dickens_etal2000} and still reasonably close to the cometary value within a factor $\\sim 2$ \\citep[in Hale-Bopp, HC$_3$N is observed with a ratio HC$_3$N/HCN $\\sim$ 0.1,][]{Bockelee-morvan_etal2000}. However, the observed \\hcccn/CN ratio of both Herbig Ae and T\\,Tauri disks is far from the value observed in L134N \\citep[HC$_3$N/CN $\\sim$ 1,][]{Dickens_etal2000}, suggesting that the later ratio is more affected by photo-dissociation.\\\\ These observations can be used to place the detectability of complex molecules in perspective with ALMA. Our typical 3$\\sigma$ limit is of order 50 mJy.km.s$^{-1}$ for a single line (Table \\ref{tab:fluxes}), but we gain $\\sqrt{3}$ by the multiplex factor due to wide frequency coverage. By comparison, ALMA reaches 1 mJy in 1 hour of integration time for 1 km.s$^{-1}$ resolution at 90 GHz, i.e.1$\\sigma$ flux sensitivity of 1.4 mJy.km.s$^{-1}$. ALMA should greatly improve the detectability of such molecules, but imaging at high ($\\leq 0.5''$) resolution would require many (more than 10) hours. \\begin{figure} \\centering \\includegraphics[width=8.5cm]{fig3.eps} \\caption{ Predicted column densities for \\hcccn and CCS using Nautilus. Left panel: observed and modeled surface densities. Right panel: Cumulative column densities (from mid-plane to atmosphere) in function of $Z/H$ (left axis) and Z (right axis) at 300\\,AU. Top panel: MWC\\,480, medium panel: LkCa15, Bottom panel: GO Tau and DM Tau. The (cumulative) column densities (right panel), whose asymptotic values are half of the surface densities (left panel).} \\label{fig:chemistry} \\end{figure} \\begin{figure} \\centering \\includegraphics[width=8.5cm]{fig3a.eps} \\caption{Abundances of \\hcccn\\ (solid line) and CCS (dashed line) in the vertical direction at 300 AU in the four protoplanetary disks. The black and grey lines represents the abundances in the gas-phase and in the ices.} \\label{fig:abun} \\end{figure}" }, "1207/1207.3085_arXiv.txt": { "abstract": "We present 1.4 GHz catalogs for the cluster fields Abell 370 and Abell 2390 observed with the Very Large Array. These are two of the deepest radio images of cluster fields ever taken. The Abell 370 image covers an area of 40$'$ $\\times$ 40$'$ with a synthesized beam of $\\sim$1.7$''$ and a noise level of $\\sim$5.7 $\\mu$Jy near field center. The Abell 2390 image covers an area of 34$'$ $\\times$ 34$'$ with a synthesized beam of $\\sim$1.4$''$ and a noise level of $\\sim$5.6 $\\mu$Jy near field center. We catalog 200 redshifts for the Abell 370 field. We construct differential number counts for the central regions (radius $<$ 16$'$) of both clusters. We find that the faint (S$_{1.4{\\rm \\: GHz}}<$ 3 mJy) counts of Abell 370 are roughly consistent with the highest blank field number counts, while the faint number counts of Abell 2390 are roughly consistent with the lowest blank field number counts. Our analyses indicate that the number counts are primarily from field radio galaxies. We suggest that the disagreement of our number counts can be largely attributed to cosmic variance. ", "introduction": "Radio surveys probe active galactic nuclei (AGNs) and star-forming (SF) galaxies. SF galaxies produce non-thermal radio continuum through synchrotron emission from supernova remnants. For these sources, the 1.4 GHz luminosity is found to be an accurate indicator of the star formation rate \\citep{condon92}. AGNs also emit synchrotron radiation, which is ultimately powered by accretion onto supermassive black holes (SMBHs). Observed tight correlations between SMBH mass and bulge properties \\citep{kormendy95,ferrarese00,tremaine02} plus various theoretical considerations \\citep{granato04,croton06,bower06} have led to the hypothesis that AGNs regulate star formation. Thus, radio surveys are an important tool to constrain how stars and SMBHs evolve and interact as a function of cosmic time. There is still debate over the faint radio population, especially regarding the differential source counts. AGNs are found to dominate the counts at high flux densities, while SF galaxies are thought to emerge at lower flux densities \\citep{condon89}. However, the exact composition --- and the counts themselves --- are debated below S$_{1.4\\:{\\rm GHz}}$ $=$ 100 $\\mu$Jy. The disagreement in faint radio counts from one survey to the next has been attributed mostly to instrumental and analysis effects, not to cosmic variance \\citep{condon07}. For example, for the heavily studied Hubble Deep Field-North (HDF-N), three different groups have derived faint-end number counts, with each subsequent study finding them to be incrementally higher \\citep{richards00,biggs06,morrison10}. However, studies such as \\citet{biggs06}, which present catalogs for three deep radio fields, have shown that similar instrument configurations and analysis methods still result in faint counts that are inconsistent with merely Poisson variation. It is important to obtain as many deep radio images as possible to help resolve this issue, especially since cosmic variance could be an important factor. In this paper, we present deep radio observations of two heavily-studied cluster fields. Radio surveys of cluster fields have been vital to many areas of study. For example, they have improved our understanding of cluster members by finding evidence for a population of dust-obscured SF cluster galaxies that were previously classified as post-starburst galaxies based on optical spectra. With the help of high-resolution near-infrared (NIR) and optical imaging, \\citet{smail99} interpreted radio emission from these objects as an indication of ongoing star formation. Radio surveys of cluster fields have also been useful in the study of cluster evolution. For example, \\citet[][]{morrison99} found that the population of low-luminosity radio sources in clusters rapidly increases with redshift (0.02 $\\leq$ $z$ $\\leq$ 0.41). This can be interpreted as an extension of the Butcher-Oemler effect \\citep[][]{butcher84}, since the majority of low-luminosity radio sources are found to be blue SF galaxies \\citep[][]{morrison03}. Radio data are key to multiwavelength studies. Radio emission is unobstructed by dust, which avoids a major source of bias found in UV and optical studies. With the addition of far-infrared (FIR) data, radio data can be used to indicate the dominant emission mechanism. The radio luminosity is observationally found to be tightly correlated with the FIR power for SF galaxies locally \\citep{helou85,condon92} and at high redshift \\citep{appleton04,ivison10,mao11}. Any significant departure from the FIR-radio correlation is an indication of AGN activity. In addition, the positional accuracy of radio surveys can be used to pinpoint counterparts at other wavebands. In the submillimeter, single dish telescopes have very poor resolution, and the unambiguous identification of counterparts is not possible without additional information. From the FIR-radio correlation, we know that FIR luminous SF galaxies, such as submillimeter galaxies (SMGs), are correspondingly luminous in radio emission. Thus, radio data can be used to identify SMG counterparts. \\citet{barger00}, \\citet{ivison02}, and \\citet{chapman03} found that $\\sim$60\\% of bright ($>$2 mJy) SMGs had radio counterparts. This feature allowed \\citet{chapman05} to use radio positions to target spectroscopically a large sample of bright SMGs, establishing the redshift distribution for this population. Bright SMGs are predominantly massive, dust-obscured, SF galaxies at a median redshift of $\\sim$2.2 \\citep{chapman05,alexander05}. Submillimeter observations are a redshift independent probe (due to a negative $K$-correction: 1 $<$ $z$ $<$ 8) of dust-reprocessed UV light. However, the positive $K$-correction of the radio synchrotron emission results in faint observed 1.4 GHz fluxes for high-redshift objects. Thus, a high submillimeter-to-radio flux measurement is an indication of a high-redshift source (Wang et al.\\ 2007, 2009; \\citealt{danner08}; Capak et al.\\ 2008, 2011; Schinnerer et al.\\ 2008; Daddi et al.\\ 2009a,b; Coppin et al.\\ 2009, 2010; Riechers et al.\\ 2010; \\citealt{knudsen10}). This extreme population of very massive sources at very high redshifts has been the topic of intense study and has led many authors to suggest the need for significant modifications to models of galaxy evolution \\citep[e.g., ][]{granato04,baugh05,hopkins05,lukic07}. Deep radio surveys continue to be an important tool in the discovery of these objects. Blank-field submillimeter surveys with the Submillimeter Common-User Bolometer Array \\citep[SCUBA; ][]{holland99} first resolved the bright SMGs that account for $\\sim20\\%-30\\%$ of the 850 $\\mu$m extragalactic background light \\citep[e.g., ][]{barger98,hughes98,barger99,eales99}. These surveys cannot reach the sensitivities required to detect directly the dominant population of $<2$ mJy sources because of confusion noise resulting from the coarse resolution of SCUBA. In order to detect this population, one must observe fields with massive cluster lenses to take advantage of both gravitational amplification by the lens and reduced confusion noise \\citep[][]{smail97,cowie02,knudsen08}. Deep radio images of cluster fields are therefore valuable for identifying the counterparts and determining the properties of this important faint SMG population. Such data will also be very useful for helping to interpret \\textit{Herschel} \\citep[][]{egami10} and SCUBA2 (C.-C. Chen et al.\\ 2012, in preparation) observations. In this paper, we construct deep radio catalogs of the cluster fields Abell 370 (A370) and Abell 2390 (A2390), which we have reduced and analyzed in a similar fashion. We derive the number counts for the two fields, and we discuss the influence of the cluster and the importance of cosmic variance on these counts. In our A370 catalog we include redshifts for 200 radio sources that we obtained from the literature, from unpublished work, and from our own observations. We note that, in addition to copious ancillary data, these clusters also have excellent lens models \\citep[][]{kneib93,kneib02,richard10}. Our radio catalogs in combination with the public data already available for these extensively studied fields will aid in both cluster and field galaxy studies. \\begin{figure*}[!t] \\begin{raggedright} \\includegraphics[bb=12bp 140bp 1000bp 1100bp,clip,scale=0.27]{f1a}\\includegraphics[bb=1127bp 140bp 1950bp 970bp,clip,scale=0.27]{f1b} \\par\\end{raggedright} \\caption{\\textbf{\\textit{Left}}: Contours of constant rms noise overlaid on the 40$'$$\\times$40$'$ A370 image. Contours levels are 6.5, 8.0, 12.0, and 20.0 $\\mu$Jy and are located approximately 6$'$, 11$'$, 16$'$, and 20$'$ from field center. The image has a 1$\\sigma$ rms noise level of $\\sim$5.7 $\\mu$Jy near field center and a 1.8$''$ $\\times$ 1.6$''$ synthesized beam.\\textbf{\\textit{ Right}}: Contours of constant rms noise overlaid on the 34$'$$\\times$34$'$ A2390 image. Contours levels are 5.5, 6.5, 8.0, 12.0, and 20.0 $\\mu$Jy and are located approximately 8$'$, 10$'$, 12$'$, 16$'$, and 19$'$ from field center. The image has a 1$\\sigma$ rms noise level of $\\sim$5.6 $\\mu$Jy near field center and a 1.4$''$ $\\times$ 1.3$''$ synthesized beam. } \\end{figure*} \\begin{figure*}[!t] \\begin{centering} \\includegraphics[bb=15bp 9bp 1550bp 1524bp,clip,scale=0.3]{f2} \\par\\end{centering} \\caption{The 40$'$$\\times$40$'$ A370 image. Field center is located at $02^{h}39^{m}32^{s}$, $-01^{\\circ}35'07''$ in J2000 coordinates. This is offset by $\\sim$5$'$ from the cluster center at $02^{h}39^{m}50.5^{s}$, $-01^{\\circ}35'08''$. We observe no radio halos or relics; however, smaller array configurations would be more sensitive to these large structures.} \\end{figure*} \\begin{figure*}[!t] \\begin{centering} \\includegraphics[bb=15bp 10bp 1550bp 1524bp,clip,scale=0.3]{f3} \\par\\end{centering} \\caption{The 34$'$$\\times$34$'$ A2390 image. Field center is located at $21^{h}53^{m}36^{s}$, $+17^{\\circ}41'52''$ in J2000 coordinates. This is only offset by 20$''$ from the cluster's central cD galaxy at $21^{h}53^{m}37^{s}$, $+17^{\\circ}41'44''$. We observe no radio halos or relics; however, smaller array configurations would be more sensitive to these large structures.} \\end{figure*} We describe our observations and data reduction in Section 2. In Section 3, we describe our source extraction and cataloging. In Section 4, we compare our catalog to large-area survey results for bright sources. In Section 5, we construct differential number counts for the central regions (radius $<$ 16$'$) of both clusters. We discuss the derived counts in Section 6 and summarize the paper in Section 7. Throughout this paper, our adopted cosmology is H$_{0}$ $=$ 70 km s$^{-1}$ Mpc$^{-1}$, $\\Omega_{M}$ $=$ 0.3, $\\Omega_{\\Lambda}$ $=$ 0.7. ", "conclusions": "The A370 number counts \\begin{figure*}[!t] \\begin{centering} \\includegraphics[bb=55bp 55bp 710bp 530bp,clip,angle=180,scale=0.5]{f13_color}\\caption{Illustration of the effect on our derived number counts of applying radial cluster masks. Symbols with error bars indicate the cluster counts with no mask applied (full 16$'$ radius field). The dashed lines indicate the counts with a radial area of 6$'$ from the cluster centers masked. The solid lines indicate the counts with a radial area of 12$'$ from the cluster centers masked. The red dotted line indicates the counts with spectroscopically confirmed A370 cluster members excluded. On average, the A370 counts are reduced, bringing them in closer agreement to blank field results. The A2390 counts are consistently low. Even taking into account the greater Poisson error of the more restrictive masks, the counts in the two cluster fields are found to be inconsistent with one another in the 110 - 350 $\\mu$Jy flux range.} \\par\\end{centering} {\\footnotesize (A color version of this figure is available in the online journal.)} \\end{figure*} form an upper envelope to the blank field number counts. This could potentially be explained by an over-density of cluster galaxies superimposed on the $z$ $\\sim$ 1 field galaxies. In addition to the galaxy cluster at $z$ = 0.375, Keenan et al. (2012) found galaxy over-densities at $z$ $\\sim$ 0.18 and $z$ $\\sim$ 0.25. We clearly see evidence for the $z$ $\\sim$ 0.25 over-density in the redshift distribution (see Figure 12). $ $ The A2390 number counts form a lower envelope to the blank field number counts, which is more surprising. This might indicate that the galaxy cluster at $z$=0.228 does not significantly contribute to a sparsely populated field of 37 $<$ S$_{1.4\\:{\\rm GHz}}$ $<$ 3000 $\\mu$Jy galaxies. In the following, we investigate disentangling the cluster galaxy population from the field galaxy population using two techniques. 1) \\textit{Redshift mask}: For the A370 field, we re-derive the number counts with all the spectroscopically confirmed cluster members excluded. We lack sufficient spectroscopic data to perform a similar analysis on A2390. 2) \\textit{Radial masks}: We re-derive the number counts for both fields with the cluster center masked. In other words, we re-derive the number counts within a range of annuli to eliminate the majority of cluster objects. As we reduce the influence of the cluster, we note the effect that has on the number counts. \\subsection{Redshift Mask} Our large sample of A370 redshifts allows us to re-derive the number counts after excluding known cluster members. To construct the A370 number counts, we used the 529 sources with 37 $<$ S$_{1.4{\\rm \\: GHz}}$ $<$ 3000 $\\mu$Jy that lie within 16$'$ of the field center. Of these, 167 (32\\%) have known redshifts. In Figure 12, we show the distribution of these redshifts. We define cluster members as any object with $z$$_{A370}\\pm$0.025. Rather than a physical cluster scale, we base the redshift range on a conservative estimate of the redshift error. Using this criteria, we identify 35 objects as cluster members, and we exclude them from the re-derived counts. In Figure 12, we use positive-slope diagonal lines to denote the cluster members. We note that our spectroscopic completeness declines as the 1.4 GHz flux decreases and as the projected distance from the cluster center increases. Given these biases and the proximity of the cluster members compared to the average field galaxy at $\\left\\langle {\\rm z}\\right\\rangle $ $\\sim$ 1, we expect to have succeeded in excluding more than 32\\% of the cluster members. We show the re-derived number counts in Figure 13 as a red dotted line. We find that the A370 counts are lowered by a weighted average of 5.7\\%$\\pm$6.5\\%. The estimated error in our weighted average calculation (6.5\\%) results from propagating the Poisson error of our measured counts. Although not corrected for, excluding spectroscopic cluster members will also remove the cluster's effective volume. This will reduce the survey's effective area and therefore increase the re-derived counts. This effect is small (estimated to be on the order of a few percent), since the effective volume of our radio field is large when compared to the effective volume of the cluster. Given our spectroscopic completeness of 32\\%, we expect the systematic uncertainty in our estimate of cluster influence to be dominated by any unidentified cluster members. Thus, we regard the reported 5.7\\%$\\pm$6.5\\% reduction in our re-derived counts as a lower limit. \\subsection{Radial Masks} We also re-derive the number counts with cluster center radial masks. We apply masks to eliminate all galaxies within 2$'$, 4$'$, 6$'$, 8$'$, 10$'$, and 12$'$ from the cluster center. We obtain an estimate of the effectiveness of removing cluster members from both fields using radial masks by considering the cluster evolution study of \\citet[][hereafter M99]{morrison99}. M99 investigated the radial distribution of low-luminosity (10$^{22.3}$$<$ L$_{1.4\\:{\\rm GHz}}$ $<$ 10$^{23}$ W Hz$^{-1}$) radio galaxies (LLRGs) for a $z$ $<$ 0.25 and a $z$ $\\sim$ 0.4 cluster sample. The low-redshift sample was composed of 76 cluster LLRGs, while the z$\\sim$0.4 sample was composed of 43 cluster LLRGs. From M99's derived cumulative radial density profile (his Figure 6.18), we see that a 1.9 Mpc (6$'$ radius) A370 mask should exclude $\\sim$82\\% of the cluster's LLRG population, while a 3.7 Mpc (12$'$ radius) mask should exclude $>$97\\%. The 1.3 Mpc (6$'$) and 2.6 Mpc (12$'$) A2390 masks should exclude $\\sim$82\\% and $\\sim$97\\% of the cluster's LLRGs, respectively. LLRGs correspond to galaxies with flux densities between 41 - 206 $\\mu$Jy for the A370 cluster and 129 - 648 $\\mu$Jy for the A2390 cluster. The radial distributions of higher luminosity radio sources are found to be more centrally compact. Thus, we expect our masks to reduce significantly the population of all S$_{1.4\\:{\\rm GHz}}>$129 $\\mu$Jy ($>$41 $\\mu$Jy for A370) cluster sources. There are uncertainties in this estimate, since individual clusters will have deviations from this averaged radial profile. Additionally, the M99 study notes a lack of confirmed outer cluster members in the $ $$z$ $\\sim$ 0.4 sample, which may bias this radial profile. As a check on these percentages, we note how many spectroscopically confirmed cluster members remain after applying the A370 radial masks. Figure 12 shows the redshift histogram before and after the application of the 6$'$ (histogram filled with negative-slope diagonal lines) and the 12$'$ (filled histogram) masks. Our 6$'$ and 12$'$ cluster masks remove 74\\% and 97\\% of the spectroscopically identified cluster members, respectively. This estimate, while roughly consistent with the M99 result, suffers from our spectroscopic incompleteness, which gets worse with increasing radius and decreasing flux. We conclude that our 6$'$ radial masks should significantly reduce cluster influence on the derived number counts. We adopt the 12$'$ mask result as our best estimate of the total cluster influence. In Figure 13, we show the number counts after excluding radial areas of 6$'$ (\\textit{dashed}) and 12$'$ (\\textit{solid}) from the cluster centers. These angular radii, respectively, correspond to 1.9 and 3.7 Mpc at $z$$_{A370}$ = 0.375 and 1.3 and 2.6 Mpc at $z$$_{A2390}$ = 0.228. The A370 field has a 5.0 Mpc (16$'$) radial extent, while the A2390 field has a 3.5 Mpc (16$'$) radial extent. Excluding radial areas of 6$'$ and 12$'$ from the A370 cluster center reduces the number counts by a weighted average of 3.5\\%$\\pm$7.4\\% and 7.6\\%$\\pm$10.7\\%. If we also exclude any additional spectroscopically identified cluster members not in the masked regions, then these percentages increase by 1\\%. Excluding radial areas of 6$'$ and 12$'$ from the A2390 cluster center reduces the number counts by an average of 0.7\\%$\\pm$8.8\\% and 1.1\\%$\\pm$15.1\\%. In Figure 13, we can see that the re-derived number counts in A2390 are consistently low in the 110 to 350 $\\mu$Jy range, even when compared to the lowest blank field results. Our best estimate indicates that the A370 cluster influences the derived number counts on the 10$\\%$ level. The A2390 cluster does not appear to alter significantly the derived number counts. However, the large Poisson error of our measurements weakens this conclusion. \\subsection{Comparison to Coma Cluster Counts} Coma is a local ($z$=0.0231) cluster of galaxies, while our clusters are at redshifts of $z$=0.228 and $z$=0.375. In cluster environments, the fraction of LLRGs decreases with decreasing redshift (M99). Thus, evolution effects may be significant and no direct comparison is possible. However, we may compare our results with the Coma cluster counts merely to check for consistency given this known evolution with redshift. Specifically, our finding that $z$$\\sim$0.3 cluster LLRGs are not a major component in the derived number counts would be at odds with a local measurement showing a significant LLRG contribution. \\citet[][]{miller09} surveyed two $\\sim$0.5 deg$^{2}$ 1.4 GHz VLA observations covering the core and southwest region of the Coma cluster. They presented number counts from 0.110 mJy (their 5$\\sigma$ limit) to 100 $ $mJy. Their counts have only been corrected for areal coverage. Thus, their faint counts ($\\lesssim$ 0.230 mJy) should be regarded as lower limits. At the Coma cluster's redshift LLRGs would correspond to the flux density range 16-82 mJy. No overdensity is noted in this range. Moreover, they find that their cluster counts are consistent with blank field counts (we also show this in Figure 11). The only excess of sources is in the $\\sim$1-5 mJy range, and this is not attributed to a cluster effect. We conclude that our result, indicating that field galaxies make up the vast majority of sources in our cluster fields, is consistent with the local counts found in the Coma cluster field. \\subsection{Gravitational Lensing Bias} Gravitational lensing affects number counts by boosting observed source flux and by magnifying the source plane. This lensing bias occurs over a relatively small area compared to our r $=$ 16$'$ fields. Given reasonable estimates for the cluster masses and background source redshifts, we estimate that beyond 2 Mpc there is no significant lensing bias. We believe our estimate of cluster influence does not require any modifications due to lensing considerations. \\subsection{Cosmic Variance} Our results are consistent with the faint number counts being primarily determined by field radio galaxies. We performed similar reduction and extraction procedures on both cluster fields. Further, the lack of A2390 sources is already seen in the bright radio flux bins, and these should not suffer significantly from incompleteness. For these reasons, we investigate the significance of cosmic variance. We estimate the effect of cosmic variance with the method of \\citet{somerville04}. Assuming a correlation length of 5 Mpc (consistent with the values given by \\citet{overzier03} for NVSS mJy sources) and a redshift interval of $z$ = 1 $\\pm$ 0.5, we estimate $\\sim$12\\% rms variance due to large-scale structure (see \\citealt{simpson06} for a similar calculation). Since there is no clear consensus on the `true' faint number counts, we choose to determine if our derived cluster number counts for the two fields can be brought into agreement given reasonable assumptions. This avoids having to develop some ad hoc method of averaging together faint number counts from radio surveys that have different analysis methods. An accurate comparison of our derived number counts depends on the accuracy of our derived error bars. While the estimated error does take into account Poisson error and the variance of our Monte Carlo simulations, it does not account for uncertainties inherent to the completeness simulations. For example, we determined a simulated object's major axis by randomly sampling from OM08's size distribution. Any inaccuracy in this distribution could alter our derived completeness corrections and thus, our derived number counts. We adopt 1.5$\\sigma$ error bars to account for these additional concerns. Assuming a 10\\% reduction in the A370 number counts to account for the cluster's influence, we find that our counts can be brought to within 1 - 2 sigma of each other. This suggests that cosmic variance can explain the disagreement seen in our derived number counts." }, "1207/1207.4200_arXiv.txt": { "abstract": "GRB 120422A is a nearby ($z = 0.283$) long-duration GRB (LGRB) detected by Swift with $E_{\\gamma,iso} \\sim 4.5 \\times 10^{49}$ erg. It is also associated with the spectroscopically-confirmed broad-lined Type Ic SN 2012bz. These properties establish GRB 120422A/SN 2012bz as the sixth and newest member of the class of subluminous GRB/SNe. Observations also show that GRB 120422A/SN 2012bz occurred at an unusually large offset ($\\sim$8 kpc) from the host galaxy nucleus, setting it apart from other nearby LGRBs and leading to speculation that the host environment may have undergone prior interaction activity. Here we present spectroscopic observations using the 6.5m Magellan telescope at Las Campanas. We extract spectra at three specific locations within the GRB/SN host galaxy, including the host nucleus, the explosion site, and the ``bridge\" of diffuse emission connecting these two regions. We measure a metallicity of log(O/H) + 12 = 8.3 $\\pm$ 0.1 and a star formation rate per unit area of 0.08 M$_{\\odot}$ yr$^{-1}$ kpc$^{-2}$ at the host nucleus. At the GRB/SN explosion site we measure a comparable metallicity of log(O/H) + 12 = 8.2 $\\pm$ 0.1, but find a much lower star formation rate per unit area of 0.01 M$_{\\odot}$ yr$^{-1}$ kpc$^{-2}$. We also compare the host galaxy of this event to the hosts of other LGRBs, including samples of subluminous LGRBs and cosmological LGRBs, and find no systematic metallicity difference between the environments of these different subtypes. ", "introduction": "\\footnotetext[2]{CASA, Department of Astrophysical and Planetary Sciences, University of Colorado 389-UCB, Boulder, CO 80309, USA; \\texttt{Emily.Levesque@colorado.edu}} \\footnotetext[3]{Einstein Fellow} \\footnotetext[4]{Harvard-Smithsonian Center for Astrophysics, 60 Garden St., Cambridge, MA 02138, USA} Recent work on LGRBs at $z < 1$ has suggested a connection between LGRBs and low-metallicity host environments. Their host galaxies, on average, fall below the luminosity-metallicity and mass-metallicity relations for star-forming galaxies out to $z \\sim 1$ (e.g. Stanek et al.\\ 2006; Levesque et al.\\ 2010a,b; Mannucci et al.\\ 2011). However, the physical mechanism driving this apparent metallicity trend is still poorly understood. LGRBs do not appear to be exclusive to low-metallicity environments, with several super-solar host galaxies and explosion sites for LGRBs (e.g. Levesque et al.\\ 2010b,c). There is also no apparent correlation between host metallicity and gamma-ray energy release for LGRBs (Levesque et al.\\ 2010e), a result that is at odds with previous predictions of LGRB progenitor models (MacFadyen \\& Woosley 1999). However, it is possible that our current picture of these objects is oversimplified. There appears to be evidence for multiple sub-classes of LGRBs. Detailed studies of nearby LGRBs have revealed a subset of these events with unusually low gamma-ray energies and luminosities ($E_{\\gamma,iso} \\lesssim 10^{50}$ erg and $L \\lesssim 10^{49}$ ergs s$^{-1}$, e.g. Kulkarni et al.\\ 1998, Soderberg et al.\\ 2006a, Guetta \\& Della Valle 2007). These subluminous events, which dominate the $z \\lesssim 0.3$ LGRB population, are thought to be much more frequent that the higher-luminosity ($E_{\\gamma,iso}$ $\\sim 10^{52}$) cosmological LGRBs detected at higher redshifts. Each subluminous LGRB is also associated with a spectroscopically identified supernova (see Woosley \\& Bloom 2006 for a review). However, supernova associations are not restricted to only subluminous LGRBs: GRB 030329/SN 2003dh ($z = 0.168$) and GRB 091127/SN 2009nz ($z = 0.49$) are both associated with spectroscopically-confirmed Ic-BLs despite having ``cosmological\" luminosities (Stanek et al.\\ 2003, Berger et al.\\ 2011), and a number of other more distant bursts have shown late-time photometric rebrightenings in their afterglow lightcurves from associated SNe (e.g. Bloom et al.\\ 2002a, Soderberg et al.\\ 2005, 2006b; Cano et al.\\ 2011). Conversely, two subluminous LGRBs (060505 and 060614)have also been observed that show {\\it no} evidence of any associated supernovae (Fynbo et al.\\ 2006) although the classification of these bursts and their connection to the general LGRB population is still uncertain (e.g. Gal-Yam et al.\\ 2006, Ofek et al.\\ 2007, Zhang et al.\\ 2007, Th\\\"{o}ne et al.\\ 2008). It is currently unclear whether subluminous LGRBs represent a phenomenologically-distinct subclass within the broader LGRB sample (see Cobb et al.\\ 2006, Zhang et al.\\ 2012). The sixth and newest member of this potential subclass of subluminous GRB/SNe, GRB 120422A, was detected by the {\\it Swift} Burst Alert Telescope on 12 April 22 at 07:12:03 UT (Troja et al.\\ 2012). Prompt emission observations determined a duration of $T_{90} \\sim 5$ s, while early follow-up observations measured a redshift of $z = 0.283$ based on Mg II absorption in the optical afterglow of the GRB as well as nebular emission features from the presumed host galaxy, SDSS J090738.51+140108.3 (Schulze et al.\\ 2012, Tanvir et al.\\ 2012). Subsequently an associated Ic-BL supernova, SN 2012bz, was spectroscopically confirmed by Wiersema et al.\\ (2012) and found to be very similar to other Ic-BLs associated with LGRBs (Melandri et al.\\ 2012). The total isotropic energy of the burst was measured to be $E_{\\gamma, iso} \\sim 4.5 \\times 10^{49}$ erg, with a peak energy of $\\sim$53 keV, marking it as subluminous compared to the general LGRB population (Schulze et al.\\ 2012, Zhang et al.\\ 2012). GRB 120422A/SN 2012bz is unique among nearby LGRBs due to its localization at an unusually large offset from the center of its host galaxy - Tanvir et al.\\ (2002) measure a projected physical offset of $\\sim$8 kpc, much larger than the median offset measured in the sample of Bloom et al.\\ (2002b). Such an offset is one of the largest observed for an LGRB, which are typically localized in the brightest and bluest regions of their hosts (Bloom et al.\\ 2002b, Fruchter et al.\\ 2006). This suggests that GRB 120422A occurred in a star-forming region near the outskirts of the host, similar to other events such as GRB 980425 and GRB 990705 (e.g. Christensen et al.\\ 2008, Bloom et al.\\ 2002b); however, the absence of clearly identified spiral arms in this host has led to speculation that the star-forming region hosting this burst may have been produced by an interacting system (Tanvir er al.\\ 2012, Perley et al.\\ 2012, Sanchez-Ramirez et al.\\ 2012). Here we present spectroscopy of several locations within the GRB 120422A/SN 2012bz host galaxy. We discuss the observations and describe the reduction and analysis applied to these spectra in \\S2. Based on these we derive ISM properties within the host (\\S3) and place GRB 120422A/SN 2012bz in context with the larger LGRB host environment population, considering the implications for our understanding of LGRBs and the subluminous LGRB subclass (\\S4). Throughout this work we adopt the standard cosmological parameters $H_0=71$ km s$^{-1}$ Mpc$^{-1}$, $\\Omega_m=0.27$, and $\\Omega_\\Lambda=0.73$. ", "conclusions": "The association of GRB 120422A/SN 2012bz with its host galaxy is robust (see Schulze et al.\\ 2012). However, it is clear from our spectra, particularly the weak emission features and SFR at the GRB/SN Site, that the explosion environment is very weakly star-forming relative to the rest of the host (unlike, for example, GRB 100316D, which was localized near the strongest star-forming region in its host; Levesque et al.\\ 2011). It is worth noting that several spectral features, most notably [OII] $\\lambda$3727, [OIII] $\\lambda$5007, and H$\\alpha$ (see Figure 2), show signs of blue-shifted or red-shifted asymmetries in emission, with these deviations from a standard line profile becoming the most pronounced at the GRB/SN Site. Such asymmetries could be indicative of outflows and inflows of ionized gas with distributed opaque clouds (e.g. Kewley et al.\\ 2001), and could suggest that the host galaxy has undergone some prior merger or interaction that is still impacting the dynamics of the host regions examined here. Similar explanations have been proposed for other LGRB hosts with disturbed morphologies (e.g. Wainwright et al.\\ 2007, Starling et al.\\ 2011). However, it is important to note that such asymmetries can also be attributed to multiple star-forming regions or structure within the nebula (e.g. Wiersema et al.\\ 2007). A proper examination of the nature of these asymmetries, and their implications for interaction activity or star-forming region components, will require higher-resolution spectroscopy and comparisons with deeper multi-band host images. In Figure 3 we plot the existing sample of LGRB host galaxies on a luminosity-metallicity (L-Z) diagram, comparing them to contours from the L-Z relation for star-forming SDSS galaxies from Tremonti et al.\\ (2004). We also plot samples of Ic-BL host galaxies, from both Modjaz et al.\\ (2008, 2011) and Sanders et al.\\ (2012), to examine whether there is any clear environmental distinction between Ic-BLs without LGRBs and those accompanied by LGRBs. Finally, we include data on the relativistic Ic-BL SN 2009bb from Levesque et al.\\ (2010d). Data for previous LGRB hosts comes from Levesque et al.\\ (2010a,b), Levesque et al.\\ (2011), and references therein. From archival SDSS global photometry of the host galaxy (DR8; $g = 21.16 \\pm 0.09$ and $r = 20.60 \\pm 0.09$), and the corrections of Blanton \\& Roweis (2007), we determine $M_B = -19.4 \\pm 0.2$ for the GRB 120422A/SN 2012bz host. From the comparison in Figure 3, there appears to be no clear distinction in metallicity among LGRBs based on an event's classification as a subluminous burst - both subluminous LGRBs and cosmological LGRBs have an average log(O/H) + 12 = 8.2 $\\pm$ 0.1. As a whole, this comparison shows that any differences within the LGRB sample at $z < 1$ cannot be discerned based on metallicities. However, this does not rule out the possibility that other burst properties (blastwave velocity, X-ray fluence, etc.) may reveal a fundamental difference in the internal properties driving the production of LGRBs in these different sub-classes. It should also be noted that these $z < 1$ GRB host studies consider only O abundances, while many single-star progenitor models depend on line-driven winds and are therefore strongly dependent on heavier element abundances which may be enhanced in GRB hosts (see Chen et al.\\ 2007). Conversely, many binary progenitor models for LGRBs do not have a strong metallicity dependence (e.g. Fryer \\& Heger 2005, Podsiadlowski et al.\\ 2010), suggesting that a distinction between different LGRB progenitor channels may not be discernible based purely on metallicity. The comparison between LGRB hosts and Ic-BL hosts is more complex. The Ic-BL sample from Modjaz et al.\\ (2008) has an average metallicity of log(O/H) + 12 = 8.6 $\\pm$ 0.1, while the Sanders et al.\\ (2012) sample has a lower average metallicity of log(O/H) + 12 = 8.2 $\\pm$ 0.1 (the cause of the disagreement between these two samples is not yet understood). The LGRB hosts and Sanders et al.\\ (2012) Ic-BL hosts shown here also fall below the general L-Z relation for star-forming galaxies form SDSS. The physical explanation driving this offset is unclear. Mannucci et al.\\ (2011) suggest that this may be attributable to a fundamental relation between metallicity, SFR, and stellar mass in star-forming galaxies, arguing that LGRBs simply occur preferentially in environments with higher SFRs. However, Kocevski \\& West (2011) find that this relation is not sufficient to explain such an offset. From our examination of the GRB 120422A/SN 2012bz host environment, we find that this galaxy is a fairly typical LGRB host, with a low metallicity measured at both the Nucleus and the GRB/SN Site. Including this newest host within the larger sample of LGRB host galaxies, we find that there is no difference in metallicity between the subluminous and cosmological LGRB host samples. In addition, the distance of GRB 120422A/SN 2012bz from the bright star-forming nucleus of its host ($\\sim$8 kpc) marks this LGRB and host environment as unique; combined with asymmetries in the emission features of the host and the lack of a clear spiral arm component in the host region, this could be indicative of prior merger or interaction activity in the host. Future work on both this and other nearby spatially-resolved LGRB hosts will allow us to further probe the nature of these galaxies. Studies of additional properties such as ionization parameter, stellar population age, and star formation history, as well as dynamical studies that can explore potential merger activity, will all be valuable in characterizing the key environmental parameters that drive progenitor formation and energetic properties for LGRBs. \\\\ EML is supported by NASA through Einstein Postdoctoral Fellowship grant number PF0-110075 awarded by the Chandra X-ray Center, which is operated by the Smithsonian Astrophysical Observatory for NASA under contract NAS8-03060. The Berger GRB group at Harvard is supported by the National Science Foundation under Grant AST-1107973. Partial support was also provided by a Swift AO7 grant number 7100117. The paper includes data gathered with the 6.5 meter Magellan Telescopes located at Las Campanas Observatory, Chile. We thank the support staff at Las Campanas for their hospitality and assistance. This paper utilized data from the Gamma-Ray Burst Coordinates Network (GCN) circulars and SDSS Data Release 8. Funding for SDSS-III has been provided by the Alfred P. Sloan Foundation, the Participating Institutions, the National Science Foundation, and the U.S. Department of Energy Office of Science. The SDSS-III web site is http://www.sdss3.org/. This work was made possible in part by collaborations and discussions at the Aspen Center for Physics, supported by NSF grant 1066293." }, "1207/1207.0615_arXiv.txt": { "abstract": "{The length of the asteroseismic timeseries obtained from the \\textit{Kepler} satellite analysed here span 19 months. \\textit{Kepler} provides the longest continuous timeseries currently available, which calls for a study of the influence of the increased timespan on the accuracy and precision of the obtained results.} {We aim to investigate how the increased timespan influences the detectability of the oscillation modes, and the absolute values and uncertainties of the global oscillation parameters, i.e., frequency of maximum oscillation power, $\\nu_{\\rm max}$, and large frequency separation between modes of the same degree and consecutive orders, $\\meandnu$.} {We use published methods to derive $\\nu_{\\rm max}$ and $\\meandnu$ for timeseries ranging from 50 to 600 days and compare these results as a function of method, timespan and $\\meandnu$.} {We find that in general a minimum of the order of 400 day long timeseries are necessary to obtain reliable results for the global oscillation parameters in more than 95\\% of the stars, but this does depend on $\\meandnu$. In a statistical sense the quoted uncertainties seem to provide a reasonable indication of the precision of the obtained results in short (50-day) runs, they do however seem to be overestimated for results of longer runs. Furthermore, the different definitions of the global parameters used in the different methods have non-negligible effects on the obtained values. Additionally, we show that there is a correlation between $\\nu_{\\rm max}$ and the flux variance.} {We conclude that longer timeseries improve the likelihood to detect oscillations with automated codes (from $\\sim$60\\% in 50 day runs to $>$ 95\\% in 400 day runs with a slight method dependence) and the precision of the obtained global oscillation parameters. The trends suggest that the improvement will continue for even longer timeseries than the 600 days considered here, with a reduction in the median absolute deviation of more than a factor of 10 for an increase in timespan from 50 to 2000 days (the currently foreseen length of the mission). This work shows that global parameters determined with high precision - thus from long datasets - using different definitions can be used to identify the evolutionary state of the stars.} ", "introduction": "Many breakthrough results for red-giant (G-K) stars have been presented using data obtained by the CoRoT \\citep{baglin2006} and NASA \\textit{Kepler} \\citep{borucki2010} missions. These results include statistical ensemble studies of global oscillation parameters, i.e., frequency of maximum oscillation power, $\\nu_{\\rm max}$, mean frequency separation between modes of the same degree and consecutive orders, $\\meandnu$, small frequency separations between modes of different degree, $\\ell$, amplitudes and visibilities of the oscillations, and tests of scaling relations \\citep[e.g.,][]{deridder2009,hekker2009,bedding2010,huber2010,hekker2011pub,huber2011,mosser2012}. Additionally, it has been possible to determine stellar parameters such as masses and radii \\citep{kallinger2010corot,kallinger2010}. In addition to these results, asteroseismic investigations into the granulation \\citep{mathur2011}, red giants in clusters \\citep{basu2011,hekker2011clus,stello2011mem,stello2011ampl} and red giants in eclipsing binaries \\citep{hekker2010ecb} have been performed, as well as detailed investigations into the internal structure of single stars \\citep[e.g.,][]{dimauro2011,jiang2011,baudin2012}. The \\textit{Kepler} results referred to are based on timeseries with a near regular cadence of either 29.4 min or 58.85 s and a timespan ranging from $\\sim$30 days up to more than 1.5 yr. These are the first datasets from space-based telescopes with such long timespan and high fill ($\\gtrsim$90\\%) and frequency resolution ($\\approx$ 0.019~$\\mu$Hz). Underpinning much of this work is the ability to determine global oscillation parameters and the uncertainties in these values. It is reasonable to ask if there are now enough data available and whether there are any gains to be obtained from observing individual stars for longer periods. In this paper we address the precision and reliability of the determination of some of the global seismic parameters. There are other areas where there is a clear need for data of longer duration because the features detected in the power spectra are narrow and hence barely resolved even by the current datasets. In particular, we highlight the detection of g-p mixed modes \\citep{beck2011}. The observed mean period spacings appear to have different values for stars that burn only H (in a shell) and those that also burn He in the core \\citep{bedding2011,mosser2011mm}, hence the period spacing can be used to distinguish between different evolutionary states in which red giants are observed using the characteristics of their frequency spectra. Another method to distinguish between different evolutionary phases is based on the difference in frequency dependence of radial modes \\citep{kallinger2012}. Furthermore, recently, the timeseries obtained with \\textit{Kepler} have become long enough to study rotational splitting of the oscillation modes, which led to the detection of differential rotation in red giants \\citep{beck2012}. In this work, we use the 19 months of data available from Q0 to Q7 to investigate how the increased timespan influences the detectability of the oscillation modes, and the absolute values and uncertainties of the global oscillation parameters, $\\nu_{\\rm max}$ and $\\meandnu$. These are important in several ways. Knowing the dependence of the precision on data duration is a guide for observing strategies, and for the determination of those secondary parameters that are derived from the primary global oscillation parameters, such as stellar mass and radius. Furthermore, it is crucial to be able to estimate the proportion of false negatives and false positives for population studies. Also, for detailed modelling of individual oscillation frequencies $\\nu_{\\rm max}$ turned out to be of great diagnostic potential \\citep{gruberbauer2012}. We will include in our considerations the impact of other relevant parameters such as the observed height-to-background ratio of the oscillation excess. This work is a follow-up of \\citet[][hereafter paper I]{hekker2011comp} on the red giants and \\citet{verner2011} on solar-type stars, in which results obtained with different methods have been compared and validated. Paper I described the comparison of global oscillation parameters extracted from about four month of \\textit{Kepler} data using different methods. From this comparison, it was concluded that 1) the results from the different methods agree for most stars within a few percent; 2) at least five methods (out of the seven tested) obtained results for 92\\% of stars for $\\nu_{\\rm max}$ within the range of 50~$\\mu$Hz to 170~$\\mu$Hz, and this percentage decreased to 69\\% when all stars with $\\nu_{\\rm max}$ covering the complete frequency range, i.e., 0 -- 283.4\\,$\\mu$Hz (the Nyquist frequency) were included; 3) the scatter due to realization noise, originating from the stochastic nature of the oscillations, is non-negligible and can be at least as important as the internal uncertainty of the results due to the method used, but this depends on the frequency of maximum oscillation power, $\\nu_{\\rm max}$, and on the methods. In case a model is used to describe the variation of $\\Delta\\nu$ with frequency the results are less sensitive to realization noise than others; 4) the influence of the obtained value of $\\meandnu$ is less dependent on the frequency range over which it is computed than is the case for solar-type stars. A theoretical follow-up study to explain the latter has been performed by \\citet{hekker2011dnu}. ", "conclusions": "In this work we investigated the impact of the length of the timeseries on the precision and accuracy of the determined global oscillation parameters $\\nu_{\\rm max}$ and $\\meandnu$ of red giants. We used \\textit{Kepler} light curves spanning about 600 days and divided them in short runs of 50, 100, 200 and 400 days. All these runs have been analysed using automated methods. The oscillation detection rate has been compared with predictions and the resulting values for the global oscillation parameters have been compared as a function of method, run length, $\\meandnu$ of the oscillations. From this study we find that: \\begin{itemize} \\item For 95\\% of the stars consistent global oscillation parameters are obtained from 600 day timeseries with different methods. For the remaining 5\\%, there were good reasons for the lack of consistency. \\item Using the observational methods we find more than 95\\% (of the consist results of 600 day data) or more reliable detections of oscillations in timeseries of 400 days or longer. \\item Current predictions of the detectability of oscillations are based on the amplitudes and predict that in the majority of the cases the likelihood to detect oscillations are above 90\\% for both the long and short runs. However, most observational algorithms use the regularity in the power spectrum to detect the oscillations and the regularity has reduced sensitivity for shorter runs. \\item The precision of the determined global oscillation parameters increases with increasing timeseriess and the trends suggest that this continues for even longer timeseries than investigated here. From the extrapolation of fits to the median absolute deviations a reduction of more than a factor of 10 for an increase in timespan from 50 to 2000 days (the currently foreseen length of the mission) is foreseen. Thus, there are real advantages to be gained from working with even long timeseries than considered here. We note that the universal pattern is already effective for short datasets. \\item The distributions of the offsets - difference between results of short runs with respect to the result obtained with the same method on the 600-day long timeseries - divided by the quoted uncertainties show that the quoted uncertainties have a tendency to be overestimated, which is in general more severe for longer datasets. However, this does depend on the method. % \\item We find that 50 day timeseries are not long enough to be certain to pick up more than 90\\% of the oscillations with the currently employed methods. \\item When comparing different methods it is clear that the differences due to different definitions are non-negligible. This difference is a function of the evolutionary state of the stars and this could be used to determine the evolutionary state. \\item The different strengths, definitions and sensitivity to realization noise of the different methods indicate that the simultaneous use of more methods is likely to be profitable. \\end{itemize} Additionally, we propose and justify a new method to estimate the frequency of maximum oscillation power from variance in the timeseries. We show that the dependence of the flux variance on $\\nu_{\\rm max}$ is also a function of evolutionary phase. The effectiveness of this method does not depend on the data duration nor on the location of the peak of the spectrum -- always assuming that the necessary data detrending is not attenuating the oscillations signal. We recommend that this method be used in conjunction with the methods described here as an additional independent constraint to detect the oscillations." }, "1207/1207.7345_arXiv.txt": { "abstract": "Solar flare emissions in the chromosphere often appear as elongated ribbons on both sides of the magnetic polarity inversion line (PIL), which has been regarded as evidence of a typical configuration of magnetic reconnection. However, flares having a circular ribbon have rarely been reported, although it is expected in the fan--spine magnetic topology involving reconnection at a three-dimensional (3D) coronal null point. We present five circular ribbon flares with associated surges, using high-resolution and high-cadence \\ha\\ blue wing observations obtained from the recently digitized films of Big Bear Solar Observatory. In all the events, a central parasitic magnetic field is encompassed by the opposite polarity, forming a circular PIL traced by filament material. Consequently, a flare kernel at the center is surrounded by a circular flare ribbon. The four homologous jet-related flares on 1991 March 17 and 18 are of particular interest, as (1) the circular ribbons brighten sequentially, with co-spatial surges, rather than simultaneously, (2) the central flare kernels show an intriguing ``round-trip'' motion and become elongated, and (3) remote brightenings occur at a region with the same magnetic polarity as the central parasitic field and are co-temporal with a separate phase of flare emissions. In another flare on 1991 February 25, the circular flare emission and surge activity occur successively, and the event could be associated with magnetic flux cancellation across the circular PIL. We discuss the implications of these observations combining circular flare ribbons, homologous jets, and remote brightenings for understanding the dynamics of 3D magnetic restructuring. ", "introduction": "\\label{introduction} The elongated structures of solar flare emissions, as the name ``flare ribbons'' implies, have long been observed at \\ha\\ and UV wavelengths \\citep[e.g.,][]{zirin88}. A pair of flare ribbons residing in opposite magnetic polarities run parallel to the magnetic polarity inversion line (PIL) lying between them, and the ribbons separate from each other in the direction perpendicular to the PIL. Such a configuration and its kinematics has been regarded as evidence of the classical two-dimensional-like (2D) magnetic reconnection model called the CSHKP model \\citep{carmichael64,sturrock66,hirayama74,kopp76}, in which two ribbons form as the energy release from a series of coronal X-points along an arcade of loops produces bright flare emissions at the loop footpoints, and the ribbon separation is resulted from the successive reconnections of higher coronal arcades. Nevertheless, actual flares take place in a more complicated three-dimensional (3D) structure, hence the standard model, despite of its general applicability \\citep{hudson11}, may not explain some features of the flare ribbon evolution, such as the frequently observed expansion of ribbons along the PIL before the perpendicular separation \\citep[e.g.,][]{fletcher04}. In relation to magnetic topology, chromospheric flare ribbons are generally situated at locations intersected by separatrices dividing domains of distinct connectivity \\citep[e.g.,][]{mandrini91}, or quasi-separatrix layers (QSLs) possessing strong connectivity gradients \\citep[e.g.,][]{demoulin97}. This is because that intense current sheets are preferentially built up at separatrices/QSLs \\citep[e.g.,][]{lau90,aulanier05}, along which reconnection-accelerated particles can precipitate into the lower atmosphere. It is known that separatrices can be typically generated by magnetic null points, which are common structures in the corona due to the mixed-polarity nature of the photospheric fields \\citep{schrijver02}. Magnetic field lines associated with a 3D coronal null point usually display a fan-spine configuration \\citep{lau90,torok09}, where the dome-shaped fan portrays the closed separatrix surface and the inner and outer spine field lines in different connectivity domains pass through the null point. The footpoint of the inner spine has a magnetic polarity opposite to those of the fan, which forms a circular PIL. Magnetic reconnection can be induced in such single null points, as the fan/spines deviate from the null when subject to shearing or rotational perturbations \\citep[e.g.,][]{pontin07a,pontin07b}. It is then expected that as a result of null-point reconnection, flare emissions at the footpoints of the fan field lines would constitute a closed circular flare ribbon, and that the spine-related flare footpoint would be a compact source. Surprisingly, there seems to be very few events reported in the literature that assume the shape of circular ribbons in \\ha\\ \\citep{sundara94} or UV \\citep{ugarte07,su09} wavelengths. In retrospect, the most comprehensive study of circular ribbon flares was carried out by \\citet[][hereafter M09]{masson09} and \\citet{reid12} for a confined C8.6 flare on 2002 November 16 using 1600~\\AA\\ UV continuum images from \\trace. Several observational features are prominent in this event. First, three ribbons were present, including two elongated ribbons inside and outside the circular ribbon, respectively. Second, the circular ribbon brightened sequentially in the counterclockwise direction. Third, the appearance of the outside, remote ribbon had a \\sm30~s delay relative to the main ribbons. A potential field extrapolation revealed that these ribbons indeed seem to map the photospheric intersections of the fan and spine field lines, which stem from a coronal null point (see Figure~\\ref{f1}). With the aid of 3D MHD numerical simulations, the authors further suggested that the fan and spine separatrices are embedded in larger QSLs. The extended shape of QSLs surrounding the singular spine field lines is consistent with the observed elongated spine ribbons. Importantly, field lines can undergo slipping and slip-running reconnection within the QSLs \\citep{aulanier06}. Their results that field lines closer to the null would reconnect first and that the slipping motion is toward the null could then account for the counterclockwise propagation of the circular ribbon emission. The implied sequential occurrence of slipping/slip-running reconnection within the fan and null-point reconnection involving the outer spine may explain the delayed brightening of the remote ribbon corresponding to a second phase of flare emission. By flipping the outer spine field lines in \\citetalias{masson09} to open outward, one could arrive at an axisymmetrical null-point and fan-spine topology (see Figure~\\ref{f1}), which \\citet{pariat09} used to model solar polar jets. Jets in the solar atmosphere can appear in a variety of forms that may be interrelated \\citep[e.g.,][]{chae99}, such as the cool ejections (surges) in \\ha\\ \\citep[e.g.,][]{schmieder84} and the hot ejections (jets) in EUV/UV \\citep[e.g.,][]{alex99} and soft X-rays (SXRs) \\citep[e.g.,][]{shibata92}. In fact, the above configuration can be simply developed by flux emergence into a unipolar region such as the polar coronal holes, a scenario typically assumed for jets \\citep[e.g.,][]{shibata92}. In the model of \\citet{pariat09}, the imposed twisting motion within the fan circle at the photospheric boundary builds up magnetic stress, until an ideal instability sets in to cause interchange reconnection at the null, driving massive, high-speed jets. The numerical investigation was then extended in \\citet[][hereafter P10]{pariat10} by tilting the outer spine to break the axisymmetry and applying a constant stress to both closed- and open-connectivity domains (see Figure~\\ref{f1}). Under these prescribed perturbations, it was then shown that 3D null-point topology can naturally produce successive, homologous jets, which are frequently observed \\citep[e.g.,][]{chifor08b}. The axes of the modeled homologous jets rotate in the same direction that is opposite to the boundary driving direction. The simulations also demonstrated that the twist originally stored in the closed domain (e.g., emerging fluxes) can be transferred to the reconnected, open field lines, which explains the observed unwinding motion and helical structure in jets \\citep[e.g.,][]{canfield96,liu11,liu+wei11}. Moreover, the authors pointed out that even though there should form a circular ribbon related to the fan and another ribbon corresponding to the inner spine, the properties of flare ribbons on the surface may be hard to predict, due to the dynamic nature of the 3D reconnection at the null point. To our knowledge, jet-associated circular ribbon flares have not been reported before. In this paper, we present five flares including four jet-related events, the ribbons of which all have a circular shape. It is interesting that three homologous flares are also accompanied by a remote ribbon. Our events thus combine the key observational features on which the models of \\citetalias{masson09} and \\citetalias{pariat10} are based, incorporating circular flare ribbons, homologous jets, and remote brightenings. The uncovering of these observations benefits from studying a sample portion of the historical films of the Big Bear Solar Observatory (BBSO), which was established and previously directed by Professor Hal Zirin. The plan of this paper is as follows: in Section~\\ref{data}, we describe the data sets and reduction procedure. In Section~\\ref{results}, we present the main results of data analysis and discuss their implications. Major findings are summarized in Section~\\ref{summary}. ", "conclusions": "\\label{summary} In this paper, we have taken advantage of the high spatio-temporal resolution BBSO film observations that have recently been digitized to study five major flares characterized with circular chromospheric ribbons and surges. In particular, four flares originated from the same active region exhibit homologous properties and three of them are accompanied by remote brightenings. Circular ribbon flares were rarely reported before but are an important observational signature predicted as a result of the 3D null-point reconnection in the fan-spine magnetic topology. The high-quality \\ha\\ blue wing images allow us to study the event evolution in great detail in relation to the different types of 3D reconnection as proposed by \\citetalias{masson09} and \\citetalias{pariat10}. It is remarkable that different from the confined C-class event in \\citetalias{masson09}, our ejective M- and X-class flares are clearly associated with surges; moreover, the model of \\citetalias{pariat10} produces homologous jets, but does not accommodate remote brightenings as observed in our events (see Table~\\ref{table1}). Therefore, the present observations point to a comprehensive picture combining the related topological magnetic field structure of \\citetalias{masson09} and \\citetalias{pariat10} (Figure~\\ref{f1}). Our results can be summarized as follows. \\begin{enumerate} \\item The main attention is focused on the four homologous flares and surges on 1991 March 17--18 (events 1--4). By tracking the ribbon motion, we derive the magnetic flux change rates in the opposite polarities $\\dot\\phi_{+-}$, which are temporally correlated with the non-thermal flare emissions but are unbalanced. A caveat is that our method is only applicable for events with a 2D-like configuration. In fact, the existence of a 3D structure comprising the coronal null point and the related fan-spine magnetic structure are inferred from the unambiguous circular shape of the outer ribbon and the potential field extrapolation \\citepalias{masson09}, as well as the drifting surge \\citepalias{pariat10}. \\item In the event 1, the well-observed sequential brightening of the fan ribbon k2 along the circle and the elongated shape of the inner spine ribbon k1 strongly indicate the presence of extended QSLs embedding the separatrices and the slipping/slip-running reconnection \\citepalias{masson09}. It is evidently observed that k1 undergoes a ``round-trip'' motion, presumably toward and then away from the coronal null point. As surges are produced through the null-point reconnection involving the outer spine \\citepalias{pariat10}, a consistent observation could be that the surge begins the abrupt ejection after k1 reaches the turning point. We suggest that the intriguing motion of k1 may be due to the succession of slipping reconnection, first toward the null and then after the null reconnection, away from the null \\citepalias{masson09}. \\item Remote brightenings are observed in the events 2--4, and are co-temporal with a second phase of flare non-thermal emissions about \\sm1~minute after the main energy release. The delay might be related to the time of field slipping toward the null \\citep{reid12}. The fact that the surge ejection is after the maximum of the remote brightening could indicate a quasi-instantaneous change of magnetic topology. These observations are the first signature that may indicate the opening and closing of the outer spine in the corona during a flare. \\item The successive occurrence of slipping/slip-run reconnection and null-point reconnection proposed by \\citetalias{masson09} and \\citet{reid12} is also manifested as two well separated phases of circular ribbon flare and ejective surge as observed in the event 5. The flaring source region experiences flux cancellation across the circular PIL, which may also contribute to the triggering of the null-point reconnection \\citep{fletcher01}. \\end{enumerate} It is important to understand how the 3D null-point system may be perturbed to be subjective to the subsequent reconnection by analyzing the magnetic field evolution \\citepalias{masson09,pariat10}. This aspect requiring long-term magnetic field observation is out of the scope of this study. While the paucity of circular ribbon flares was suggested to be due to the possible uncommon presence of a single null point or the pronounced asymmetry of the fan \\citepalias{masson09}, we emphasize that our survey in progress using the rich film data base has revealed a few more circular ribbon flares with the associated surges and/or remote brightenings. These events plus the continued advanced observations from the \\Sdo\\ have promise to shed further insights on the dynamics of magnetic reconnection in the null-point topology." }, "1207/1207.3346_arXiv.txt": { "abstract": "We have assembled a large-area spectroscopic survey of giant stars in the Sagittarius (Sgr) dwarf galaxy core. Using medium resolution ($R \\sim 15,000$), multifiber spectroscopy we have measured velocities of these stars, which extend up to 12 degrees from the galaxy's center (3.7 core radii or 0.4 times the King limiting radius). From these high quality spectra we identify 1310 Sgr members out of 2296 stars surveyed distributed across 24 different fields across the Sgr core. Additional slit spectra were obtained of stars bridging from the Sgr core to its trailing tail. Our systematic, large area sample shows no evidence for significant rotation, a result at odds with the $\\sim 20$ km s$^{-1}$ rotation required as an explanation for the bifurcation seen in the Sgr tidal stream; the observed small ($\\le 4$ km s$^{-1}$) velocity trend along primarily the major axis is consistent with models of the projected motion of an extended body on the sky with no need for intrinsic rotation. The Sgr core is found to have a flat velocity dispersion (except for a kinematically colder center point) across its surveyed extent and into its tidal tails, a property that matches the velocity dispersion profiles measured for other Milky Way dwarf spheroidal (dSph) galaxies. We comment on the possible significance of this observed kinematical similarity for the dynamical state of the other classical Milky Way dSphs in light of the fact that Sgr is clearly a strongly tidally disrupted system. ", "introduction": "The Sagittarius (Sgr) dwarf galaxy is a satellite of the Milky Way (MW) with an apparent dwarf spheroidal (dSph) morphology that is, however, distinct in at least two ways. First, the relative proximity of Sgr to the Sun ($d \\approx 24$ kpc) makes its brightest stars quite accessible to study with even relatively modest aperture telescopes. Thus, Sgr is open to much more straightforward and extensive examination than its much more distant, counterpart MW satellites making Sgr, in principle, a highly approachable, local laboratory for dSph galaxies. Second, Sgr represents the most certain and vivid example of a satellite enduring tidal disruption by the MW, with the highest known surface brightness tidal tails among MW dSph satellites. Indeed, not only does no other MW dSph have anywhere near as obvious tidal debris, whether other dSphs are showing any evidence of visible mass loss is still debated for most MW satellites. For this reason, Sgr is often relegated to a ``special case'' category among MW satellites and investigators of the properties of ``classical dSphs'' often avoid Sgr because of an uncertainty and wariness of how it fits into the larger dSph context. Obviously prudence dictates care in such matters, but the potential advantages of a much more accessible test laboratory for dSph dynamics makes it worth understanding whether the second of the above two Sgr properties (tidal disruption) truly nullifies the usefulness of the first (proximity) in terms of adopting Sgr as a valid prototype for other dSph satellites. An early analysis of the internal dynamics of Sgr was provided by Ibata et al. (1997). This impressively assembled early sample of radial velocities numbered some 410 stars using the CTIO/ARGUS and AAT/AUTOFIB spectrographs, and it remains the only fully published wide-spread detailed analysis of the dynamical properties of the Sgr galaxy main body. While the area covered provides some strength to look at the larger Sgr main body, the limitation of most fields being along the main axis limits further analysis, as does the poor velocity precision of the AAT data ($\\epsilon_{RV} \\sim 12$ km s$^{-1}$). However, many analyses of the dynamics and $M/L$ of this system specifically make use of the single dispersion ($\\sim11$ km s$^{-1}$) derived for Ibata et al.'s database of 114 stars located at the nucleated (perhaps unrepresentative) center. Recently, a first analysis of a new AAT data set was presented by, Pe\\~narrubia et al. (2011), based on 1805 Sgr members by RV in 6 fields along the major axis and one side of the minor axis of Sgr reaching to $\\pm 4^{\\circ}$ along the major axis. This data analysis and modeling focuses primarily on whether rotation is found in Sgr. An improved kinematical analysis of the Sgr inner core and attempting to isolate it chemically and kinematically from M54 was conducted by Bellazzini et al. (2008). They analyzed 1152 VLT/FLAMES and Keck/DEIMOS spectra using RVs and CaT metallicities determinations to isolate the Sgr core population from that of M54. While this provides an excellent analysis of the central core of Sgr ($\\sigma_{v} = 9.6 \\pm 0.4$ km s$^{-1}$), the limited area, all within 9' of core, provides just one data point compared to the larger main body of Sgr. Of course, understanding the Sgr system --- among the nearest known examples of a minor merger --- is sufficient cause for systematic exploration of its core, whether or not it can be considered a suitable representative of the dSph morphological class in general. Some recent evidence makes a compelling case that the primary difference between Sgr and other MW satellites may be principally in its life phase within what may be a universal evolutionary course for dwarf galaxies. In general, the star formation histories of dwarf galaxies in the local universe are similar for most of cosmic time, regardless of their present morphological type, with the latter mostly a function of variations in the influence of external mechanisms (like ram pressure stripping or tidal effects) within the past Gyr or so and that give rise to a strong morphology-density relationship (Skillman et al. 2003; Grebel 2004; Weisz et al. 2011). According to models by Mayer et al.\\ (2001), tidal stirring via close passages around large host galaxies can transform star forming, gas-rich, dwarf irregular (dIrr) systems with small rotating disks into classical, moribund, gas-depleted, pressure-supported dSphs. The mechanism driving such a metamorphosis is the formation and decline of a central bar, incited by the dynamical interaction of the satellite with the gravitational potential of the host, particularly at pericenter passages. Certainly it is clear that at least some of the present ``classical'' MW dSphs (e.g., Carina, Leo I), though now devoid of gas, were forming stars in the not too distant past (the past few Gyrs; Smecker-Hane et al.\\ 2009; Gallart et al.\\ 2007) --- just like Sgr \\citep{siegel07}. Interestingly, Carina and Leo I may be the two classical MW dSphs apart from Sgr with the strongest evidence for tidal disruption \\citep*[Carina:][ Leo I: Sohn et al.\\ 2007; Mateo et al.\\ 2008]{IH95,kuhn96, majew00, majew04, munoz08}. Ironically, despite Sgr's proximity, less is known about the distribution of properties of stars across the Sgr core than for other satellites of the MW, for a variety of potential reasons --- e.g., because of (1) the tendency of observers to avoid it on grounds that it is perceived to be unusual, as mentioned above, (2) the fact that Sgr has only been relatively recently discovered (Ibata et al.\\ 1994), as well as, probably most importantly, (3) its intrinsic size at its distance from the Sun means that Sgr spans a large angle on the sky. Even with large field of view, multifiber spectroscopy it is simply difficult to cover the enormous size of the Sgr core (with a semi-major axis length of almost 30 degrees). Nevertheless, some basic knowledge about the Sgr core is, of course, in place. Ibata et al.\\ (1994, 1995) made the first maps showing the extended and elongated nature of the Sgr core and determined that it was comparable in size and luminosity to the largest of the MW dSph satellites, the Fornax system. These authors also showed that, like Fornax, the Sgr system has its own retinue of globular clusters, including M54, a large globular cluster that is either the very heart, or {\\it at} the very heart, of the dSph (see discussion by Bellazzini et al.\\ 2008; Monaco et al. 2005; but also cf. Siegel et al. 2011). Other globular clusters in the Galactic halo farther from the Sgr core may also be debris from the original satellite (see review by Law \\& Majewski 2010b). Majewski et al.\\ (2003; hereafter PAPER I) assembled the first all-sky view of the Sgr dwarf galaxy core along with its tidal tails by identifying constituent M giant stars within the final release Two Micron All-Sky Survey (2MASS; Skrutskie et al.\\ 2006) point source catalog. M giants, which are copious in the relatively metal rich ([Fe/H] as high as $-$0.4; Layden \\& Sarajedini 2000) Sgr populations, are separable from M dwarfs with $J$,$H$,$K$ photometry (Bessell \\& Brett 1988), and in any case the former well outnumber the latter to 2MASS magnitude limits (e.g., $K_s \\sim 13-14$). The 2MASS M giant sample, which is more robust to the differential reddening effects that have made optical studies of Sgr core difficult, yielded the first ``clean'' view of the entire main Sgr body. From the distribution of the large population of M giants, the morphology of the Sgr core could be more accurately fit, and Majewski et al.\\ showed that it could be described by a King profile but with, at the largest radii, a break in the King profile due to the presence of its tidal tails. Apart from the fact that Sgr shows a small density cusp at its very center, which may or may not (see Monaco et al. 2005) be due to the density enhancement provided by the superposition of M54, the overall King+power law break profile of Sgr looks identical to those previously observed in other ``classical'' Galactic dSphs (see Fig.\\ 13 below); however, the source of these profile breaks in other dSph systems have been debated as due either to prodigious tidal mass loss or as bound populations expected for the high inferred $M/L$ in these systems (e.g., Munoz et al.\\ 2006, 2008; Sohn et al. 2007; Koch et al. 2007; McConnachie, Pe\\~narrubia, \\& Navarro 2007). It is commonly suggested (e.g., Kroupa 1997; Kleyna et al.\\ 1999) that discrimination of these two models should be possible from the radial velocity (RV) dispersion trend with radius. Sgr represents the closest, most accessible dSph in the class of those exhibiting King+break profiles, but it is also the one example where the break is from {\\it certain} tidal disruption. Thus it provides a Rosetta Stone by which the dynamical markers of this process may be empirically established. In the present paper we undertake the first large, wide-field {\\it and} high resolution ($R\\sim15,000$) spectroscopic survey to sample across the larger area encompassing the core body of the Sgr system. With this survey, we explore the velocity distribution of 2MASS-selected K and M giant stars across the face of this core, from its very center out to the trailing tidal tail and along the major, minor and other symmetry axes of the system. Thus, with this survey we provide kinematical information for Sgr that is comparable in extent to datasets existing for other Milky Way dSphs. This enables us to put Sgr in a dynamical context with these other systems, and provide a baseline for comparison to a dSph known to be undergoing severe tidal disruption. Our study focuses on two primary observables to provide this context: (1) Within the distribution of radial velocities (\\S3) we look for signatures of rotation and show that Sgr, like other dSphs, has no significant rotation signature (\\S4). (2) The velocity dispersion profile of Sgr is shown to be more or less flat, which is the same trend found in all other dSphs (\\S5). Thus, in these two properties, Sgr is found not to be unusual compared to other dSphs, despite its current state of enduring extreme mass loss (see, e.g., PAPER I; Chou et al. 2007). We comment on the potential significance of this apparent ``normality'' of Sgr as a dSph system for the potential state of other dSph systems in \\S6.2. ", "conclusions": "\\subsection{Summary of Sgr Core Properties Found Here} We have conducted the first large scale systematic survey of the Sgr dSph main body with broad azimuthal coverage. Using over 1200 member stars, we are able to investigate evidence for effects that would be seen over large areas, such as rotation and global velocity dispersion trends. A complementary study of the metallicity distribution across Sgr using these same spectra will be included in a future contribution. The primary results from this paper may be summarized as follows: 1. The Sgr central velocity dispersion is found to be $9.9 \\pm 0.7$ km s$^{-1}$, in agreement with the dispersion ($9.6 \\pm 0.4$ km s$^{-1}$) found by Bellazzini et al. (2008) after removing stars associated with M54 (e.g., Figure~\\ref{fig:vGSR_R_Re}, \\S 4). This central velocity dispersion is not atypical among the ``classical'' dSph galaxies (see Table~\\ref{Tab:dwarfs}). 2. Within the distribution of radial velocities, we find no signatures of rotation along the minor axis of Sgr. With a detailed investigation of trends along the major axis, we find evidence for at most a small trend ($\\le 4$ km s$^{-1}$ deg$^{-1}$), though even this small effect can be explained without needing intrinsic rotation, as shown by various adopted models of an extended Sgr system moving along its orbit. Thus, we confirm that Sgr is like other dSphs in having no significant rotation ($v_{rot}/\\sigma(v) \\gtrsim 1 $) signature along any axis (e.g., Figures~\\ref{fig:angle}, \\S 4). \\subsection{Consistency with the Tidal Stirring Model} In a companion study by {\\L}okas et al. (2010a), $N$-body simulations were generated to explain both the elongated shape and velocity characteristics (as found in the present study) of the Sgr core. Of critical significance to the model described there was that the present study found no significant rotation curve across the face of the dSph, which supports the notion that Sgr is bar-like with the major axis almost perpendicular to the line of sight, not a disk seen almost edge-on, which would exhibit a significant rotational signal. The bulk of the small observed rotation is simply due to projection effects, and the insignificant intrinsic rotation of the system points toward a prolate shape for the galaxy, a configuration supported by radially-orbiting stars. A disk-like configuration was recently employed by Pe\\~narrubia et al. (2010) as an intriguing way to produce a bifurcated leading arm (as a means to explain the two Sgr streams seen in the Sloan Digital Sky Survey data --- Belokurov et al. 2006; Yanny et al. 2009), but the 20 km s$^{-1}$ level of rotation expected to be seen today in such a model is clearly precluded by the present observations, a point also made by Pe\\~narrubia et al. (2011) from observations they made to test their model (over a smaller area than covered here). On the other hand, as shown in {\\L}okas et al. (2010a), the data presented here are consistent with the model of tidal stirring (Mayer et al. 2001; Klimentowski et al. 2007, 2009; Kazantzidis et al. 2011) proposed there to account for the shape and dynamics of the Sgr core. According to the proposed scenario, dwarf galaxies like Sgr did indeed start out as disky dwarfs systems (embedded in extended dark matter halos), but in the presence of the tidal field of a larger host galaxy like the MW, these systems are transformed into spheroids via formation and subsequent shortening of a bar, and through this process the stellar motions are altered from ordered to random. Depending on what phase of this dynamical transformation the dwarf galaxy is viewed, it can be highly elongated (e.g., just after bar formation) to very spherical (e.g., much further evolved). The best fitting model for the Sgr case explored by {\\L}okas et al. places the presently observed phase just beyond the second perigalacticon, intermediate between the phase of bar formation at first perigalacticon (a phase where Sgr would have tidal arms too short) and the third perigalacticon (a phase where Sgr would be spherical). This is the first model to explain both the very large ellipticity of Sgr as well as the observed kinematics of its stars, but presupposes that the current Sgr orbit is a recent product of dynamical friction from a larger orbit having an apocenter exceeding 100 kpc on which Sgr was deposited after cosmological infall into the MW halo. This orbital evolution could occur if the initial mass of the system was large --- similar to that of the Large Magellanic Cloud, a system itself that has probably only recently been accreted (e.g., Besla et al. 2007) and that presently contains a bar, a morphology consistent with the early phases of the proposed evolutionary scenario. Coincident with the orbital erosion from dynamical friction, the Sgr system must have shed a significant amount of mass, mostly dark matter, but also probably baryons, to bring it to its presently smaller size; a large, LMC-mass progenitor would be required to achieve a significant amount of orbital erosion (Colpi et al. 1999; Jiang \\& Binney 2000; Taffoni et al. 2003). That the LMC may be an apt model for the Sgr progenitor is supported by the similarity in extended star formation histories (e.g., Siegel et al. 2007; Harris \\& Zaritsky 2009) and detailed chemical evolution (Chou et al. 2010). Indeed, as we reiterate from {\\L}okas et al. (2010a), the primary difference in the present appearance of the LMC and Sgr may well have to do with timing --- i.e., the phase of dynamical evolution dictated by the relative orbital sizes and times they have been bound to the MW.\\footnote{For example, Sgr differs from the LMC in being presently devoid of gas, but the former clearly had gas as recently as $\\sim$0.75 Gyr ago to create its youngest stellar population (Siegel et al. 2007). That last star formation episode may well have depleted Sgr's gas reservoir, and/or other processes --- such as ram pressure stripping or supernovae blowout --- may have contributed. All three processes --- (1) gas depletion, (2) ram pressure stripping, and (3) supernovae blowout --- are conceivably accelerated in the case of Sgr because of (1) more frequent and intense tidal shocking to produce starbursts at its smaller orbital radius, (2) the higher density of the ram pressure medium at this orbital radius, or (3) the weakened binding energy for the more dynamically evolved, tidally diminished Sgr core.} \\subsection{Is the Disrupting Sagittarius Galaxy a Dwarf Spheroidal Exception or Rosetta Stone?} In the previous section we made a connection of the Sgr system to dwarf irregular/dwarf spiral galaxies via dynamical evolution models with tidal stirring and through observed similarities in star formation and enrichment history. But how does Sgr fit within the context of the dwarf spheroidal galaxies with which it is normally morphologically classified, and what does it potentially tell us about the origin and evolution of these systems? More than seventy years after their discovery, there still remains no universally agreed upon, complete picture of the nature of dwarf spheroidal (dSph) galaxies. Hodge (1964a, 1961b, 1962, 1966) and Hodge \\& Michie (1969) first invoked tidal effects to explain the diffuse structure of at least some dSphs. On the other hand, ever since the work of Aaronson (1983) to derive the first radial velocity dispersion ($\\sigma_{v}$) of a dSph, the notion that they are laden with dark matter (DM) has emerged as the standard paradigm to explain the observed structure and dynamics of these systems. In the last decade, enormous efforts have been plied to collecting data on dSphs to their greatest radial extent. Evidence for very distended structure has been reported for the classical (i.e., the most luminous, not including the ``ultrafaints\") MW dSphs (e.g., Majewski et al.\\ 2000; Martinez-Delgado et al.\\ 2001; Palma et al.\\ 2003; Westfall et al.\\ 2006; Mu\\~noz et al.\\ 2006; Sohn et al. 2006; some of these data are collected in Figure~\\ref{fig:all_dense}). These observed morphologies might be used to infer either that dSphs are being tidally disrupted, with the stripped stars accounting for the extended structures (the general view of the above cited sources; see also {\\L}okas et al. 2008), or that the dSphs are {\\it very} large, fully internally-bound, massive structures (e.g., Gilmore et al. 2007; Wu 2007).\\footnote{In a study of the Sculptor dSph, Coleman et al. (2005a) suggest yet another interpretation of the extended structure as one relating to variations in the distribution of stellar populations; nevertheless these authors do not rule out the possibility that a small extratidal component may exist in this system. Meanwhile, Coleman et al. (2005b) identify an apparent excess of stars around the Fornax system as due to the remains of another, smaller satellite that previously merged with Fornax.} The latter view has often been bolstered by the advent of spectroscopic studies that have collected hundreds, and in some cases thousands, of individual radial velocities over much of the luminous extent of dSphs. These new data sets have revealed flat or just slowly rising/declining velocity dispersion profiles all the way to the limits of where data exist (e.g., Westfall et al.\\ 2006; Sohn et al.\\ 2006; Mu\\~noz et al.\\ 2006; Walker et al.\\ 2007 --- some of these data are summarized in Fig.~\\ref{fig:all_disp}), a behavior that might be reasonably well explained by the assumption that dSphs live inside much larger DM halos in dynamical equilibrium (e.g., Kleyna et al. 2002; {\\L}okas 2002; {\\L}okas et al. 2005; Maschenko et al. 2005; Gilmore et al. 2007; Koch et al. 2007; Wu 2007; Walker et al. 2007) and resulting in dramatically rising mass-to-light ratios with radius and large inferred total satellite masses. However, as argued by Mu\\~noz et al. (2005, 2006), the extent of some of the mapped dSphs, as verified through RV-membership studies of extreme outlying stars in the luminosity profiles, imply very large linear dimensions, mass-to-light ratios, and total masses if all of the identified outer stellar members are assumed to be bound; for example, in the case of Carina, Mu\\~noz et al. obtain $M/L>16,000$, $M_{dSph}=7.2\\times10^{9}$ and $R_{tidal}\\sim4$\\,kpc to keep the outermost RV-member star bound. While there is currently debate about whether $M/L$ ratios this large make sense for the classical MW dSphs --- e.g., whether they sit in the largest DM subhalos seen in simulations or in systematically smaller halos (Ferrero et al. 2011; Boylan-Kolchin et al. 2011) --- the above-implied dimensions for Carina rival those of the Magellanic Clouds and are not likely to hold universally for all classical MW dSphs based on expected subhalo mass spectra (e.g., Moore et al. 1999). It seems even less likely that these dimensions would hold, just by chance, for the first few examples of dSphs that have received the most radially extensive observational attention. Alternatively, the observed flat velocity dispersion profiles up to and beyond the observed King limiting radii, $r_{lim}$, of these systems might also be taken as further proof for a more universal tendency for the classical MW dSphs, while containing DM to be sure, to be experiencing tidal disruption of their luminous parts as well. It is now well established that the MW halo contains extensive stellar substructure from hierarchically merged ``subhalos'', and dSph galaxies are prime candidates for the progenitors of these substructures. Moreover, Sohn et al.\\ (2006); Mu\\~noz et al.\\ (2008); {\\L}okas et al. (2008, 2010b), among others, have successfully produced $N$-body simulations of DM-filled, tidally disrupting, mass-follows-light satellites that well match the observed properties of the Carina and Leo I dSphs --- two of the three dSph systems (apart from Sgr) with the most extensive radial coverage in their velocity mappings. Similar models are also seen to work for Fornax, Sculptor and Sextans ({\\L}okas 2009). Here we lend further support to this morphological and evolutionary paradigm from a decidedly phenomenological point of view. Figure ~\\ref{fig:all_dense}, a synthesis of previous structural studies (see references above), clearly demonstrates the similarity of the Sgr {\\it morphology} to that of other dSphs --- i.e., an inner King profile with, at the largest radii, an extended, slower declining ``break\\footnote{That is, ``breaking away'' from the King profile.} population'' --- whereas the present study demonstrates quite vividly that the Sgr dwarf galaxy, {\\it a DM-dominated system that nonetheless is demonstrably undergoing tidal disruption}, also has observed kinematics strongly resembling those of the other classical MW dSph galaxies. Of course, the lack of any significant rotation in Sgr is commensurate with the kinematics of other dSphs. But, in addition, as Figure~\\ref{fig:all_disp} demonstrates (data from Mu\\~noz et al 2005, 2006; Walker et al.\\ 2007; and Sohn et al.\\ 2006), the velocity dispersion profile for Sgr as a function of radius for the data presented in this study also strongly imitates those for other MW dSphs. Just like the velocity dispersion profiles of the other classical MW dSphs, that of Sgr remains more or less flat to large radii, including that part of the system where it is clearly dominated by tidal debris. As may be seen from the compilation of properties of the classical MW dSph systems in Table~\\ref{Tab:dwarfs}, Sgr is at one extreme within this group in terms of its distance, luminosity, linear size, ellipticity and derived mass, which might merit it being considered ``exceptional'' among the group. However, while Sgr is the most luminous and most massive dSph, it is not significantly brighter or more massive than Fornax, and its central velocity dispersion is similar to those of Sculptor, Draco, Ursa Minor, and Leo I and less than that of Fornax. Sgr also has an extended star formation history, like Fornax, Carina, Sculptor, Sextans, Leo I and Leo II. Meanwhile, we would argue that Sgr's large ellipticity and half-light radius may well be a function of its presently small Galactocentric distance: As discussed above (\\S 6.2) the ellipticity may be due to the recent infall of the system to a small orbit that induced tidal stirring, whereas the large present size may be due to Sgr recently experiencing catastrophic tidal disruption and mass loss, as argued by Chou et al. (2007). Indeed, guided by the remarkable morphological and dynamical similarity of Sgr to the other MW dSph systems as well as its compatibility with the tidal stirring model, we conclude that Sgr may sometimes appear to be a ``dSph outlier'' simply because it is presently in a more flashy, but intermediate phase of a more universal evolutionary sequence from a disky, star-forming, ``LMC-like\" state to a more spherical, staid, dwarf spheroidal state long past active star formation, like the Draco system. As shown by {\\L}okas et al. (2010a, 2011), the exact timescale for a dSph to follow this evolution is a function of its mass and orbital size, with the latter, of course, a function of the former through dynamical friction and the zero-age of the evolution set by the epoch of initial infall into the MW. In the case of Sgr, its present relative ``flashiness'' is driven by its currently small orbit (which amplifies the tidal stirring effect), proximity to perigalacticon (the most dramatic phase of both tidal stirring and mass loss), and its having recently formed significant stellar tidal tails (again, a function of the stronger tidal force experienced on the smaller orbit). If star formation history and morphology present ersatz timestamps of this ``universal'' evolutionary sequence (admittedly crude timestamps, because of the vagaries of viewing perspective, orbital shape, stochasticity of star formation histories, and initial conditions), then perhaps the less bar-like ({\\L}okas et al. 2012) and recently star forming Fornax system may be another system in a somewhat earlier phase of transformation, while Carina and Leo I may be in a somewhat later phase. This proposed paradigm of a universal evolutionary track for tidally-stirred, infalling galaxies is explored in more detail, as a function of more parameters and including all Local Group galaxies in {\\L}okas et al. (2011). Our primary point here is that because Sgr offers a clear example {\\it already found in nature} of a disrupting dwarf galaxy system containing DM that shares so many properties --- e.g., density profile, lack of rotation, $\\sigma_v$ profile --- with other classical, DM-filled MW dSphs that it is worth considering that the paradigms of not only tidal influence (e.g., stirring), but tidal {\\it disruption}, may apply more ubiquitously to these other (classical) dSphs and that perhaps these other objects are also suffering varying degrees of tidal disruption, along the Sgr paradigm. In fact, as found by {\\L}okas et al. (2010a), morphological transformation and mass loss take place together. Thus, we contend that tidal disruption, as in the case of Sgr, remains a viable explanation for the observed kinematics and structure at large radii of these other dSphs (as already suggested by e.g., Kuhn et al. 1996; Majewski et al. 2000; Palma et al.\\ 2003; Westfall et al.\\ 2006; Mu\\~noz et al.\\ 2006; Sohn et al. 2007; Mateo et al. 2008; {\\L}okas et al. 2008 and modeled as DM-dominated but tidally influenced systems by, e.g., Sohn et al. 2007; Mu\\~noz et al. 2006, 2008; {\\L}okas et al. 2008). Indeed, with Sgr a widely-accepted representative of the tidal debris model, but with good cases in hand also for the Carina and Leo I systems (modeled well by even simple mass-follows-light DM models), one may well question the need to have a second explanation --- i.e., large, extended DM halos to keep the observed extended stellar distributions bound and dynamically hot --- to explain the outer structure of the other classical MW dSphs. In the least, if there are two mechanisms responsible for the outer structure and dynamics of dSphs, our results here show that dynamical studies are hard pressed to discriminate them. Moreover, our analysis of the kinematics of the Sgr system demonstrate clearly that flat velocity dispersion profiles do {\\it not} prove the existence of extended DM halos around dSphs. On the other hand, tidal tails, seen morphologically and not just inferred from kinematics, prove when extended DM halos are not present. Thus, one clear means by which we might hope to distinguish between structural/dynamical models --- i.e., extended DM halos versus unbound tidal debris --- for specific dSphs is to build more extensive and refined morphological maps capable of detecting (subtle) tidal tails. A more sensitive and systematic search for evidence of tidal tails around the Galactic dSph population would help build tighter empirical constraints on their true masses and help answer the questions of whether dSphs typically live in truly huge ($10^{10-11}$ M$_{\\odot}$) DM subhalos or smaller ones and what this implies for the efficiency of galaxy formation on these scales (Ferrero et al. 2011; Boylan-Kolchin et al. 2011). In conclusion, in contrast to the way Sgr is often portrayed in studies of MW dSphs (see Section 1), we argue that it may not be an anomaly, but rather a Rosetta Stone of dSph galaxies and their evolution in a MW-like environment. Perhaps Sgr's only overriding exceptional quality is that it is the MW dSph presently in the most dramatic stages of a universal evolutionary path that includes tidal harassment and tidal disruption." }, "1207/1207.6425_arXiv.txt": { "abstract": "We consider the implications of Lorentz-invariance violation (LIV) on cosmogenic neutrino observations, with particular focus on the constraints imposed on several well-developed models for ultra-high energy cosmogenic neutrino production by recent results from the Antarctic Impulsive Transient Antenna (ANITA) long-duration balloon payload, and Radio Ice Cherenkov Experiment (RICE) at the South Pole. Under a scenario proposed originally by Coleman and Glashow, each lepton family may attain maximum velocities that can exceed $c$, leading to energy-loss through several interaction channels during propagation. We show that future observations of cosmogenic neutrinos will provide by far the most stringent limit on LIV in the neutrino sector. We derive the implied level of LIV required to suppress observation of predicted fluxes from several mainstream cosmogenic neutrino models, and specifically those recently constrained by the ANITA and RICE experiments. We simulate via detailed Monte Carlo code the propagation of cosmogenic neutrino fluxes in the presence of LIV-induced energy losses. We show that this process produces several detectable effects in the resulting attenuated neutrino spectra, even at LIV-induced neutrino superluminality of $(u_{\\nu}-c)/c \\simeq 10^{-26}$, about 13 orders of magnitude below current bounds. ", "introduction": " ", "conclusions": "" }, "1207/1207.2492.txt": { "abstract": "{The number and spatial distribution of confirmed quasi-stellar objects (QSOs) behind the Magellanic system is limited. This undermines their use as astrometric reference objects for different types of studies.} {We have searched for criteria to identify candidate QSOs using observations from the VISTA survey of the Magellanic Clouds system (VMC) that provides photometry in the $YJK_\\mathrm{s}$ bands and $12$ epochs in the $K_\\mathrm{s}$ band.} {The $(Y-J)$ versus $(J-K_\\mathrm{s})$ diagram has been used to distinguish QSO candidates from Milky Way stars and stars of the Magellanic Clouds. Then, the slope of variation in the $K_\\mathrm{s}$ band has been used to identify a sample of high confidence candidates. These criteria were developed based on the properties of $117$ known QSOs presently observed by the VMC survey.} {VMC $YJK_\\mathrm{s}$ magnitudes and $K_\\mathrm{s}$ light-curves of known QSOs behind the Magellanic system are presented. About $75$\\% of them show a slope of variation in K$_\\mathrm{s}$ $>10^{-4}$ mag/day and the shape of the light-curve is in general irregular and without any clear periodicity. The number of QSO candidates found in tiles including the South Ecliptic Pole and the 30 Doradus regions is $22$ and $26$, respectively, with a $\\sim20$\\% contamination by young stellar objects, planetary nebulae, stars and normal galaxies. } {By extrapolating the number of QSO candidates to the entire VMC survey area we expect to find about $1\\,200$ QSOs behind the LMC, $400$ behind the SMC, $200$ behind the Bridge and $30$ behind the Stream areas, but not all will be suitable for astrometry. Further, the $K_\\mathrm{s}$ band light-curves can help support investigations of the mechanism responsible for the variations. } ", "introduction": "\\label{intro} The astrometric accuracy and the photometric sensitivity of observations made with VISTA is sufficiently good that we expect data from the VISTA Magellanic Clouds survey (VMC; Cioni et al. \\cite{cio11} hereafter Paper I) can be used to derive proper motions of the Magellanic Clouds (MCs). Such proper motion studies require a reference grid of bright, distant, non-moving, point like objects. Quasi-stellar objects (QSOs) provide such a grid (e.g. Kallivayalil et al. \\cite{kal06}; Costa et al. \\cite{cos11}). QSOs are point-like sources believed to be powered by accretion onto black holes in the centre of distant galaxies. QSOs can also be used as background sources to examine the composition of the MC interstellar medium along the line of sight (e.g. Redfield et al. \\cite{red06}; van Loon et al. \\cite{loo09}), and are important for studies of galaxy formation and evolution. The density of QSOs with $i<19$ mag is $\\sim 11$ per deg$^2$ as estimated from the Sloan Digital Sky Survey (SDSS) Quasar Catalogue (Schneider et al. \\cite{sch10}) but discovery of candidate QSOs behind the MCs is complicated by the necessity to distinguish candidate QSO\u00d5s from the dense stellar content of the MCs themselves. Candidate QSOs must then be observed spectroscopically to confirm which are true QSOs, and to make this process efficient the sample of candidate QSOs should be as clean as possible. Selection of candidate QSOs behind the MCs has been greatly improved by long-term multi-epoch observations, as part of micro-lensing projects such as the MAssive Compact Halo Objects (MACHO -- Alcock et al. \\cite{alc00}) and the Optical Gravitational Lensing Experiment (OGLE -- Udalski et al. \\cite{uda92}), and using data at various wavelengths, from ultraviolet (UV) to infrared (IR), that permit better removal of stellar objects (young stellar objects, planetary nebulae, hot and red stars) from QSO candidate samples. Methods used to identifying candidate QSOs include X-ray (Shanks et al. \\cite{sha91}) or radio emission (Schmidt \\cite{sch68}), mid-IR (Stern et al. \\cite{ste05}) and near-IR colours (see below). Flux variations probably associated with the accretion disc have also been used (Hook et al. \\cite{hok94}). QSO candidates are spectroscopically confirmed on the basis of optical and UV ionic emission lines, from which their redshifts are measured (e.g. Vanden Berk et al. \\cite{vbe01}). The most recent such investigation by Koz{\\l}owski et al. (\\cite{koz12}) focused on the southern edge of the Large Magellanic Cloud (LMC) and is relatively complete for objects with $I< 19.2$ mag, with a candidate sample based on Spitzer Space Telescope mid-IR colours, X-ray emission and/or optical variability. Spectra of their sample quadrupled the number of confirmed QSOs behind the LMC. Currently there are $360$ known QSOs behind the Magellanic system of which $233$ are behind the LMC, $100$ behind the Small Magellanic Cloud (SMC) and $27$ behind the inter-cloud region including the Magellanic Bridge, and many more QSO candidates awaiting follow-up observations to establish their nature (e.g. Koz{\\l}owski \\& Kochanek \\cite{koz09}; Kim et al. \\cite{kim12}). However these objects cover a limited area compared to the extent of the whole MC system being surveyed in the VMC survey. VMC detects sources as faint as $K_\\mathrm{s}=23.4$ mag (AB) with S/N$=5$, corresponding to the luminosity of sources below the old main-sequence turnoff in the LMC which occurs at $I\\sim 22$ mag. The $YJK_\\mathrm{s}$ VMC survey, which is multi-epoch in $K_\\mathrm{s}$, has the potential to considerably enlarge the parameter space for the search of QSOs, behind the Magellanic system, both in terms of sensitivity and spatial distribution. Near-IR criteria to select QSOs have been proposed previously. Kouzuma \\& Yamaoka's (\\cite{kou10}) criteria were based on 2MASS $JHK$ bands. A series of papers have described the $K$ excess ($K$X) method (Warren, Hewett \\& Foltz \\cite{war00}) using UKIDSS $JK$ bands and the SDSS $g$ band (e.g. Maddox et al. \\cite{mad12}; Mortlock et al. \\cite{mor12}). Findlay et al. (\\cite{fin12}) based their selection on single epoch VISTA $ZYJ$ data. Our aim is to establish VMC $YJK_\\mathrm{s}$ bands selection criteria using the known QSOs behind the Magellanic system which have already been observed, and investigate their utility in identifying new candidate QSOs in the areas being surveyed. The first stage of our method uses the $(Y-J)$ vs $(J-K_\\mathrm{s})$ diagram and the second stage uses $K_\\mathrm{s}$ variability. Section \\ref{data} describes the VMC data for the known QSOs behind the Magellanic system that have been covered so far. Section \\ref{results} uses VMC\u00d5s $YJK_\\mathrm{s}$ photometry of these QSOs to establish colour criteria for isolating them from Milky Way (MW) and MC objects, and to examine the $K_\\mathrm{s}$ variability of the QSOs over a baseline of up to $12$ epochs over two years. Section \\ref{discussion} discusses the reliability of our method and Sect. \\ref{conclusions} presents our conclusions. The Appendix provides further information on the known QSOs and their $K_\\mathrm{s}$ band light-curves. The influence of reddening and of photometrically selected non-QSOs on our selection criteria is also discussed in there. ", "conclusions": "\\label{conclusions} A large sample of QSOs across the Magellanic system is of fundamental importance as an astrometric reference source for investigations of the proper motion, of the interstellar and intergalactic medium along the line of sight and to address the nature of the extra-galactic sources themselves. We present criteria for selecting candidate QSOs based solely on the multi-epoch near-IR photometry from the VMC survey developed from the properties of $117$ known QSOs. Most of them occupy two neighbouring regions of the $(Y-J)$ vs. $(J-K_\\mathrm{s})$ diagram, one is dominated by QSOs with a star-like appearance and the other by QSOs with a galaxy-like appearance in the VMC data. The analysis of $K_\\mathrm{s}$ multi-epochs shows that $\\sim75$\\% of the QSOs have a slope of variation $>10^{-4}$ mag/day across a time range up to $600$ days. The combination of $YJK_\\mathrm{s}$ colours and $K_\\mathrm{s}$ variability criteria, with appropriate data quality flags, has been used to define samples of candidate QSOs with a $\\sim20$\\% contamination by non-QSOs on average, but the contamination is null in the region dominated by objects with a star-like appearance. In LMC tile $6\\_6$, including the 30 Doradus region, and $8\\_8$, including the South Ecliptic Pole region, we find $22$ and $26$ QSO candidates with photometric errors $<0.1$ mag and quality flags $=0$ in each VMC wave band, brighter than $19.32$, $19.09$ and $18.04$ mag in the $Y$, $J$ and $K_\\mathrm{s}$ band, respectively, and with a $K_\\mathrm{s}$ slope of variation $>10^{-4}$ mag/day. The selection of QSOs from the VMC data is in excellent agreement with the sample of high confidence QSOs by Kim et al. (\\cite{kim12}). PNe represent the major source of contamination in the regions of the colour-colour diagram dominated by star-like QSOs while normal galaxies and YSOs are the major contaminants in the region dominated by galaxy-like QSOs. The sub-set of sources with a slope of variation in the $K_\\mathrm{s}$ band $>10^{-4}$ mag/day is not influenced by PNe and red stars. The sub-set of sources with photometric errors $<0.1$ mag and quality flags $= 0$ in each VMC wave band is very small. Since known QSOs may be de-blended and have quality flags $=16$ the latter represents a reliable whilst not complete sample of candidates. VMC magnitudes and $K_\\mathrm{s}$ light-curves for known QSOs are provided in a table and in the appendix. The full sample of known QSOs comprises $332$ QSOs from the literature and $28$ newly confirmed QSOs by Kamath et al. (in prep.). VMC data on a fully observed tile in the outer LMC disc ($8\\_8$) have been used to estimate the number of QSO candidates behind the entire system as covered by the VMC survey. This is of the order of $1\\,200$ behind the LMC, $400$ behind the SMC, $200$ behind the Bridge and $30$ behind the Stream areas. The VMC survey is the most sensitive near-IR survey of the Magellanic system to date providing for the first time near-IR counterparts to many QSOs. Spectroscopic observations of the candidates identified here will support an extension of the selection to faint magnitudes exploiting the whole VMC range available." }, "1207/1207.2481_arXiv.txt": { "abstract": "Using the OSIRIS instrument installed on the 10.4-m Gran Telescopio Canarias (GTC) we acquired multi-color transit photometry of four small ($R_{p} \\aplt 5$ $R_{\\oplus}$) short-period ($P \\aplt 6$ days) planet candidates recently identified by the $Kepler$ space mission. These observations are part of a program to constrain the false positive rate for small, short-period $Kepler$ planet candidates. Since planetary transits should be largely achromatic when observed at different wavelengths (excluding the small color changes due to stellar limb darkening), we use the observed transit color to identify candidates as either false positives (e.g., a blend with a stellar eclipsing binary either in the background/foreground or bound to the target star) or validated planets. Our results include the identification of KOI 225.01 and KOI 1187.01 as false positives and the tentative validation of KOI 420.01 and KOI 526.01 as planets. The probability of identifying two false positives out of a sample of four targets is less than 1\\%, assuming an overall false positive rate for $Kepler$ planet candidates of 10\\% \\citep[as estimated by][]{morton11}. Therefore, these results suggest a higher false positive rate for the small, short-period $Kepler$ planet candidates than has been theoretically predicted by other studies which consider the $Kepler$ planet candidate sample as a whole. Furthermore, our results are consistent with a recent Doppler study of short-period giant $Kepler$ planet candidates \\citep{santerne12}. We also investigate how the false positive rate for our sample varies with different planetary and stellar properties. Our results suggest that the false positive rate varies significantly with orbital period and is largest at the shortest orbital periods ($P < 3$ days), where there is a corresponding rise in the number of detached eclipsing binary stars (i.e., systems that can easily mimic planetary transits) that have been discovered by $Kepler$. However, we do not find significant correlations between the false positive rate and other planetary or stellar properties. Our sample size is not yet large enough to determine if orbital period plays the largest role in determining the false positive rate, but we discuss plans for future observations of additional $Kepler$ candidates and compare our program focusing on relatively faint $Kepler$ targets from the GTC with follow-up of $Kepler$ targets that has been done with warm-$Spitzer$. ", "introduction": "\\label{intro} The $Kepler$ space mission has discovered, to date, 61 transiting planets as well as over 2,000 planet candidates and 2,000 eclipsing binary stars \\citep{batalha12,prsa11,slawson11}.\\footnote{Up to date catalogs can be found at http://kepler.nasa.gov} With such a vast number of planet candidates, it can be difficult to decide which to focus follow-up efforts on. Recent studies have tried to address this issue by estimating the false positive rate for the $Kepler$ sample. Based on the list of 1,235 $Kepler$ planet candidates published by \\citet{borucki11b}, it has been predicted that as many as $\\sim$95\\% of these candidates are true planets \\citep{morton11}. However, previous studies have not taken into account how different subsets of $Kepler$ targets may have different false positive rates. For example, there is a rapid rise in the number of detached eclipsing binary stars that have been discovered by $Kepler$ at orbital periods of less than $\\sim$ 3 days, and such systems can mimic planetary transits \\citep{prsa11, slawson11}. This suggests that there may be corresponding changes in the false positive rate with orbital period. Because the probability of observing a transit event increases as the orbital period of the planet decreases, many of $Kepler$'s planet candidates have short periods. Thus it is necessary to be cautious when estimating false positive rates over the whole $Kepler$ sample. While observational studies with warm-$Spitzer$ support predictions of low false positive rates over the entire $Kepler$ sample \\citep{desert12}, biases in target selection can affect observationally-constrained false positive rates (see \\S\\ref{discuss} for further discussion). We also note that a recent imaging study has found that nearly 42\\% of their sample of 98 $Kepler$ planet candidate hosts has a visual or bound companion within 6 arcseconds of the target star (Lillo-Box et al., in preparation). Such studies emphasize the need for follow-up imaging to exclude blend scenarios imitating planet candidates or contaminated transit depths. The false positive scenarios we consider in this paper are those that result from stellar eclipsing binaries that are either in the background (or, in rare cases, foreground) or bound to the target star and are not always easy to identify with $Kepler$ due to the flux from the different stars being blended together within $Kepler$'s photometric aperture ($\\sim$6 arcsec). As discussed by, e.g., \\citet{colon11}, different techniques can be used to eliminate many, but not all, blends. Multi-color transit photometry is an efficient method for recognizing blends that cannot be spatially resolved, as measuring the transit depth in different bandpasses (i.e. the transit color) allows one to test the planet hypothesis. This is possible since the magnitude of the transit color changes as long as the blended stars have significantly different colors. \\citet{colon11} presented multi-color transit photometry of a $Kepler$ target, KOI 565.01, that was first announced by \\citet{borucki11a} to be a super-Earth-size planet candidate but later recognized as a likely false positive due to measurements of a centroid shift away from the location of the target on the CCD during transit \\citep{borucki11b}. In \\citet{colon11}, we used near-simultaneous multi-color observations acquired using the narrow-band tunable filter imaging mode on the Optical System for Imaging and low Resolution Integrated Spectroscopy (OSIRIS) installed on the 10.4-m Gran Telescopio Canarias (GTC) to confirm that KOI 565.01 is indeed a false positive, as we both resolved a stellar eclipsing binary $\\sim$15 arcsec from the target and measured a color change in the ``unresolved'' target+eclipsing binary system. Thus, \\citet{colon11} demonstrated the capability of the GTC/OSIRIS for efficient vetting of planet candidates via its capabilities for near-simultaneous multi-color photometry within a single transit event. In this paper, we present observations of four $Kepler$ planet candidates specifically selected to have small radii and short orbital periods ($R_{p} \\aplt 5$ $R_{\\oplus}$ and $P \\aplt 6$ days). Measuring the false positive rate for this extreme subset of planet candidates will allow us to test if there is a correlation between the false positive rate and different planetary and stellar properties. As in \\citet{colon11}, the observations presented here were acquired with the GTC/OSIRIS. However, these observations used broadband filters in lieu of the narrow-band tunable filters in order to collect more photons and to obtain greater wavelength coverage (and thereby probe greater color differences). In \\S\\ref{targets} we discuss our target selection criteria, and we discuss the corresponding observations for our four targets in \\S\\ref{obs}. In \\S\\ref{reduction} and \\S\\ref{analysis} we describe our data reduction and light curve analysis procedures, and we present results for each target in \\S\\ref{results}. We include a discussion of our results and how they relate to the distribution of eclipsing binaries that have been discovered by $Kepler$ as well as previous estimates of the false positive rate for the $Kepler$ sample in \\S\\ref{discuss}. In particular, we discuss theoretical estimates from \\citet{morton11} and observational constraints from studies by \\citet{desert12} and \\citet{santerne12}. Finally, we summarize our results and conclusions in \\S\\ref{conc}, and we also discuss our plans for future observations of additional $Kepler$ targets with the GTC. ", "conclusions": "\\label{conc} We acquired multi-color transit photometry of four small ($\\sim$2.5$-$4.9 R$_{\\oplus}$), short-period ($\\sim$0.37$-$6.0 days) $Kepler$ planet candidates with the GTC. Based on the transit color, we identified two candidates as false positives. For two, we find further evidence supporting the planet hypothesis, consistent with {\\it validated} planets. We also remind the reader of KOI 565.01, which was a planet candidate observed in a similar fashion and that was also found to be a false positive (albeit it was selected from the first KOI catalog published by Borucki et al. 2011a; Col{\\'o}n \\& Ford 2011). While we find a high false positive rate (2/4 or 3/5, if we include KOI 565) in our small sample, we caution that this is likely not representative of the entire sample of $Kepler$ planet candidates, due to the small number of targets we observed and the specific properties of these candidates (e.g. the orbital period and size). Nevertheless, our results demonstrate the importance of considering these properties when evaluating the false positive probability of specific systems. While our findings seem to contradict the theoretical estimates from \\citet{morton11}, the low false positive rate that they estimate is based on the assumption that all candidates had passed preliminary false-positive vetting metrics based on $Kepler$ photometry and astrometry. Thus, if we consider the obviously V-shaped transits for KOI 225.01 and KOI 1187.01 to imply a false positive nature for these KOIs, then according to \\citet{morton11} the low false positive probabilities for these targets are not accurate. The false positive rate for our sample is also much larger than the observational constraints from \\citet{desert12} that predict that the false positive rate is much less than 10\\%. This is likely partly a result of different targets being probed by the different studies. {The recent study by \\citet{santerne12} further supports this idea, as they found a $\\sim$35\\% false positive rate for short-period giant $Kepler$ planet candidates}. We plan to continue observing small, short-period KOIs with the GTC in order to improve the sample size of our study. The observations presented here, as well as future observations with the GTC, greatly complement follow-up of KOIs done with warm-$Spitzer$ as well as other observatories. As we can expect the number of $Kepler$ planet candidates to continue to increase, we can use results from all such studies to pinpoint which targets are the best to follow-up in order to maximize the science output from the $Kepler$ mission." }, "1207/1207.0351_arXiv.txt": { "abstract": "We study the evolution of galactic bars and the link with disk and spheroid formation in a sample of zoom-in cosmological simulations. Our simulation sample focuses on galaxies with present-day stellar masses in the $10^{10-11}$\\,M$_{\\odot}$ range, in field and loose group environments, with a broad variety of mass growth histories. In our models, bars are almost absent from the progenitors of present-day spirals at $z>1.5$, and they remain rare and generally too weak to be observable down to $z \\approx 1$. After this characteristic epoch, the fractions of observable and strong bars raise rapidly, bars being present in 80\\% of spiral galaxies and easily observable in two thirds of these at $z \\leq 0.5$. This is quantitatively consistent with the redshift evolution of the observed bar fraction, although the latter is presently known up to $z \\approx0.8$ because of band-shifting and resolution effects. Our models hence predict that the decrease in the bar fraction with increasing redshift should continue with a fraction of observable bars not larger than 10-15\\% in disk galaxies at $z>1$. Our models also predict later bar formation in lower-mass galaxies, in agreement with existing data. We find that the characteristic epoch of bar formation, namely redshift $z \\approx 0.8-1$ in the studied mass range, corresponds to the epoch at which today's spirals acquire their disk-dominated morphology. At higher redshift, disks tend to be rapidly destroyed by mergers and gravitational instabilities and rarely develop significant bars. We hence suggest that the bar formation epoch corresponds to the transition between an early ``violent'' phase of spiral galaxy formation at $z\\geq 1$ and a late ``secular'' phase at $z \\leq 0.8$. In the secular phase, the presence of bars substantially contributes to the growth of the (pseudo-)bulge, but the bulge mass budget remains statistically dominated by the contribution of mergers, interactions and disk instabilities at high redshift. Early bars at $z>1$ are often short-lived, while most of the bars formed at $z\\leq1$ persist down to $z=0$, late cosmological gas infall being necessary to maintain some of them. ", "introduction": "Bars are one of the most frequently and easily quantified substructures in spiral galaxies, and are hence often used as a tracer of galaxy evolution. Most spiral galaxies today contain a central bar, although with largely variable amplitudes \\citep{block02,whyte02}. Spiral arms are equally common in optical light, but are much weaker in the near-infrared light that more closely traces the stellar mass distribution, and the strength of bars is generally easier to quantify independently of the imaging sensitivity. Both the formation of bars and their time evolution are connected to the baryonic and dark matter properties of their host galaxies, and their mass assembly history. Bars form spontaneously in stellar disks that are sufficiently massive and dynamically cold to be gravitationally unstable, with typical Toomre stability parameters $Q \\approx 1.5-2.0$ \\citep{toomre63, combes-sanders, gerin, combes-elmegreen}, and can also be amplified by dynamical friction on the dark matter halo \\citep{2002ApJ...569L..83A}. The gaseous content can also trigger bars: gas helps to form outer spiral arms, which can remove angular momentum from the inner regions and strengthen a bar seed \\citep{BCS05}. Once formed, bars evolve through the exchange of angular momentum with the dark matter halos \\citep{1985MNRAS.213..451W,deb-sel2000}, as well as with stellar and gaseous disks \\citep{1993A&A...268...65F,BC02,BCS05}. This can reinforce bars when they lose angular momentum, but can also weaken and destroy them in the opposite case \\citep{BCS05}. Bars can also be weakened or destroyed by the growth of central mass concentrations \\citep{norman,pfenniger,hasan}, however central concentrations with low-enough mass and/or low-enough mass growth rates could have little effect on real bars \\citep{2005MNRAS.363..496A}. If bars are formed in conditions where they are intrinsically short-lived, sufficient accretion of external gas onto the disk could enable their survival or re-formation \\citep{BC02}. In addition, galaxy interactions could in theory trigger bar (re-)formation \\citep{gerin,1998ApJ...499..149M,2004MNRAS.347..220B}, although observations do not show a clear environmental dependence of the bar in disk galaxies \\citep{vdbergh,2009A&A...495..491A,2011MNRAS.410L..18B,2012ApJ...745..125L}. Hence, the fraction of barred galaxies, and the redshift evolution of this fraction, are fundamental {\\em tracers} of the evolution history of galaxies: this indicates when disks became sufficiently massive and self-gravitating to be bar-unstable, and whether the conditions for bars being long-lived or re-formed were met. Furthermore, bars can directly drive structural evolution of their host galaxies. They trigger radial gas flows and may provide gas to nuclear disks and central black holes \\citep{2000ApJ...529...93K,2002ApJ...567...97L,2004ApJ...607..103L}. They can also thicken through vertical resonances, leading to the formation of pseudo-bulges -- i.e. , bulges with relatively low concentration and substantial residual rotation \\citep[e.g.,][]{1999AJ....118..126B,kormendy-kennicutt,martinez-v06}. \\medskip In the nearby Universe, the bar fraction in disk galaxies is very high. Depending on classification techniques, the fraction of strong bars in the optical light is at least $50\\%$ \\citep{2008ApJ...675.1194B}. Optical classifications reveal roughly one third of strongly barred galaxies, one third of weakly or moderately barred galaxies, and one third of optically unbarred galaxies \\citep{1991trcb.book.....D}. In the near-infrared, where weak bars are not obscured by dust and more easily distinguished from spiral arms, the bar fraction is at least $80\\%$ \\citep{2000AJ....119..536E,block02,whyte02,2007ApJ...657..790M}. The first searches for bars at redshift $z>0.5$ found a very low bar fraction \\citep[e.g.,][]{1999MNRAS.308..569A}, possibly because of small number statistics. Their work also illustrated the difficulties to identify bars at high redshift: the observed optical light traces the ultraviolet emission, in which bars are harder to detect, even locally. Near-infrared data revealed a number of barred galaxies at $z \\geq 0.7$ \\citep{2003ApJ...592L..13S,2004ApJ...612..191E,2004ApJ...615L.105J}. The first sample large enough to robustly quantify the redshift evolution of the bar fraction without being affected by resolution and band-shifting bias up to $z \\approx 0.8$, was studied by \\cite{2008ApJ...675.1141S} in the COSMOS field. These observations indicate that {\\em the bar fraction drops by a factor of about three from $z=0$ to $z=0.8$}. This result holds both for all observable bars and for strong bars separately, and using either visual classifications or quantitative estimates of the bar strength. \\citet{2008ApJ...675.1141S} also found a downsizing-like behaviour for bar formation, i.e. more massive galaxies tend to get barred at higher redshifts. This trend can explain why previous studies, such as \\cite{2004ApJ...615L.105J}, using shallower data targeted to more massive systems, observed higher bar fractions -- but still consistent with a declining bar fraction \\citep[see also][]{2004ApJ...612..191E}. \\medskip In this paper, we study the evolution of bars in a sample of cosmological zoom-in simulations of 33 galaxies with present-day stellar masses ranging from $1\\times 10^{10}$ to $2\\times 10^{11}$\\,M$_{\\sun}$, in field and loose group environments. The simulation technique and structural evolution of these galaxies (bulge and disk fractions, angular momentum evolution) were presented in \\citet{martig12}. The paper is organized as follows. In Section~\\ref{Sec:simu_analysis} we present our simulations and methods for the identification of bars and morphological analysis. In Section~\\ref{Sec:redshift_evolution} we analyze the redshift evolution of the bar fraction in the whole sample and in disk-dominated galaxies\\footnote{most galaxies in this sample are disk-dominated spirals at $z=0$, but a larger fraction is spheroid-dominated at $z>1$}, using quantitative measurements of the bar strength. Our main result, the emergence of bars along the cosmic time that traces the epoch of thin disk formation subsequent to the growth of spheroids and thick disks, is presented in Section~\\ref{Sec:bars_vs_thin_disks}. Section~\\ref{Sec:bar_lifetime} studies the lifetime of bars and its dependence on external gas accretion. In Section~\\ref{Sec:bars_vs_bulge}, we quantify the contribution of bars in the late growth of bulges, comparing to unbarred galaxies. Finally, we discuss and summarize our results in Sections~\\ref{Sec:discussion} and~\\ref{Sec:summary}, respectively. ", "conclusions": "\\label{Sec:discussion} In our models, bars emerge in present-day spirals at $z\\sim 1$, being rare and only weak at higher redshift, and rapidly becoming ubiquitous at lower redshift. We suggested from our simulations that this corresponds to the transition between an early ``violent'' phase down to $z \\simeq 0.8-1$ and a late ``secular'' phase at lower redshifts. Indeed, the morphology of the progenitors of today's spirals is rapidly evolving and uncorrelated with their final structure in the proposed ``violent'' phase at $z>1$ (see \\citealt{martig12}), with mostly the thick disk and stellar spheroids forming in this phase, while after $z<1$ the thin disk grows and the structural parameters such as the bulge and disk fractions become well correlated with their final values at $z=0$. The influence of the dark matter halo can have additional effects on the evolution of bars (see e.g., \\citealt{2006ApJ...648..807B,2010MNRAS.406.2386M}, for the effects of the shape of halo and \\citealt{2007MNRAS.375..460W,2008ApJ...679..379S}, for the effects involving more general halo properties). As the dark matter halo evolves with redshift, it affects the evolution of the entire galaxy, in particular the bar, which in turn influences the halo itself. Bar formation can thus be reinforced or delayed depending on exact halo properties. Such effects should be resolved by our simulations, but the key epoch of bar formation appears to correspond mostly to the evolution of baryonic properties in our analysis. We show on Figure~\\ref{fig:images_z_evolution} three representative examples at $z>1$: one, which is disk-dominated at $z>1$ but where violent disk instabilities (including giant clumps) destroy the early disk into a thick disk and spheroid, another one, which is spheroid-dominated with a major merger at $z\\sim 2$, and a third one, which is also spheroid dominated with several minor mergers at $z\\simeq1-2$. This is illustrative of the ``violent phase'' at $z>1$. Indeed, it is known from other work that high-redshift disks have high gas fractions and are violently unstable with giant clumps and transient features, but do not frequently develop bars \\citep[e.g.,][]{BEE07, CDB10} and this instability destroys any thin disk that would have started to grow, while major mergers that can reform some disk components \\citep[e.g.,][]{robertson06} but mostly convert disks into spheroids even in high-redshift conditions \\citep{bournaud11}. Hence, in the two examples on Figure~\\ref{fig:images_z_evolution} and in the majority of our sample, no massive thin disk can stabilize before redshift one and develop a substantial bar. On the other hand, mergers with mass ratios larger than 5:1 are almost absent from our sample after $z=1$ and diffuse gas accretion occurs at much lower rate, and the rate of stellar bulge growth also drops. Thus a massive thin disk can form and start to dominate the mass distribution, as probed by the formation of spiral arms between $z=1$ and $z=0.5$ in the face-on images shown on Figure~\\ref{fig:images_z_evolution}. This thin disk grows secularly down to $z=0$ and generally gets barred by $z \\approx 0.5$ -- strongly in the two first cases, weakly in the third one. These three cases illustrate the transition between an early violent phase with frequent mergers and disk instabilities, and a late secular phase dominated by slower mass infall, and the fact that this transition occurs when the massive thin disk forms and is traced by the emergence of bars. Overall, the epoch of bar formation in our simulations probes the epoch at which spiral galaxies have formed the bulk of their disk, stellar halo and thick disk, and start to be dominated (in stellar mass) by their final thin disk with only slow (secular) evolution down to $z=0$. \\begin{figure*}[t] \\centering \\includegraphics[width=1.0\\textwidth]{multi_images_36_50.eps}\\\\ \\vskip 0.5cm \\includegraphics[width=1.0\\textwidth]{multi_images_106_50.eps}\\\\ \\vskip 0.5cm \\includegraphics[width=1.0\\textwidth]{multi_images_47_50.eps} \\caption{\\label{fig:images_z_evolution} Examples of morphological evolution from $z=2$ down to $z=0$ for three simulated galaxies. Stellar density maps (face-on $50 \\times 50$ kpc$^2$ projections) are shown for $z=2$, $1$, $0.5$ et $0$. Galaxies evolve rapidly during the ``violent'' phase at $z\\gtrsim1$ when they frequently undergo phases of violent disk instabilities (top panels) and major mergers (middle panels). The morphology at $z=0$ and $z\\gtrsim1$ are uncorrelated with early spheroids being possible progenitors of today's spirals (bottom panels). Once galaxies enter the late ``secular'' phase at $z<1$, their structural parameters become more tightly correlated with the final disk morphology. The color coding of the projected density maps is the same as in Figure~\\ref{fig:lin_fit}.} \\end{figure*} \\medskip In observations, the bar fraction also decreases with increasing redshift, and although it could be probed in detail only up to $z \\simeq 0.8$ it is in close agreement with our models so far. In the following we discuss whether the emergence of bars at $z\\simeq 0.8-1$, if confirmed observationally, could also correspond to the epoch at which spiral galaxies acquire their final morphology and start being dominated by their thin rotating stellar disk. We are thus led to wonder whether a transition between a violent phase of galactic assembly, with rapid episodes of merging and violent instabilities building a thick disk and spheroids, could be followed by a calm secular phase of thin disk growth and evolution, with the typical transition at $z\\simeq 0.8-1$. Chemical properties of disk galaxies could be consistent with such two-phase assembly of their components. Thick disks are ubiquitous around spiral galaxies \\citep{dalcanton-bernstein02,seth05}, without star formation and no or almost no young stellar populations \\citep{yoachim-dalcanton06, ibata09}. The Milky Way thick disk contains no or almost no star younger than 8~Gyr \\citep{gilmore1985, reddy2006}. This is consistent with the occurrence of events disrupting any cold rotating disk into a thick disk only at redshifts higher than $z\\sim 1$, while the thin disk wouldn't have been significantly disrupted/thickened after this epoch. The enhanced $\\alpha$ element abundances in the thick disk and central bulge \\citep{lecureur07, zoccali07} suggest that their star formation occurred mostly in brief events (not longer than a few $10^8$~yr), and that the formation of stars that belong to the present-day thin disk occurred at later epochs with longer timescales \\citep[see also][]{chiappini}. These brief events could have been mergers, as well as violent instabilities (giant clumps) in rapidly-accreting galaxies, that grow both a bulge and a thick disk \\citep{BEM09} over short timescales, while a thin disk component can form and remain stable after these phases with lower gas infall rates. \\medskip Direct searches for disks, through gas kinematics and/or optical and near-infrared spectroscopy, are also consistent with $z \\simeq 1$ as the typical redshift for the emergence of modern thin spiral disks, along with their bars. Near-infrared spectroscopy of a $z\\sim3$ sample of star-forming galaxies (AMAZE and LSD, \\citealt{maiolino10,2011A&A...528A..88G}), with masses typical for the progenitors of Milky Way-like galaxies and present-day spirals, reveals only a minor fraction of rotating disks, the majority of these galaxies being mergers or irregular systems dominated by high velocity dispersions. A similar survey of star-forming galaxies in a comparable mass range at a median redshift $z\\simeq1.2$ (MASSIV, \\citealt{2012A&A...539A..92E}) finds about 40-50\\% of rotating disk galaxies. At $z\\simeq 0.6-0.8$, the IMAGES survey finds a majority of rotating disk galaxies \\citep{yang08} in a sample that still covers masses typical for the progenitors of today's spiral galaxies in the $10^{10-11}\\,M_{\\odot}$ mass range. This survey furthermore suggests that many of these $z\\simeq 0.6-0.8$ galaxies have formed their disks only recently after undergoing violent events such as major mergers \\citep{2009A&A...507.1313H} and that they will undergo only slow evolution of their global properties, such as their Tully-Fisher relation, down to $z=0$ \\citep{puech10}. Morphological studies are also consistent with the emergence of modern thin disks around redshift $z\\sim1$. The Hubble Ultra Deep Field sample studied by \\citet{elmegreen07,elmegreen09} at $z>1$ is dominated by irregular morphologies corresponding to major merger and interactions and ``clumpy'' unstable disks, which are typically forming thick disks, bulges and stellar haloes \\citep{BEM09}, not thin spiral disks. Some cold spiral disks are found in this sample but their fraction is quite low before $z=1$. At $z\\simeq0.7$, the situation is largely different as the fraction of clumpy irregular disks drops steadily and stable spiral disks rapidly become more numerous (still for progenitors of present-day $10^{10-11}\\,M_{\\odot}$ galaxies). For somewhat more massive galaxies, \\citet{2007ApJS..172..434S} found that a large fraction of massive disks are in place around redshift one, but substantially fewer than at lower redshift. \\medskip Hence, existing observations are consistent with our suggestion that the emergence of galactic bars at z$\\simeq0.8-1$ traces the transition between an early ``violent'' phase during which stars that belong to the modern thick disk, bulge and halo are formed in systems that do not have a permanent disk-dominated structure, and a late ``secular'' phase of thin disk growth and evolution with ubiquitous bars and limited pseudo-bulge growth. In this picture, the downsizing of bar formation (\\citealt{2008ApJ...675.1141S} and Section~4.2) could correspond to the later termination of the violent phase and later disk stabilization in lower-mass galaxies. This would be consistent with the fact that both merging activity and violent disk instabilities should persist at lower redshift for lower mass galaxies \\citep[][and references therein]{bournaud12}. It is also possible that bars grow more rapidly in more massive systems once their cold, thin disk is stabilized \\citep{elmegreen-bars07}. The observed evolution of the bar fraction so far is consistent with our model, but further confirmation could be obtained by confirming the drop in the bar fraction at $z\\sim1$ and above, with almost only weak bars (strength $\\leq$ 0.2) being present at $z > 1$ for the mass range studied here. We studied a sample of cosmological zoom-in simulations of 33 Milky Way-mass galaxies in field and loose group environments from $z=2$ down to $z=0$. The method used to determine the presence of a bar is based on the decomposition of the stellar surface density profiles into Fourier components. We also analyzed the disk/spheroid structure of the modeled galaxies using the S\\'{e}rsic index of the surface density profile. Our main results are as follows: \\begin{enumerate}[i)] \\item The \\textit{total} bar fraction among spiral galaxies declines with increasing redshift. It drops from almost $90\\%$ at $z=0$ to about $50\\%$ at $z\\simeq1$ and to almost zero at $z\\simeq2$. The fraction of \\textit{observable} and \\textit{strong} bars declines from about $70\\%$ at $z=0$ to $10-20\\%$ at $z\\simeq1$ and to zero at $z=2$. This result holds for galaxies with mass range of $2\\times 10^{9}-8\\times 10^{10}$\\,M$_{\\sun}$ at $z\\sim1$ and $4\\times 10^{5}-1\\times 10^{10}$\\,M$_{\\sun}$ at $z\\sim2$, i.e. typical progenitors of Milky Way-like spirals. However, the bar fraction could remain higher at $z\\sim2$ for more massive galaxies, if the downsizing of bar formation observed in our sample still holds for higher masses/redshifts. \\item The epoch of bar formation traces the epoch of the emergence of final disk of spiral galaxies. This corresponds to the termination of an early ``violent'' phase at $z>1$, characterized by frequent mergers, violent disk instabilities and rapidly evolving structure, forming thick disks, bulges and stellar halos. It is followed by a ``secular'' phase at $z<0.8$, dominated by the slower growth and evolution of modern thin disks and limited bulge growth at late times. The $z=0.8-1.0$ transition epoch is for the mass of typical Milky Way progenitors, and tends to move to higher redshift for more massive systems. \\item We find that there is only a minor contribution of bars in the late growth ($z<1$) of (pseudo-)bulges in spiral galaxies. This late growth is dominated by continued cosmic infall and minor mergers rather than by bars. \\item Finally, early bars (formed at $z>1$) are often short lived and may reform several times. Bars formed below $z\\sim1$ are found to persist down to $z=0$, some of them being intrinsically short-lived but maintained by late cosmological gas infall. \\end{enumerate} If confirmed observationally, the scarcity of significant bars at $z \\geq 1$ would indicate, according to our models, that present day spirals and Milky Way-like galaxies have formed and stabilized their modern thin spiral disk only relatively late in their growth history, typically at $z \\simeq 0.8-1$. At earlier times, they would be mostly forming their spheroidal components (bulges, halos) and thick disk, under the effect of both hierarchical merging and violent instabilities in rapidly-accreting systems. The continuation of this violent phase with mergers, rapid cold gas accretion and disk instabilities down to lower redshift in lower mass galaxies \\citep[e.g.,][]{bournaud12} could then explain a ``downsizing''-like behaviour for bar formation." }, "1207/1207.0810_arXiv.txt": { "abstract": "{ We present a fast likelihood method for including event-level neutrino telescope data in parameter explorations of theories for new physics, and announce its public release as part of \\iclike. Our construction includes both angular and spectral information about neutrino events, as well as their total number. We also present a corresponding measure for simple model exclusion, which can be used for single models without reference to the rest of a parameter space. We perform a number of supersymmetric parameter scans with IceCube data to illustrate the utility of the method: example global fits and a signal recovery in the constrained minimal supersymmetric standard model (CMSSM), and a model exclusion exercise in a 7-parameter phenomenological version of the MSSM. The final IceCube detector configuration will probe almost the entire focus-point region of the CMSSM, as well as a number of MSSM-7 models that will not otherwise be accessible to e.g. direct detection. Our method accurately recovers the mock signal, and provides tight constraints on model parameters and derived quantities. We show that the inclusion of spectral information significantly improves the accuracy of the recovery, providing motivation for its use in future IceCube analyses. } ", "introduction": "\\label{intro} Despite ongoing efforts, we have yet to identify dark matter. One of the most promising candidate classes is the weakly-interacting massive particle (WIMP) \\cite{Bergstrom00,Bertone05,BertoneBook}, so named because WIMPs interact with standard model (SM) particles only via the weak nuclear force. WIMPs are appealing because the expected cosmological density of a thermal relic with a weak-scale annihilation cross-section is the same order of magnitude as the observed value \\cite{WMAP7}. WIMPs are typically sought via three observational channels: direct WIMP-nucleon scattering (e.g.~\\cite{CerdenoGreen10,Pato11,CoGeNTAnnMod11,CRESST11,XENON100,CDMSLowE12,COUPP12}), production at accelerators (e.g.~\\cite{White07,Bai10,Goodman10,Fox12a,Fox12b,CMS12a,Frandsen12}) and indirect detection of SM products of WIMP self-annihilation (e.g.~\\cite{Pamelapositron,CompositeAlt,LATDwarfComposite,MAGICSegue,VERITASSegue,IceCube09,IceCube09_KK}). A promising version of indirect detection is to use neutrino telescopes such as IceCube \\cite{IC40DM}, ANTARES \\cite{Lim09} and SuperKamiokande \\cite{SuperK11} to search for high-energy neutrinos produced by annihilation of WIMPs in the core of the Sun or Earth \\cite{Press85,Krauss:1985ApJ,Freese86,Krauss86,Gaisser86}. Due to their weak interactions with nuclei, such WIMPs would have scattered on solar nuclei and become gravitationally bound to the solar system, eventually returning to scatter repeatedly and settle (in the case of the Sun) to the solar core \\cite{Gould87b} (see also \\cite{Scott09} for a review of this process and impacts of WIMPs on stellar structure, and \\cite{Gould91, Lundberg04, Peter09a, Peter09, Sivertsson12} for additional corrections due to the influence of planets). WIMPs arise in many proposed extensions to the SM. Standard neutrino telescope analyses \\cite{IceCube09,IceCube09_KK,ICRC2011ic86,IC40DM,Lim09,SuperK11,Kumar12} take an effective view of WIMP interactions, placing limits on WIMP-nucleon scattering cross-sections as a function of the WIMP mass, by assuming a certain annihilation cross-section and final state. These limits are difficult to translate into actual particle models, where the annihilation cross-section and branching fractions may take on a range of different values for any given WIMP mass and nuclear-scattering cross-section. To properly interpret limits on neutrino fluxes in terms of the parameters (defined at e.g. the Lagrangian level) of a theory for new physics, it becomes necessary to compare the observed neutrino flux with the predicted neutrino signal for each individual point in the parameter space of the theory. Having translated the flux limits into direct constraints on the parameters of a theory, it is then also possible to compare and combine the sensitivities of multiple experiments, even if they probe entirely different sectors of the theory (e.g. neutrino and accelerator searches). Within the realm of theories for new weak-scale physics, this `global fit' approach has so far been applied mostly to supersymmetry (SUSY) \\cite{Baltz04,SFitter,Allanach06,Ruiz06,Trotta08,Fittino,Abdussalam09a, Abdussalam09b,Mastercode12}, in the form of the minimal supersymmetric standard model (MSSM). Recent analyses have included new data from the Large Hadron Collider (LHC) \\cite{MastercodeHiggs,SuperBayeSHiggs,Gfitter11}, direct \\cite{BertoneLHCDD,Akrami11DD,MastercodeXENON100,SuperBayeSXENON100} and indirect detection \\cite{Scott09c,Ripken11,Fittino12}. Whilst great care needs to be taken over the details of the statistical methods employed \\cite{Akrami09,SBSpike,SBcoverage,Akrami11coverage,Strege12}, these analyses have proven a clear success, pointing the way to a future of closer comparison between astronomical and terrestrial experiments. To date global fits have not included neutrino telescope data. The ability of future incarnations of IceCube to detect WIMP annihilation in the constrained MSSM (CMSSM) has been studied \\cite{Trotta09,Nightmare,Roszkowski12}, based on the number of observed events only and simple estimates of the instrumental sensitivity. Other authors have looked at predicted neutrino rates in 2D slices through MSSM parameter spaces (e.g. \\cite{Feng01,Barger02,Baer04,Ellis09,Ellis11}), or from sets of models in non-statistical random scans of various incarnations of the MSSM (e.g. \\cite{Bergstrom98b,Wikstrom09,Cotta12}). Here we present a fast likelihood method for including full event-level data from neutrino telescopes in global fits and related analyses. In particular, our formulation allows the directions and energy estimators associated with each event to be included in the final unbinned likelihood calculation, which can then be employed as a likelihood component in a global fit. The inclusion of spectral information in neutrino searches for WIMPs has been of particular interest recently \\cite{Rott11,Allahverdi12}. As a by-product of our approach, we also present a rigorous but simple exclusion measure for individual models, based just on the observed number of events in a neutrino telescope. We give full details of the input data required, and explicit examples using real data from the IceCube neutrino telescope. The IceCube data and simulations we employ are described in Section \\ref{icecube}. These data have been made publicly available on the IceCube webserver \\cite{filesloc} and in \\iclike.\\footnote{\\href{www.darksusy.org}{www.darksusy.org}} The likelihood formalism, outlined in Section \\ref{likelihood}, is implemented and also available in \\iclike. Our example global fits and MSSM scans are detailed in Section \\ref{examples}. We summarise our method and examples in Section \\ref{conclusions}. ", "conclusions": "\\label{conclusions} We have constructed an unbinned likelihood formalism for including full event-level IceCube data in parameter explorations of theories for new physics. The likelihood function includes information about the number, direction and spectral characteristics of neutrino events, and is fast to calculate. We also constructed a simple associated measure for model exclusion, which is even faster to calculate and can be used for single models, without any reference to other parts of a parameter space. We performed a number of example model scans using our likelihood construction within the MSSM. We carried out global fits to the CMSSM with actual 22-string data, and with the 22-string effective area rescaled to represent the final detector configuration. Existing 22-string data has little impact on the CMSSM, but the final detector configuration will robustly exclude the majority of the focus point region if IceCube finds no evidence for WIMP annihilation. We carried out a mock signal recovery in the CMSSM, showing that our method accurately recovers a benchmark point, and constrains model parameters very well. In the process, we showed the utility of including spectral information in IceCube searches for dark matter. Finally, we carried out a simple example model-exclusion analysis in the MSSM-7, showing that the 86-string configuration of IceCube will test a number of models that cannot be probed by direct detection experiments in the near future. Our likelihood construction is implemented in \\iclike, and freely accessible to the community. It will also be made available in a future release of \\textsf{SuperBayeS}. The data and simulations we have used for this study, including 22-string IceCube event lists \\cite{IceCube09}, are freely available on the web \\cite{filesloc}, and in \\iclike." }, "1207/1207.5432_arXiv.txt": { "abstract": "Vacuum ultraviolet light sensitive photomultiplier tubes directly coupled to liquid xenon are being used to efficiently detect the 178 nm scintillation light in a variety of liquid xenon based particle detectors. Good knowledge of the performance of these photomultipliers under cryogenic conditions is needed to properly characterize these detectors. Here, we report on measurements of the quantum efficiency of Hamamatsu R8520 photomultipliers, used in the XENON Dark Matter Experiments. The quantum efficiency measurements at room temperature agree with the values provided by Hamamatsu. At low temperatures, between 160K and 170K, the quantum efficiency increases by $\\sim5-11$\\% relative to the room temperature values. ", "introduction": "\\label{} Liquid Xenon (LXe) detectors are currently a favored choice for rare event search experiments, from dark matter direct detection to neutrinoless double beta decay~\\cite{Aprile:2009dv}~\\cite{Danilov:2000}~\\cite{Ackerman:2011}~\\cite{Abe:2009}. In particular, the XENON100 dark matter experiment~\\cite{Aprile:2011}\\cite{Aprile:2011hi} uses 242 Hamamatsu R8520-06-Al photomultiplier tubes (PMTs)~\\cite{hamamatsu} to detect the scintillation light produced by Weakly Interacting Massive Particles (WIMPs) as they scatter off Xe nuclei. These compact metal-channel PMTs were specifically designed to work at 177 K, close to the temperature of LXe, and at a pressure up to 5 atm. They have a quartz window and a bialkali photocathode, for high sensitivity and low dark current in the UV regime \\citep{Hamamatsu_handbook}. Recently, a new version of the same PMT (R8520-406) with improved sensitivity at 178 nm, the scintillation wavelength of LXe, was produced by Hamamatsu and procured for an upgrade of the XENON100 experiment. The quantum efficiency (QE) is one of the most important characteristics of a PMT. The QE is defined as the ratio between the number of photoelectrons emitted from the photocathode and the number of incident photons. The energy absorbed from the incident photons is transferred to the valence band of the photocathode. Since the photoemission process has a probabilistic nature, not all of the electrons that absorb energy from the incident photons are emitted as photoelectrons~\\cite{Hamamatsu_handbook}. A quantity related to the QE is the radiant sensitivity ($\\mathrm{Sk}$), used to express the spectral characteristics and especially the relationship between the photocathode response and the incident light wavelength. The radiant sensitivity is defined as the photoelectric current generated by the photocathode, I$_\\mathrm{K}$, divided by the incident radiant flux at a given wavelength, L$_\\mathrm{P}$: \\begin{equation} \\mathrm{Sk} = \\frac{\\mathrm{I_K}}{\\mathrm{L_P}}\\\\ \\label{rad_sensitivity} \\end{equation} The relationship between the QE and the $\\mathrm{Sk}$ is given by~\\cite{Hamamatsu_handbook} \\begin{equation} \\mathrm{QE} = \\frac{h\\cdot c}{\\lambda \\cdot e}\\cdot \\mathrm{Sk} \\simeq \\frac{1240}{\\lambda}\\cdot \\mathrm{Sk}\\cdot 100 ~[\\%]\\\\ \\label{qe_sk} \\end{equation} where $h$ is Planck's constant, $c$ is the speed of light in vacuum, $\\lambda$ is the wavelength of the incident light in nm, $e$ is the electron charge, and $\\mathrm{Sk}$ is the radiant sensitivity in A/W. The radiant sensitivity is often provided by the manufacturer for each PMT. More generally, the blue sensitivity index (sometimes called cathode blue sensitivity), $\\mathrm{Sk_{b}}$, which is the photoelectric current generated from the photocathode with a blue filter interposed in the same setup, is specified. However, this value is usually measured only at room temperature. Hence, a quantitative study of the QE as a function of temperature is necessary for an improved understanding of LXe experiments. ", "conclusions": "\\label{} We have designed a system to measure the QE of R8520-406 PMTs, which have special bialkali photocathodes capable of providing a good performance at temperatures down to 160~K. Four PMTs were measured with this setup. Their absolute QEs at room temperature were measured by comparing their response to that of a calibrated photodiode, and the values obtained match those quoted by the manufacturer within the uncertainties. The QE of the PMTs was also measured for temperatures down to 160 K, the temperature of LXe at which these PMTs are usually operated. A relative increase at low temperatures of up to $\\sim$10\\% is measured for all four PMTs. This constitutes an important measurement for detectors based on LXe scintillation since it will help to more accurately determine the number of photons observed from the number of photoelectrons detected. The setup can be adopted to enable testing of other types of PMTs being considered for next generation experiments based on liquid noble elements. \\vspace{1cm}" }, "1207/1207.4434_arXiv.txt": { "abstract": "We investigate the level of fine-tuning of neutralino Dark Matter below \\hbox{200~GeV} in the low-energy phenomenological minimal supersymmetric Standard Model taking into account the newest results from XENON100 and the Large Hadron Collider as well as all other experimental bounds from collider physics and the cosmological abundance. We find that current and future direct Dark Matter searches significantly rule out a large area of the untuned parameter space, but solutions survive which do not increase the level of fine-tuning. As expected, the level of tuning tends to increase for lower cross-sections, but regions of resonant neutralino annihilation still allow for a band at light masses, where the fine-tuning stays small even below the current experimental limits for direct detection cross-sections. For positive values of the supersymmetric Higgs mass parameter $\\mu$ large portions of the allowed parameter space are excluded, but there still exist untuned solutions at higher neutralino masses which will essentially be ruled out if XENON1t does not observe a signal. For negative $\\mu$ untuned solutions are not much constrained by current limits of direct searches and, if the neutralino mass was found outside the resonance regions, a negative $\\mu$-term would be favored from a fine-tuning perspective. Light stau annihilation plays an important role to fulfill the relic density condition in certain neutralino mass regions. Finally we discuss, in addition to the amount of tuning for certain regions in the neutralino mass--direct detection cross-section plane, the parameter mapping distribution if the allowed model parameter space is chosen to be scanned homogeneously (randomized). ", "introduction": "\\indent Recently, the XENON100 collaboration has released new results after analyzing 225 live days of data taking. Limits on the spin-independent elastic Dark Matter-nucleon cross-section, $\\sigma^{\\rm SI}$, have been increased by a factor of roughly four with \\hbox{$2.0\\times10^{-9}$~pb} as the minimal value of the upper limit on $\\sigma^{\\rm SI}$ at a Dark Matter particle mass of \\hbox{55~GeV}~\\cite{xenon2012}. This leads to further tests for dark matter models. Furthermore, the ATLAS and CMS collaborations have presented their analysis of more than \\hbox{5 fb${}^{-1}$} of \\hbox{7~TeV}, also including \\hbox{8~TeV}, data and claimed close to 5 local sigma level the existence of a Higgs boson with a mass of approximately \\hbox{125~GeV}~\\cite{ATLASseminarHiggs,CMSseminarHiggs}. This fact fits very well to the minimal supersymmetric standard model (MSSM) because its prediction of the lightest Higgs boson mass is, when the LEP limit is taken into account, between \\hbox{115 - 135~GeV} depending on the supersymmetric parameters, see~\\hbox{\\it e.g.}~\\cite{Ellis:1990nz,Ellis:1991zd,Okada:1990vk,Haber:1990aw,% Drees:1991mx,Degrassi:2002fi,Buchmueller:2009fn}. The existence of Dark Matter is supported by various cosmological observations such as gravitational effects on visible matter in the infrared and gravitational lensing of background radiation. Its total abundance, that has important implications for the evolution of the Universe, has been precisely measured by the WMAP collaboration~\\cite{Komatsu:2010fb} during the last decade. This requires that a different kind of matter beyond the Standard Model (SM) of particle physics must be postulated. One of the most popular and most intensive studied candidate is the so-called weakly interacting massive particle (WIMP) that may constitute most of the matter in the Universe. Cosmology provides therefore a good motivation for Supersymmetry (SUSY), since the MSSM possesses a natural WIMP candidate as the lightest supersymmetric particle (LSP) is stable due to $R$-parity conservation~\\cite{Goldberg:1983nd,Ellis:1983ew} (for reviews see \\hbox{\\it e.g.}~\\cite{Jungman:1995df,Bergstrom:2000pn}). SUSY~(for reviews we refer to~\\cite{Nilles:1983ge,Haber:1984rc,Martin:1997ns}) has moreover the ability to solve the famous hierarchy problem by introducing superpartners with opposite spin statistics to each SM particle such that the loop contributions from superpartners cancel exactly and the weak scale is stabilized. Since SUSY must be broken, however, these cancellations are not exact and the non-discovery of SUSY particles pushes the breaking scale further up. This separation of the weak scale and of the SUSY breaking scale raises the question how easily this stability can be maintained. We apply therefore in this paper a measure of naturalness~\\cite{Ellis:1986yg,Barbieri:1987fn}, which was used for electroweak symmetry breaking, to the Dark Matter sector and study the level of fine-tuning. There is a series of studies on Dark Matter in the framework of simplified variants of the MSSM, the so-called constrained MSSM (CMSSM), which possess universal supersymmetry breaking mass parameters at the grand unification scale (for example~\\cite{Farina:2011bh,Buchmueller:2011sw,Buchmueller:2011ab}). Due to the existence of the grand unification condition on the gaugino masses in these models, there exists a LEP limit on the lightest neutralino mass: they must be heavier than \\hbox{$46$~GeV}. According to reference~\\cite{Buchmueller:2011ab} the lightest neutralino mass must be larger than about \\hbox{$200$~GeV} at \\hbox{$95$~\\% C.L.} after taking into account all relevant experimental constraints. Instead of this restricted class of models non-universal gaugino models within the framework of the phenomenological MSSM (pMSSM) (see for example~\\cite{AbdusSalam:2009qd,Sekmen:2011cz,Arbey:2011un,AlbornozVasquez:2012px,CahillRowley:2012rv}) have gained much attention. In the pMSSM low-energy input parameters are used with no high-energy relations between them. These models were used to explain the possible annual modulation signals of DAMA/LIBRA~\\cite{Bernabei:2010mq} and CoGeNT~\\cite{Aalseth:2010vx} (\\hbox{\\it e.g.}~\\cite{Hooper:2002nq,Bottino:2002ry,Dreiner:2009ic,Kuflik:2010ah,% Feldman:2010ke,Vasquez:2010ru,Fornengo:2010mk,Calibbi:2011ug,Arbey:2012na}), as well as the excess of nuclear recoil events reported by CRESST~\\cite{Angloher:2011uu}. This would be interpreted in terms of Dark Matter with a mass between roughly \\hbox{10~GeV} and \\hbox{30~GeV} and spin-independent cross-section of order $10^{-4}-10^{-7}$ pb. However, it was shown that light neutralino Dark Matter scenarios consistent with DAMA/LIBRA, CoGeNT and CRESST within the pMSSM are disfavored by LHC constraints~\\cite{Calibbi:2011un}. In addition there are discussions about the validity and natural consistency of these signals~\\cite{Kopp:2011yr} and XENON100~\\cite{xenon2012,Aprile:2011hi,Aprile:2010um} (see also~\\cite{Angle:2007uj}) as well as CDMS~\\cite{Ahmed:2010wy}, since these experiments have excluded these ``would be'' Dark Matter signals anyway. We assume therefore that Dark Matter has so far not been detected and ask how natural or fine-tuned the left-over parameter space is. Specifically we study in detail the not so well investigated neutralino mass range less than \\hbox{200~GeV}, taking into account all collider, cosmological and flavor constraints including the recent results of LHC Higgs researches as well as flavor studies. Over this complete mass region we find valid scenarios that may have escaped every experiment so far. We will especially show that it is possible to fulfill the muon anomalous magnetic moment condition for positive gaugino masses and a negative supersymmetric Higgs mass parameter, the $\\mu$-term. This article is organized as follows: in the next section, we define the fine-tuning measures and fix our notation of the neutralino sector. Then, in section~\\ref{sec:analysis} the method of our numerical analysis and the SUSY parameter space is discussed. Our results will be presented in section~\\ref{sec:results} including discussions about the annihilation mechanisms for neutralinos, the mapping of the level of fine-tuning into the direct detection cross-section plane, the direct detection cross-section and its dependence on the sign of the $\\mu$-term, how the muon anomalous magnetic moment can be obtained correctly with a negative $\\mu$-term, and lastly about functional fine-tuning and the parameter mapping distribution. Finally, we conclude in section~\\ref{sec:conclusions}. ", "conclusions": "\\label{sec:conclusions} \\indent In this paper we have analyzed the naturalness of neutralino Dark Matter in the non-universal gaugino model within the framework of the minimal supersymmetric extension of the Standard Model. We have taken into account all cosmological (upper and lower bound on the relic density), collider and flavor constraints including the results of XENON100 (2012), LHC data on the Higgs mass, ${\\rm Br}(B_s\\to\\mu^+\\mu^-)$ and pseudo-Higgs searches. Hereby the soft supersymmetry breaking terms are parameterized by 11 independent free parameters, that we have chosen such that the lightest supersymmetric particle is the lightest neutralino with a mass smaller than 200~GeV. We studied from the Dark Matter perspective how much fine-tuning is needed to reach a certain point in the $m_{\\widetilde{\\chi}} - \\sigma^{\\rm SI}$ plane. Therefore we use a parameter fine-tuning measure (see equation~(\\ref{eq:finetuning})) which was used before in order to study the naturalness of the Higgs -- SUSY breaking scale separation. We first presented in figure 1 the contribution to the Dark Matter abundance $\\Omega h^2$ for different annihilation mechanisms as a function of $m_{\\widetilde{\\chi}}$. Demanding that neutralinos provide the right amount of Dark Matter, we restrict the further scans to cases where $\\Omega h^2 \\in [0.089,0.136]$. We also investigated the dominant neutralino annihilation mechanisms (figure~\\ref{fig:relic}) and have shown their arrangement in the $m_{\\widetilde{\\chi}}-\\sigma^{\\rm SI}$-plane (figure~\\ref{fig:mechanism_xenonplane}). In case of a signal in a direct detection experiment the most important annihilation mechanism can be deduced and we would know which channel at the LHC is promising for production of neutralino Dark Matter. To avoid the limits of the direct detection results of XENON100 (2012), we showed that light stau annihilation of neutralinos in the early Universe play a special role, not only in the mass range of light neutralinos, $\\lesssim30$~GeV, but also between $ \\simeq [60, 80]$~GeV, which has been missed so far in other studies. It is important to differently parameterize the soft SUSY masses of the left- ($m_{{\\widetilde l}_L}$) and right-handed ($m_{{\\widetilde l}_R}$) sleptons. Note that we scanned the input parameter space in such a way that we obtain the fine-tuning for every point in the $m_{\\tilde\\chi} - \\sigma^{\\rm SI}$ plane which is accessible. This implies that the density of points has no meaning, but that the envelope implies that certain areas cannot be reached by any input parameter. With this method we showed in figure~\\ref{fig:pos_nog2} how the electroweak fine-tuning maps into the direct detection plane and found that a great part of untuned regions is already excluded by the current XENON100 (2012) limit when the supersymmetric Higgs mass parameter, the $\\mu$-term, is positive. A general trend for higher fine-tuning for smaller $\\sigma^{\\rm SI}$ is then visible. This trend can easily be understood, since the cross-section departs more and more from its natural value set by the generic scale. Thus, future direct detection experiments will push the amount of the electroweak fine-tuning further up. The only exception occurs for neutralino masses that allow for resonant annihilations, especially near the $Z$- and $h$-boson resonances. This last statement is also valid for a negative $\\mu$-term and the electroweak fine-tuning near the $Z$- and $h$-resonance stays small independent of the value of $\\sigma^{\\rm SI}$. Additionally, due to cancellations between contributions from light and heavy Higgs exchanges the direct detection cross-section gets shifted to smaller values, such that the XENON100 (2012) exclusion limit is fulfilled easily. A negative $\\mu$-term is therefore favored from a fine-tuning perspective (see figure~\\ref{fig:neg_nog2}). In our analytical study (see figure~\\ref{fig:explaining_sigma}) we have discussed the reason why the negative value of the $\\mu$-term allows for these cancellations in the spin-independent cross-section, and showed that the combination of input parameters $(M_1 + \\mu \\sin2\\beta)$ and the sign of the wino component are responsible for decreasing $\\sigma^{\\rm SI}$ to values that can be as low as $10^{-20}$ pb (see also figure~\\ref{fig:neg_nog2}). Since these cancellations might be viewed as an instance of tuning, we reevaluated the fine-tuning by adding a measure of ``equation-tuning'' and found that scenarios with $\\sigma^{\\rm SI} \\lesssim 10^{-15}$ pb are always unbearably tuned (figure~\\ref{fig:ft_sigma}). We find the scenario with the lowest fine-tuning, independent of the sign of $\\mu$, at $m_{\\widetilde\\chi} \\approx$ \\hbox{84~GeV} and $\\sigma^{\\rm SI} \\approx 2.0 \\times$ \\hbox{10$^{-9}$~pb} just below the new limit. Even for a negative $\\mu$-term we were able to get the correct positive pull for the muon anomalous magnetic moment, $a_\\mu$, to correctly deviate from the Standard Model (figure~\\ref{fig:explaining_g2}). This has been thought to be very difficult, but is possible due to the bino--higgsino--right-handed smuon loop which contributes to $a_\\mu$ dominantly when both gaugino masses are positive ($M_1>0$ and $M_2>0$) and $m_{{\\widetilde l}_L} \\gg m_{{\\widetilde l}_R}$. If the latter condition is fulfilled staus are generally light ($ \\lesssim 400$ GeV) and help to respect the cosmological abundance of Dark Matter by light stau annihilation in the complete mass region of the neutralino. It should be stressed that there is an easy way to satisfy $a_\\mu$, namely, if additional parameters for the smuon masses, \\hbox{\\it i.e.} $m_{{\\widetilde \\mu}_{L, R}}$, are introduced. In this case we can avoid the connection between the relic density and the anomalous muon magnetic moment and fulfill both conditions without any doubt. Therefore, the case of a negative sign of the $\\mu$-term is equally important and should be investigated more carefully in future studies. Note that the density of points is meaningless except in figure~\\ref{fig:probability}, since we do not assign a probability measure, but determine only the amount of tuning required to reach a certain point in the $m_{\\widetilde{\\chi}} - \\sigma^{\\rm SI}$ plane. The envelope implies, however, that these points cannot be reached. In this context it is interesting to note that the cross-section of the neutralino annihilating into two photons and into a pair of photon and $Z$-boson is loop suppressed and is therefore much smaller than the requirement from the claimed \\hbox{130~GeV} gamma-ray line in the Fermi-LAT data. Thus, our models can not explain this ``evidence''. Besides, very light neutralino scenarios consistent with the DAMA/LIBRA, CoGeNT, CRESST experiments cannot be explained in the pMSSM especially due to the limits on ${\\rm Br}(B_s \\rightarrow \\mu^{+} \\mu^{-})$ and on the pseudo-Higgs mass-$\\tan \\beta$-plane. Finally, in section~\\ref{sec:probability} we have discussed in addition the parameter mapping distribution of our models into the $m_{\\widetilde{\\chi}} - \\sigma^{\\rm SI}$ plane (figure~\\ref{fig:probability}). We found that the $Z$- and $h$-boson resonant areas become the preferred regions to detect neutralino dark matter if the $\\mu$-term is positive. For negative $\\mu$ another important region is formed by light stau annihilation, that has avoided direct searches so far. Note that taking into account the branching ratio of the decay $B_s \\to \\mu^+ \\mu^-$~\\cite{LHCb:2012ct} does not change our discussion and results because we are in the decoupling regime and $\\tan\\beta$ is not too high. There are only few models that do not satisfy the lower limit of ${\\rm Br}(B_s \\to \\mu^+ \\mu^-)$ at 95 \\% C.L. Furthermore, the strong bounds from LHC for light generation squarks and gluino masses~\\cite{Aad:2012fqa} do not affect the main conclusion of our discussions, since its contributions to the direct detection and pair-(co)annihilation cross-sections are typically subdominant. Note added: After the completion of this work, two papers appeared which have studied the importance of the $\\mu$-term sign for the direct detection cross-section within the frame work of MSSM~\\cite{Cheung:2012qy} and NMSSM~\\cite{Perelstein:2012qg}, respectively. In section~\\ref{sec:cross-section} we have discussed the suppression of $\\sigma^{\\rm SI}$ in the region (so-called ``blind spot'') where a particular combination of SUSY parameter $M_1 + \\mu \\sin 2\\beta$ is small." }, "1207/1207.1816_arXiv.txt": { "abstract": "We present the {\\it AKARI} near-infrared (NIR; 2.5--5 $\\mu$m) spectroscopic study of 36 (ultra)luminous infrared galaxies [(U)LIRGs] at $z=0.01-0.4$. We measure the NIR spectral features including the strengths of 3.3 $\\mu$m polycyclic aromatic hydrocarbon (PAH) emission and hydrogen recombination lines (Br$\\alpha$ and Br$\\beta$), optical depths at 3.1 and 3.4 $\\mu$m, and NIR continuum slope. These spectral features are used to identify optically elusive, buried AGN. We find that half of the (U)LIRGs optically classified as non-Seyferts show AGN signatures in their NIR spectra. Using a combined sample of (U)LIRGs with NIR spectra in the literature, we measure the contribution of buried AGN to the infrared luminosity from the SED-fitting to the {\\it IRAS} photometry. The contribution of these buried AGN to the infrared luminosity is 5--10\\%, smaller than the typical AGN contribution of (U)LIRGs including Seyfert galaxies (10--40\\%). We show that NIR continuum slopes correlate well with {\\it WISE} [3.4]--[4.6] colors, which would be useful for identifying a large number of buried AGN using the {\\it WISE} data. ", "introduction": "Since the {\\it Infrared Astronomical Satellite} \\citep[{\\it IRAS};][]{neu84} first opened the all-sky view of far-infrared universe, a large number of luminous infrared galaxies (LIRGs; 10$^{11} \\leqq L_{\\rm IR}$(8--1000 $\\mu$m) $<~$10$^{12}~L_{\\odot}$) and ultraluminous infrared galaxies (ULIRGs; $L_{\\rm IR} \\geqq$ 10$^{12}~L_{\\odot}$) have been identified and studied extensively (see \\citealt{san96,lon06,soi08} for review). In the local Universe, many of them are interacting systems between gas-rich disk galaxies [e.g., \\citealt{kim02,vei02,wan06,kav09,hwa10a}; but see \\citealt{elb07,lot08,ide09,kar10,kar11} for high-$z$ (U)LIRGs]. They may evolve into quasars and then into intermediate-mass elliptical galaxies \\citep[e.g.,][]{san88a,kor92,gen01,tac02,das06,vei09b,rot10,haa11}. Their enormous infrared luminosity comes from cool dust primarily heated by young stars [i.e., star formation (SF)] and hot dust heated by supermassive black holes (SMBHs) rapidly accreting matter [i.e., active galactic nuclei (AGN)]. Their contribution to the infrared luminosity density increases with redshift \\citep[e.g.,][]{lef05,mag09,got11}. Therefore, the study of (U)LIRGs allows us to better understand galaxy-galaxy interactions, starburst-AGN connection, and cosmic star formation history. To identify the energy sources of galaxies (i.e., SF vs. AGN), the optical line ratios sensitive to the photoionization source have been used \\citep[e.g.,][]{bal81,vei87,kew06}. Based on this method, the optical spectral types for a large number of (U)LIRGs are also determined \\citep[e.g.,][]{vei95,vei99a,kew01,got05,cao06,hou09,lee11}. However, the optical spectral classification can be uncertain, in particular for (U)LIRGs, due to the difficulty of detecting dust-enshrouded AGN. In this case, near-infrared (NIR) spectroscopy, less affected by dust extinction, is a more efficient tool. There are several ground-based $K$-band (1.9--2.4 $\\mu$m) spectroscopic surveys searching for obscured AGN signatures in (U)LIRGs: the presence of broad Pa$\\alpha$ emission centered at (rest-frame) 1.875 $\\mu$m or of high-excitation coronal line [\\ion{Si}{6}] at 1.963 $\\mu$m \\citep[e.g.,][]{vei97,vei99b,mur99,mur01}. The ground-based $L$-band (2.8--4.2 $\\mu$m) spectroscopy was also used to constrain their energy source \\citep[e.g.,][]{ima00,ima03,ima06,ima11,ris06,ris10,san08}. The SF-dominated (U)LIRGs show a strong emission line at 3.29 $\\mu$m attributed to polycyclic aromatic hydrocarbons (PAHs), whereas AGN-dominated (U)LIRGs show a relatively PAH-free continuum attributed to larger-sized hot dust grains \\citep[e.g.,][]{moo86,ima00}. The strong absorption features at 3.05 and 3.4 $\\mu$m by H$_{2}$O ice-covered dust grains and bare carbonaceous dust, respectively, are found in (U)LIRGs with buried AGN. However, these absorptions are weak or absent in normal SF (U)LIRGs where energy sources and dust are often spatially well mixed \\citep[e.g.,][]{ima00,ima03}. For a similar reason, NIR continua of AGN-dominated (U)LIRGs can be much redder than those of SF-dominated (U)LIRGs \\citep[e.g.,][]{ris06}. Thanks to the wide wavelength coverage (2.5--5 $\\mu$m) and high sensitivity of the Infrared Camera (IRC) on-board the {\\it AKARI} space telescope \\citep[][]{mur07,ona07}, the $L$-band diagnostic could be applied to more distant and faint galaxies \\citep[e.g.,][]{ima08,ima10a}. The AGN detection rate from $L$-band spectroscopy appears to be larger than that from $K$-band spectroscopy (roughly 70\\% vs. 20\\% in ULIRGs) because the $K$-band diagnostic detects only obvious AGN. These NIR spectroscopic studies of (U)LIRGs suggest that there are many optically elusive buried AGN and that the buried AGN fraction increases with increasing infrared luminosity. In this study, we analyze the {\\it AKARI} NIR spectra of 36 (U)LIRGs\\footnote{Although 5 out of 36 galaxies do not satisfy the definition of (U)LIRGs (i.e., $L_{\\rm IR} <$ 10$^{11}~L_{\\odot}$), they are simply referred as (U)LIRGs in this study. We do not include them when we compare our results with those in previous studies.} mainly from the cross-correlation between the {\\it IRAS} and Sloan Digital Sky Survey \\citep[SDSS;][]{yor00}. Combining our new NIR spectroscopic data with those in \\citet{ima08,ima10a}, we investigate the NIR properties of a large sample of (U)LIRGs. The structure of this paper is as follows. The target selection is explained in Section 2. Observations and data reduction are described in Section 3, and the method of NIR spectral analysis is given in Section 4. Our findings are discussed and summarized in Sections 5 and 6, respectively. Throughout this paper, we adopt flat $\\Lambda$CDM cosmological parameters: $H_{0}$ = 75 km s$^{-1}$ Mpc$^{-1}$, $\\Omega_{M}$ = 0.3, and $\\Omega_{\\Lambda}$ = 0.7. ", "conclusions": "\\subsection{AGN Diagnostics} \\subsubsection{AGN signature in the NIR spectra}\\label{agnsig} Emission at 3.3 $\\mu$m is a prominent feature in (U)LIRG NIR spectra. It probably originates from the reprocessing of ultraviolet (UV) radiation by PAH molecules. The contribution of Pf$\\delta$ emission line at a similar wavelength is usually negligible. If (U)LIRGs host AGN, the EW$_{\\rm 3.3 PAH}$ is suppressed because hot dust emission from the AGN increases the NIR continuum flux level and X-ray photons may destroy PAH molecules \\citep[e.g.,][]{smi07}. A red NIR continuum (i.e., a high value of continuum slope) as well as strong absorption features at 3.1 and 3.4 $\\mu$m from H$_{2}$O ice-covered dust grains inside molecular clouds and bare carbonaceous dust in the diffuse interstellar medium, respectively \\citep[see][]{dra03}, indicate the presence of highly obscured compact sources. While many of these absorbed sources with red continuum slopes have been shown to harbor buried AGN \\citep[e.g.,][]{ris06, san08, ima10a}, deeply buried starbursts can also produce the same spectral signatures \\citep[e.g.,][]{des07, spo07, vei09a} and in general all that is required is a warm, highly obscured heating source. The small PAH equivalent width condition is useful for identifying weakly obscured AGN, while the large continuum slope and optical depth conditions are efficient in detecting highly obscured AGN. Therefore, all of these diagnostics are necessary to detect as many obscured AGN as possible. Following the criteria in \\citet{ima10a}, we regard EW$_{\\rm 3.3 PAH} <$ 40 nm, $\\Gamma~(F_\\lambda \\varpropto \\lambda^{\\Gamma}) > -$1, $\\tau _{3.1} >$ 0.3, and $\\tau _{3.4} >$ 0.2 as AGN signatures, and classify sources satisfying at least one of these conditions as NIR AGN-detected galaxies (see column 10 in Table \\ref{table3}). In the results, we find 19 AGN out of 36 (U)LIRGs based on these criteria. Among these AGN, there are 13 sources with EW$_{\\rm 3.3 PAH} <$ 40 nm (68\\%), five sources with $\\Gamma > -$1 (26\\%), five sources with $\\tau _{3.1} >$ 0.3 (26\\%), and three sources with $\\tau _{3.4} >$ 0.2 (16\\%). Note that some sources satisfy more than one criterion. Therefore, EW$_{\\rm 3.3 PAH} <$ 40 nm is the primary criterion to select AGN. A similar trend is also seen in the Imanishi sample (56, 43, 45, and 10\\% for EW$_{\\rm 3.3 PAH}$, $\\Gamma$, $\\tau _{3.1}$, and $\\tau _{3.4}$, respectively). We count the number of sources with AGN signatures among our sample (U)LIRGs in bins of optical spectral type and of infrared luminosity. These results are summarized in Table \\ref{table4} together with those from the Imanishi sample. The NIR AGN detection rates for our sample and the Imanishi sample are on average 53\\% and 51\\%, respectively, and agree well in each bin. % \\citet{ima08,ima10a} found that the AGN signature in NIR spectra is more often detected in optical AGN-like and in more luminous infrared galaxies. Our sample appears to follow these trends, as shown in Figure \\ref{fig-agndet}. However, it is not conclusive with our data alone because of large uncertainties, in particular, for the NIR AGN detection rate depending on optical spectral type. In the combined sample of 180 (U)LIRGs with $L_{\\rm IR} \\geqq 10^{11} L_{\\odot}$, the NIR AGN detection rate depends on optical spectral type as follows: 36\\% for SF (U)LIRGs, 55\\% for composite (U)LIRGs, and 66\\% for Seyfert 2 (U)LIRGs. Note that most of previously classified LINERs are called composites in this study because we adopted the selection criteria of \\citet{kew06} rather than those of \\citet{vei87} that were used in \\citet{ima08,ima10a}. There are two Seyfert 1 (U)LIRGs in the combined sample, and both of them show AGN signatures in their NIR spectra. The total NIR AGN detection rates for LIRGs and ULIRGs are 29\\% and 65\\%, respectively. If we select (U)LIRGs without any priority to non-Seyfert galaxies, the NIR AGN detection rates for (U)LIRGs increase slightly [e.g., 30\\% for LIRGs and 70\\% for ULIRGs based on the (U)LIRG sample in \\citet{yua10}]. When we consider non-Seyfert galaxies with AGN signature in the NIR to be optically elusive buried AGN, the buried AGN fraction (i.e., the number ratio of buried AGN to non-Seyferts) increases with infrared luminosity: 24\\% for (U)LIRGs with $L_{\\rm IR} = 10^{11}$--$10^{12} L_{\\odot}$, 60\\% for (U)LIRGs with $L_{\\rm IR} = 10^{12}$--$10^{12.3} L_{\\odot}$, and 84\\% for (U)LIRGs with $L_{\\rm IR} = 10^{12.3}$--$10^{13} L_{\\odot}$. The higher AGN fraction in more luminous infrared galaxies has been reported in previous studies based on the data from not only {\\it AKARI} but also the {\\it Infrared Space Observatory} \\citep[{\\it ISO}; e.g.,][]{lut98,rig99,tra01} and {\\it Spitzer Space Telescope} \\citep[e.g.,][]{des07,ima09a,val09,vei09a,pet11}, suggesting an important role of AGN in increasing infrared luminosity. As expected, we find many non-Seyfert (U)LIRGs with AGN signatures in their NIR spectra [buried AGN; 49\\% (55/113)]. On the other hand, a substantial number of Seyfert (U)LIRGs do not show the AGN signatures in the NIR [33\\% (14/43)]. The apertures used in the optical spectroscopy (our sample: 3\\arcsec\\ diameter fiber; Imanishi sample: 2\\arcsec\\ wide slit and 2 kpc extraction) are smaller than those of NIR spectroscopy (covers the entire galaxy size). The line measurements in the optical spectra for our sample are aperture-corrected following the method in \\citet{hop03}, but those for the Imanishi sample are not. The small aperture used for the optical spectral classification may miss the extended emission associated with star formation, and hence be more sensitive to weak, central AGN. This helps to understand the existence of Seyferts without AGN signatures in the NIR. We find no difference in the redshift distribution between Seyferts with and without AGN signatures in the NIR spectra. The different aperture size between the optical and NIR spectroscopy seems partially responsible for the disagreement of spectral types. Regardless of the aperture effect, if the hot dust emission from AGN is very weak because of a tiny covering of dust around the AGN, such AGN are clearly visible in the optical spectrum. However, they would be not distinguishable from SF-dominated galaxies based on the NIR diagnostics \\citep[see][]{nar10}. In contrast, when AGN are really heavily obscured, both optical and NIR diagnostics are less powerful. Then the observations at other wavelengths are necessary to detect them (e.g., X-ray: \\citealt{bau10,ten10}; mid/far-infrared: \\citealt{far07,spo07,vei09a,hat10,elb11}; but see also \\citealt{elb10,hwa10b}; radio: \\citealt{saj08,ima10b}). In Figure \\ref{fig-bpt}, we present the optical AGN diagnostic diagram based on [O{\\sc iii}]$\\lambda5007$/H$\\beta$ and [N{\\sc ii}]$\\lambda6584$/H$\\alpha$ line ratios, and NIR AGN diagnostic diagram based on EW$_{\\rm 3.3 PAH}$ and continuum slope. In panel (a), the optical AGN (Seyfert$+$LINER) without the NIR AGN signature have low [O{\\sc iii}]/H$\\beta$ ratios compared to those with AGN signatures both in optical and in NIR spectra. On the other hand, in panel (b), the NIR properties of optical AGN without the NIR AGN signature are not significantly different from those of non-AGN both in optical and in NIR spectra. \\subsubsection{AGN contribution to the infrared luminosity}\\label{agncont} To measure the contribution of buried AGN to the infrared energy budget of (U)LIRGs, we use the infrared spectral energy distribution (SED) templates and fitting routine of \\citet[][]{mul11}, {\\small DECOMPIR}\\footnote{http://sites.google.com/site/decompir}. These templates consist of one AGN and five host-galaxy SEDs. For the host-galaxy SEDs, {\\it Spitzer} mid-infrared spectra of starburst galaxies are extrapolated to the far-infrared using {\\it IRAS} photometry. These host-galaxy SEDs are grouped into five categories, referred to as `SB1' through `SB5', in terms of their overall shape and relative strength of PAH features. For the AGN SED, the intrinsic SEDs of AGN-dominated sources are derived after subtracting suitable host-galaxy components from the observed SEDs, and these SEDs are averaged. The infrared SEDs of AGN show a large spread, mainly dependent on dust distribution around AGN (i.e., smooth vs. clumpy torus structures). However, the different AGN SEDs do not significantly change the resulting AGN contribution to the infrared luminosity in (U)LIRGs \\citep[see][]{mul11,poz12}. Based on this SED fitting with sparse photometric data points such as {\\it IRAS}, \\citet[][]{mul11} found that the intrinsic AGN luminosities measured are actually correlated with those from high-resolution mid-infrared observations of the AGN cores. We apply this routine to 69 (U)LIRGs with S/Ns $>$ 3 at all four {\\it IRAS} bands (12, 25, 60, and 100 $\\mu$m) in the combined sample, and fit their SEDs five times with AGN and one of host-galaxies by allowing renormalization of these two templates. We choose the best-fit solution with the lowest $\\chi^2$ value, computing the AGN contribution to the infrared (8--1000 $\\mu$m) luminosity from this template set. Figure \\ref{fig-sed} represents example SEDs with the best-fit AGN and host-galaxy templates. The AGN contribution in (U)LIRGs ranges from 0\\% to 69\\%, and is on average 6--8\\% in LIRGs and 11--19\\% in ULIRGs. Because the combined sample preferentially includes non-Seyferts, the AGN contribution in this study seems to be small compared to other studies (e.g., LIRGs: $\\sim$10\\% in \\citealt{pet11}; ULIRGs: $\\sim$20\\% in \\citealt{far03}; 35--40\\% in \\citealt{vei09a}; $\\sim$25\\% in \\citealt{nar10}; 15--20\\% in \\citealt{ris10}). Figure \\ref{fig-agncont} shows the correlation of the AGN contribution with the presence of AGN signature in the NIR, optical spectral type, and infrared luminosity. Not surprisingly, the AGN contribution is higher in (U)LIRGs with AGN signature in the NIR than in those without AGN signature. The AGN contribution clearly increases with increasing infrared luminosity \\citep[see also][]{don12}, similar to the trend of buried AGN fraction in Figure \\ref{fig-agndet} (b). The AGN contribution for Seyfert (U)LIRGs is slightly larger than those for SF and composite (U)LIRGs. \\subsubsection{AGN diagnostics with {\\it WISE} data}\\label{agnwise} Recently, the {\\it Wide-field Infrared Survey Explorer} \\citep[{\\it WISE};][]{wri10} opens up the opportunity to probe mid-infrared properties (3.4, 4.6, 12, and 22 $\\mu$m) for a large sample of galaxies with excellent sensitivity. We use the {\\it WISE} all-sky survey source catalog\\footnote{ http://wise2.ipac.caltech.edu/docs/release/allsky/} to identify the {\\it WISE} counterparts of the (U)LIRGs observed with {\\it AKARI} IRC and (U)LIRGs at 0.01 $< z <$ 0.4 in the SDSS DR7 \\citep{hwa10a} within 3\\arcsec . In Figure \\ref{fig-wise} (a), we compare the continuum slope $\\Gamma$ with {\\it WISE} [3.4]--[4.6] color \\citep[Vega magnitude system;][]{jar11} for the {\\it AKARI} (U)LIRGs. We overplot the expected {\\it WISE} [3.4]--[4.6] colors from simple power-law continuum models as a function of continuum slope (solid line). The continuum slope $\\Gamma$ and {\\it WISE} [3.4]--[4.6] color show a tight correlation (Spearman rank correlation coefficient = 0.82; the probability of obtaining the correlation by chance = 1.03$\\times10^{-30}$) with some offset and scatter around the expected relation. Most sources have small values of continuum slope compared to the expectation because the presented slopes (fitting range: 2.7--4.8 $\\mu$m) are affected by the spectrum at $<$3.4 $\\mu$m where the stellar population contribution is large (see \\citealt{lee10}). The scatter may come from the contamination by the 3.3 $\\mu$m PAH feature. For the galaxies at $z <$ 0.13, the presence of 3.3 $\\mu$m PAH emission makes [3.4]--[4.6] color bluer than the color without the PAH emission. On the other hand, if the galaxies are at 0.26 $< z <$ 0.55, the presence of 3.3 $\\mu$m PAH emission makes [3.4]--[4.6] color redder than the color without the PAH emission. The amount of change in colors depends on redshift and on the strength of PAH emission. From the experiment with the spectra of our sample, we find that the presence of PAH emission can change the [3.4]--[4.6] color by $\\pm$0.2 mag, consistent with the scatter in Figure 9 (a). We find that the AGN selection criterion based on the continuum slope (i.e., $\\Gamma$ $> -$1) in this study is roughly equivalent to that of [3.4]--[4.6] $> 0.8$ suggested by \\citet{ste12} \\citep[see also][]{ass10, jar11}. If (U)LIRGs with AGN signatures in the NIR spectra are regarded as genuine AGN (filled symbols), the {\\it WISE} color criterion selects AGN with 72\\% completeness (51 out of 71 genuine AGN satisfy the {\\it WISE} color criterion) and 76\\% reliability (51 out of 67 objects which satisfy the {\\it WISE} color criterion are genuine AGN). Figure \\ref{fig-wise} (b) shows the {\\it WISE} [3.4]--[4.6] colors versus {\\it IRAS} flux density ratios between 25 and 60 $\\mu$m (hereafter {\\it IRAS} 25--60 $\\mu$m colors) for the SDSS (U)LIRG sample. The {\\it IRAS} 25--60 $\\mu$m color is known to be associated with nuclei activity in infrared-luminous galaxies \\citep[e.g.,][]{deg85,san88b,nef92,vei09a,lee11}. AGN-dominated galaxies show warm {\\it IRAS} 25--60 $\\mu$m colors ($f_{25}/f_{60} \\geqq 0.2$), while SF-dominated galaxies show cool colors ($f_{25}/f_{60} < 0.2$). The composite and SF galaxies show similar distributions in this domain, but the AGN are significantly different. The median colors of AGN differ from those of non-AGN with significance levels of 6.8$\\sigma$ and 7.4$\\sigma$ in the {\\it IRAS} and {\\it WISE}, respectively. Interestingly, there are a substantial number of SF (U)LIRGs in the lower-right corner. It seems real even if we consider large uncertainties associated with the {\\it IRAS} colors. They have warm dust emission without hot dust emission, therefore appearing to be heavily obscured AGN. As a result, the {\\it WISE} [3.4]--[4.6] color is a good tracer of AGN-heated hot dust emission, but may not be sufficient to detect heavily obscured AGN. \\subsection{Comparison between Optical and Infrared Properties}\\label{comp} \\subsubsection{Star formation rate indicators}\\label{sfr} The total infrared continuum, 3.3 $\\mu$m PAH emission, and recombination lines including H$\\alpha$ and Br$\\alpha$ of galaxies are useful indicators of star formation rate (SFR) \\citep[][]{ken98}. In Figure \\ref{fig-sfr} (a--b), we compare these SFR indicators: infrared luminosity versus (a) H$\\alpha$ and (b) 3.3 $\\mu$m PAH luminosities. The H$\\alpha$ luminosities are extinction-corrected using the Balmer decrement and \\citet{cal00} extinction curve. The expected relationships are overplotted between these parameters (dotted lines; hereafter SF galaxy sequences): $L_{\\rm H\\alpha}/L_{\\rm IR}=10^{-2.25}$ from the empirical relation in \\citet{ken98} and $L_{\\rm 3.3 PAH}/L_{\\rm IR}=10^{-3}$ from the observations of \\citet{mou90} and \\citet{ima02}. These panels show that there are large offsets and scatters between the data and the expected relationships, even for SF (U)LIRGs without any AGN signature [pure SF (U)LIRGs; large filled symbols]. The offset from the SF galaxy sequence is larger in ULIRGs than in LIRGs, suggesting that H$\\alpha$ and 3.3 $\\mu$m PAH emissions are more depressed in ULIRGs. There could be several reasons for the strong depression of H$\\alpha$ and 3.3 $\\mu$m PAH emission in ULIRGs. (1) The amount of dust extinction could be systematically underestimated in ULIRGs. To check this effect, we plot the observed line ratios of Br$\\alpha$/H$\\alpha$ and H$\\alpha$/H$\\beta$ in panel (c). The pure SF (U)LIRGs appear to have larger Br$\\alpha$/H$\\alpha$ ratios by considering the extinction curve of \\citet{cal00} with a typical $R_{\\rm V}$ value of 4.05. This can imply the need of a larger $R_{\\rm V}$ value for the extinction correction in (U)LIRGs [e.g., $R_{\\rm V}$ = 5.01 denoted by the arrow]. Some studies also suggest the need of flatter/grayer extinction curves (i.e., large $R_{\\rm V}$ values) in ULIRGs, which may be attributed to supernovae-driven large-size dust grains (e.g., \\citealt{kaw11}; \\citealt{shi11}; see also \\citealt{boq12}). However, because of large scatter in the data and different aperture size between the optical and NIR spectroscopy, it should be checked with a more extensive data set in future studies. (2) Even for ULIRGs without any AGN signature in optical and NIR observations, there could still be hidden AGN that play a role in the relative depression of line emission. To remove contamination of hidden AGN, it is necessary to search for AGN at other wavelengths. However, it is expected that the contribution of buried AGN is not significant as discussed in Section \\ref{agncont}. (3) The line emission from ULIRGs may be intrinsically weak. Regardless of the presence of AGN, the star formation in ULIRGs produces large infrared continuum emission because a larger fraction of stellar UV photons is absorbed by dust inside star-forming regions under the strong radiation field in ULIRGs \\citep[][]{abe09}. The PAH emission in ULIRGs may be depressed in the sense that intense radiation fields do not produce photo-dissociation regions that are necessary for PAH emission, or lead to the destruction of PAH carriers \\citep[see][]{voi92,smi07,vei09a}. It is still unclear whether the ionization state of grains is actually related to infrared luminosity \\citep[e.g.,][]{des07,smi07,vei09a,ima10a,pet11}. \\subsubsection{Dust extinction as a function of infrared color}\\label{extinction} In Figure \\ref{fig-av}, we present the Br$\\alpha$/H$\\beta$ line ratio of (U)LIRGs as a function of {\\it IRAS} 25--60 $\\mu$m color. There is an anti-correlation between the two parameters with Spearman rank correlation coefficient = $-$0.67 and the probability of obtaining the correlation by chance = 0.02. This supports that warmer sources are less extinguished than cooler sources \\citep[see][]{vei99a,vei09a,kee05,ima08}. When using Balmer decrement values instead, such a correlation becomes weaker since the optical lines do not trace well the buried sources." }, "1207/1207.4034.txt": { "abstract": "We study the wave optics features of gravitational microlensing by a binary lens composed of a planet and a parent star. In this system, the source star near the caustic line produces a pair of images in which they can play the role of secondary sources for the observer. This optical system is similar to the Young double-slit experiment. The coherent wave fronts from a source on the lens plane can form diffraction pattern on the observer plane. This diffraction pattern has two modes from the close- and wide-pair images. From the observational point of view, we study the possibility of detecting this effect through the Square Kilometer Array (SKA) project in the resonance and high magnification channels of binary lensing. While the red giant sources do not seem satisfy the spatial coherency condition, during the caustic crossing, a small part of source traversing the caustic line can produce coherent pair images. Observations of wave optics effect in the longer wavelengths accompanied by optical observations of a microlensing event provide extra information from the parameter space of the planet. These observations can provide a new basis for study of exoplanets. ", "introduction": "Gravitational lensing is caused by the bending of light rays due to the gravitational effect of a foreground mass. Depending on the distribution of mass on the lens plane and on the relative distances of the lens and source from the observer, multiple images or distortion in the source shape can be formed. In the case of star-star lensing inside the Milky Way, the separation between images is less than few milliarcseconds and the images are unresolvable for the ground based telescopes. This type of gravitational lensing is termed gravitational microlensing. Einstein (1936) derived the gravitational lensing equation, but it was decades until the first gravitational lensing was observed in 1979. The source of this lensing was a quasar and observations were performed in radio frequencies \\cite{Walsh79}. A few years later Paczy\\'ski proposed studying the MACHO (Massive Astrophysical Compact Halo Objects) population in the Galactic halo by the method of gravitational microlensing \\cite{pa1986}. His suggestion was to observe stars in the Large and Small Magellanic Clouds, counting the number of microlensing events and measuring their transit times (Einstein crossing time). Based on this observation one can measure the contribution of MACHOs to the mass of Galactic halo. In addition to dark matter studies, another interesting astrophysical application of microlensing was suggested by Moa~\\&~Paczynski (1991), namely the use of gravitational microlensing to aid in the discovery of exoplanets. Microlensing effect due to single or multiple lenses have been studied mainly using geometric optics. An important study of wave optics features of gravitational lensing was undertaken by Ohanian (1983), who investigated the magnification of a radio point source when a galaxy acts as a gravitational lens. He showed that wave optics smoothes singular features of the light curve at the position of caustic lines. In another work, Jaroszy\\'ski and Paczy\\'ski (1995) studied the caustic crossing of Quasar Q2237+0305 by a galaxy composed of individual stars. By studying the diffraction images of this system, they could put limit on the size of the quasar. Wave optics observation of gravitational lensing inside the Milky Way also have astrophysical applications, for example studying the limb darkening of small sources like white dwarfs \\cite{zabel diff}. Recently, Heyl (2010,2011a,2011b) discussed the possibility of detecting wave optics signals in microlensing light curves with a single substellar lens. In this work our aim is to extend the application of the wave optics to the conventional method of extra solar planet detection by gravitational microlensing. Here we assume a binary lens composed of a lensing star and a planet. The crossing of the caustic lines of this system by the source star produces high magnification in the light curve. Moreover, owing to the small separation of the images on the lens plane, the gravitational lensing system resembles a multiple slit optical system in the astronomical scales. With a coherent condition for the wave fronts on the lens plane, the result would be a diffraction pattern on the observer plane. We study the applications of this method in both the resonance and high magnification channels of the exoplanet detection. Observations of the contrast in the fringes and transit time of the fringes enable us to break degeneracy between the lens parameters. We also study the possibility of observing the wave optics features of binary microlensing using the future Square Kilometer Array (SKA) project. In section \\ref{Ampl1}, we introduce wave optics formalism in gravitational lensing and calculate the wave optics light curve for a binary lens system. In section \\ref{near} we carry out semi-analytic calculations of the wave optics feature for microlensing near the caustic lines and study the temporal and spatial coherency conditions. We also numerically compute wave optics light curves and compare them with the results of geometric optics. In section \\ref{det} we discuss the possibility of detecting microlensing wave optics signals by a binary lens in which one of the lenses is a planet. Our study uses observation in radio or micrometer wavelengths and future observations with SKA. We also discuss possibility of degeneracy breaking between the lens parameters in the resonance and high-magnification channels of exoplanet detection. Conclusion and a summary are given in section \\ref{conc}. ", "conclusions": "\\label{conc} In this work we studied the wave optics effect of gravitational microlensing by a binary lens composed of a lens star and a companion planet. The lensing effect of a planet during the caustic crossing produces close images, a suitable configuration for the images on the lens plane to generate diffraction pattern on the observer plane. This effect is an example of Young's double-slit experiment in the astronomical scales. We derived the wave optics features of a binary lens, showing that it depends only on the third-order derivatives of the Fermat potential $\\phi_{222}$ and $f = 2k R_s$. We take red giants and super-giants as the source stars of gravitational microlensing toward the Galactic bulge, as there is a natural selection bias for observing this type of source stars. We suggested using SKA future project for the observation of the wave optics signals in the light curve. In this observational program, radio observatories accompany the microlensing follow-up telescopes in the visual bands. These two observations at long and short wave lengths can provide a complimentary program for studying binary microlensing events to break the degeneracy of binary systems. We discussed the problem of spatial coherency of sources in binary lensing and showed that only the part of the source that crosses the caustic line contributes to the formation of close-pair images. While an extended giant star may have no spatial coherency, the spatial coherency condition holds for the small part of the source crossing the caustic line. We studied the observability of the wave optics parameters in a Monte Carlo simulation by fitting the simulated microlensing light curves with the theoretical wave optics light curve. Out of the two channels for the detection of exoplanets by microlensing, namly (i) the high-magnification channel and (ii) the resonance channel, we showed that the wave optics observation is in favour of resonance binary microlensing events. The extra information from the wave optics light curve enable us to solve for the lens parameters with better accuracy. Our analysis has shown that the use of radio telescopes for observations of planetary microlensing events will open a new window for studies of exoplanet ." }, "1207/1207.3678_arXiv.txt": { "abstract": "I present recent high-resolution submillimeter and millimeter observations of molecular gas and dust in some mergers, luminous galaxy nuclei, and possible mergers. Such observations tell us the behavior and properties of interstellar medium in merger nuclei. For example, the gas sometimes makes a mini disk around the remnant nucleus, feeds starburst and/or a massive black hole there, hides such a power source(s) by enveloping it, and is blown out by the embedded power source. Even when the power source is completely enveloped and hidden we can still constrain its physical parameters and nature from high-resolution (sub)millimeter observations. The observables include gas motion such as rotation (hence dynamical mass) and inflow/outflow, luminosity and luminosity density of the embedded nucleus, and mass, temperature, density, chemical composition, and (sometimes unusual) excitation conditions of gas. ", "introduction": "Among various ways to observationally study mergers, high-resolution observations at submillimeter and millimeter wavelengths can shed unique light on the inner working and evolution of mergers. This is partly because the above-mentioned observations provide detailed spatial and kinematical information of molecular gas that is the dominant interstellar medium in the inner regions of mergers. It is also because star formation and active nuclei, the main energy-generating mechanisms in mergers, rely on the interstellar molecular gas for raw material or fuel. We can examine this important ISM at the very region where star formation is most active or a massive black hole is being fueled, owing to recent progress in the high-resolution observing capabilities at the (sub)millimeter wavelengths. I show below the power of high-resolution (sub)millimeter observations in three areas related to mergers. First, compact luminous nuclei are often, though not exclusively, seen in major mergers. Submillimeter imaging of these nuclei at subarcsecond resolution can reveal hidden properties of young AGNs or starbursts in these nuclei. Examples of such nuclei in Arp 220 and NGC 4418 are presented. Second, outflow of molecular gas from mergers can be found thorough high-resolution observations of (sub)millimeter molecular lines. Such outflows can be driven by extreme starburst or AGN in mergers and likely affect the evolution of these activities. Examples of such outflows are presented for Arp 220 and NGC 3256. Third, subtle effects of a minor galaxy merger may be detectable with high resolution observations of molecular gas. Observations of NGC 4418 and M83 are addressed as possible cases where minor merger may have played a role. ", "conclusions": "I have shown examples of high-resolution (sub)millimeter observations of mergers and possible mergers and illustrated that such observations can uncover the effect of galaxy merging to the dynamics of the interstellar medium and the energy-generating activities in the system. The latter effect is caused by the fueling to the activities through gas concentration and also by quenching of the activities through gas dispersal via outflows. There is no doubt that similar observations of mergers are going to lead us to deeper understanding of galaxy evolution thorough mergers. The arrival of ALMA is particularly encouraging for the future of such observational studies." }, "1207/1207.7217_arXiv.txt": { "abstract": "The present work purports to identify candidate carriers of the UIBs. This requires a procedure for the computation of the emission spectrum of any given candidate. The procedure used here consists in exciting the carrier into a state of internal vibration, waiting until the system has reached dynamic equilibrium and, then, monitoring the time variations of the overall electric dipole moment associated with this vibration. The emission spectrum is shown to be simply related to the FT of these variations. This procedure was applied to more than 100 different chemical structures, inspired by the exhaustive experimental and theoretical analyses of Kerogens, the terrestrial sedimentary matter, which is known to be mainly composed of C, H, O, N and S atoms. From this data base, 21 structures were extracted, which fall in 4 classes, each of which contributes preferentially to one of the main UIBs. Summing their adequately weighted spectra delivers an emission spectrum which indeed exhibits the main UIB features (allowing for computational errors inherent in the chemical simulation code). By changing the weights, it is possible to change the relative band intensities so as to mimic the corresponding (moderate) changes observed in the sky. The defects of the present simulation are discussed, and directions for improvement are explored. ", "introduction": "The quest for a realistic model of the structure, composition and excitation of the carriers of UIBs (Unidentified Infrared Bands) has been going on for more than three decades now. An essential requirement is, for such a model, to correctly reproduce the spectral features of UIB emission, which, by now, are very well documented in a large variety of astronomical sites. Among the latest publications, note the exhaustive analysis of a rich collection of spectra from star-forming regions (Smith et al. 2007). Despite some variations from object to object, there are enough similarities that it has been possible to partition most of the observed spectra into a small number of classes (see, for instance, Peeters et al. 2002, 2004a; van Diedenhoven et al. 2004), thus allowing systematic comparisons to be made with laboratory or theoretical spectra. Most present day candidate model carriers are big molecules of widely different sizes (see Kwok and Sandford 2008, Tielens 2008, Kwok 2009), the size being limited, of course, by experimental or computational constraints,or even by theoretical considerations. While the experimental production and spectral analysis of single macromolecules have proven to be daunting, modern computing capabilities make it easier to use theory or chemical simulation for modeling purposes. The most usual computational way of determining the IR (InfraRed) spectrum of a given structure (molecule or grain) is to perform a Normal Mode Analysis (NMA; see Wilson et al. 1955). This delivers all the different modes of atomic vibration of infinitesimal amplitude, characteristic of the structure. For a grain made up of N chemically bound atoms, there are 3N-6 such modes, each characterized by its frequency and the directions and velocities of the excursion of each atom from its equilibrium position along the 3 space coordinates. Further treatment yields the mode \\emph {ir intensity} (or \\emph{line intensity}, or \\emph{integrated band intensity}), which is proportional to the \\emph {absorbance} at the corresponding frequency, the quantity commonly measured in the laboratory. NMA can be performed using chemical simulation codes based on different theories and approximations: Molecular Mechanics (MM+, Ambers, BIO+, etc.) and quantum chemical codes: Semi-empirical (AM1, PM3, etc.), \\emph{ab initio} (e.g. DFT) methods, which differ widely in accuracy and time consumption. These are available with commercial packages such as HyperChem, Spartan, Gaussian, etc. As early as 1992, Brenner and Barker \\cite{bre} warned that ``unambiguous identification of the IS (InterStellar) emitter cannot be established just on the basis of matching the emission frequencies to laboratory absorption spectra\". For emission spectra differ in many respects from absorption spectra: the features of the former are generally shifted and broadened relative to those of the latter, because of anhormonicity effects, which set in as soon as the grain is excited; correlatively, they depend on the excitation process (photon absorption, chemical reaction, etc.). Brenner and Barker themselves computed the emission of excited benzene and naphtalene based on a) the thermal equilibrium assumption (statistical distribution of energy among molecular states), b) Einstein's spontaneous emission coefficients, as deduced from measured (relative) absorption coefficients at a few characteristic spectral frequencies of the two molecules (see also Allamandola et al. 1989). Intrinsic bandwidth and underlying emission continuum were assumed. Several off-springs of this procedure have been used (e.g. Verstraete et al. 2001, Boersma et al. 2010, 2011); the state of the art in this domain is clearly summarized by Bauschlicher et al. \\cite{bau}. Addressing the same issue, Papoular \\cite{pap01}, \\cite{pap12b} proposed a distinct procedure which bypasses NMA altogether. Based on the unversal relation between emission, electric dipole moment and frequency, he showed that the emission spectrum of an excited molecule is simply related to the energy content and overall dipole moment of the structure. The proposed procedure hence consists in monitoring the variations in time of the energy content and overall dipole moment of a selected model structure after it has been excited. The FT (Fourier Transform) of these variations deliver the emission spectrum. As the specific anharmonicity of the molecule, and all its effects, are reflected in its internal motions, it also affects both the energy and dipolar spectra, and, hence, the emission spectrum. As a consequence, the \\emph{absorption line} spectrum of NMA is replaced by an \\emph{emission band} spectrum. No assumptions are added to shift or broaden the bands artificially. Plateaus arise naturally from overlapping band wings. If the excitation energy is high enough, overtone and combination bands automatically emerge from the analysis. The anharmonicity is generally built in the dynamics code used for chemical simulation. In semi-empirical codes, it is generated by a system of parameters tailored by analogy with measured properties of structures made of the particular atomic species and bonds of interest. The PM3 code used here is best fit to simulate hydrocarbon molecules. Semi-empirical and Molecular Mechanics commercial codes were used to demonstrate relaxation after excitation (see Papoular 2001 and 2002), as well as anharmonicity effects: line shift, mode-locking, combination frequencies, Fermi-resonance (see Papoular 2006) and line broadening (Papoular 2012). Both DFT (Density Functional Theory) and Semi-empirical quantum chemistry methods deliver accurate Normal Mode intensities and frequencies. However, DFT methods are basically designed to deal with \\emph{equilibrium states}, so they are not directly applicable to the analysis of dipole moment variations. To my knowledge, even Time-Dependent DFT (TDDFT) has been applied only to adiabatic transitions between equilibrium states, not to \\emph{dynamics}. Although, the procedure described above provides a sound theoretical transition from NMA to emission spectra and yields a degree of line broadening, this does not allow to mimic the wide mid-IR bands observed in the sky, unless very large aggregates of atoms are considered, or the emission spectra of a very large number of different medium-sized molecules are summed up. The former option is excessively computer-time consuming, and is not pursued further here. As for the latter, it is preferably started with a survey of the Normal Mode spectra (easier to compute!) of structures built up with the most common chemically active atoms: C, H, O, N and S (CHONS for short). Even with this restriction, the number of possible combinations is daunting. Papoular \\cite{pap10} therefore sought cues in kerogens and coals, whose spectra have features in common with the UIBs, and which have been abundantly modeled using laboratory analysis and chemical codes (e.g. Durand 1980, Behar and Vandenbroucke 1986, Carlson 1992, Speight 1994). It turned out that some types of structures have spectral lines falling mostly within one or two UIBs. Changing slightly the structure or composition slightly displaces the spectral lines, thus helping to build a model band. More than a hundred normal mode spectra were computed and distributed over eight families of structures, each contributing mainly to one or two spectral regions. By combining these spectra in the right proportions, it was possible to synthesize spectra bearing some resemblance with UIB spectra. These computational experiments led to the selection of a number of structural types exhibiting spectral features which appeared to be promising for our present purposes. In the present work, the \\emph{emission} spectra of many more structures of these types were computed according to the procedure outlined above. Each was given a weight and their weighted sum compared with UIB spectra. The details and outcome of this endeavour are given below. Section 2 lays down the theoretical basis of the proposed method of spectral emission calculation; it is complemented by Appendix A. The details of its implementation in the computational procedure are given in Sec. 3; Sec. 4 illustrates the 4 types of chemical structures that were included, by displaying one of each type, all the others being collected in Appendix B; Sec. 5 assembles the corresponding spectra in 4 graphs, one for each type, and compares the resulting overall emission spectrum with an observed galactic spectrum. Section 6 discusses the defects of this simulation and Sec. 7 explores possible improvements. ", "conclusions": "Clearly, more work is needed to obtain a satisfactory fit of model to observed spectra. The straightforward part might consist in including many more structures similar to those of the 4 mentioned classes. Exploring larger structures would help, as diversity increases with size, but this is more challenging. Also, that would call for more accurate and efficient codes and machines. Studying the transition from molecules to grains, however, cannot be avoided if important issues are to be addressed, such as the emission continuum, the graphitization at the origin of the 2175 \\AA{\\ } extinction band and the evolution from the aromatic 3.3-$\\mu$m to the aliphatic 3.4-$\\mu$m band. Even in the realm of medium-sized molecules, it should be possible to consider more tightly knitted structures than those studied here, which are essentially 2D. By contrast, taking on ionized molecules at present seems very laborious and uncertain with available simulations codes for molecular dynamics. From the strictly chemical point of view, it would be interesting to understand why and under what circumstances, structures similar to those described in this work would come to be privileged. Laboratory experiment could help advance these issues. In particular, the IR emission of a gaseous assembly of medium-sized molecules could be collected and spectrally analyzed at different pressures and temperatures: this may provide both a test of emission theories and an understanding of molecule-to-grain transition." }, "1207/1207.1431_arXiv.txt": { "abstract": "Over the last years both cosmic-ray antiproton measurements and direct dark matter searches have proved particularly effective in constraining the nature of dark matter candidates. The present work focusses on these two types of constraints in a minimal framework which features a Majorana fermion as the dark matter particle and a scalar that mediates the coupling to quarks. Considering a wide range of coupling schemes, we derive antiproton and direct detection constraints using the latest data and paying close attention to astrophysical and nuclear uncertainties. Both signals are strongly enhanced in the presence of degenerate dark matter and scalar masses, but we show that the effect is especially dramatic in direct detection. Accordingly, the latest direct detection limits take the lead over antiprotons. We find that antiproton and direct detection data set stringent lower limits on the mass splitting, reaching 19\\% at a 300 GeV dark matter mass for a unity coupling. Interestingly, these limits are orthogonal to ongoing collider searches at the Large Hadron Collider, making it feasible to close in on degenerate dark matter scenarios within the next years. ", "introduction": "\\par Complementarity between different signals has always been highlighted as the preferred way forward in the search for weakly interacting massive particles (WIMPs) -- see \\cite{BertoneReview,BergstromReview,BertoneBook,BergstromMulti} for general-purpose reviews. However, WIMP searches were for a long time an extremely data-starved field of research with experiments lagging far behind theoretical predictions. That is no longer the case. With large amounts of data pouring in from direct, indirect and collider dark matter (DM) searches, complementarity studies are now possible at an unprecedented level of detail \\cite{Zheng:2010js,Bertone:2011nj,Bertone:2011pq,Strege:2011pk,Yu:2011by}. Direct searches, in particular, have developed tremendously over the last decade benefiting from the use of various targets and techniques to measure WIMP-induced nuclear recoils. Presently, the situation is rather blurry: while experiments such as DAMA/LIBRA \\cite{DAMA2008,DAMA2008b,DAMA2010}, CoGeNT \\cite{cogent,cogentannualmod} and CRESST \\cite{CRESST2011} report excess of nuclear recoil events and also evidence for annual modulation in the case of DAMA/LIBRA and CoGeNT, other collaborations -- including XENON10/100 \\cite{Xenon10,XENONSD,Xenon100,Xenon1002,Xenon100_2012}, CDMS \\cite{CDMS2009,cdms10,Ahmed:2012vq}, SIMPLE \\cite{SIMPLE10a,SIMPLE10b,SIMPLE11} or COUPP \\cite{COUPP11,COUPP12} -- find null results. Leaving aside this controversy, that has been discussed at length in the literature \\cite{ChangIso,HooperCoGeNT,FengIso,Arina2011,Schwetz:2011xm,Kopp:2011yr}, it is clear that the whole array of different targets used boasts huge sensitivities to both spin-dependent (SD) and spin-independent (SI) WIMP-nucleus scattering. Together with direct detection, indirect searches via antiprotons have provided valuable hints in shaping the phenomenology of viable WIMP models. In fact, the exquisite data on the cosmic-ray antiproton flux collected by experiments such as BESS \\cite{Orito:1999re,Abe:2008sh,BESSPolarII} or PAMELA \\cite{Adriani:2010rc} fall nicely on top of the expectations from cosmic-ray spallations in the Galaxy. This means in practice that dark matter annihilations or decays cannot yield copious fluxes of antiprotons, and correspondingly the coupling to quarks is much constrained \\cite{Donato:2008jk,Kappl:2011jw}. Notice that this is precisely the coupling that drives WIMP-nucleus scattering in underground detectors, which makes antiprotons and direct searches particularly suitable for complementarity studies. \\par Several complications arise when exploring the complementarity between two or more observables. Firstly, if we are to draw sound conclusions, all relevant uncertainties must be treated carefully. In the case at hand, the unknowns are sizable and different in nature: on the one hand, antiproton searches are prone to uncertainties on the galactic dark matter profile \\cite{Navarro:2003ew,Gao:2007gh,Diemand:2008in,Navarro:2008kc} and on cosmic-ray propagation \\cite{SM98,Maurin:2001sj,Delahaye07,dragon2,Putze1,Trotta:2010mx,Evoli:2011id}; on the other hand, direct searches suffer from the lack of knowledge on the local dark matter density \\cite{CaldwellOstriker,Gates1995,UllioBuckley,BelliFornengo,CatenaUllio,Weber:2009pt,SaluccilocalDM,paperDMlocal,papermuLens,Bovy:2012tw} and local velocity distribution \\cite{Smith:2006ym,Xue2008,Sofue2008,Reid2009,Bovy2009,McMillan2009,Lisanti:2010qx} as well as from nuclear uncertainties \\cite{Pumplin:2002vw,EllisDD}. Secondly, combining signals such as antiprotons and direct detection requires the specification of an underlying particle physics framework. For instance, the author of references \\cite{Lavalle:2010yw,Lavalle:2011fm} points out the usefulness of antiproton measurements to constrain light WIMPs able to accommodate the hints of signal in DAMA/LIBRA and CoGeNT. This claim is, though, dependent on how dark matter couples to quarks, especially the light quarks abundant in nucleons that get struck in direct detection experiments. A definitive conclusion is therefore necessarily model-dependent. This highlights the difficulty in pursuing complementarity studies and deriving conclusive results. \\par Now, model building imposes hardly any boundaries on the phenomenology of dark matter candidates. Nevertheless, one can learn a lot by focussing on relatively simple realisations. For example, it has been recently emphasized \\cite{Hisano:2011um,Garny:2011cj,Garny:2011ii,Asano:2011ik} that the presence of mediator particles degenerate in mass with the WIMP leads to enhanced signals in both antiprotons and direct detection. In the case of the former, degeneracy boosts the contribution of $2\\to3$ processes to the annihilation yields \\cite{Bergstrom:1989jr,Flores:1989ru,Drees:1993bh,Garny:2011cj,Garny:2011ii,Asano:2011ik}, while in the latter it induces almost-resonant scattering in SD and SI interactions \\cite{Hisano:2011um}. Besides, all this proceeds at mass degeneracies below the trigger of searches at the Large Hadron Collider (LHC). To the best of our knowledge, the complementarity between antiprotons and direct searches has not been conveniently explored in this important case of degenerate mass states. The present work is precisely devoted to fill that gap, and is a first step in exploring realistically the interplay between direct and indirect searches within mass-degenerate WIMP frameworks. In particular, we show explicitly how data already at hand effectively constrain the presence of degenerate mass states, and how a reasonable experimental effort over the coming years will be able to close in on these scenarios. As we shall see, collider searches are orthogonal to antiproton and direct detection, and hence a multi-disciplinary approach is highly desirable for the near future of dark matter searches. \\par We start by introducing the main ingredients of our particle physics framework in Section \\ref{sec:model}. Then, Sections \\ref{sec:Indirect} and \\ref{sec:DD} are devoted to the formalism behind antiproton and direct detection constraints, respectively. Taking special care in modelling all relevant uncertainties, we derive in Section \\ref{sec:res} the limits imposed by the latest SD and SI direct detection searches as well as cosmic-ray antiproton data, before concluding in Section \\ref{sec:conc}. ", "conclusions": "\\label{sec:conc} \\par We have shown in this work that the latest data on cosmic antiprotons and direct detection place useful constraints on the phenomenology of mass-degenerate dark matter scenarios. In particular, we have considered a minimal framework featuring a Majorana fermion as dark matter that couples to light quarks via a scalar close in mass, encompassing e.g.~a simplified model with bino-like neutralino and squark as lightest and next-to-lightest supersymmetric particles. This setup allows for a direct comparison of scattering rates on nuclei with dark matter annihilation rates in our Galaxy, the dominant channel being $\\chi\\chi\\to q\\bar q g$ for quarks of the first and second generation. The derived constraints on coupling to quarks suffer from sizeable astrophysical and nuclear uncertainties, but it is nevertheless clear that antiprotons lag significantly behind direct detection, a fact that can be attributed mainly to the extreme sensitivity of underground searches to mass degeneracy. Fine degeneracies are conclusively discarded by current direct detection data. This is precisely the range that escapes detection at collider searches. Accordingly, we find that the interplay between antiprotons, direct and collider searches will be of crucial importance in closing in on simple mass-degenerate dark matter models over the coming years. Further work is needed to study the implications of this complementarity in the framework of more complicated particle physics realisations. \\vspace{0.5cm} {\\it Acknowledgements:} The authors thank Junji Hisano for helpful discussions and Jose Manuel Alarc\\'on, Germano Nardini and Pat Scott for useful comments. This work has been partially supported by the DFG cluster of excellence ``Origin and Structure of the Universe'' and by the DFG Collaborative Research Center 676 ``Particles, Strings and the Early Universe''. S.V.~also acknowledges support from the DFG Graduiertenkolleg ``Particle Physics at the Energy Frontier of New Phenomena''." }, "1207/1207.4352_arXiv.txt": { "abstract": "Supplemental material. Contains expanded figures. ", "introduction": " ", "conclusions": "" }, "1207/1207.3744_arXiv.txt": { "abstract": "A new method for solving the relativistic inverse stellar structure problem is presented. This method determines a spectral representation of the unknown high density portion of the stellar equation of state from a knowledge of the total masses $M$ and radii $R$ of the stars. Spectral representations of the equation of state are very efficient, generally requiring only a few spectral parameters to achieve good accuracy. This new method is able, therefore, to determine the high density equation of state quite accurately from only a few accurately measured $[M,R]$ data points. This method is tested here by determining the equations of state from mock $[M,R]$ data computed from tabulated ``realistic'' neutron-star equations of state. The spectral equations of state obtained from these mock data are shown to agree on average with the originals to within a few percent (over the entire high density range of the neutron-star interior) using only two $[M,R]$ data points. Higher accuracies are achieved when more data are used. The accuracies of the equations of state determined in these examples are shown to be nearly optimal, in the sense that their errors are comparable to the errors of the best-fit spectral representations of these realistic equations of state. ", "introduction": "\\label{s:introduction} The standard stellar structure problem consists of determining the structure of a star by solving the coupled gravitational and hydrodynamic equations with an assumed equation of state for the stellar matter. The solutions to the standard problem determine the various observable properties of the stars with a given equation of state like their total masss $M$, their total radii $R$, etc. The inverse stellar structure problem determines what equation of state is required to produce stellar models having a given set of macroscopic obserables. The goal of this paper is to find efficient and robust methods of solving this inverse stellar structure problem. The method developed here is based on the use of spectral expansions to represent the equation of state. The values of the spectral coefficients in these expansions are fixed in this method by matching stellar models based on these equations of state to observed values of the masses and radii of the stars. Once fixed, these coefficients determine the equation of state that represents the (approximate) solution to the inverse stellar structure problem. For non-rotating stars in general relativity theory, the simplest version of the stellar structure equations were first derived by Oppenheimer and Volkoff~\\cite{Oppenheimer1939}, \\begin{eqnarray} \\frac{dm}{dr} &=& 4\\pi r^2\\epsilon, \\label{e:OVm}\\\\ \\frac{dp}{dr} &=& - (\\epsilon+p)\\frac{m+4\\pi r^3 p}{r(r-2m)}, \\label{e:OVp} \\end{eqnarray} where $m(r)$ represents the mass contained within a sphere of radius $r$; $p(r)$ is the pressure; and $\\epsilon(p)$ is the total energy density of the fluid. Solving these equations with a given equation of state is the standard relativistic stellar structure problem. Once an equation of state $\\epsilon=\\epsilon(p)$ is specified, these equations determine a one parameter family of stellar models, $m=m(r,p_c)$ and $p=p(r,p_c)$, where $p_c$ is the value of the pressure at the center of the star $r=0$. These solutions then determine various macroscopic properties of the stars, including their outer radii $R(p_c)$ where $p[R(p_c),p_c]=0$, and their total masses $M(p_c)=m[R(p_c),p_c]$. These macroscopic properties are (at least in principle) observable. The standard stellar structure problem can be thought of as a map that takes the equation of state [a curve in energy density -- pressure space $\\epsilon=\\epsilon(p)$], into a curve in the space of macroscopic observables, e.g. $[M(p_c),R(p_c)]$. The inverse stellar structure problem consists of finding the inverse of this map~\\cite{Lindblom1992}, i.e. determining the equation of state of the stellar matter from a knowledge of some information about the macroscopic structures of the stars (like their masses $M$ and radii $R$). The solution to this problem, like the solutions to many inverse problems, is less straightforward than the solution to the standard problem. The inverse stellar structure problem is probably more relevant for practical relativistic astrophysics, however, than the standard problem. The highest density part of the equation of state in neutron stars, for example, is not well known. Matter in this state is well beyond the reach of laboratory experiments, so there is no independent way of directly measuring its properties, including its equation of state. Numerous attempts have been made to model this matter theoretically, but even today there is no consensus among theoreticians. Predictions of the energy density for a typical neutron-star central pressure, for example, often differ by an order of magnitude. Therefore, the standard stellar structure problem for neutron stars is not terribly useful. In contrast, the inverse problem may provide an important tool for learning about high density nuclear matter. Numerous high quality measurements of the masses of neutron stars are now available~\\cite{Lattimer2007}, and a few (fairly imprecise and model-dependent) radius measurements are starting to become available as well~\\cite{Steiner2010, Galloway2012, Guver2012a, Guver2012b}. In principle then, the inverse stellar structure problem should (eventually) allow us to measure the high density equation of state of neutron-star matter directly. This measurement would provide important information about nuclear interactions that can not be obtained at present in any other way. One naive approach to solving the inverse stellar structure problem for neutron stars would simply be to match their observed properties, e.g. their $[M,R]$ data, with models of those stars based on different micro-physical models of the dense material in their cores. In this approach the model equation of state whose stellar models best matches the data would be declared the observed neutron-star equation of state. This approach would clearly be ideal if there were wide consensus on exactly what the high density core material is, and if there were a reasonably simple model for this material that were known, up to a few undetermined parameters that could be fixed by these observations. Unfortunately the wide diversity of ``realistic'' neutron-star equations of state in the literature, suggest that (in the near term at least) this approach is not likely to be effective or conclusive. A more practical variation of this approach uses some knowledge about the expected properties of neutron-star matter in an intermediate range of densities, and a more empirical description of the equation of state for larger densities~\\cite{Steiner2010, Steiner2012}. This approach is more promising, but the proposed model equations of state of this type have many free parameters that must all be fit by the observational data. Since these data are likely to be sparse for some time, we take a different approach here. Our goal is to find efficient and robust methods for solving the inverse stellar structure problem that use no prior knowledge of the high density micro-physics at all. A (somewhat impractical) method for solving the inverse stellar structure problem that uses no information about the micro-physics of the high density equation of state was given in the literature about 20 years ago~\\cite{Lindblom1992}. This traditional method can be summarized as follows. The total masses $M$ and radii $R$ of all of the stars associated with a particular equation of state are assumed to be known. The equation of state is also assumed to be known up to some pressure $p_i$ with corresponding energy density $\\epsilon_i=\\epsilon(p_i)$. Let $M_i=M(p_i)$ and $R_i=R(p_i)$ denote the mass and radius of the star whose central pressure is $p_c=p_i$. Now choose another point, $[M_{i+1},R_{i+1}]$, along the mass-radius curve, with a slightly larger central pressure. The outer layers of this new star are composed of low pressure material, $p\\leq p_i$, where the equation of state is known. The stellar structure equations, (\\ref{e:OVm}) and (\\ref{e:OVp}), can therefore be solved in this outer region starting at the surface of the star, $r=R_{i+1}$, where $p(R_{i+1})=0$ and $m(R_{i+1})=M_{i+1}$, by integrating inward toward $r=0$. This integration can be continued until the point $r=r_{i+1}$ where $p(r_{i+1})=p_i$ and the known equation of state ends. This integration determines the radius $r_{i+1}$ and the mass $m_{i+1}=m(r_{i+1})$ of a ``small'' core of high pressure material, $p\\ge p_i$ where the equation of state is not yet known. If this core is small enough, the stellar structure equations can be solved there as a power series expansion about $r=0$. The coefficients in this expansion are functions of the central density $\\epsilon_{i+1}$ and pressure $p_{i+1}$ of this little core. Since the mass and radius of this core are known, $m_{i+1}$ and $r_{i+1}$, this power series can be ``inverted'' to determine $\\epsilon_{i+1}$ and $p_{i+1}$~\\cite{Lindblom1992}. This new point $[\\epsilon_{i+1},p_{i+1}]$ provides a small extension of the equation of state beyond $[\\epsilon_i,p_i]$. Iterating these steps then determines a sequence of closely-spaced points along the high density portion of the equation of state curve. This traditional solution to the inverse stellar structure problem is unfortunately very impractical. A large number of points $[M_i,R_i]$ are needed from the mass-radius curve to achieve modest accuracy in the calculation of the corresponding points $[\\epsilon_i,p_i]$ along the equation of state curve. Since $[M_i,R_i]$ points are very difficult to measure (at least for neutron stars) it is unlikely that there will ever be enough data to use this traditional method to determine the unknown high density part of the neutron-star equation of state. This paper proposes a rather different approach to the solution of the inverse stellar structure problem, an approach that can be very effective even when only a small number of $[M_i,R_i]$ data points are available. The equation of state in this new approach is expressed as a parametric equation, e.g. $\\epsilon=\\epsilon(p,\\gamma_k)$, instead of a table of values $[\\epsilon_i,p_i]$ . The parameters $\\gamma_k$ are adjusted to give the best-fit approximation to a particular equation of state model. Parametric representations of this sort, based on spectral expansions, have been shown to be extremely efficient at representing the high density portions of ``realistic'' neutron-star equations of state~\\cite{Lindblom10}. Only a few non-vanishing $\\gamma_k$ are generally needed to achieve 1\\% accuracy in most cases. The basic idea of this new method for solving the inverse stellar structure problem is to choose the equation of state parameters $\\gamma_k$ by minimizing the differences between the masses and radii of real neutron stars, $M_i$ and $R_i$, with those based on the parametric model equation of state, $M(p_c,\\gamma_k)$ and $R(p_c,\\gamma_k)$. Once the $\\gamma_k$ are fixed, the parametric equation $\\epsilon=\\epsilon(p,\\gamma_k)$ then provides an approximate solution of the inverse stellar structure problem. Spectral expansions typically converge exponentially as the number of terms in the expansion are increased (for smooth functions). These approximate solutions to the inverse stellar structure problem are therefore expected to converge to the exact equation of state as the number of data points $[M_i,R_i]$ and the number of parameters $\\gamma_k$ fixed by this method are increased. The remainder of this paper presents details on how to implement this new spectral approach to the solution of the inverse stellar structure problem, along with practical tests of its accuracy and efficiency. Section~\\ref{s:SpectralEOS} reviews the particular spectral representation of the equation of state used in the solution presented here. Section~\\ref{s:FixingSpectralParameter} describes how the spectral parameters $\\gamma_k$ are fixed by matching to the given $[M_i,R_i]$ data points. Section~\\ref{s:TestingSpectralInversionMethod} presents a series of numerical tests of the accuracy and efficiency of this new method. Mock $[M_i,R_i]$ data computed from a collection of 34 ``realistic'' neutron-star equations of state are used as input in these tests. These tests show, for example, that the resulting spectral equation of state agrees with the exact to within a few percent (on average) when only two $[M_i,R_i]$ data points are used. Higher accuracies are (generally) achieved when more data points are used. Section~\\ref{s:Discussion} discusses some of the limitations of the numerical tests presented here, and proposes several ways that the basic method developed here might be extended and improved. Some of the more complicated technical details needed to implement this method are described in two Appendices. Appendix~\\ref{s:AppendixA} describes how to evaluate the derivatives of $M(h_c,\\gamma_k)$ and $R(h_c,\\gamma_k)$, with respect to the parameters $h_c$ and $\\gamma_k$. Appendix~\\ref{s:AppendixB} describes the interpolation method used here to bridge the gaps between points in the exact ``realistic'' equation of state tables. ", "conclusions": "\\label{s:Discussion} In summary, we have developed a new method for solving the relativistic inverse stellar structure problem based on the construction of a spectral expansion of the unknown high density part of the equation of state of the star. The results of our numerical tests of this new method, described in Sec.~\\ref{s:TestingSpectralInversionMethod}, are quite impressive. Using only two $[M_i,R_i]$ data points, this new method can determine the entire high density part of the neutron-star equation of state with errors (on average) of just a few percent. The addition of more data points (generally) results in higher accuracy approximations. We also show that $N$-parameter spectral approximations to the equation of state determined in this way are almost as accurate as the best possible $N$-parameter spectral approximations. This is quite remarkable. It shows that macroscopic mass-radius measurements are strongly correlated to the properties of the equation of state, and such measurements should therefore allow us (eventually) to measure the high density part of the neutron-star equation of state with great precision. A close inspection of the results from the various tests summarized in Table~\\ref{t:TableI} reveals a number of anomalies that merit further study. For example, the error measure $\\Delta_{N_{\\gamma_k}}$, defined in Eq.~(\\ref{e:DeltaDef}), is expected to decrease as the number of spectral parameters $N_{\\gamma_k}$ is increased, i.e. that $\\Delta_{N_{\\gamma_k}}\\geq\\Delta_{N_{\\gamma_k}+1}$. This seems to be true for most of our tests, but there are also a number of exceptions in Table~\\ref{t:TableI}. The equation of state FPS, for example, has $\\Delta_2=0.0048$, $\\Delta_3=0.0061$, $\\Delta_4=0.0096$, and $\\Delta_5=0.0048$. What is going on? Such a sequence of errors would be consistent, for example, with the idea that this particular equation of state is not well represented by these low order spectral expansions, i.e. that these expansions in this case are not yet in the convergent regime. This does not seem to be the case however since the optimal spectral fits to the FPS equation of state do appear to be convergent with these same numbers of spectral parameters, cf. Table II of Ref.~\\cite{Lindblom10}. Another (more likely) explanation of the anomalous results found in Table~\\ref{t:TableI} is that the minima of $\\chi^2(h_c,\\gamma_k)$ found by the Levenberg-Marquardt algorithm for these cases are just local minima and not the desired global minima. An interesting area for further research on this problem, therefore, will be to explore the use of more robust numerical methods for finding global minima of complicated non-linear functions like $\\chi^2(h_c,\\gamma_k)$. Another interesting direction for future research on this problem will be to explore how robust this kind of solution to the inverse stellar structure problem will be when applied to more realistic $[M_i,R_i]$ data sets. The data used here were idealized in two important ways. First, the mock $[M_i,R_i]$ data used here were supplied with very high precision. Real astrophysical measurements of these quantities will have significant errors. How will measurement errors influence the accuracy of the equation of state that is constructed by these techniques? Second, the mock $[M_i,R_i]$ data used here were chosen to cover uniformly the astrophysically relevant range of neutron-star masses. Real astrophysical measurements will not be distributed in such an orderly way. How will the accuracy of the implied equation of state be affected by different, presumably less ideal, data distributions? The version of the inverse stellar structure problem studied here is based on the use of mass $M$ and radius $R$ measurements to determine the high density part of the equation of state. These are not the only macroscopic properties of neutron stars that could potentially be measured. It is not too difficult to imagine for example that the moment of inertias or the tidal Love numbers might be more easily observable for some types of observations. Another interesting direction for future study will therefore be to explore the use of other measurement data, say the mass and Love number (which could be measured using gravitational wave observations of neutron-star mergers), as input for solving the inverse stellar structure problem using the spectral methods developed here." }, "1207/1207.4487_arXiv.txt": { "abstract": "The Visible Integral field Replicable Unit Spectrograph (VIRUS) is an array of at least 150 copies of a simple, fiber-fed integral field spectrograph that will be deployed on the Hobby-Eberly Telescope (HET) to carry out the HET Dark Energy Experiment (HETDEX). Each spectrograph contains a volume phase holographic grating as its dispersing element that is used in first order for $350 < \\lambda \\mathrm{(nm)} < 550$. We discuss the test methods used to evaluate the performance of the prototype gratings, which have aided in modifying the fabrication prescription for achieving the specified batch diffraction efficiency required for HETDEX. In particular, we discuss tests in which we measure the diffraction efficiency at the nominal grating angle of incidence in VIRUS for all orders accessible to our test bench that are allowed by the grating equation. For select gratings, these tests have allowed us to account for $>90$\\% of the incident light for wavelengths within the spectral coverage of VIRUS. The remaining light that is unaccounted for is likely being diffracted into reflective orders or being absorbed or scattered within the grating layer (for bluer wavelengths especially, the latter term may dominate the others). Finally, we discuss an apparatus that will be used to quickly verify the first order diffraction efficiency specification for the batch of at least 150 VIRUS production gratings. ", "introduction": "\\label{sec:intro} % The upcoming Hobby-Eberly Telescope Dark Energy eXperiment (HETDEX; Ref. \\citenum{Hill08a}) will amass a sample of $\\sim0.8$ million \\lya\\ emitting galaxies (LAE) to be used as tracers of large-scale structure for constraining dark energy and measuring its possible evolution from $1.9 < z < 3.5$. To carry out the 120 night blind spectroscopic survey covering a 420 square degree field (9 Gpc$^{3}$), a revolutionary new multiplexed instrument called the Visible Integral field Replicable Unit Spectrograph (VIRUS; Ref. \\citenum{Hill12a}) is being constructed\\cite{Tuttle12} for the upgraded 9.2 m Hobby-Eberly Telescope (HET\\footnote{The Hobby-Eberly Telescope is operated by McDonald Observatory on behalf of the University of Texas at Austin, the Pennsylvania State University, Ludwig-Maximillians-Universit\\\"{a}t M\\\"{u}nchen, and Georg-August-Universit\\\"{a}t Goettingen.}; Ref. \\citenum{Hill12b}). The VIRUS spectrograph array consists of at least 150 copies (with a goal of 192) of a simple fiber-fed integral field spectrograph and for the first time has introduced industrial-scale replication to optical astronomical instrumentation. The spectrographs are mechanically built into unit pairs and are fed by dense-pack fiber bundle integral field units (IFU) with 1/3 fill factor, each consisting of 448 fiber optic elements with a core diameter of 266 $\\mu$m (1.5\\arcsec\\ on the sky). Thus, each spectrograph images 224 fibers. At least 75 IFUs will be arrayed on the 22\\arcmin\\ diameter focal plane of the upgraded HET, yielding $\\sim33000$ individual spectra per exposure. Each spectrograph consists of a double-Schmidt optical design with a volume phase holographic (VPH) diffraction grating at the pupil between a $f$/3.33 folded collimator and a $f$/1.25 cryogenic camera. The spectral coverage of VIRUS is $350 < \\lambda \\mathrm{(nm)} < 550$ at $R = \\lambda / \\Delta\\lambda \\approx 700$ for measuring the baryonic acoustic oscillation via the \\lya\\ emission of star-forming galaxies from $1.9 < z < 3.5$. Fig. \\ref{fig:VIRUS} shows a rendering of VIRUS and its optical design. \\begin{figure}[t] \\begin{center} \\begin{tabular}{c} \\includegraphics[width=0.95\\textwidth]{f1.png} \\end{tabular} \\end{center} \\caption[example] { \\label{fig:VIRUS} \\textit{a}) A rendering of the upgraded HET showing the large enclosures mounted on either side of the telescope structure that contain VIRUS spectrographs. The green cables extending from the prime-focus instrument package to the enclosures are large bundles of fiber optics. \\textit{b}) Close view of two enclosures, each containing an 8$\\times$3 array of VIRUS units (48 total spectrographs). \\textit{c}) Section view of a single VIRUS pair. \\textit{d}) A ray trace of the VIRUS optical system.} \\end{figure} The VIRUS concept has already been proven by the Mitchell Spectrograph (formerly known as VIRUS-P; Ref. \\citenum{Hill08b}), a single prototype VIRUS spectrograph that has been in use at the McDonald Observatory 2.7 m Harlan J. Smith telescope since 2007. The instrument has excellent throughput in the blue down to 350 nm ($\\sim30$\\%, excluding the telescope and atmosphere). For VIRUS, similar or better throughput will be essential to keep the number of LAE detections sufficiently high for achieving the goals of HETDEX. At the bluest wavelengths, a limiting optical component is the VPH diffraction grating that is used as the instrument's dispersing element. Ref. \\citenum{Adams08} discusses the performance of VPH gratings developed for the Mitchell Spectrograph, which at that time pushed the technology to the highest diffraction efficiency achieved at 350 nm ($\\sim60$\\%); for VIRUS, we desire even higher diffraction efficiency of $\\sim70$\\% at 350 nm (see $\\S$\\ref{sec:gratingspec}). In order to tune the VPH layer fabrication prescription for achieving such high efficiency and maintaining it across the instrument's spectral coverage, we must understand how a given diffraction grating distributes the incoming light into various diffraction orders and where losses (i.e., scattering and/or absorption) occur. In addition, a challenge that is currently specific to VIRUS is achieving consistency in the required high standard of performance for large batches of $>150$ gratings. In this paper, we discuss methods for testing and verifying the suite of $> 150$ VPH diffraction gratings for VIRUS. We begin in $\\S$\\ref{sec:gratingspec} by describing the production design of the gratings. In $\\S$\\ref{sec:performance}, we discuss the methods we have used to date in measuring the diffraction efficiency and highlight the measurements of select prototype gratings that have helped lead us toward an acceptable fabrication prescription. In $\\S$\\ref{sec:tester}, we present the design of a new test apparatus that will allow for the efficient verification of the diffraction efficiency specification for the entire suite of VIRUS gratings. Finally in $\\S$\\ref{sec:conclusions}, we outline the general conclusions of the paper. ", "conclusions": "" }, "1207/1207.3979_arXiv.txt": { "abstract": "{After years of modest optical activity, the quasar-type blazar \\object{4C 38.41} (\\object{B3 1633+382}) experienced a large outburst in 2011, which was detected throughout the entire electromagnetic spectrum, renewing interest in this source.} {We present the results of low-energy multifrequency monitoring by the GASP project of the WEBT consortium and collaborators, as well as those of spectropolarimetric/spectrophotometric monitoring at the Steward Observatory. We also analyse high-energy observations of the {{\\it Swift}} and {{\\it Fermi}} satellites. This combined study aims to provide insights into the source broad-band emission and variability properties.} {We assemble optical, near-infrared, millimetre, and radio light curves and investigate their features and correlations. In the optical, we also analyse the spectroscopic and polarimetric properties of the source. We then compare the low-energy emission behaviour with that at high energies.} {In the optical--UV band, several results indicate that there is a contribution from a quasi-stellar-object (QSO) like emission component, in addition to both variable and polarised jet emission. In the optical, the source is redder-when-brighter, at least for $R \\ga 16$. The optical spectra display broad emission lines, whose flux is constant in time. The observed degree of polarisation increases with flux and is higher in the red than the blue. The spectral energy distribution reveals a bump peaking around the $U$ band. The unpolarised emission component is likely thermal radiation from the accretion disc that dilutes the jet polarisation. We estimate its brightness to be $R_{\\rm QSO} \\sim 17.85$--18 and derive the intrinsic jet polarisation degree. We find no clear correlation between the optical and radio light curves, while the correlation between the optical and $\\gamma$-ray flux apparently fades in time, likely because of an increasing optical to $\\gamma$-ray flux ratio.} {As suggested for other blazars, the long-term variability of 4C 38.41 can be interpreted in terms of an inhomogeneous bent jet, where different emitting regions can change their alignment with respect to the line of sight, leading to variations in the Doppler factor $\\delta$. Under the hypothesis that in the period 2008--2011 all the $\\gamma$-ray and optical variability on a one-week timescale were due to changes in $\\delta$, this would range between $\\sim 7$ and $\\sim 21$. If the variability were caused by changes in the viewing angle $\\theta$ only, then $\\theta$ would go from $\\sim 2.6 \\degr$ to $\\sim 5 \\degr$. Variations in the viewing angle would also account for the dependence of the polarisation degree on the source brightness in the framework of a shock-in-jet model. } ", "introduction": "The Compton Gamma Ray Observatory (CGRO) launched in 1991 revealed 271 $\\gamma$-ray sources, one fourth of which were identified as blazars \\citep{har99}, i.e.\\ active galactic nuclei (AGNs) showing the most extreme properties. They are indeed characterised by violent activity at almost all frequencies on different timescales, from long-term flux changes to intraday variability (IDV), by high radio to optical polarisation, and by the superluminal motions of radio knots. The observations can be explained by assuming that their emission is relativistically beamed, which occurs if the emitting plasma jet, produced by a supermassive black hole fed by an accretion disc, is directed toward us. The polarised, low-energy emission (from the radio to the optical--X-ray band) is likely synchrotron radiation produced by relativistic electrons in the jet, while the high-energy emission (from the X-ray to the $\\gamma$-ray frequencies) is usually interpreted as the result of inverse-Compton scattering of low-energy photons off the same relativistic electrons. Whether these low-energy photons come from the jet itself (synchrotron self-Compton, or SSC, models) or from the AGN environment (external Compton, or EC, models), and in the latter case whether they come from the accretion disc, the broad line region, or the dusty torus, is still a matter of debate. Blazars include flat spectrum radio quasars (FSRQs) and BL Lacertae objects. Frequently FSRQs exhibit quasi-stellar-object (QSO) like broad emission lines, as well as a blue and unpolarized continuum that is presumed to be the signature of the thermal emission from the accretion disc, the so-called ``big blue bump\" \\citep[e.g.][]{wil92,pia99,rai07b,dam09}. In contrast, BL Lacertae objects may have by definition at most weak lines, even if sometimes these sources challenge their classification \\citep[e.g.][]{ver95}. Researchers are trying to understand the structure of blazars and the mechanisms behind their emission and variability by analysing the multifrequency behaviour of these objects extended over the broadest possible energy range with data sampling and quality that is as high as possible \\citep[see e.g.][]{mar10,jor10,agu11a,agu11b}. The international collaboration known as the Whole Earth Blazar Telescope (WEBT)\\footnote{\\tt http://www.oato.inaf.it/blazars/webt/} was created in 1997 for this purpose, and involves tens of optical, radio, and near-infrared observatories. The WEBT campaigns have produced low-frequency light curves of extraordinary sampling, and these results have often been analysed in conjunction with high-frequency data from satellites \\citep[see e.g.][and references therein]{vil06,rai08a,rai08c,lar08,boe09,vil09a,rai09}. Renewed interest in blazars occurred with the launch of the new-generation $\\gamma$-ray satellites Astrorivelatore Gamma a Immagini Leggero (AGILE) in 2007, and particularly with that of {\\it Fermi} (formerly GLAST) in 2008, which operates in survey mode. To acquire low-energy data to compare with the high-energy observations of AGILE and {\\it Fermi}, in 2007 the WEBT started the GLAST-AGILE Support Program \\citep[GASP; see e.g.][]{vil08,vil09b,dam09,rai10,rai11,dam11}. Since 2008, a ground-based monitoring programme has been running at the Steward Observatory \\citep{smi09}, providing support for the {\\it Fermi} gamma-ray telescope. It uses the 2.3 m Bok and 1.54 m Kuiper telescopes with the SPOL spectropolarimeter \\citep{sch92} and provides publicly available spectropolarimetry, spectrophotometry, and calibrated broad-band flux measurements for about 40 blazars\\footnote{\\tt http://james.as.arizona.edu/$\\sim$psmith/Fermi/}. In this paper, we study the multifrequency behaviour of one blazar, namely the optically violent variable (OVV) FSRQ 4C 38.41 (1633+382) at redshift $z=1.814$, which belongs to the target list of both the GASP and the Steward Observatory monitoring programmes. The period of observations presented here includes a big optical outburst in 2011 \\citep{rai11_atel} that was also detected at $\\gamma$-ray energies by {\\it Fermi} \\citep{szo11}, and in the X-ray and UV bands by the {\\it Swift} satellite, providing an excellent opportunity to study the source variability over most of the electromagnetic spectrum. 4C 38.41 had already been observed in $\\gamma$-rays by the Energetic Gamma Ray Experiment Telescope (EGRET) instrument onboard CGRO several times, and from these data its emission was found to vary significantly on a day timescale \\citep{mat93}. The maximum flux detected by EGRET above 100 MeV was $(107.5 \\pm 9.6) \\times 10^{-8} \\rm \\, ph \\, cm^{-2} \\, s^{-1}$ in mid September 1991\\footnote{\\tt http://cossc.gsfc.nasa.gov/cossc/egret/}. AGILE observed 4C 38.41 several times. In particular, a preliminary analysis of the AGILE Gamma Ray Imaging Detector (GRID) data between 2009 December 1 and 2010 November 30 yielded a $\\gamma$-ray flux $F_{\\rm E>100\\,MeV} = (28 \\pm 5) \\times 10^{-8} \\rm \\, ph \\, cm^{-2} \\, s^{-1}$. This flux is consistent with the value reported in the Second {\\it Fermi} Large Area Telescope (LAT) Catalog \\citep{nol12}. A refined analysis of the AGILE/GRID data, along with a search for possible short-term variability, will be presented in Vercellone et al.\\ (2012, in preparation). Multifrequency observations of 4C 38.41 during the $\\gamma$-ray flares observed by {\\it Fermi} in 2009--2010, including Very Long Baseline Array (VLBA) images, was presented by \\citet{jor11}. They conclude that high states at $\\gamma$-ray energies are due to interaction between a disturbance travelling down the jet and the 43 GHz VLBI core. \\begin{figure*} \\centering \\includegraphics{ottico.eps} \\caption{Optical and near-infrared light curves of 4C 38.41 in 2008--2011 built with GASP-WEBT data (blue dots), {\\it Swift}-UVOT data (red crosses), and data from the Steward Observatory (green plus signs). The UVOT $v$-band points have been shifted by $-0.1$ mag to match the ground-based data. } \\label{ottico} \\end{figure*} ", "conclusions": "We have presented the results of a huge observing effort on the FSRQ 4C 38.41, carried out by the GASP-WEBT and collaborators from 2007.5 to 2011.9 and by the Steward Observatory from 2008.8 to 2012.1. Earlier data were also collected, so that the optical and radio light curves cover in total about 17 years. Moreover, we also analysed the UV and X-ray data acquired in 2007--2011 by the {\\it Swift} satellite and the $\\gamma$-ray data taken in 2008--2011 by {\\it Fermi}. Light curves from the near-IR to the UV band, spanning a factor $\\sim 11$ in frequency, show a quite impressive correlation, presumably because this whole range is in the upper part of the synchrotron bump. In the near-IR-to-UV spectral range, the presence of a QSO-like emission contribution in addition to the synchrotron emission from the jet, is revealed by several findings: \\begin{itemize} \\item the maximum flux-variation amplitude decreases with increasing frequency; \\item the optical colour shows a redder-when-brighter trend; \\item the optical spectrum reveals unpolarised broad emission-lines with constant flux; \\item the optical polarisation is higher in the red than in the blue; \\item the SEDs display a bump peaking around the $u$ band in faint states. \\end{itemize} This unpolarised emission component is likely thermal emission from the accretion disc. Thermal AGN signatures have apparently been found in several blazars based on different information (see e.g. \\citealt{per08} for a review, and \\citealt{ghi09}), but this is maybe the first time so much observational evidence has been collected for a single object. Furthermore, we have been able to estimate the brightness of the unpolarised component, $R_{\\rm QSO} \\sim 17.85$--18, and to correct the observed degree of polarisation for its dilution effect to obtain the polarisation of the jet emission. This still shows a dependence on brightness, which is thus an intrinsic dependence. The analysis of the radio light curves confirms a scenario where radiation of increasing wavelength is emitted in progressively larger and more external jet regions. Optical and radio light curves do not show a general correlation, at least on month--year timescales. One possible explanation, which has already been proposed for other blazars, is that the optical and radio emissions come from different jet zones that have variable orientations with respect to the line of sight. Optical and radio flares would thus result when the corresponding emitting regions are more closely aligned, with a consequent enhancement of the Doppler factor. We found a general correlation between the flux variations in the optical band and at $\\gamma$-ray energies, as the optical and $\\gamma$-ray periods of increased activity coincide. However, the shape of the $\\gamma$-ray outbursts differs from that of the optical ones, and the flux ratios as well as the time lags between variations at the two frequencies change with time. The correlation was strong during 2009, but became weaker in 2011, likely because of an increased activity in the optical band. We analysed broad-band SEDs of 4C 38.41 built with contemporaneous data in different brightness states. A careful spectral analysis of both the X-ray data from {\\it Swift} and $\\gamma$-ray data from {\\it Fermi} was also performed to obtain a reliable spectral shape in the high-energy part of the SED. All the selected epochs show a strong Compton dominance, even when the jet emission is weak and the unpolarised component clearly emerges in the optical band. We finally discussed a geometrical interpretation of the flux and polarisation variability, which is able to fairly account for the observational data. In this view, at least the long-term flux variations can be ascribed to changes in the Doppler boosting factor produced by changes in the viewing angle. In particular, the weekly $\\gamma$-ray and optical light curves would imply changes in $\\delta$ and $\\theta$ in the ranges of $\\sim 7$--21 and $\\sim 2.6\\degr$--5\\degr, respectively. When using these values in the framework of a shock-in-jet model, where the direction of the wave front, and hence the polarisation degree, changes according to the viewing angle, we could also account for the trend of polarisation with brightness." }, "1207/1207.4234.txt": { "abstract": "{}{}{}{}{} % 5 {} token are mandatory \\abstract{}{}{}{}{} % \\abstract % context heading (optional) % {} leave it empty if necessary %{General relativistic emission lines and images produced by compact objects, such as accretion material around black holes, are of great importance in probing the effects of extreme gravity environments.} % {The relativistic emission lines produced by thin accretion disks and tori around Kerr black holes are an important diagnostic tool for probing the effects of extreme gravity % near the event horizon. Accretion disks in high luminosity AGN are geometrically thick and require a three-dimensional (3D) treatment of the radiative transfer (RT) of the emission from near the black hole event horizon.} % aims heading (mandatory) % {We present 3D covariant radiative transfer calculations of the emission from these thick accretion disks (tori). } {We aim to construct a general relativistic radiative transfer formulation, applicable to particles with or without mass in astrophysical settings, wherein ray-tracing calculations can be performed for arbitrary geodesics for a given space-time geometry. % methods heading (mandatory) The relativistic radiative transfer formulation is derived from first principles: conserving particle number and phase-space density. The formulation is covariant, and transfer calculations are conducted along particle geodesics connecting the emitters and the observer. The geodesics are determined through the space-time metric, which is specified beforehand. Absorption and emission in the radiative transfer calculations are treated explicitly. The particle-medium interaction is evaluated in the local inertial frame, co-moving with the medium. Relativistic, geometrical and optical depth effects are treated self-consistently within an integral covariant framework. % results heading (mandatory) We present a self-consistent general relativistic radiative transfer formulation with explicit treatment of emission and absorption. The formulation is general and is applicable to both particles with mass and without mass. The presence of particles has two major effects: firstly the particle bundle ray is no longer along the null geodesic, and secondly the intensity variation along the particle bundle ray is reduced by an aberration factor. The radiative transfer formulation can handle 3D geometrical settings and structured objects with variations and gradients in the optical depths across the objects and along the line-of-sight. Such scenarios are applicable in calculations of photon emission from complex structured accretion flows around black holes and neutrino emission from remnant neutron tori in neutron-star mergers. We apply the formulation and demonstrate radiation transfer calculations for emission from accretion tori around rotating black holes. We consider two cases: idealised optically thick tori that have a sharply defined emission boundary surface, and structured tori that allow variations in the absorption coefficient and emissivity within the tori. We show intensity images and emission spectra of the tori obtained in our calculations. Our findings in the radiative transfer calculations are summarised as follows. (i) Geometrical effects, such as lensing-induced self-occulation and multiple-image contribution are much more significant in accretion tori than geometrically thin accretion disks. (ii) Optically thin accretion tori show emission line profiles distinguishable from the profiles of lines from optically thick accretion tori and lines from optically thick geometrically thin accretion disks. (iii) The line profiles of the optically thin accretion tori have a weaker dependence on the viewing inclination angle than those of the optically thick accretion tori or accretion disks, especially at high viewing inclination angles. (iv) Limb effects are present in accretion tori with finite optical depths, due to density and temperature stratification within the tori. We note that in accretion flows onto relativistic compact objects, gravitationally induced line resonance can occur. This resonance occurs easily in 3D flows, but not in 2D flows, such as a thin accretion disk around a black hole.} % conclusions heading (optional), leave it empty if necessary ", "introduction": "The X-ray emission observed in active galactic nuclei (AGN) and black hole binaries is believed to be powered by accretion of material onto black holes \\citep{Salpeter1964, Lynden-Bell1969, Shakura1973}. The accreting hot plasmas rotating around the black hole form a disk or a torus. It has been suggested that dense remnant neutron tori can also be formed around compact objects, which could be a very massive neutron star or a black hole, after two neutron stars merge \\citep[see][]{Shibata2003, Baiotti2008, Rezzolla2010}. In such systems, the influence of curved space-time is significant. It affects the radiative transport of particles in the accretion flow as well as the hydrodynamics of the flow itself \\citep[see][]{Novikov1973}. Emission from accretion disks around compact objects has been investigated for several decades now. Emission lines from geometrically thin accretion disks around gravitating objects are expected to have two peaks \\citep{Smak1969}. The peaks correspond to emission from the two parts of the disk, which have opposite projected line-of-sight velocities. Double-peaked optical lines have been observed in a variety of binary systems \\citep[e.g.\\ black hole X-ray binaries,][]{Johnston1989, Marsh1994, Soria1999, Wu2001}. Double-peaked optical lines are also seen in a small fraction of AGN \\citep{Puchnarewicz1996, Eracleous2003, Strateva2006}. These double-peaked lines can be explained in a Newtonian framework as described in \\cite{Smak1969} \\citep[see also][]{Horne1986}. Double-peaked lines have also been observed in the X-ray spectra of accreting black holes. Broad asymmetric double-peaked Fe K$\\alpha$ lines were found in the spectra of a number of AGN \\citep[e.g.\\ MCG -6-30-15,][]{Tanaka1995}. The X-rays of AGN are believed to originate from regions very close to the central black hole, where the accretion flow is highly relativistic and the gravity is strong. The emissions from different parts of the accretion flow are therefore boosted differentially. Various relativistic effects also cause additional differential broadening and distortion of any line emission. As such, the emission lines from the inner regions of relativistic accretion disks around black holes have a very broad profile, with an extended red wing and a pronounced blue peak \\citep{Cunningham1975, Reynolds1999, Fabian2000}. Black holes with faster spins would give rise to relativistic accretion lines with a broader red wing, because the inner boundary of an accretion disk around a maximally rotating black hole can extend very close to the black hole event horizon. % Moreover, studies have also show that gravitational lensing produces multiple images, % and this allows self-occulation to occur, thus further modifying the total emission and the profiles of the lines % \\citep{Fuerst2004, Beckwith2005}. % % These line profiles are usually interpreted as emission from material circulating with high % velocities in the inner accretion disk close to the event horizon of the central black % hole. These effects arise due to line photons having to climb out of the deep % gravitational well of the central black hole. By modelling the observed emission line profiles the viewing inclination % of the accretion disk/torus, the black hole spin parameter and the spatial distribution of % the line emission may be constrained \\citep{McClintock2011}. The emission from accreting gas and outflows in the vicinity of black holes is subject to Doppler shifts, lensing, gravitational time-dilation, and other relativistic dynamical effects. There have been numerous calculations of relativistic (photon) lines from accretion disks and tori around black holes \\citep[e.g.][]{Cunningham1975, Gerbal1981, Fabian1989, Stella1990, Kojima1991, Bao1992, Fanton1997, Reynolds1999, Fabian2000, Fuerst2004, Beckwith1_2004, Beckwith2005, Cadez2006, Schnittman2006, Fuerst2007, Dexter2009, Sochora2011, Vincent2011, Wang2012}. The three most common methods for calculating relativistic line profiles are (i) the transfer function method \\citep[e.g.][]{Cunningham1975, Fabian2000}, (ii) the elliptic function method \\citep[e.g.][]{Dexter2009} and (iii) the direct geodesic integration method \\citep[e.g.][]{Fuerst2004, Schnittman2006, Fuerst2007, Anderson2010, Vincent2011}. The transfer function method and the elliptic function method are efficient for the calculation of emission from thin axisymmetric optically thick accretion disks but are not applicable to systems that lack the appropriate geometry and symmetry. The direct geodesic integration method is a brute force approach, and less restrictive in this context, compared to the other two methods. It works well with any three-dimensional (3D) accretion flow, e.g. time-dependent accretion flows from numerical relativistic hydrodynamic simulations, and can also handle opacity variations within the system. In most of these relativistic calculations, the focus was on the investigation of line broadening due to relativistic effects. While relativistic ray-tracing of photons in strong gravity and the corresponding calculations of emission line profile broadening have been investigated in various astrophysical settings for decades, there have been only a few studies, often in restricted settings, of the opacity effects due to self-absorption, emission and scattering within the accretion flows and along the line of sight \\citep[e.g.][]{Zane1996, Fuerst2006, Wu2006, Wu2008, Dolence2009}. Covariant radiative transfer calculations of emission from accretion flows in more general settings with an explicit treatment of absorption, emission and scattering are lacking. There is also no corresponding covariant radiative transport formulation applicable to both relativistic particles without mass (photons) and with mass (e.g.\\ neutrinos), in a general astrophysical setting in the present literature. (Note that particles with mass do not follow null geodesics, and the relativistic formulation and the corresponding ray-tracing need to be modified.) In this work we present the general covariant formulation for radiative transport of relativistic particles. The radiative transfer equation is derived from the Lorentz-invariant form of the conservation law. It explicitly includes emission and absorption processes. The formulation is not restricted to general relativistic radiative transfer of photons. It is general and can be applied to both relativistic particles without mass (e.g.\\ photons) and with mass (e.g.\\ relativistic electrons and neutrinos). The radiative transfer recovers its conventional form in the Newtonian limit. We carried out demonstrative radiative transfer calculations of emission from model accretion tori around rotating black holes with different optical thicknesses. Our calculations show the convolution of geometrical and optical depth effects, such as limb darkening and brightening caused by optical-depth and emissivity variations, multiple image contribution, and absorption and self-occultation induced by gravitational lensing, which are characteristic of 3D structured flows in strong gravity environments. The paper is organised as follows. In Sec. 2 the covariant radiative transfer equation is derived and its solution discussed. In Sec. 3 we show the construction of two torus models, one with a specific sharp emission boundary surface, and another with a stratified density and temperature structure. In Sec. 4 we present the results of radiative transfer calculations for the torus models. We briefly discuss the astrophysical implications of our calculations and the possible occurrence of gravitationally induced line resonance in 3D accretion flows near black holes. %__________________________________________________________________ ", "conclusions": "We have derived the radiative transfer equation from first principles, conserving both particle number and phase-space density. The equation is thus manifestly covariant, and it is applicable in arbitrary 3D geometrical settings and for any pre-defined spacetime metric. The more general form of the equation explicitly considers a covariant particle flux with mass. We have found that ind addition to modifying the geodesic trajectories that connect the observer and the emitting regions, the presence of particle mass in the radiative transfer equation introduces a mass-dependent aberration, which reduces the intensity gradient along the ray. In the zero-mass limit this general radiative transfer equation recovers its original form for the massless particles. We carried out demonstrative numerical general relativistic radiative transfer calculations with a ray-tracing algorithm constructed from the formulation. Different 3D accretion tori around rotating black holes with different geometrical aspects, physical structures, emission properties, and optical depth variations were considered. We demonstrated that radiative transfer calculations based on the formulation are able to deal with the complexity in the various combinations and convolution of relativistic, geometrical, physical and optical effects. Our calculations clearly showed the significant role that structures and optical depth, and their gradients, together with geometrical and relativistic factors, play in shaping the emission properties of these 3D relativistic flows in the vicinity of rotating black holes. The calculations also showed the presence of limb effects in 3D objects with finite optical depths. We note that gravitationally induced line resonance can occur in 3D accretion onto a compact object. This phenomenon is not present in 2D planar objects, such as geometrically thin accretion disks, where the radiation can escape from the disk surface to free space without additional absorption or re-emission. % \\end{enumerate} % \\begin{appendix} % %" }, "1207/1207.2839_arXiv.txt": { "abstract": "We present the first orbit--integrated self force effects for an IMRI or EMRI source, specifically the effects of its conservative piece on the orbit and on the waveform. We consider the quasi--circular motion of a particle in the spacetime of a Schwarzschild black hole, find the orbit and the corresponding gravitational waveform, and discuss the importance of the conservative piece of the self force in detection and parameter estimation. We also show the effect of the conservative piece of the self force on gauge invariant quantities, specifically $u^t$ as a function of the angular frequency $\\Omega$. 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. ", "introduction": " ", "conclusions": "" }, "1207/1207.3074_arXiv.txt": { "abstract": "{Understanding the relation between the star formation rate (SFR) and stellar mass (\\mstar) of galaxies, and the evolution of the specific star formation rate (sSFR=SFR/\\mstar) .} {We examine the dependence of derived physical parameters of distant Lyman break galaxies (LBGs) on the assumed star formation histories (SFHs), their implications on the SFR--mass relation, and we propose observational tests to better constrain these quantities.} {We use our SED-fitting tool including the effects of nebular emission to analyze a large sample of LBGs from redshift $z \\sim 3$ to 6, assuming five different star formation histories, extending thereby our first analysis of this sample (de Barros et al.\\ 2012, paper I). In addition we predict the IR luminosities consistently with the SED fits. } {Compared to ``standard\" SED fits assuming constant SFR and neglecting nebular lines, models assuming variable SFHs yield systematically lower stellar masses, higher extinction, higher SFR, higher IR luminosities, and a wider range of equivalent widths for optical emission lines. Exponentially declining and delayed SFHs yield basically identical results. Exponentially rising SFHs yield similar masses, but somewhat higher extinction than exponentially declining ones. We find significant deviations between the derived SFR and IR luminosity from the commonly used SFR(IR) or SFR(IR+UV) calibration, due to differences in the SFHs and ages. Models with variable SFHs, favored statistically, yield generally a large scatter in the SFR--mass relation. We show the dependence of this scatter on assumptions of the SFH, the introduction of an age prior, and on the extinction law. We show that the true scatter in the SFR--mass relation can be significantly larger than inferred using SFR(UV) and/or SFR(IR), if the true star formation histories are variable and relatively young populations are present. We show that different SFHs, and hence differences in the derived SFR--mass relation and in the specific star formation rates, can be tested/constrained observationally with future IR observations with ALMA. Measurement of emission lines, such as \\ha\\ and \\Oii, can also provide useful constraints on the SED models, and hence test the predicted physical parameters. } {Our findings of a large scatter in the SFR--mass relation at high-z and an increase of the specific star formation rate above $z \\ga 3$ (paper I) can be tested observationally. Consistent analysis of all the observables including the rest-frame UV to IR and emission lines are required to establish more precisely true SFR values and the scatter in the SFR--mass relation. } ", "introduction": "One of the important, recent results of multi-wavelength galaxy surveys is the finding of a well-defined relation between the star formation rate (SFR) and stellar mass of galaxies -- now often called the main sequence of star forming galaxies -- and the evolution of this relation with redshift, at least from the nearby Universe up to redshift $z \\sim 2$ \\citep{noeskeetal2007,daddietal2007,elbazetal2007,rodighieroetal2011}. The observed redshift evolution corresponds to an increase of the typical specific star formation rate (sSFR) by more than an order of magnitude from $z \\sim 0$ to 2, indicating a higher star formation activity of galaxies in the past. Recent analysis of Lyman break galaxies (LBGs), combining often HST and Spitzer observations, have tried to examine whether a SFR--mass relation is already in place at higher redshift, and how the specific star formation rate of galaxies behaves at $z \\sim$ 4--7 \\citep[e.g.][]{stark09,schaerer&debarros2010,labbeetal2010,mclureetal2011,gonzalezetal2011}. The first studies have found a relation similar to the $z \\sim 2$ SFR--mass relation, which would indicate no evolution of the sSFR beyond redshift $\\ga 2$ -- a plateau in sSFR -- \\citep[e.g.][]{stark09,gonzalezetal2011}. This result appears difficult to reconcile with most theoretical models, which successfully explain the behavior of star formation properties from $z \\sim$ 0 to 2, but predict a continuing rise of the specific star formation rate towards higher redshift \\citep{boucheetal2010,duttonetal2010,weinmannetal2011,daveetal2011}, although others propose alternatives \\citep[e.g.]{krumholz&dekel2012}. More recent work, including different spectral energy distribution (SED) models, or allowing for a revision of dust attenuation have shown that the observed sSFR at $z \\sim$ 5--7 could be higher than previously thought \\citep{schaerer&debarros2010,debarros2011,bouwens2011_beta,yabeetal2009,dBSS12}. If all galaxies at $z \\ga 2$ obey a mass--SFR relation with star formation rate increasing with galaxy mass and its normalization remains basically unchanged with redshift, it is evident that their star formation rate must increases with time. The apparent, small scatter in the SFR--mass relation, and the value of the exponent $\\alpha$ close to unity in this relation SFR $\\propto \\mstar^\\alpha$ found by various studies \\citep[e.g.][]{daddietal2007,elbazetal2007, gonzalezetal2011} has led several authors to conclude that high redshift galaxies must have gone through a phase of quasi-exponential growth \\citep[e.g.][]{renzini2009,marastonetal2010}. Earlier studies have independently advocated rising star formation histories (SFHs) for high redshift galaxies from hydrodynamic simulations and semi-analytical models \\citep{finlatoretal2007,finlatoretal2010,finlatoretal2011}. Arguments in favor of rising SFHs, at least on average, have been put forward by \\citet{finkelsteinetal2010,papovichetal2011} to explain the evolution of the UV luminosity function. If representative of the star formation history of individual galaxies, such star formation histories would indeed by quite different from exponentially declining or constant SFHs, most frequently assumed in the literature to model/analyze the observed SED of distant galaxies. Assumptions on the star formation history affect the physical parameters of galaxies derived from SED fits. This is therefore generally treated as a free parameter. See e.g.\\ \\citet{reddy2010,marastonetal2010,wuytsetal2011,reddy2012} for examples on $z \\ga 2$ star forming galaxies. Determinining star formation histories and their associated timescales is also important, as this may provide constraints on different modes of star formation and on feedback processes at high redshift, recognized as key features for our understanding of galaxy evolution. For example, star formation may proceed on different timescales if regulated by cold-accretion or by star formation feedback or triggered by mergers \\citep[cf.][]{khochfar&silk2011,wyithe&loeb2011}. These points already illustrate the interest and importance of determining the star formation histories and typical timescales of star formation at high $z$. We here wish to discuss and reexamine these issues on the basis of an analysis of a large sample of LBGs covering redshifts from $\\sim 3$ to 6, and to present observational tests which should be able to distinguish different star formation histories. One of main arguments often invoked to argue for rising star formation histories is the small scatter of the SFR--mass relation. However, it should be recognized that generally the SFR is derived from observables -- the UV or IR luminosity -- which depend on relatively long ($\\ga$ 100 Myr) timescales, and which are assumed to be at an equilibrium value, which is only reached after this timescale and for constant SFR. With these assumptions entering e.g.\\ the commonly used SFR(UV) or SFR(IR) calibrations or \\citet{kennicutt1998}, it is natural that the ``observational\" scatter is smaller than the true scatter in the current SFR, if typical ages are less than 100 Myr and/or the timescale shorter than this. Especially at high redshift, where timescales are shorter \\citep[e.g.\\ the dynamical timescale decreases with $(1+z)^{-3/2}$, cf.][]{wyithe&loeb2011} one should therefore carefully (re)examine the SFR--mass relation and its ``tightness\" using consistent diagnostics, and with all the available observational constraints including on age, star formation timescales and histories. In any case, the conclusion that rising SFHs are favored is {\\em a priori} inconsistent with the assumptions on the SFH for the determination of the SFR--mass relation and its scatter, where most studies simply assume a constant star formation rate. Although this has been re-examined for $z\\sim 2$ samples \\citep{marastonetal2010,wuytsetal2011,reddy2012} this obviously calls for a revision at higher redshifts. Most SED studies of LBGs at $z \\ga 3$ have assumed constant star formation rates, or exponentially declining SFHs to determine the physical parameters of these galaxies \\citep[e.g.][]{egamietal2005,SP05,eylesetal2005,vermaetal2007,yabeetal2009,stark09,leeetal2010,schaerer&debarros2010,gonzalezetal2011}, some of them imposing no dust attenuation, motivated by the blue UV slopes observed at high redshift. Notable examples are the work of \\citet{finlatoretal2007} who analyzed 6 galaxies at $z>5.5$ with different SFHs, including rising ones taken from they hydro-dynamic simulations. The physical parameters they derive are consistent with those using simple parametrised histories, albeit with reduced uncertainties. Their study also shows that SED fits with rising star formation histories yield a higher attenuation than inferred assuming constant SFR, a result not yet appreciated enough, which we also find in \\cite{dBSS12} and in this paper. Most recently, other groups have also analyzed high-$z$ samples with rising SFHs \\citep[e.g.][]{Curtis-Lake2012,Gonzalez2012}. However, dust extinction is, for example, not treated consistently with the star formation history in the approach of \\cite{Gonzalez2012}. Another drawback of most earlier studies is that the contribution of nebular emission (most emission lines) is not taken into account, an effect which can significantly alter the ages, masses, and other physical parameters derived from SED fits, as demonstrated by \\cite{schaerer&debarros2009, schaerer&debarros2010}. Indeed, there is now clear evidence for the presence of nebular lines affecting the broad-band photometry of high-$z$ star-forming galaxies (Lyman alpha emitters and LBGs in particular), as discussed e.g.\\ in \\citet{schaerer&debarros2011}. The best demonstration comes from LBGs with spectroscopic redshifts between 3.8 and 5, among which a large fraction shows a clear excess in the 3.6 \\micron\\ filter with respect to neighboring filters (K and 4.5 \\micron), as shown by \\cite{shimetal2011}. At these redshifts \\ha\\ is located in the 3.6 \\micron\\ filter, whereas very few lines are expected at 4.5 \\micron, showing that the observed 3.6 \\micron\\ excess is naturally explained by strong \\ha\\ emission. Given these limitations of published SED studies and the interesting results obtained from our work on a small sample of $z \\sim$ 6--8 LBGs \\citep{schaerer&debarros2010}, we have recently undertaken an extensive study of the physical parameters of a large sample of $\\sim 1400$ LBGs at $z \\sim$ 3--6, using our state-of-the art photometric redshift and SED fitting model including nebular emission, and considering a range of different star formation histories. Among the numerous detailed results obtained in this work \\citep[][hereafter dBSS12]{dBSS12} we find in particular from the preferred models: 1) a large scatter in the SFR--mass relation for the preferred, variable star formation histories, 2) higher dust attenuation than obtained from models assuming SFR=const and from standard methods using the UV slope, and 3) a higher sSFR than commonly obtained, and a rising sSFR with redshift. Our models therefore reconcile the observationally-inferred specific star formation rate with predictions from cosmological models predicting a continuous rise of the sSFR with redshift \\citep{boucheetal2010,duttonetal2010,weinmannetal2011,daveetal2011,krumholz&dekel2012}. The results from dBSS12 have also other important implications, e.g.\\ on the typical star formation timescales at high redshift, and on the cosmic star formation rate density. Given the importance of these findings, it is of interest to carry out additional independent tests of the models and to provide further constrains on the ages, extinction, and star formation histories of high redshift galaxies. In this paper we present predictions allowing such tests. Furthermore, we extend the study of dBSS12 by exploring other star formation histories, not considered in our previous paper. Indeed, while dBSS12 adopted the fixed, rising star history of \\citet{finlatoretal2011} obtained from the hydrodynamic simulations, we here explore exponentially rising SFHs and so-called ``delayed\" histories with variable timescales. These histories were previously applied to the SED fits of galaxies at lower redshift, e.g.\\ at $z \\sim 2$ by \\citet{marastonetal2010,wuytsetal2011,reddy2012} and found to be preferred over other simple star formation histories. First, we examine the effect of the different SFHs on the derived physical parameters of LBGs. We then present the implications these different model assumptions have on the SFR--mass relation. After that we present consistent predictions for the infrared luminosity of these galaxies, based on their SED fits and the assumption of energy conservation, abandoning the standard SFR--\\lir\\ conversions. We show in particular that such a consistent prediction yields a smaller scatter in the observables (\\lir) than in the current SFR. The true star formation rate of LBGs at $z \\ga 3$ could thus very well show a large scatter around a ``main sequence\" even if the UV and/or IR luminosities show a small scatter. Our models also show that star formation histories can, to some extend, be distinguished from measurement of IR luminosities, which will become possible with ALMA, since different amounts of UV attenuation are expected for different SFHs. We also present the predicted strength of some selected emission lines, which can be used to test our models and the different star formation histories. Our paper is structured as follows. The observational data and the method used for SED modelling are described in Sect.\\ \\ref{s_models}. The dependence of the physical parameters on the assumed star formation histories is discussed in Sect.\\ \\ref{s_phys}. Our general predictions for the IR emission are presented in Sect.\\ \\ref{s_ir}. Specific predictions for $z \\sim 4-6$ LBGs and ALMA are given in Sect.\\ \\ref{s_alma}. The predicted strengths of optical emission lines are shown in Sect.\\ \\ref{s_lines}. We discuss our results in Sect.\\ \\ref{s_discuss}, and Sect.\\ \\ref{s_conclude} summarises our main conclusions. We adopt a $\\Lambda$-CDM cosmological model with $H_{0}$=70 km s$^{-1}$ Mpc$^{-1}$, $\\Omega_{m}$=0.3 and $\\Omega_{\\Lambda}$=0.7. ", "conclusions": "\\label{s_conclude} Following up on our earlier detailed study \\citep{dBSS12} of a large sample of LBGs from redshift $z \\sim$ 3 to 6 located in the GOODS-South field, using for the first time an SED fitting tool including the effects of nebular emission on the synthetic photometry, we have examined the impact of different star formation histories (SFHs) on the derived physical parameters of these galaxies, on the SFR--mass relation, on different SFR indicators (UV and IR), on the expected dust extinction and the corresponding IR luminosity, and on the expected strengths of emission lines such as \\ha\\ and \\Oii. To do so, we have carried out SED fits for five different SFHs including exponentially rising and so-called delayed SFHs, plus the three histories already considered in dBSS12 (see Table \\ref{t_sfh} and Fig.\\ \\ref{fig_sfh}). Metallicity is also treated as a free parameter, and we have examined the effect of two different extinction/attenuation laws (Calzetti and SMC). The usual physical parameters, stellar mass, SFR, age, and attenuation are derived from the SED fits to the broad band photometry reaching from the U band to 8 \\micron, using Monte Carlo simulations to derive their median values (and the detailed probability distribution function, generally not discussed here). We have also computed consistently the predicted IR luminosities, \\lir, for all galaxies, assuming energy-conservation, i.e.\\ that all the radiation absorbed by dust is reemitted in the IR. Finally, the predicted IR luminosities have been translated to flux predictions in various IR bands, assuming modified black body spectra. The \\lir\\ predictions allow us in particular to examine in a consistent way the effects of variable SFHs and ages on this observable quantity, showing thus significant departures from results assuming inconsistent SFR(IR) or SFR(UV) calibrations. Our main results concerning the impact of star formation histories on the physical parameters of LBGs, exemplified to a sample of 705 LBGs at $z \\sim 4$ (B-drop galaxies), can be summarized as follows (see Sect.\\ \\ref{s_phys}): \\begin{itemize} \\item Compared to commonly adopted SED fits assuming constant SFR, no nebular emission and an age prior of $t>50$ Myr, models with exponentially declining SFHs, nebular lines and no age constraint yield younger ages, lower stellar masses, higher current SFR, higher specific star formation rates (sSFR=SFR$/\\mstar$), and higher dust extinction ($A_V$), as already shown in \\citet{dBSS12}. Exponentially declining SFHs yield overall the best fits (in terms of \\ki2 ) for the majority of LBGs. Based on the available SED constraints it is, however, difficult to distinguish different SFHs, although various arguments (in particular the distribution of emission line strengths inferred from broad-band photometry for $ z\\sim$ 3.8--5 galaxies) favor clearly variable, i.e.\\ non constant, histories (cf.\\ dBSS12). \\item Assuming delayed star formation histories one obtains basically identical physical parameters (and fit qualities) as for exponentially declining SFHs. \\item Rising star formation histories with variable timescales imply generally a similar stellar masses, and comparable or somewhat higher dust extinction than models assuming declining SFHs. The latter leads to the highest star formation rates and to similar or higher IR luminosities as for declining histories. \\item Overall ``standard\" models assuming constant and neglecting lines predict systematically higher stellar masses, lower extinction, lower SFR, lower IR luminosities, and more narrow range of equivalent widths for optical emission lines than all the other star formation histories considered here. \\end{itemize} Combining these different physical parameters we obtain the following (Sect.\\ \\ref{s_ir}): \\begin{itemize} \\item We find significant deviations between the derived SFR and IR luminosity from the commonly used SFR(IR) or SFR(IR+UV) calibration of \\citet{kennicutt1998}. Such differences naturally arise, due to differences in the derived ages and in the adopted star formation histories. In most cases (i.e.\\ for most galaxies and SFHs) we find that the Kennicutt relation will underestimate the true, current SFR derived from the SED fits (Sect.\\ \\ref{s_iruv}). Consistent SED studies including also the IR are therefore necessary, if the SFHs and ages may differ from those assumed in standard SFR calibrations. \\item A large scatter is found in the SFR--mass relation for models with declining and delayed SFH and no age prior \\cite[cf.][]{dBSS12}. The same also hold for models with rising star formation histories. The scatter is reduced when a minimum age (e.g.\\ $t>50$ Myr) is adopted. Even in this case, models with rising SFHs can show a large scatter, since high SFRs are found to the high(er) extinction. \\item A large scatter in the SFR--mass relation does not necessarily imply the same scatter in the \\lir\\ (or ($\\lir+\\luv$)--mass relation and vice versa, when the IR luminosity is computed consistently from the chosen SED model (i.e.\\ accounting for age and SFH effects). The same also applies to SFR(UV)--mass diagrams where the UV luminosity and the corresponding standard SFR conversion is used. We suggest that the true scatter in the SFR--mass relation obtained in this way may indeed be underestimated, if the true star formation histories are variable on relatively short timescales. Indeed such SFH variations can reproduce more successfully features related to emission lines, such as the observed 3.6 \\micron\\ excess in $z \\sim$ 4--5 LBGs. They may also be more relevant to at higher redshift, where the dynamical timescales decrease with $(1+z)^{-3/2}$. \\end{itemize} Our consistent predictions of IR luminosities (and fluxes) show that different SFHs lead to significantly different amounts of reddening and hence to different IR/UV luminosity ratios. Measurements of IR luminosities of individual LBGs or statistical samples of such galaxies can be used to distinguish different SFHs, and hence also different specific SFRs predicted by such models. ALMA observations will thus be able to provide independent constraints on behavior of the sSFR at high redshift, and on the scatter in the SFR--mass relation. We show predictions for the IR luminosities of B-drop and i-drop galaxies for different star formation histories and as a function of UV magnitude. The typical/median \\lir\\ is predicted to be $\\lir \\sim 10^{10 \\ldots 11}$ \\lsun\\ for LBGs with absolute UV magnitudes of $M_{\\rm UV} \\sim -22$ to $-19$. Finally we also show the predicted strengths (equivalent widths) of the \\ha\\ and \\Oii\\ emission lines. Again, different star formation histories naturally lead to different EW distributions, which can in principle be used to constrain the SFHs. Our models predict on average higher equivalent widths in low mass galaxies, in agreement with currently available observations at $z<3$, and a clear anti-correlation of \\whalpha\\ with the specific SFR. Our predictions should in particular provide new tests using IR observations with ALMA and/or measurements of (rest-frame) optical emission lines to obtain a better insight on the star formation histories of high redshift LBGs, on the behaviour of the SFR--mass relations and on the evolution of the specific SFR with redshift." }, "1207/1207.3845_arXiv.txt": { "abstract": "We present a formulation for multigroup radiation hydrodynamics that is correct to order $O(v/c)$ using the comoving-frame approach and the flux-limited diffusion approximation. We describe a numerical algorithm for solving the system, implemented in the compressible astrophysics code, CASTRO. CASTRO uses an Eulerian grid with block-structured adaptive mesh refinement based on a nested hierarchy of logically-rectangular variable-sized grids with simultaneous refinement in both space and time. In our multigroup radiation solver, the system is split into three parts, one part that couples the radiation and fluid in a hyperbolic subsystem, another part that advects the radiation in frequency space, and a parabolic part that evolves radiation diffusion and source-sink terms. The hyperbolic subsystem and the frequency space advection are solved explicitly with high-order Godunov schemes, whereas the parabolic part is solved implicitly with a first-order backward Euler method. Our multigroup radiation solver works for both neutrino and photon radiation. ", "introduction": "\\label{sec:intro} Numerical simulations are a useful tool for many radiation hydrodynamic problems in astrophysics. However, numerical modeling of radiation hydrodynamic phenomena is very challenging for a number of reasons. To achieve a highly accurate description of the system usually requires a multidimensional treatment with high spatial resolution. It is also desirable to handle the radiation frequency variable properly, because radiation transport is a frequency-dependent process. Moreover, the numerical algorithms for radiation hydrodynamics must be stable and efficient. Core-collapse supernovae are such complex phenomena \\citep{Bethe90, KotakeST06, JankaLM07, Janka12} that neutrino radiation hydrodynamics simulations have been indispensable for helping understand the explosion mechanism. Various algorithms have been developed to solve neutrino radiation hydrodynamics equations. Because of the complexities involved in the problem and the limited computer resources available, various tradeoffs have to be made. Many of these algorithms are 1D \\citep[e.g.,][]{BowersWilson82, Bruenn85, MezzacappaBruenn93, BurrowsYP00, LiebendorferMM04, OConnorOtt10}, or have used the so-called ``ray-by-ray'' approach, which does not truly perform multi-dimensional transport \\citep[e.g.][]{BurrowsHF95, RamppJanka02, BurasRJK06, TakiwakiKS12}. The two-dimensional code of \\citet{Millerwm93} is a gray flux-limited diffusion (FLD) code that does not include order $O(v/c)$ terms, where $v$ is the characteristic velocity of the system. The 2D multigroup multiangle code of \\citet{LivneBW04} does not include all order $O(v/c)$ terms, and in particular it lacks direct coupling among radiation groups. The two-dimensional algorithm of \\citet{HubenyBurrows07} is based on a mixed-frame formulation of a multigroup two-moment system that is correct to order $O(v/c)$, but it remains to be implemented in a working code. \\citet{SwestyMyra09} have developed a fully coupled 2D multigroup FLD code based on a comoving frame formulation that is correct to order $O(v/c)$. Besides the grid-based methods, smoothed particle hydrodynamics (SPH) methods have also been applied to multi-dimensional neutrino radiation hydrodynamics simulations of core-collapse supernovae \\citep[e.g.,][]{HerantBH94, FryerWarren02, FryerRW06}, but the SPH scheme has not yet been extended to multigroup. Recently \\citet{AbdikamalovBO12} developed a Monte Carlo method for neutrino transport, but hydrodynamics and radiation transport are not yet coupled in the scheme. In this paper, we present a numerical algorithm for solving multi-dimensional multigroup photon and neutrino radiation hydrodynamics equations in the comoving frames that are correct to order $O(v/c)$. Radiation quantities that are of interest are the radiation energy density, radiation flux, and radiation pressure tensor. It is more accurate to evolve both the radiation energy density and the radiation flux especially when the radiation is highly anisotropic and optically thin. However, the two-moment approach is very expensive in terms of computer time and memory for multi-dimensional multigroup radiation hydrodynamics simulations. Thus, we have adopted the flux-limited diffusion approach \\citep{AlmeWilson73} for computational efficiency. In the FLD approach, only the radiation energy density needs to be evolved over time, and the radiation flux is derived through a diffusion approximation. Furthermore, we use an analytic closure for the radiation pressure tensor. We have adopted a multigroup method, in which the radiation frequency is discretized into multiple groups. The equations we solve are similar to those in \\citet{SwestyMyra09}. However, we have neglected inelastic scattering of neutrinos on matter, which is currently treated as elastic scattering. Since the total cross section for the latter at the neutrino energies near and outside the neutrinospheres are $\\sim 20$--100 times smaller than the cross sections for scattering off nucleons and nuclei, not including inelastic effects at this stage has a minor effect on the qualitative behavior of collapse, bounce, and shock dynamics. Nevertheless, the approach of \\citet{ThompsonBP03} to handle inelasticity and energy redistribution can be incorporated into our scheme. Another difference between our scheme and the scheme of \\citet{SwestyMyra09} is that we have developed a Godunov scheme in which radiation is fully coupled into a characteristic-based Riemann solver for the hyperbolic part of the system. Many multigroup neutrino transport codes use a fixed Eulerian grid. However, it is difficult to achieve sufficient resolution in multi-dimensional simulations with this simple approach. To efficiently achieve high resolution, we use an Eulerian grid with block-structured adaptive mesh refinement (AMR). The AMR algorithm we use has a nested hierarchy of logically-rectangular variable-sized grids and refines simultaneously in both space and time. AMR techniques have been successfully used in multi-dimensional multigroup photon radiation transport \\citep{RAGE,crash}. However, to the best of our knowledge, this is the first time that AMR has been employed in a multi-dimensional multigroup neutrino solver. Besides AMR, the Arbitrary Lagrange Eulerian (ALE) method is an alternative approach for mesh refinement. We note that ALE has been successfully employed in multi-dimensional multigroup neutrino transport \\citep{LivneBW04,BurrowsLD07,OttBDL08} and photon transport \\citep{HYDRA}. We have implemented our algorithm in the compressible astrophysics code, CASTRO. This is the third paper in a series on the CASTRO code and its numerical algorithms. In our previous papers, we describe our treatment of hydrodynamics, including gravity and nuclear reactions \\citep[][henceforth Paper I]{CASTRO}, and our algorithm for flux-limited gray radiation hydrodynamics based on a mixed-frame formulation \\citep[][henceforth Paper II]{CASTRO2}. Here, we describe our algorithm for flux-limited multigroup radiation hydrodynamics based on a comoving-frame formulation. In \\S~\\ref{sec:RHD}, we present the multigroup radiation hydrodynamics equations that CASTRO solves. We show that the system can be divided into a hyperbolic subsystem (\\S~\\ref{sec:hyper}), a subsystem of advection in frequency space (\\S~\\ref{sec:fspace}), and a parabolic subsystem for radiation diffusion (\\S~\\ref{sec:para}). Analytic results for the mathematical characteristics of the hyperbolic subsystem are also presented in \\S~\\ref{sec:hyper}. The hyperbolic and the frequency space advection steps are solved by an explicit solver, whereas the parabolic subsystems is solved by an iterative implicit solver. We describe the explicit solvers in \\S~\\ref{sec:explicit}, followed by a description of the implicit solver in \\S~\\ref{sec:implicit}. We also discuss acceleration schemes that can greatly improve the convergence rate of the iterative implicit solver (\\S~\\ref{sec:accel}). In \\S~\\ref{sec:tests}, we present results from a series of test problems. A scaling test is shown in \\S~\\ref{sec:performance} to demonstrate the scaling behavior of the multigroup group radiation solver. Finally, the results of the paper are summarized in \\S~\\ref{sec:sum}. ", "conclusions": "\\label{sec:sum} In this paper, we have presented a multi-dimensional multigroup radiation hydrodynamics solver that is part of the CASTRO code. Block-structured AMR is utilized in CASTRO. In this paper, we focus attention on the single-level algorithms because the AMR algorithms have been presented in Papers I, II, and references therein. Our multigroup radiation hydrodynamics solver is based on a comoving frame formulation that is correct to order $O(v/c)$, and uses a FLD approximation. Our mathematical analysis shows that the system we solve contains a hyperbolic subsystem that is the usual hydrodynamics system modified by radiation, a frequency space advection part that accounts for the Doppler shift of radiation due to the motion of matter, and a parabolic diffusion part. We have presented the mathematical characteristics of the hyperbolic subsystem. The eigenvalues and eigenvectors we have obtained are useful for characteristic-based Godunov schemes. We also described our treatment of the frequency space advection part. An implicit solver involving two nested iterations is presented. We have also presented two acceleration schemes that improve the convergence rate of the implicit solver. Extensive testing is presented demonstrating the accuracy and stability of CASTRO to solve multigroup radiation hydrodynamics problems. We have presented a series of photon radiation test problems that cover a wide range of regimes. In addition to photon radiation problems, we have demonstrated the applicability of CASTRO in core-collapse supernova simulations. Recently, \\citet{Lentz_etal12} have argued, via a series of 1D simulations, that general relativistic effects and inelastic scattering of neutrinos on matter have significant effects on the numerical modeling of core-collapse supernova explosions. In contrast, \\citet{NordhausBAB10} have concluded that spatial dimensionality has far more impact than general relativistic effects or microphysics such as inelastic scattering. Our 1D simulations are consistent with the assessment of \\citet{Lentz_etal12} on the effects of inelastic scattering. However, our 2D simulation supports the assessment of \\citet{NordhausBAB10} on the importance of spatial dimensionality. Our simulations have shown that multidimensional effects (e.g., convective motion in the region behind the shock and exterior to the nascent neutron star) appear less than 10 milliseconds after bounce, and 2D results are profoundly different from 1D results. Thus, one should be cautious about assessment of various effects such as inelastic scattering that are based on 1D simulations that last several hundred milliseconds past bounce like the ones presented in \\citet{Lentz_etal12}. The main advantages of the CASTRO code are the efficiency due to the use of AMR and the accuracy due to the coupling of radiation force into the Riemann solver. Three-dimensional multigroup neutrino radiation hydrodynamics simulations of core-collapse supernovae with a good resolution using CASTRO can be carried out, if not now, in the near future with faster computers. In the strong scaling study with a total of 36 radiation groups and a cell size of $0.6\\:\\mathrm{km}$ on the finest level (Section~\\ref{sec:performance}), it took about 0.8 hours to advance one coarse time step (about $0.1\\:\\mathrm{ms}$) on 12288 cores of Hopper. Thus, it would take 800 hours on 12288 cores of Hopper to evolve to $100\\:\\mathrm{ms}$. Admittedly, it is still very expensive to perform long-duration 3D simulations, but it is possible to study the first $100\\:\\mathrm{ms}$ after bounce and investigate the development of turbulent convection region via 3D simulations." }, "1207/1207.5464_arXiv.txt": { "abstract": "A dataset of 126,501 spiral galaxies taken from Sloan Digital Sky Survey was used to analyze the large-scale galaxy handedness in different regions of the local universe. The analysis was automated by using a transformation of the galaxy images to their radial intensity plots, which allows automatic analysis of the galaxy spin and can therefore be used to analyze a large galaxy dataset. The results show that the local universe (z$<$0.3) is not isotropic in terms of galaxy spin, with probability $P<\\sim5.8\\cdot10^{-6}$ of such asymmetry to occur by chance. The handedness asymmetries exhibit an approximate cosine dependence, and the most likely dipole axis was found at RA=132$^o$, DEC=32$^o$ with 1$\\sigma$ error range of 107$^o$ to 179$^o$ for the RA. The probability of such axis to occur by chance is $P<1.95\\cdot10^{-5}$ . The amplitude of the handedness asymmetry reported in this paper is generally in agreement with Longo, but the statistical significance is improved by a factor of 40, and the direction of the axis disagrees somewhat. ", "introduction": "\\label{introduction} Assuming that the universe is isotropic, it is expected that in a sufficiently large sector of the universe the number of galaxies that rotate clockwise will be roughly equal to the number of galaxies that rotate counterclockwise. However, recent evidence suggests that the local universe does not follow that expected balance, and show that the ratio between clockwise and counterclockwise galaxies in some regions is significantly different than 1:1, introducing galaxy handedness asymmetry. Longo \\citep{Lon11} analyzed the sense of rotation of 15,158 spiral galaxies imaged by Sloan Digital Sky Survey (SDSS) and showed cosmic parity violation and a possible cosmic dipole axis. The analysis was based on a relatively small dataset of galaxies classified manually by four undergraduate students, so that the number of galaxies in each 30$^o$ Right Ascension (RA) slice ranged between 0 to 3512 galaxies. Redshift of the galaxies was smaller than 0.085. The experiment was done using a web-based user interface that mirrored half of the galaxy images, so that the analysis was not biased by certain preferences of the human readers who classified them. A dipole axis was detected with RA=217$^o$, DEC=32$^o$ with a probability of occurring by chance of 7.9$\\cdot10^{-4}$ \\citep{Lon11}. Asymmetry of clockwise and counterclockwise galaxies was also observed in the Galaxy Zoo dataset \\citep{Lin10}. However, since the classification was done manually by a static graphical user interface, the asymmetry could be attributed to certain preferences of the Galaxy Zoo participants who used the web-based system. The sense of rotation of a galaxy is a gross metrics that is relatively easy to measure, and it is not affected by atmospheric effects or hardware inaccuracy. However, analyzing large datasets of spiral galaxies manually requires significant labor, and might be biased by human errors or reader preferences. In this study a computer analysis was used to determine the handedness of spiral SDSS galaxies, and the results were analyzed in the light of the Cosmological Principle and the cosmological isotropy assumption. ", "conclusions": "\\label{conclusion} The results of the experiment show that galaxy handedness asymmetry is different in different RA ranges within z$<$0.3, and the asymmetries exhibit a cosmological dipole axis. According to these observations, the local universe is not isotropic, meaning that the observed physical characteristics of the universe are different in different directions of observation. These results are in agreement with \\citep{Lon11}. The most likely dipole axis was found at RA=132$^o$, DEC=32$^o$. This axis is somewhat different from the most likely dipole axis reported by \\citep{Lon11} at RA=217$^o$, DEC=32$^o$ with an uncertainty of $\\sim$35$^o$, but the difference is within the error ranges. The limited range of declination in the SDSS dataset as well as the absence of data in certain right ascension ranges did not allow comparing all opposite hemispheres for handedness asymmetry for the purpose of fully profiling the characteristics of the handedness asymmetry in the local universe. It should be noted that the experiment is limited to the local universe, although the range of z$<$0.3 exceeds far beyond the scale of a galaxy supercluster. Higher-resolution analysis of the cosmic parity violation in the more distant universe will be possible when the Large Synoptic Sky Survey (LSST) starts operating, providing a much larger galaxy dataset and seeing deeper space. The experiment is based solely on data acquired by Sloan Digital Sky Survey, and under the assumption that these raw data are not biased. A possible source of bias might be the automatic galaxy detection algorithm that is part of the SDSS pipeline, that might have a preference to galaxies of some certain directionality. This, however, conflicts with the observation that some of the opposite hemispheres have opposite asymmetries. In any case, even in the case that such bias exists, it is expected to be consistent for all SDSS data, so that the asymmetry values might be shifted, but the asymmetry profile itself should not change. Another possible weakness is that galaxies from different RA were taken at different times, and any change made to SDSS or its processing algorithms during that time could potentially lead to a biased dataset, and different galaxy asymmetries at different redshifts. It should also be noted that the SpecObj view and the Galaxy Zoo dataset contain bright objects, and are not random samples of the galaxies in DR7. A bias in the selection of these objects can also affect the results. Clearly, this study does not offer explanations neither to the nature of asymmetry at the cosmological scale, nor to the observation that the asymmetry changes at different regions. %" }, "1207/1207.2558_arXiv.txt": { "abstract": "Cosmological parameters from WMAP 7 year data are re-analyzed by substituting a pixel-based likelihood estimator to the one delivered publicly by the WMAP team. Our pixel based estimator handles exactly intensity and polarization in a joint manner, allowing to use low-resolution maps and noise covariance matrices in $T,Q,U$ at the same resolution, which in this work is 3.6$^\\circ$. We describe the features and the performances of the code implementing our pixel-based likelihood estimator. We perform a battery of tests on the application of our pixel based likelihood routine to WMAP publicly available low resolution foreground cleaned products, in combination with the WMAP high-$\\ell$ likelihood, reporting the differences on cosmological parameters evaluated by the full WMAP likelihood public package. The differences are not only due to the treatment of polarization, but also to the marginalization over monopole and dipole uncertainties present in the WMAP pixel likelihood code for temperature. The credible central value for the cosmological parameters change below the 1 $\\sigma$ level with respect to the evaluation by the full WMAP 7 year likelihood code, with the largest difference in a shift to smaller values of the scalar spectral index $n_S$. ", "introduction": "The anisotropy pattern of the cosmic microwave background (CMB) is a treasure for understanding the costituents of our Universe and how it evolved from the Big Bang. Under the assumption of isotropy and Gaussianity of CMB fluctuations, the power spectra of intensity and polarization anisotropies include all the compressed information on our Universe through the determination of the cosmological parameters. There has been a tremendous improvement in the estimate of cosmological parameters driven by the increasingly better quality of CMB data, mainly due to the full sky observations in temperature and polarization by the Wilkinson Microwave Anisotropy Probe (WMAP), (see \\cite{Larson:2010gs,Komatsu:2010fb} and references therein) and to the small angular scales measurements by QUaD in polarization \\citep{QUAD}, by the South Pole Telescope \\citep{SPT,SPT2011,SPT2011_2} and the Atacama Cosmology Telescope \\citep{ACT,ACT2010} in temperature. {\\sc Planck} will lead to a drastic improvement of CMB full sky maps in temperature and polarization, leading to an eagerly expected improvement in cosmological parameters with uncertainties at the percent level \\citep{bluebook}. A joint likelihood analysis in temperature and polarization is one of the accepted methods in securing the scientific expectations of observational achievements in terms of cosmological parameters. Although the likelihood could be written exactly in the map domain under the Gaussian hypothesis, its computation is almost prohibitive already at the resolution of 2 degrees, whereas cosmological information is encoded in the temperature and polarization power spectra up to the angular scales of the order of few arcminutes, where the Silk damping suppress the CMB primary anisotropy spectrum. It is now commonly accepted to use an hybrid approach which combines a pixel approach at low resolution with an approximated likelihood based on power spectrum estimates at high multipoles (see \\cite{BJK,verdeetal,HL} for some of these approximations). Since the three year release of the full polarization information, the WMAP team adopted such a hybrid scheme approach, which has been suggested independently in \\cite{Efstathiou04,Slosar04,ODwyer04,Efstathiou06}. At a first appearance of the three year data, the WMAP team adopted a pixel approach on HEALPIX \\citep{gorski} resolution $N_{\\rm side}=8$ \\footnote{The number of pixels in a map is given by $N_{\\rm pix} = 12 N_{\\rm side}^2$, i.e. 768 for $N_{\\rm side}=8$ and 3072 for $N_{\\rm side}=16$.} temperature and polarization maps, and considered the high-$\\ell$ approximated likelihood to start at $\\ell =13$ in temperature and $\\ell=24$ in polarization and temperature-polarization cross-correlation for the determination of cosmological parameters in \\cite{Spergel:2006hy}. The WMAP team treats separately temperature and polarization as explained in \\cite{Page:2006hz} and \\cite{Hinshaw:2006ia}, by using the approximation that the noise in temperature is negligible. As a consequence, the WMAP likelihood code includes either $(Q,U)$ and the temperature-polarization cross-correlation in the same sub-matrix. It was then shown by \\cite{Eriksenetal2007} that by increasing the resolution of the temperature map to HEALPIX $N_{\\rm side}=16$ and therefore the multipole of transition to high-$\\ell$ approximated likelihood in temperature from $\\ell=12$ to $\\ell=30$, the mean value for the scalar spectral index $n_s$ shifted to higher values by a 0.4 $\\sigma$. The asymmetric handling of the low-resolution temperature map at $N_{\\rm side}=16$ and polarization at $N_{\\rm side}=8$, became the final treatment of the three year data release. This low-$\\ell$ likelihood aspect in the WMAP hybrid approach has not changed since the final release of the WMAP 3 year data to the current WMAP 7 year one. In this paper we wish to perform an alternative determination of the cosmological parameters from WMAP 7 public data, substituting the WMAP low-$\\ell$ likelihood approach with a pixel based likelihood code which treats $T,Q,U$ at the same HEALPIX resolution $N_{\\rm side}=16$ connected to the standard WMAP high-$\\ell$ package. In this analysis we therefore increase the resolution of polarization products digested by the pixel base likelihood from $N_{\\rm side}=8$ to $N_{\\rm side}=16$, in analogy with what done by \\cite{Eriksenetal2007} for temperature only. The WMAP 7 year foreground cleaned $(Q,U)$ maps, covariance matrices and masks at the resolution $N_{\\rm side}=16$ are also publicly available at http://lambda.gsfc.nasa.gov: therefore, all data used in this paper are made available by the WMAP team. The paper is organized as follows. In Section II we briefly describe the WMAP hybrid approach to the likelihood, with particular care to the low multipole part. In Section III we describe our pixel approach, implemented in the BoPix code. We then present in Section IV the cosmological parameters obtained by using our alternative pixel approach in place of the WMAP one for a $\\Lambda$CDM scenario. In Section V we extend our investigations to other cosmological models. In Section VI we draw our conclusions. ", "conclusions": "\\label{conclusions} We have performed an alternative estimate of the cosmological parameters from WMAP 7 year public data, by substituting the WMAP 7 low-$\\ell$ likelihood with a pixel likelihood code which treats $(T,Q,U)$ at the same resolution without any approximation. We have used this code at the HEALPIX resolution $N_{\\rm side}=16$ on foreground cleaned public data, therefore increasing the resolution of the pixel based polarization products used in our extraction of the cosmological parameters with respect to the WMAP standard one. We have consistently increased the transition multipole from $\\ell=24$ to $\\ell=31$ for the high-$\\ell$ WMAP 7 year temperature-polarization cross-correlation likelihood and included the marginalization over the nuisance parameter $A_{\\rm SZ}$. With this setting we have found estimates for the cosmological parameters consistent with those obtained by the full WMAP 7 year likelihood package, although for some parameters the differences are of half $\\sigma$ or more. These differences between the two low-$\\ell$ likelihood treatments we find are larger than the WMAP 7 yr likelihood uncertainties from tests on simulations reported in \\cite{Larson:2010gs}; however, we need to keep in mind that our differences between two likelihood treatments are reported for {\\em real} data, with WMAP 7 year beam/points source corrections and various marginalizations taken fully into account, differently from the simulation analysis performed in \\cite{Larson:2010gs}. The difference between the two best-fit $C_\\ell^{TT}$ for $\\Lambda$CDM found by the two alternative likelihood treatments show a maximum of $4\\%$ around at $\\ell \\sim 10$ and oscillate with an amplitude below $1\\%$ for $\\ell > 100$ \\footnote{We have checked that either the difference between the two best-fit $C_\\ell$ or between the estimates of the cosmological parameters decrease when the nuisance parameter $A_{\\rm SZ}$ is set to zero in both alternative likelihood treatments. The net effect of the variation of this foreground parameter, which is unconstrained by the data, is to increase the differences between the estimates of the cosmological parameters from the two likelihood treatments for the $\\Lambda$CDM model.}. A $5\\%$ percent difference is found in the two best-fits for the lensing power spectrum, whereas smaller differences are found for temperature-polarization cross-correlation and polarization power spectra. We have shown how part of the discrepancy, but not all, can be ascribed to the monopole/dipole marginalization used in the WMAP temperature likelihood and described in \\cite{Slosar04}. On restricting to the $\\Lambda$CDM model the most important difference is for the scalar spectral index $n_S$, which decrease to 0.956 from the value 0.968 we obtain with the full WMAP 7 yr likelihood code, i.e. a decrease of 0.86 $\\sigma$. This different value for $n_S$ would increase the evidence against the Harrison-Zeldovich spectrum from WMAP 7 yr data. This difference for $n_S$ is consistent with the one between the two best-fit $C_\\ell$ and depend only partially from the threshold multipole from which the high-$\\ell$ TE likelihood starts. Other previous alternative likelihood treatments also reported the most important discrepancy for the scalar spectral index \\citep{Eriksenetal2007,Rudiordetal2007a}. A smaller value for $n_S$ with respect to the estimate by the full WMAP 7 year likelihood code, always within 1 $\\sigma$, is then seen in all the extension of $\\Lambda$CDM considered here. No major changes are found for the 95 \\% credible intervals for the tensor to scalar ratio and for the running of the scalar spectral index. A slight degradation has been found for the 95 \\% credible interval on the neutrino mass. The case of cosmological birefringence has been taken as a sensitive test for the two alternative likelihoods, whose most relevant difference is the treatment of polarization on large scales. A slight difference on the posterior of the polarization angle $\\alpha$ has been found when only low resolution data are used, whereas the results are fully consistent when the high-$\\ell$ TB data are added to both likelihoods." }, "1207/1207.7182_arXiv.txt": { "abstract": " ", "introduction": "The solar atmosphere is host to a variety of highly energetic events such as coronal mass ejections (CMEs), flares, prominences, and coronal heating. The huge reservoir of energy in the solar atmosphere is related to the magnetic field originating in the convection zone, where shearing motion of field line foot points can stretch and twist the anchored field \\citep{P79}. A variety of magnetohydrodynamic (MHD) waves are generated in the corona due to convective motion in the photosphere. The convective shearing motion not only produces \\alf, fast or, slow magnetoacoustic waves, but also brings topologically disparate parts of the magnetic configuration closer together resulting in the formation of current sheets. All in all, a tiny fraction of convective energy carried by the waves to higher altitude may suffice to heat the coronal plasma to high temperatures. This simple physical picture of wave excitation, propagation and the ensuing coronal heating is attractive as it ties the heat transport in the solar corona and heliosphere to the ultimate source of energy -- shearing photospheric convective motions. Therefore the investigation of wave propagation and concomitant heating of coronal plasma in the framework of MHD has been a popular topic in solar physics \\citep{P87, P97a, P97b, P98, H99, GP04, A06}. Most solar atmospheric heating, with the possible exception of flares, takes place in the chromosphere. The chromosphere extends up to about nine pressure scale heights, i.e. about $2$\\,Mm above the photospheric surface. The lower chromosphere is threaded by strong ($\\sim$ kiloGauss) vertical flux tubes located in the network regions where they are observed as bright points. These tubes, which have a low filling factor ($< 1 \\%$) near the foot point in the photosphere, expand with increasing height to fill about $15 \\%$ of the lower chromosphere (altitude $\\sim 1 \\mbox{Mm}$, where CaII emission features are observed in H and K lines) before filling the entire atmosphere and forming a canopy in the chromosphere. The quiet solar internetwork region is also magnetised, with patches of concentrated kG magnetic field \\citep{D09, SA10} and an order of magnitude smaller field everywhere else. Observation suggests that localised active regions that emit $\\sim 80 \\%$ of the coronal radiative loss near solar maximum contain plasma of chromospheric origin \\citep{A01}. This raises the possibility that the same mechanism that transports mechanical energy from the convection zone to the chromosphere to sustain its heating rate also supplies the energy needed to heat the corona and accelerate the solar wind. However, the number of plasma particles in the partially ionised solar atmosphere, particularly below $\\lesssim 2.5 \\mbox{Mm}$ is small \\citep{VAL81}. As a result the plasma particles are not {\\it frozen} in the field owing to frequent collision with the neutral hydrogen \\citep{MS56}. Therefore, the ideal MHD description which is valid only for the fully ionised fluid is unsuitable to describe the low temperature photosphere-chromosphere where non--ideal MHD effects such as Hall, ambipolar and Ohm diffusion dominate \\citep{PW06, P08a, P08b, K12, SK12, PW12a}. All in all, there is no unified ideal MHD like framework to describe the weakly ionised and weakly magnetised photosphere and fully ionised and highly magnetised corona with weakly ionised and highly magnetised chromosphere sandwiched in-between. The inclusion of neutral dynamics not only destroys the economy and simplicity of the single fluid MHD description of fully ionised plasma but the very concept of a flux tube may be difficult to define in the multi--fluid framework. Furthermore, high frequency \\alf waves may not survive the collision-dominated photosphere and chromosphere \\citep{G04, LAK05, VPP07, AHL07, VPPD07, G11}. This could have been anticipated on the grounds that in the photosphere ($\\lesssim 500 \\,\\mbox{km}$) and chromosphere ($\\lesssim 2500 \\,\\mbox{km}$) the plasma number density is much smaller than the neutral number density and thus, high frequency (with respect to the ion-neutral collision frequency) MHD waves are severely damped by collisional dissipation of the wave energy. A way out of this difficulty is to retain the MHD momentum equation for the ionized component and include the effect of collisions with the neutrals by modifying the induction and energy equations to incorporate a conductivity tensor \\citep{E04, L06}, an approach often employed to study the lower ionosphere of the Earth. However, the derivation of the \\emph{time--independent} conductivity tensor neglects time derivatives in the electron and ion momentum equations, i.e. $d_{e\\,,i}/dt \\sim \\omega_{e, i} \\ll \\omega_{ce, ci}$, requiring that the dynamical response frequencies of the ions and electrons, $\\omega_{e, i}$, are much smaller than their respective gyro-frequencies, $\\omega_{ce, ci}$. Further, the relative ion--neutral drift is assumed much smaller than the centre of mass ion--neutral velocity \\citep{MK73}. Therefore, neglecting plasma inertia, a linear relationship between the electric field $\\E$ and plasma current $\\J$ can be easily derived $\\E = \\bmath{\\sigma}\\cdot \\J$ where $\\bmath{\\sigma}$ is the \\emph{time-independent} conductivity tensor. However, the MHD equation of motion assumes $\\omega_{i} \\sim \\omega_{ci}$. In fact the ion carries the inertia of the fluid. Therefore, on the one hand, the \\emph{time--independent} conductivity tensor implies $\\omega_{e, i} \\ll \\omega_{ce, ci}$, on the other hand, the MHD momentum equation requires $\\omega_{i} \\sim \\omega_{ci}$. Clearly, investigation of the collisional effects by merely modifying the induction and energy equation in the MHD framework is highly unsatisfactory. Any realistic model of the solar atmosphere must reflect two basic observational facts: (a) the magnetic field distribution on the solar surface is not continuous but is organised into network and internetwork elements. Whereas the network field ($\\gtrsim \\mbox{kG}$) is predominantly vertical and organised into flux tubes (diameter $\\lesssim 100\\,\\mbox{km}$) located in the intergranular lanes, the internetwork field $(\\mbox{few}\\,\\mbox{G}-\\mbox{kG})$ in the interior of supergranule cells is primarily horizontal \\citep{H09, L08} \\footnote{The internetwork field may instead be isotropic \\citep{SA11}, but see also \\cite{S12}.}; (b) the plasma in the photosphere--chromosphere is weakly ionised (with fractional ionisation, i.e.\\ the ratio of the electron and neutral number densities, $X_e = n_e / n_n \\sim 10^{-4}$, VAL81). Both the active and quiet phases of the solar atmosphere are highly dynamic and consist of convectively driven vortices and flows on a variety of spatial and temporal scales \\citep{B08, W09, Ba10, B10}. Most vortices are small ($\\lesssim 0.5\\,\\mbox{Mm}$) with average size $\\sim 241 \\pm 25\\,\\mbox{km}$ and typical lifetime $\\sim 3-5\\,\\mbox{min}$, although large vortices $\\sim 20\\,\\mbox{Mm}$ with lifetime $\\gtrsim 20\\,\\mbox{min}$ have also been observed \\citep{At09}. The bright points associated with the vortex motion in the intergranular lane moves with typical speed $\\lesssim 2\\,\\mbox{km} / \\mbox{s}$ \\citep{W09}. Magnetic fields appear to play a crucial role in mediating vortex motion in the photosphere and chromosphere \\citep{S12}. Rotation has been invoked in the past to explain the stability of flux tubes \\citep{S84}. Models of spicules also invoke rotating flux tubes \\citep{KS97}. Numerical simulations of solar convection display turbulent vortex flows at intergranular lanes \\citep{Z93, SN98}. Vorticity generation near the boundaries of granules has also been seen in numerical simulation of the photosphere \\citep{N09, M10}. The formation of small-scale, intergranular vortices suggests that vorticity is formed due to the interaction of photospheric plasma with the ambient magnetic field in intergranular lanes \\citep{M11, S11}. High-resolution simulations including the effects of non-ideal MHD show that the Hall effect generates out-of-plane velocity fields with maximum speed $\\sim 0.1\\,\\mbox{km} / \\mbox{s}$ at the interface layers between weakly magnetized light bridges and neighbouring strong field umbral regions \\citep{C12}. To summarise, both observation and numerical simulation points to the presence of shear flow at various spatial scales in the solar photosphere. Large scale shear flow acts as a source of free energy in the solar plasma that can easily destabilise waves. For example, shear driven Kelvin-Helmholtz instability (KHI), which converts shear flow energy into vortex kinetic energy, is invoked to explain the dynamical structures in the solar atmosphere \\citep{ J93, Kl04, Sr10, Z10, OF11, F11}. The KHI possibly acts as a triggering mechanism for large scale solar transient phenomena such as solar flares, CMEs, and associated eruptions \\citep{Sva12}. The generation of the highly dynamical structures observed by the Solar Dynamical Observatory (SDO) is most likely due to the KHI \\citep{F11}. Although the solar atmosphere may be susceptible to KHI, the presence of a magnetic field is not conducive to this instability. For example, a magnetic field directed along the flow suppresses KHI whereas a transverse field has no effect on the instability \\citep{C61}. However, magnetic fields not only quench shear instabilities but can also facilitate them. For example, the most important instability in accretion discs, the magnetorotational instability is caused by an interplay between the angular velocity of the magnetised fluid and magnetic field \\citep{BH98}. Thus, depending on the presence of a velocity gradient, magnetic field, when well coupled to the plasma can suppress as well as drive the instability. Therefore, before dwelling upon the role of the magnetic field on the flow driven instabilities, it is pertinent to know how well is the magnetic field coupled to the surrounding matter. Magnetic field drift through weakly ionised matter in the presence of shear flow can assist waves to grow. For example, in protoplanetary discs, both Hall and ambipolar diffusion enhance the magnetorotational instability \\citep{W99, BT01, KB04, D04, WS12, PW12b}. Clearly, the drift of the magnetic field in a weakly-ionised diffusive medium provides new pathways through which shear energy can be channelled to the waves. Indeed, diffusive destabilisation of a partially ionised medium in the presence of shear flow is not unique to weakly-ionised discs but is generic \\citep{K08, PW12a}. The crucial ingredients required to excite this diffusion--shear instability are the presence of a shear flow, and favourable magnetic field topology. The ensuing instability is overtly similar to KHI, and not surprisingly, the growth rate is proportional to the shear gradient. However, unlike KHI which is hydrodynamic in nature, this is a magnetohydrodynamic instability. A detailed investigation of diffusive shear instability in the context of the solar atmosphere is carried out here, building on our previous work (\\citealt{PW12a}, hereafter PW12a), as follows. First, unlike PW12a, where only a vertical field and transverse fluctuation (vertical wave vector) is assumed, here field topology is more general and the wave vector may be oblique. Second, in PW12a the back reaction of the fluid on the magnetic field was completely ignored, whereas here it is retained and as a result the shear driven diffusive instability does not have a cut-off wavelength. Third, for vertical fields and transverse fluctuations we showed in PW12a that only Hall diffusion assists the instability, with ambipolar and Ohm diffusion (which combine as Pedersen diffusion) only able to damp waves. In the present work for a more general field topology and oblique wave vector, we shall see that both ambipolar and Hall diffusion can assist the instability. Finally, the general stability criterion for a magnetic-diffusion-dominated plasma in the presence of shear flows is presented in this work. The paper is organised in the following fashion. The basic set of equations and dispersion relation are given in Sec.~2; in subsection 2.1 we give the linearised equations in terms of the diffusion tensor and derive the general dispersion relation. In Sec.~3 the general stability criterion is described and the maximum growth rate of the instability is derived. The expressions for the maximum growth rate and critical wavelength are given in limiting cases. In Sec.~4, various limiting cases of the dispersion relation are discussed. In Sec.~5 applications of the results are outlined. Finally, in Sec.~6 we give a brief summary of the main results. ", "conclusions": "The solar photosphere is threaded by a kilogauss magnetic field concentrated in vertical flux tubes ($\\mbox{radius} \\sim 100-200\\,\\mbox{km}$) at intergranular boundaries \\citep{H09}. Similar field strengths have also been observed in the quiet solar internetwork region \\citep{SA10}, although less frequently ($\\lesssim 40\\,\\%$). Outside these kilogauss patches, theinternetwork field is much weaker ($\\sim 100\\,\\mbox{G}$). Flux tubes are often modelled as non-rotating cylindrical tubes with plasma and the magnetic pressure balance providing the required stability. In the present work, because we are interested in short radial wavelengths we have approximated flux tubes by a planar sheet where the $x$ and $y$ coordinates locally correspond to the radial and azimuthal directions locally. Although the chromosphere has a limited extent in comparison to the corona, its net radiative loss $\\sim 10^{7}\\,\\mbox{erg}\\,\\mbox{cm}^{-2}\\,\\mbox{s}^{-1}$ is $10$ times larger. Further, except for flares, most solar atmospheric heating occurs in the chromosphere \\citep{A01, G01}. A strong correlation between the core emission of calcium K and H resonance lines and the quiet sun magnetic field \\citep{S89} suggests that the origin of chromospheric heating ($\\sim 10^7\\,\\mbox{ergs}\\,\\mbox{cm}^{-2}\\,\\mbox{s}^{-1}$) is magnetic. The magnetic-diffusion-driven shear instability proposed in the present work can provide a viable mechanism for the excess chromospheric heating as the crucial ingredients required to excite this instability -- shear flow and magnetic field -- are always present in the network--internetwork region. The attractive feature of this fast growing diffusive shear instability is that all wavelengths of fluctuations are likely to be excited as there is no cut-off wavelength. Thus, as we see from Fig.~(\\ref{fig:DF2}), the network field below $1\\,\\mbox{Mm}$ is likely to be destabilised by this instability. In the internetwork elements, this instability may operate in the entire photosphere--chromosphere. The only uncertainty involved is the lack of information about the scale of shear flow gradient which can not be resolved by current observations. However, small whirlpools with size similar to terrestrial hurricanes [$\\lesssim 0.5\\,\\mbox{Mm}$ with typical lifetime $\\sim 5\\,\\mbox{min.}\\,$ on the solar surface \\cite{B08}] suggest the presence of such flows. Long lasting large scale vortices at supergranular junctions with typical lifetime $\\sim 1-2\\,\\mbox{h}$ with enhanced CaII emission have also been observed \\citep{At09}. The swirl motion in the chromosphere has been recently detected by \\cite{W09}. Clearly, observations suggest the ubiquitous presence of flow gradients in the photospheric-chromospheric plasma. Past numerical simulations also indicated the presence of vortex flows in intergranular lanes \\citep{Z93, SN98}. The typical vorticity of a vortex is $\\sim 6\\times 10^{-3}\\,\\persec$ which corresponds to rotation period $\\sim 35\\,$ minutes \\citep{B10}. Thus it would appear that the Hall instability does not have time to develop since the growth rate ($\\avp / 2 = 3\\times 10^{-3}\\,\\persec$) is very small. However, above vorticity value is limited by the upper limit in the vorticity resolution [$\\sim 4\\times 10^{-2}\\,\\persec$, \\cite{B10}]. The numerical simulation gives much higher vorticity value ($\\sim 0.1-0.2\\,\\persec$) in the photosphere-lower chromosphere (Fig.~31, \\cite{SN98}). The growth rate corresponding to $\\avp = 0.2\\,\\persec$ is one minute. For almost magnetosonic waves ($\\kz \\rightarrow 0$), the maximum growth rate of the Hall diffusion driven shear instability may become quite small. The maximum growth rate in the ambipolar diffusion dominated middle and upper chromosphere will be only one fifth of the Hall diffusion dominated case for the maximum $g = 0.25$ and $\\mu = 0.5$. Therefore, vortex motions with typical lifetime $\\gtrsim 15\\,\\mbox{min.}$ will be susceptible to the ambipolar diffusion driven shear instability. Since, vortex motions of various spatial and temporal scales are observed, it is likely that non--ideal MHD effects will play an important role in exciting low frequency turbulence in the medium. \\begin{figure} \\includegraphics[scale=0.31 ]{F4.eps} \\includegraphics[scale=0.31 ]{F5.eps} \\caption{The growth rate and most unstable wavelength is shown for $1.2\\,\\mbox{kG}$ [Fig.~\\ref{fig:fr7}(a)] and $120\\,\\mbox{G}$ [Fig.~\\ref{fig:fr7}(b)] fields. Following parameters have been used in the above figure: $\\kz = 1\\,,\\Bz = 1$ (bold line), $\\kz = \\Bz = 0.9$ (dashed line) and $\\kz = 0.4\\,,\\Bz = 0.9$ (dotted line).} \\label{fig:fr7} \\end{figure} The \\emph{non-ideal} MHD description of the photosphere-chromosphere provides several pathways through which shear energy can be channelled to the waves by magnetic field. For example in the excessively diffusive limit when $\\dv / v \\ll \\dB / B$, diffusion in tandem with the shear flow can destabilise the network-internetwork field. For a purely vertical field and vertical wavevector this limit has been discussed in detail in PW12a. In order to compare with PW12a, we briefly describe the effect of the field topology and wave orientation in the highly diffusive limit. Comparing Figs.~\\ref{fig:DF2} and \\ref{fig:fr7} we conclude that the maximum growth rate is similar in both cases. When $\\Bz = \\kz = 1$, the instability grows at a maximum rate whereas with decreasing $\\kz$ or, $\\Bz$ the growth rate diminishes. In Fig.~\\ref{fig:fr7}(b) we plot most unstable wavelength against height for the same parameters as in Fig.~\\ref{fig:fr7}(a). The wavelength is normalized against scale height calculated self--consistently using model C, VAL81. For both kG and weaker fields $\\lambda_0$ fits well within a scale height and thus the instability will grow at a maximum rate in the entire photosphere--chromosphere. However, it is only in the photosphere and lower chromosphere ($\\lesssim 1\\,\\mbox{Mm}$) where the instability will grow at a maximum rate for a kG field. In the middle and upper chromosphere the growth rate tapers off and becomes one eighth of the shear frequency. Therefore, in the strong field region, diffusive instability will be efficient in destabilizing a magnetic element in the photosphere and lower chromosphere. When the field is weak ($\\sim 100\\,\\mbox{G}$), the instability can operate in the entire photosphere-chromosphere at a maximum rate. How does a nonlinear saturated state of the diffusive instability will look like? The answer to this can be given only by numerical simulations. However, if the nonlinear results of the protoplanetary discs and star forming regions are any guide then this instability should be quite efficient in exciting the low frequency turbulence and heating of the plasma. In fact the interplay between the vortex flow and magnetic diffusion could be responsible for the entire energy budget of the solar atmosphere. The energetic events such as hard x-ray emissions ($\\sim 10^{26}\\,\\mbox{erg}\\,\\mbox{s}^{-1}$) are believed to be due to the presence of energetic electrons with energies above ($\\sim 20\\,\\mbox{keV}$) in solar flares. Although, the exact mechanism of the initiation and triggering of solar flares is not yet known it is widely believed that the flares and the associated eruptions may occur due to magnetic reconnection, i.e., the rapid dissipation of electric currents near the magnetic null points. The ensuing relaxation of the sheared magnetic field topology can give rise to the large--scale \\alf waves which may transport the energy to the chromosphere \\citep{FH08}. Thus the development of turbulence in the chromosphere via reflection and mode conversion Of the \\alf wave may lead to the cascade of energy to the short wavelengths. The stochastic acceleration of the electron by a turbulent wave spectrum produces a high energy spectrum \\citep{FH08}. This model of electron acceleration crucially depends on the \\alf wave propagation along the field line to the chromosphere. However, as we have noted in the introduction, the concept of well defined flux tubes in a highly magnetic diffusion dominated medium is unclear. However, the electron acceleration in the chromosphere may indeed take place as envisioned by the Fletcher-Hudson Model due to magnetic diffusion driven turbulence. Since magnetic diffusion could be an important agent in driving the diffusive shear instability, this could easily lead to the low frequency turbulence in the medium. However, further work is needed in this direction to support this hypothesis." }, "1207/1207.0241_arXiv.txt": { "abstract": "{The magnetic field of galaxies is believed to be produced by internal dynamo action, but can be affected by motion of the galaxy through the surrounding medium. Observations of polarized radio emission of galaxies located in galaxy clusters have revealed noticeable features of large-scale magnetic configurations, including displacements of the magnetic structures from the optical images and tails, which are possible imprints of ram pressure effects arising from motion of the galaxies through the intracluster medium.} {We present a quantitative dynamo model which attempts to describe the above effects. In contrast to the traditional problem of a wind affecting a body with a prescribed magnetic field, we investigate how a non-magnetized wind flow affects a magnetic field that is being self-excited by galactic dynamo action.} {In order to isolate the leading physical effects we exploit a simple dynamo model that can describe relevant effects. In particular, we use what is known as the 'no-$z$' approximation for the mean-field dynamo equations.}{In a suitable parametric range we obtain displacements of the large-scale magnetic field, as well as magnetic tails. However, the specific details of their locations are quite counterintuitive. The direction of displacement is perpendicular to, rather than parallel to, the wind direction. The point at which the tail emerges from the galaxy depends on details of the model. The tail is eventually directed downstream. In the simplest case the magnetic tail begins in the region where the wind decreases the total gas velocity. Any wind that penetrates the galaxy modifies the intrinsic dynamo action. These features are different from those found in ram-pressure models.} {Any determination of galactic motion through the cluster medium from observational data needs to take the effects of dynamo action into account.} ", "introduction": "The formation and subsequent evolution of large-scale galactic magnetic fields are believed to be a result of galactic physical processes, mainly induction effects of differential rotation and mirror-asymmetric galactic turbulence (e.g. Beck et al. 1996). Galaxies are surrounded by the intergalactic medium and participate in various intergalactic interactions which somehow modify the ``intrinsic'' magnetic properties of galaxies. The corresponding effects on galactic magnetic structures naturally attract interest (e.g. Williams et al. 2002). A straightforward example of this is provided by the propagation of galaxies through the medium pervading galaxy clusters. It appears quite natural to attempt to attribute the existence of various galactic tails to the effects of such propagation. Such a morphological interpretation {\\it a priori} appears natural; however it requires verification by modelling, including effects of galactic dynamo action. Contemporary radio polarimetric observations provide several examples of large-scale galactic magnetic configurations which appear possibly to be affected by such propagation effects. In particular, Vollmer et al. (2007) and We\\.zgowiec et al. (2007) observed magnetic fields of large spiral galaxies in the Virgo Cluster in order to study interactions of galaxies with the cluster environment. (Further relevant results can be found in Kantharia et al. 2008; Vollmer et al. 2010; We\\.zgowiec et al. 2011, see also Yoshida et al. 2012.) A straightforward intention of this work was to use the proximity of the Virgo Cluster to isolate effects of the high-velocity tidal interactions and the effects of ram pressure stripping by the intracluster gas. \\begin{figure} \\centering \\includegraphics[width=0.3\\textwidth]{n4535.ps} \\caption{Polarized radio emission (contours) and B--vectors of the spiral galaxy NGC~4535 in the Virgo Cluster, observed at 3.6\\,cm wavelength with the Effelsberg 100-m telescope, smoothed to 3\\arcmin\\ resolution and overlaid onto a DSS optical image (kindly provided by Marek We\\.zgowiec).} \\label{fig:n4535} \\end{figure} \\begin{figure} \\centering \\includegraphics[width=0.35\\textwidth]{n4438.ps} \\caption{Polarized radio emission (contours) and B--vectors of the spiral galaxy NGC~4535 in the Virgo Cluster, observed at 6.3\\,cm wavelength with the VLA at 18\\arcsec\\ resolution and overlaid onto a DSS optical image (from Vollmer et al. 2007).} \\label{fig:n4438} \\end{figure} Indeed, the observational results presented by We\\.zgowiec et al. (2007) demonstrate a magnetic configuration that impressively mimics the naive expectation of the effects of ram pressure arising from galactic motion through the surrounding intracluster gas. In particular, the distribution of total radio intensity in NGC~4535 at 6.3\\,cm wavelength (Fig.~5 in the above cited paper) practically coincides with the optical image while the measures of polarized emission (Fig.~6 of that paper) are markedly displaced from the optical image towards the west. This displacement has been confirmed by observations at 3.6\\,cm wavelength (We\\.zgowiec et al. 2012) and hence cannot be caused by Faraday depolarization. Possible interpretations of this displacement include: (1) the spiral magnetic field generated in NGC~4535 by a galactic dynamo is compressed on the western side by ram pressure, (2) weakening of the galactic dynamo on the eastern side by ram pressure, (3) amplification and ordering of the magnetic field by shear in the HI envelope of the galaxy (Vollmer et al. 2010). Note that NGC~4535 is observed almost face-on, so that if the asymmetry in PI is due to ram pressure effects then its motion through the ICM is expected to be at a small or moderate angle with respect to the galactic plane; in this case the direction of motion can hardly be inferred from spectroscopic data which gives radial velocities. We present as an instructive example a map of polarized intensity and magnetic vectors of NGC~4535 overlaid on the optical image of this galaxy (Fig.~\\ref{fig:n4535}). Note that the maximum of polarized intensity is located at the periphery of the optical image while the magnetic field vectors cover the entire optical image. Another example comes from the observations of NGC~4501. (Fig.~3 in We\\.zgowiec et al. (2007) shows the total intensity and Fig.~4 the polarized intensity.) Again, the total intensity practically coincides with the optical image while the polarized emission is displaced from the optical image. The displacement is less pronounced than for NGC~4535. NGC~4501 is more inclined towards the line of sight than NGC~4535, which makes the interpretation less straightforward. Displacements towards one side of the disc are also detected in other galaxies observed edge-on, e.g. in NGC~4402 (Vollmer et al. 2007) and in NGC~4388 (We\\.zgowiec et al. 2012). On the other hand, several face-on galaxies of the Virgo Cluster with large-scale spiral patterns of their magnetic fields, like NGC~4303 and NGC~4321, show no displacement at all, presumably indicating that they move ``head-on'' through the cluster medium, i.e. almost along the line of sight (We\\.zgowiec et al. 2012). Each galaxy needs a specific model of gas stripping, gas back-flow and the effects of the flow on the magnetic field, and these can explain the observations in many cases (e.g. Vollmer et al. 2009, 2012). However, these models consider only magnetic fields which are sheared and amplified by the distorted velocity field of the gas and neglect the ongoing dynamo action in the galaxy. Not surprisingly, several features of the magnetic field patterns, such as a large-scale asymmetry out of the disc plane as in NGC~4192 (We\\.zgowiec et al. 2012) or a long magnetic tail in NGC~4438 (Fig.~\\ref{fig:n4438}) remained unexplained in these papers. An interpretation of the displacement as an effect of the ram pressure interacting with galactic dynamo action needs quantitative investigation and verification; this is the aim of this paper. Our idea is to test the interpretation in the most straightforward way. We present the effect of ram pressure (in the coordinate system of the galaxy) as a non-magnetized flow of the intracluster gas which moves through the galaxy parallel to the galactic plane and interacts with the dynamo-generated magnetic field. We assume that this motion does not destroy the main drivers of galactic dynamo, i.e. differential rotation and helical interstellar turbulence, that are responsible for the generation of the large-scale magnetic field and subsequently for the polarized emission. We expect to see that the gas motion (``wind'') displaces the dynamo-generated magnetic field so that it is no longer centred on the galaxy. The effect of a wind on magnetic structures around various celestial bodies is a classical topic that originated in investigations of interactions between the solar wind and the Earth's magnetic field. Obviously, the solar wind has a significant effect on the shape of the Earth's magnetic field, but not on its generation. We consider here the problem where a flow directly affects the region of magnetic field generation. We do not consider the galactic magnetic field to be prescribed {\\it a priori}, but rather investigate how the generation process is modified by the wind. ", "conclusions": "\\label{disc} We have presented the results of modelling of the effects of ram pressure on the stationary, dynamo-generated, magnetic field configuration in spiral galaxies. We recognize a general displacement of the magnetic configuration from the galaxy centre as well as, in some cases (not always, see Fig.~\\ref{fig:wind3}), magnetic tails propagating from the galaxy into the surrounding medium. These effects are rather nontrivial and have not been recognized in previous modelling attempts (Roediger 2009, Vollmer et al. 2009, 2012). Figs.~\\ref{fig:wind1} and \\ref{fig:wind2} show that the wind displaces the centre of the magnetic field pattern in a galaxy from the nominal centre of mass. Surprisingly, this displacement is not in the downstream direction of the wind flow, but perpendicular to this direction. As the galaxy is assumed to rotate counterclockwise, the wind increases the total gas velocity on the ``north'' side of the galaxy and decreases it on the ``south'' side. It thus appears that the large-scale field is displaced towards the galaxy in the region where the gas velocity is reduced. The small-scale (turbulent) magnetic field component, related to the star-formation activity in a galaxy and hence the gas distribution, is not affected, in agreement with the observation that the total radio emission is not displaced. If the ram speed of the wind is too large, it evacuates the large-scale magnetic field from the galaxy and dynamo action is killed. The critical velocity $U^*$ at which dynamo action is still possible can be estimated as follows. Accepting for a crude estimate that the maximal linear velocity of differential rotation $V = \\Omega R$ and taking into account the definitions of $R_\\omega$ and $R_m$ we obtain \\begin{equation} U = V \\, (R_{\\rm m}/R_\\omega) \\, \\lambda. \\label{est} \\end{equation} For $R_{\\rm m} = 150$, $R_\\omega = 20$ (close to the maximal wind velocity for which dynamo still works) and $\\lambda = 0.05$, $V = 250$\\,km/sec we obtain $U^* \\approx 100$\\,km/sec. Comparing the estimate $U^* = 100$\\,km/sec with the observed velocities of galaxies relative to the intracluster medium it is necessary to take into account that the majority of our models do not include the density contrast between the interstellar medium in the galaxy (ISM) and the intracluster medium (ICM). Estimating the limiting velocity of a galaxy $W$ in the cluster medium that is just sufficient to kill dynamo action from the condition $\\rho_{\\rm ISM} U^* = \\rho_{\\rm ICM} W$, and taking $\\rho_{\\rm ISM}/\\rho_{\\rm ICM} \\approx 10^{-2}$\\,cm$^{-3} / 10^{-4}$\\,cm$^{-3} \\approx 100$ we derive $W \\approx 10^4$\\,km/sec (here $\\rho_{\\rm ISM}$ and $\\rho_{\\rm ICM}$ are typical ionized gas densities in the galactic disc and the intracluster medium, respectively). If $U > U^*$, the wind pushes the dynamo-generated magnetic field out of the region where dynamo is active and hence kills the dynamo. If $U \\ll U^*$, the wind almost does not affect the dynamo-generated magnetic field. If $U$ is slightly less than $U^*$, the main effect is a less efficient dynamo action in the ``northmost'' part of the galaxy where wind and rotation velocities reinforce. In contrast, the magnetic field is larger in the ``south'' part of the galaxy where the velocities tend to cancel and dynamo action is stronger. As a result, the magnetic configuration is shifted across rather along the wind direction. There is thus a substantial difference between the effect of a wind on a magnetic configuration generated {\\it in situ} (i.e. by dynamo action) from that on a magnetic configuration generated before the wind impacts on the galaxy. In our models the dynamo operates contemporaneously with the wind, and so the wind modifies the dynamo velocity field, and hence the dynamo properties. The wind direction and the general displacement of the main magnetic structure are mutually perpendicular. This must be into account in determinations of galactic motion through the cluster medium from observational data (cf. e.g. Pfrommer \\& Dursi 2010). The displacements obtained in the models are quite moderate. It is impossible to push the centre of the magnetic configuration far beyond the galaxy centre, or the boundary of the region occupied by the bulk of the field far beyond its original position. This is partly at least because in the absence of external dynamo sources (here alpha-effect) the field rapidly decays. Note that the magnetic vectors in the map of NGC~4535 cover the centre of the optical structure of the galaxy, i.e. the displacement of magnetic structure from the optical image is still quite modest. Another unexpected result is the location of the magnetic tail: it emerges from different regions of the disc depending on the details of the model, and then propagates into the area where the wind decreases the total velocity. The dynamo-generated field propagates due to turbulent diffusion out of the area which rotates significantly, and is caught by the wind, producing a tail. Observations (We\\.zgowiez et al. 2007) do not show the presence of magnetic tails in NGC~4535 and NGC~4501. On the other hand, the spectacular magnetic tails presented in Fig.~\\ref{fig:wind1} may be related to the tail of polarized emission in NGC~4438 (Fig.~\\ref{fig:n4438}; Vollmer et al. 2007, 2010). Our simple model does not describe gas flows which might be associated with the magnetic tails. It appears plausible that gas is also carried along in the magnetic tail. Chung et al. (2007) found seven spiral galaxies in the Virgo cluster with long, one-sided tails of neutral gas. The various features of magnetic field structures associated with the ram pressure effects obtained in the models presented above have different ranges of robustness. The magnetic structure is displaced {\\it perpendicularly} to the wind direction in all models. Magnetic tails are quite ``fragile'' and disappear when the turbulent magnetic diffusivity in the surrounding space is increased. The value of turbulent magnetic diffusivity in galactic halos and intracluster media is an important dynamo governing parameter, but is difficult to estimate (e.g. Sokoloff \\& Shukurov 1990). If we include a deceleration of the wind in the disc, the tails are smaller and its recognition becomes difficult -- see Sect.~3. The tail disappears if an $x$-dependence is imposed on the wind velocity (Sect.~3). Furthermore, the no-$z$ approximation can no longer be expected to be valid far outside of the disc region, where {\\it inter alia} the estimate of $-B/h^2$ for the $z$-diffusion may in turn be an over-estimate. Given that the dynamo does not operate in the external region, the latter uncertainty can probably be formally absorbed into the uncertainty concerning the external diffusivity. Thus the presence of a magnetic tail cannot be considered a robust feature of our simulations. A statistical investigation of magnetic tails in cluster galaxies could contribute to elucidation of the problem. Of course, further details such as elongated magnetic structures and their particular forms are even more uncertain and need further more detailed modelling. Our simulations are necessarily restricted to considering flows parallel to the galactic plane. Real flows will be inclined at arbitrary angles. We can split such a general flow into components parallel to and perpendicular to the plane. We have separately modelled the effects of a perpendicular flow (parallel to the $z$-axis) in an {\\it axisymmetric} dynamo model with dynamo effect confined to the disc region. Here the effect is more clearly an advection of the field structure in the direction of the flow, effectively \"lifting\" the field from the galactic midplane. As the field is moved out of the disc region, where the alpha-effect is assumed to operate, in the absence of dynamo action the field decays. The advection is thus naturally limited in extent. A somewhat faster flow removes the field from the dynamo-active region more rapidly than it can be regenerated, and so kills the dynamo completely. Thus we feel our results are reasonably representative of flows that are not too strongly inclined to the galactic plane. In contrast, strong flows at large angles to the disc plane will seriously weaken dynamo action. Finally, we wish to stress again that the problem under discussion is different from the classical problem of a magnetized wind affecting a body with prescribed magnetic field (this problem in the galactic context is addressed e.g. by Ruszkowski et al. 2012). The interaction between the wind and the internal velocity field, and hence the effect of the wind on dynamo action, cannot be avoided, unless the internal magnetic field is completely shielded from the external flow - which however is unrealistic. The problem is no longer that of a dynamo with a simple circular flow. Hence, the interpretation of an asymmetry in polarized radio intensity as a naive effect of ram pressure is similarly unrealistic. A compression of the internal field is possible only if the wind penetrates the galaxy, which in turn must affect dynamo action. Clearly there are plausible improvements that could be made to our no-$z$ model, but we believe our results are generic. We have investigated a magnetic field that is self-excited in a region that is directly affected by the wind, i.e. the flow generating the dynamo is altered. This problem has not attracted attention so far. In this preliminary study we have attempted to isolate the basic physical effects. We assume that the wind itself is non-magnetized and flows parallel to the galactic plane, and use simple models to describe the wind profile. These assumptions allowed us to use the simple no-$z$ approach. A more detailed modelling, including a magnetized wind inclined to the galactic plane and extension to a 3D model, and /or to go beyond a mean-field formulation of the problem, would be instructive and should be investigated in future." }, "1207/1207.2857_arXiv.txt": { "abstract": "We study the 17 January 2010 flare--CME--wave event by using STEREO/SECCHI EUVI and COR1 data. The observational study is combined with an analytic model which simulates the evolution of the coronal-wave phenomenon associated with the event. From EUV observations, the wave signature appears to be dome shaped having a component propagating on the solar surface ($\\overline{v}\\approx$~280~km~s$^{-1}$) as well as off-disk ($\\overline{v}\\approx$~600~km~s$^{-1}$) away from the Sun. The off-disk dome of the wave consists of two enhancements in intensity, which conjointly develop and can be followed up to white-light coronagraph images. Applying an analytic model, we derive that these intensity variations belong to a wave-driver system with a weakly shocked wave, initially driven by expanding loops, which are indicative of the early evolution phase of the accompanying CME. We obtain the shock standoff distance between wave and driver from observations as well as from model results. The shock standoff distance close to the Sun ($<$0.3~\\textit{\\textit{R$_\\odot$}} above the solar surface) is found to rapidly increase with values of $\\approx$0.03\\,--\\,0.09~\\textit{R$_\\odot$} which give evidence of an initial lateral (over-)expansion of the CME. The kinematical evolution of the on-disk wave could be modeled using input parameters which require a more impulsive driver ($t$\\,=\\,90~s, $a$\\,=\\,1.7~km~s$^{-2}$) compared to the off-disk component ($t$\\,=\\,340~s, $a$\\,=\\,1.5~km~s$^{-2}$). ", "introduction": "In large part, our knowledge of coronal mass ejections (CMEs) comes from coronagraph observations delivering white-light data. CMEs, as observed in white light, often exhibit a typical three-part structure, consisting of a bright rim encircling a dark cavity, mostly followed by a bright core \\citep[][]{illing85}. Therefore, by definition, a CME is a structured intensity enhancement observed in white-light. The actual process that launches the ejection is most probably connected to magnetic restructuring. This early evolution phase of a CME can often be observed in the extreme ultraviolet as well as soft X-ray data in the form of expanding loop systems \\citep[\\textit{e.g.}][]{harrison00,vrsnak04-cme}. CMEs, as they evolve and propagate away from the Sun, are able to drive magnetohydrodynamical (MHD) shocks in the corona that can be tracked by coronal type II radio bursts \\citep{gopalswamy97,magdalenic10}. The formation of the shock itself is dependent on the time--speed profile of the CME as well as on the spatial distribution of the Alfv\\'{e}n speed in the solar corona, which in turn is related to the local magnetic-field strength and density of the ambient plasma. To generate a shock, the CME needs to have a sufficiently high velocity with respect to the local Alfv\\'{e}n speed and such favorable conditions are assumed to be present in the middle corona over $\\approx$2~\\textit{R$_\\odot$} \\citep[\\textit{e.g.},][]{gopal01,mann03}. Recent studies showed that shocks driven by fast CMEs are observable in white-light data \\citep[\\textit{e.g.},][]{vourlidas03,ontiveros09,bemporad10,kim12} as well as EUV \\citep[e.g.][]{veronig10,kozarev11,ma11,gopalswamy11,cheng12}, and UV spectra \\citep[e.g.][]{raymond00,bemporad10}. The evolution of a three-dimensional dome connected to a surface shock wave is observed for the 17 January 2010 CME--flare event. It was studied in detail by \\cite{veronig10} who showed that the surface as well as the off-limb structure are part of an evolving three-dimensional wave-dome formed by a weak shock. The surface wave propagated with a mean speed of $\\approx$280~km~s$^{-1}$ whereas the upward moving part was of much higher speed of $\\approx$650~km~s$^{-1}$ \\citep{veronig10}. The difference between the speed of the upward-moving part of the wave and the on-disk signature was interpreted by \\cite{veronig10} in the following manner: the upward-moving part is driven all of the time by the outward moving CME, whereas the surface signature is only temporarily driven by the flanks of the expanding CME and then propagates freely. A recent article by \\cite{grechnev11}, studying the same event, supports the result that the dome-structure was actually a shock-driven plasma flow. \\cite{grechnev11} simulated the evolution of the shock wave from which they concluded that most likely an abrupt eruption of a filament caused the weak shock. They compare this with a blast-wave scenario during which the wave is only briefly driven. \\cite{zhao11} investigated, for the 17 January 2010 event the relation between the surface wave speed, the CME speed and the local fast-mode characteristic speed. They concluded that the observed CME front is in fact a wave phenomenon just like the EUV wave on the solar surface. In this study, we focus on the kinematical evolution of the off-disk signature of the dome-shaped wave event and add new aspects not covered by previous studies. Using observations of the SECCHI instruments EUVI and COR1 on STEREO we will show, by applying an analytical model with input parameters constrained by the observations, that the off-disk signature in fact consists of two components: a driver and a weakly shocked wave. The driver of the off-limb wave evolves from expanding loop structures and is interpreted as the CME, the observed frontal part is interpreted as the shock wave ahead. In particular, we investigate the shock offset (standoff) distance for the wave-driver system. ", "conclusions": "The 17 January 2010 is a well-observed event revealing a dome-shaped coronal wave structure. In the present study we give evidence from observational and model results that the off-disk part of the wave actually consists of a driver and a wave component. The driver is interpreted as the CME, the frontal part as a weakly shocked wave. We derive that the shock standoff distance shows a linear evolution with a rather rapid increase below 0.3~\\textit{R$_\\odot$} above the solar surface. These results may be interpreted that the initial lateral (over-)expansion of the CME which is short-lived \\citep[$\\approx$70~seconds; see][]{patsourakos10} acts as a piston driver to the shock, which leads to a rapid increase in the shock standoff distance. The piston nature of expanding CME flanks is also reflected in results from a recent study by \\cite{cheng12}. Using SDO observations, \\cite{cheng12} analyzed the structural and kinematical evolution of a CME together with the separation process of a diffuse wave front from the CME flanks. The wave decoupling from the driver, and with this the actual detection of the wave front, happens when the CME expansion slows down. Comparing our results to previous studies of off-disk waves moving in the radial direction away from the Sun, we find good agreement for several parameters. We derive for the first observable, \\textit{i.e.}, measurable standoff distance values of 0.03\\,--\\,0.06~\\textit{R$_\\odot$}. \\cite{ma11} who studied a low coronal shock wave using high spatial and temporal resolution data from SDO/AIA report for the thickness of the shocked layer $\\approx$0.03~\\textit{R$_\\odot$}. We find the first signatures of the shock at a distance of $\\approx$0.28~\\textit{R$_\\odot$} above the solar surface which is comparable to the results from \\cite{ma11} who find 0.23~\\textit{R$_\\odot$}. From the model we also derive the timing of the shock signatures to be close to the observed type II radio burst. \\cite{cheng12} refer to an almost simultaneous occurrence between a type II radio burst and the start of the separation process between wave and driver. For radio bursts, shock formation heights of $\\approx$0.2~\\textit{R$_\\odot$} are derived by, \\textit{e.g.}, \\cite{magdalenic10}. We note that the height at which the shock forms is strongly dependent on the speed profile of the driver, \\textit{i.e.}\\ CME. Peak accelerations of CMEs are found to occur at very low distances from their launch site: $<$0.5~\\textit{R$_\\odot$} above the solar surface \\cite[\\textit{e.g.}][]{temmer08,temmer10}. Using an analytical model, the kinematical profile of both components, driver and wave, can be simulated by applying model parameters that are constrained by observations. In addition, we are able to simulate the on-disk wave using a pure piston-type expansion of the driving source whereas the source of the off-disk wave behaves in the early evolution as piston then becomes more of a bow shock type. We find that the on-disk wave requires a more impulsive driver ($t$\\,=\\,90~seconds, $a$\\,=\\,1.7~km~s$^{-2}$) compared to the off-disk wave ($t$\\,=\\,340~seconds, $a$\\,=\\,1.5~km~s$^{-2}$). These results lie in between the findings from \\cite{grechnev11} who obtain that the off-limb wave was excited impulsively most probably from a filament eruption and then propagated freely, and those of \\cite{veronig10} who conclude that the upward-moving dome might have been driven all of the time. The dome-shaped wave under study evolves from an eruption plus the deformation of different loop systems. In morphology, the dome-shaped wave reminds much more of a CME bubble than of separated loop systems. In visible light, shock waves are reported as well-outlined and sharp boundaries \\citep[see][]{ontiveros09}. For the evolution of a coronal surface wave, \\cite{temmer10} reported that the wave was launched from two separate centers before it became of circular shape. We may speculate that the loop systems expand and are pushed aside due to the early evolution phase of the erupting structure. The magnetic loop structures form the ``observable envelope'' and are the first signatures of the evolving CME. The current study shows that relatively slow drivers may cause weak shock waves low in the corona. These waves are visible in white-light and may further propagate up to 1~AU. To investigate the evolution of shock standoff distances in interplanetary space, we require observations of the wave--driver system close to the Sun as well as their in-situ signatures. \\begin{acks} MT and AMV greatly acknowledge the Austrian Science Fund (FWF): V195-N16 and P24092-N16. The research leading to these results has received funding from the European Commission's Seventh Framework Programme (FP7/2007-2013) under the grant agreement n$^{\\circ}$~263252 [COMESEP] and n$^{\\circ}$~284461 [eHEROES]. \\end{acks}" }, "1207/1207.6755.txt": { "abstract": "{This review summarizes a few of the frontiers of Galactic center research that are currently the focus of considerable activity and attention. It is aimed at providing a necessarily incomplete sketch of some of the timely work being done on phenomena taking place in, or originating in, the central few parsecs of the Galaxy, with particular attention to topics related to the Galactic black hole (GBH). We have chosen to expand on the following exciting topics: 1) the characterization and the implications for the variability of emission from the GBH, 2) the strong evidence for a powerful X-ray flare in the Galactic center within the past few hundred years, and the likelihood that the GBH is implicated in that event, 3) the prospects for detecting the \"shadow\" of the GBH, 4) an overview of the current state of research on the central S-star cluster, and what has been learned from the stellar orbits within that cluster, and 5) the current hypotheses for the origin of the G2 dust cloud that is projected to make a close passage by the GBH in 2013. ", "introduction": "%% first-level sections will be auto-capitalized \\label{sect:intro} If we adopt the discovery of the nuclear stellar cluster in the infrared by \\citet{BN68} as the beginning of research focussed on the central parsecs of the Galactic center, the field is only 44 years old. In that time, the pace of discovery has been breathtaking, and it has not shown any signs of leveling off. This review is intended to illustrate this fact by offering brief descriptions of some selected samples of exciting research trends that are currently under intense investigation. It focusses on topics that are related in various ways to the supermassive Galactic black hole, often given the name of the associated radio source, Sagittarius A*. The review is not meant to be complete; many important topics are left out. For comprehensive overviews on topics not addressed here, the reader is referred to the recent review article by \\citet{GEG10} and to the proceedings of a recent international conference on this subject \\citep{MWY11}. ", "conclusions": "" }, "1207/1207.4706_arXiv.txt": { "abstract": "{ We make use of the \\Planck\\ all-sky survey to derive number counts and spectral indices of extragalactic sources -- infrared and radio sources -- from the \\Planck\\ Early Release Compact Source Catalogue (ERCSC) at 100 to 857\\,GHz (3\\,mm to 350\\micron). Three zones (deep, medium and shallow) of approximately homogeneous coverage are used to permit a clean and controlled correction for incompleteness, which was explicitly not done for the ERCSC, as it was aimed at providing lists of sources to be followed up. Our sample, prior to the 80\\,\\% completeness cut, contains between 217 sources at 100\\,GHz and 1058 sources at 857\\,GHz over about 12,800 to 16,550\\,deg$^2$ (31 to 40\\,\\% of the sky). After the 80\\,\\% completeness cut, between 122 and 452 and sources remain, with flux densities above 0.3 and 1.9\\,Jy at 100 and 857\\,GHz. The sample so defined can be used for statistical analysis. Using the multi-frequency coverage of the \\Planck\\ High Frequency Instrument, all the sources have been classified as either dust-dominated (infrared galaxies) or synchrotron-dominated (radio galaxies) on the basis of their spectral energy distributions (SED). Our sample is thus complete, flux-limited and color-selected to differentiate between the two populations. We find an approximately equal number of synchrotron and dusty sources between 217 and 353\\,GHz; at 353\\,GHz or higher (or 217\\,GHz and lower) frequencies, the number is dominated by dusty (synchrotron) sources, as expected. For most of the sources, the spectral indices are also derived. We provide for the first time counts of bright sources from 353 to 857\\,GHz and the contributions from dusty and synchrotron sources at all HFI frequencies in the key spectral range where these spectra are crossing. The observed counts are in the Euclidean regime. The number counts are compared to previously published data (from earlier \\Planck\\ results, {\\it Herschel}, BLAST, SCUBA, LABOCA, SPT, and ACT) and models taking into account both radio or infrared galaxies, and covering a large range of flux densities. We derive the multi-frequency Euclidean level -- the plateau in the normalised differential counts at high flux-density -- and compare it to {\\it WMAP}, {\\it Spitzer\\/} and {\\it IRAS\\/} results. The submillimetre number counts are not well reproduced by current evolution models of dusty galaxies, whereas the millimetre part appears reasonably well fitted by the most recent model for synchrotron-dominated sources. Finally we provide estimates of the local luminosity density of dusty galaxies, providing the first such measurements at 545 and 857\\,GHz.} ", "introduction": "\\begin{figure}[!ht] \\centering \\includegraphics[width=0.30\\textwidth,angle=90]{mask_gal_857_KQ75.ps} \\caption{Comparison between our \\Planck\\ 857\\,GHz mask (red indicates regions removed from the analysis) and the {\\it WMAP} 7-year KQ75 mask (green means removed); unlike the case for the mask employed in this paper, the {\\it WMAP} mask excludes some point sources. The background (blue) is the sky area used for our analysis. This map is a Mollweide projection of the sky in Galactic coordinates.} \\label{fig:mask857} \\end{figure} Among other advantages, all-sky multifrequency surveys have the benefit of probing rare and/or bright objects in the sky. One reason to probe bright objects is to study the number counts of extragalactic sources and their spectral shapes. In the far-infrared (FIR) the sources detected by these surveys are usually dominated by low-redshift galaxies with $z<0.1$, as found by {\\it IRAS} at 60\\micron\\ \\citep{ashby96} but a few extreme objects like the lensed F10214 source \\citep{rowan-robinson91} also appear. However, the population in the radio band is dominated by synchrotron sources (in particular, blazars) at higher redshift \\citep[see][for a recent review]{de_zotti2010}. Previous multifrequency all-sky surveys were carried out in the infrared (IR) range by the {\\it IRAS} satellite (between 12 and 100\\micron; \\citealt{neugebauer84}), and more recently by {\\it Akari} (between 2 and 180\\micron; \\citealt{murakami2007}) and {\\it WISE} (between 3.4 and 22\\micron; \\citealt{wright2010}). Early, limited sensitivity surveys were carried out in the IR and microwave range by {\\it COBE} (between 1.2\\micron\\ and 1\\,cm; \\citealt{boggess92}), and more recently in the microwave range by {\\it WMAP} (between 23 and 94\\,GHz; \\citealt{bennett2003,wright2009,massardi2009,de_zotti2010}). \\Planck's\\ all-sky, multifrequency surveys offer several advantages to all of the above. The frequency range covered is wide, extending from 30 to 857\\,GHz; observations at all frequencies were made simultaneously, reducing the influence of source variability; the calibration is uniform; and the delivered catalogue of sources used in this paper was carefully constructed and validated. \\begin{figure}[!ht] \\centering \\includegraphics[width=0.30\\textwidth,angle=90]{mask_gal_353_KQ85.ps} \\caption{Comparison between our \\Planck\\ 353\\,GHz mask (red means removed) and {\\it WMAP} 7-year KQ85 mask (green). The background (blue) is the sky area used for the analysis.} \\label{fig:mask353} \\end{figure} \\begin{figure}[!hb] \\centering \\includegraphics[width=0.50\\textwidth]{histo_hitcounts.ps} \\caption{Cumulative distribution of \\Planck\\ hit counts on the sky (here at 857\\,GHz with Nside=2048), with the corresponding integration time per sky pixel (given on the top axis). The three sky zones used in the analysis are defined at 857\\,GHz: shallow (50 to 75\\,\\% of the hit count distribution, short-spaced lines, blue); medium (75 to 95\\,\\% of the hit counts, medium-spaced lines, green); and deep (95\\,\\% and above hits, widely-spaced lines, red).} \\label{fig:hitcounts} \\end{figure} The \\Planck\\ frequency range fully covers the transition between the dust emission dominated regime (tracing star formation), and the synchrotron regime (tracing active galactic nuclei). The statistical analysis of the populations in this spectral range has never been done before. At large flux densities (typically 1\\,Jy and above), number counts from all-sky surveys of extragalactic FIR sources show a Euclidean component, i.e. a distribution of the number of source per flux density bin $S_{\\nu}$ at observed frequency $\\nu$ ($dN/dS_{\\nu}$ in Jy$^{-1}$\\,sr$^{-1}$) proportional to $S_{\\nu}^{-2.5}$ (see Eq. \\,\\ref{eq:p}). This result is in line with expectations from a local Universe uniformly filled with non-evolving galaxies \\citep{lonsdale89,hacking91,bertin97,massardi2009,wright2009}. In the radio range, the Euclidean part is modified by the presence of higher-redshift sources. At flux densities smaller than typically 0.1 to 1\\,Jy, an excess in the number counts compared to the Euclidean level is an indication of evolution in luminosity and/or density of the galaxy populations. This effect is clearly seen in deeper surveys in the FIR (e.g.~\\citealt{genzel2000,dole2001,dole2004a,frayer2006,soifer2008,bethermin2010}); in the submillimetre (submm) range (e.g.~\\citealt{barger99,blain99,ivison2000,greve2004,coppin2006,weiss2009,patanchon2009,clements2010,lapi2011}); -- and in the millimetre and radio ranges (e.g.~\\citealt{de_zotti2010,vieira2010,vernstrom2011}). \\begin{figure}[!ht] \\centering \\includegraphics[width=0.30\\textwidth,angle=90]{number_counts_ercsc_maskall_857.ps} \\caption{The three sky zones used in the analysis at 857\\,GHz: deep (red); medium (green); and shallow (blue); These are based on the 857\\,GHz hit counts.} \\label{fig:surveys857} \\end{figure} The \\Planck\\footnote{\\Planck\\ (\\url{http://www.esa.int/Planck}) is a project of the European Space Agency (ESA) with instruments provided by two scientific consortia funded by ESA member states (in particular the lead countries France and Italy), with contributions from NASA (USA) and telescope reflectors provided by a collaboration between ESA and a scientific consortium led and funded by Denmark.} all-sky survey covers nine bands between 30 and 857\\,GHz. It gives us for the first time robust extragalactic counts over a wide area of sky at these wavelengths, and the first all-sky coverage between 3\\,mm ({\\it WMAP}) and 160\\micron\\ ({\\it Akari}) -- e.g. see Table\\,1 of \\citet{planck2011-1.10}. The counts in turn give us powerful constraints on the long-wavelength spectral energy distribution (SED) of the dusty galaxies investigated e.g. by {\\it IRAS}, and on the short-wavelength SED of the active galaxies studied at radio wavelengths, e.g. by {\\it WMAP} or ground-based facilities. For \\Planck's\\ six highest frequency bands (100 to 857\\,GHz, we present here the extragalactic number counts and spectral indices of galaxies selected at high Galactic latitude and using identifications. \\Planck\\ number counts and spectral indices of extragalactic radio-selected sources were already published for the frequency range 30 to 217\\,GHz using results from the LFI and HFI instruments \\citep{planck2011-6.1}. The transition between synchrotron-dominated sources and thermal dust-dominated occurs in the crucial spectral range 200-800\\,GHz . Thus our broader frequency data allow a better spectral characterisation of sources. We use the {\\it WMAP} 7 year best-fit $\\Lambda$CDM cosmology \\citep{larson2011}, with $H_0$ = 71\\kmsMpc, $\\Omega_\\Lambda$ = 0.734 and $\\Omega_{\\rm M}$ = 0.266. \\begin{figure}[!ht] \\centering \\includegraphics[width=0.30\\textwidth,angle=90]{number_counts_ercsc_maskall_100.ps} \\caption{The three sky zones used in the analysis at 100\\,GHz: deep (red); medium (green); and shallow (blue). These are based on the 100\\,GHz hit counts.} \\label{fig:surveys100} \\end{figure} \\begin{figure}[!b] \\centering \\includegraphics[width=0.50\\textwidth,angle=0]{number_counts_ercsc_completeness.6all.ps} \\caption{Completeness (vs.~flux density of sources) coming from the Monte-Carlo runs for the ERCSC and derived for each zone. The horizontal dashed line represents our threshold for number count analysis.} \\label{fig:completeness} \\end{figure} ", "conclusions": "" }, "1207/1207.3310_arXiv.txt": { "abstract": "The \\pipe\\ is a massive, nearby, filamentary dark molecular cloud with a low star-formation efficiency threaded by a uniform magnetic field perpendicular to its main axis. It harbors more than a hundred, mostly quiescent, very chemically young starless cores. The cloud is, therefore, a good laboratory to study the earliest stages of the star-formation process. We aim to investigate the primordial conditions and the relation among physical, chemical, and magnetic properties in the evolution of low-mass starless cores. We used the IRAM 30-m telescope to map the 1.2~mm dust continuum emission of five new starless cores, which are in good agreement with previous visual extinction maps. For the sample of nine cores, which includes the four cores studied in a previous work, we derived a \\Av\\ to \\Nhd\\ factor of (1.27$\\pm$0.12)$\\times$10$^{-21}$~mag~cm$^{2}$ and a background visual extinction of $\\sim$6.7~mag possibly arising from the cloud material. We derived an average core diameter of $\\sim$0.08~pc, density of $\\sim$10$^5$~\\cmt, and mass of $\\sim$1.7~\\msun. Several trends seem to exist related to increasing core density: ({\\it i}) diameter seems to shrink, ({\\it ii}) mass seems to increase, and ({\\it iii}) chemistry tends to be richer. No correlation is found between the direction of the surrounding diffuse medium magnetic field and the projected orientation of the cores, suggesting that large scale magnetic fields seem to play a secondary role in shaping the cores. We also used the IRAM 30-m telescope to extend the previous molecular survey at 1 and 3~mm of early- and late-time molecules toward the same five new \\pipe\\ starless cores, and analyzed the normalized intensities of the detected molecular transitions. We confirmed the chemical differentiation toward the sample and increased the number of molecular transitions of the ``diffuse'' (e.g. the ``ubiquitous'' \\co, \\cdh, and \\cs), ``oxo-sulfurated'' (e.g. \\so\\ and \\chtoh), and ``deuterated'' (e.g. \\ndhp, \\cn, and \\hcn) starless core groups. The chemically defined core groups seem to be related to different evolutionary stages: ``diffuse'' cores present the cloud chemistry and are the less dense, while ``deuterated'' cores are the densest and present a chemistry typical of evolved dense cores. ``Oxo-sulfurated'' cores might be in a transitional stage exhibiting intermediate properties and a very characteristic chemistry. ", "introduction": "The \\pipe\\ is a massive ($10^4$~\\msun: \\citealp{onishi99}) nearby (145~pc: \\citealp{alves07}) filamentary dark cloud located in the southern sky (Fig.~\\ref{fig:pipe}). What differentiates the \\pipe\\ from other low-mass star-forming regions such as Taurus and $\\rho$-Ophiuchus is that it is very quiescent and has a very low star-formation efficiency, only the Barnard~59 (B59) region shows star formation \\citep{forbrich09,brooke07,roman09,roman11}. The cloud harbors more than one hundred low-mass starless dense cores in a very early evolutionary stage \\citep{muench07,rathborne08}. Thermal pressure appears to be the dominant source of internal pressure of these cores: most of them appear to be pressure confined, but gravitationally unbound \\citep{lada08}. Only the B59 region shows a significant non-thermal contribution to molecular line widths that could be caused by outflows feedback \\citep{duartecabral12}. Through simulations of an unmagnetized cloud compatible to the \\pipe, \\citet{heitsch09} predicted pressures lower than those required by \\citet{lada08}. This result suggests that an extra source of pressure, such as magnetic fields, is acting. In fact, \\citet{franco10} found that most of the \\pipe\\ is magnetically dominated and that turbulence appears to be sub-Alfv\\'enic. \\citet{alves08} have distinguished three regions in the cloud with differentiated polarization properties, proposed to be at different evolutionary stages (Fig.~\\ref{fig:pipe}). B59, with low polarization degree (\\polp) and high polarization vector dispersion (\\deltapa), is the only magnetically supercritical region and might be the most evolved, the \\stem\\ would be at an earlier evolutionary stage, and finally, the \\bowl, with high \\polp\\ and low \\deltapa, would be at the earliest stage. The \\pipe\\ is, hence, an excellent place to study the initial conditions of core formation which may eventually undergo star formation. \\begin{figure*}[t] \\centering \\includegraphics[width=14cm,angle=0]{pipe_franco_bw2.ps} \\caption{Position of the observed cores plotted over the 2MASS extinction map of the \\pipe\\ \\citep{lombardi06}. Black segments represent the mean polarization vector of the region \\citep{alves08} with the scale shown on the top left corner of the figure. White boxes depict the size of the 1.2~mm continuum maps (Section~\\ref{obs_mambo} and Fig.~\\ref{fig:mambo} of both \\paper\\ and the present work). The dashed lines separate the three different magnetically defined regions \\citep{alves08}. The lowest visual extinction ($A_\\texttt{v}$) corresponds to 0.5 magnitudes. The highest $A_\\texttt{v}$ is observed toward the \\bowl\\ and B59 regions, where it reaches approximately 20 magnitudes. \\label{fig:pipe} } \\end{figure*} \\begin{table}[t] \\caption{ Source List Observed in \\paper\\ and in this Work. } \\begin{tabular}{ccccc} \\hline\\hline & \\multicolumn{1}{c}{$\\alpha$(J2000)} & \\multicolumn{1}{c}{$\\delta$(J2000)} & \\multicolumn{1}{c}{$\\texttt{v}_{\\rm LSR}$$^b$} & \\multicolumn{1}{c}{} \\\\ \\multicolumn{1}{c}{Source$^a$} & \\multicolumn{1}{c}{h m s} & \\multicolumn{1}{c}{$\\degr$ $\\arcmin$ $\\arcsec$} & \\multicolumn{1}{c}{(km s$^{-1}$)} & \\multicolumn{1}{c}{Region$^c$} \\\\ \\hline Core 06 & 17 10 31.6 & -27 25 51.6 & +3.4 & B59\\\\ Core 14 & 17 12 34.0 & -27 21 16.2 & +3.5 & B59\\\\ Core 20 & 17 15 11.2 & -27 35 06.0 & +3.5 & {\\it Stem}\\\\ Core 40 & 17 21 16.4 & -26 52 56.7 & +3.3 & {\\it Stem}\\\\ Core 47 & 17 27 29.6 & -26 59 06.0 & +2.8 & {\\it Stem}\\\\ Core 48 & 17 25 59.0 & -26 44 11.8 & +3.6 & {\\it Stem}\\\\ Core 65 & 17 31 20.5 & -26 30 36.1 & +5.0 & {\\it Bowl}\\\\ Core 74 & 17 32 35.3 & -26 15 54.0 & +4.2 & {\\it Bowl}\\\\ Core 109 & 17 35 48.5 & -25 33 05.8 & +5.8 & {\\it Bowl}\\\\ \\hline \\end{tabular} $^a$ According to \\citet{lombardi06} numbering. \\\\ $^b$ \\citet{rathborne08}. \\\\ $^c$ According to \\citet{alves08} diffuse gas polarimetric properties. \\label{tab_source} \\end{table} \\citet[][hereafter \\paper]{frau10} presented the first results of a molecular line study at high spectral resolution for a sample of four cores distributed in the different regions of the \\pipe. The aim of the project was to chemically date the cores through an extensive molecular survey based in two main categories of molecules: early- and late-time \\citep[e.g.,][]{taylor96}. In addition, we mapped the 1.2~mm dust continuum emission of the cores. We found no clear correlation between the chemical evolutionary stage of the cores and their location in the \\pipe\\ and, therefore, with the large scale magnetic field. However, at core scales, there are hints of a correlation between the chemical evolutionary stage of the cores and the local magnetic properties. Recently, \\citet[][hereafter \\paperd]{frau12} have presented a 3~mm $\\sim$15~GHz wide chemical survey toward fourteen starless cores in the \\pipe. In order to avoid a density bias, we defined the molecular line normalized intensity by dividing the spectra by the visual extinction (\\Av) peak, similar to the definition of the abundance. We found a clear chemical differentiation, and normalized intensity trends among the cores related to their \\Av\\ peak value. We defined three groups of cores: ``diffuse'' cores (\\Av$\\lsim$15~mag) with emission only of ``ubiquitous'' molecular transitions present in all the cores (\\cdh, \\cthd, \\hcop, \\cs, \\so, and \\hcn), ``oxo-sulfurated'' cores (\\Av$\\simeq$15-22~mag) with emission of molecules like \\tso, \\sod, and \\ocs, only present in this group, and finally, ``deuterated'' cores (\\Av$\\gsim$22~mag), which present emission in nitrogen- and deuterium-bearing molecules, as well as in carbon chain molecules. In this paper, we further explored observationally the relationship among structure, chemistry, and magnetic field by extending the sample in five new \\pipe\\ cores, for a total of nine, and several new molecular transitions. We repeated and extended the analysis conducted in \\paper\\ for molecular (temperature, opacity, and column density estimates) and continuum (dust parameters estimates and comparison with previous maps) data. We also derived and analyzed the molecular line normalized intensities as in \\paperd. For the sake of simplicity, we omit here technical details, which are widely explained in \\paperud. ", "conclusions": "We carried out observations of continuum and line emission toward five starless cores, located on the three different regions of the \\pipe, and combined them with the observations of the four additional cores of \\paper\\ to extend the dataset to nine cores. We studied the physical and chemical properties of the cores, and their correlation following \\paperd. We also studied the correlation with the magnetic field properties of the surrounding diffuse gas following \\paper. \\begin{enumerate} \\item The \\pipe\\ starless cores show very different morphologies. The complete sample of nine cores contains dense and compact cores (6, 65, and 109; \\nhd\\gsim10$^5$~\\cmt), diffuse and elliptical/irregular ones (20, 40, 47, 48, and 74; \\nhd\\lsim5$\\times$10$^4$~\\cmt), and a filament containing the relatively dense core~14 (\\nhd$\\sim$9$\\times$10$^4$~\\cmt). The average properties of the nine cores of the sample are diameter of $\\sim$0.08~pc ($\\sim$16,800~AU), density of $\\sim$10$^5$~\\cmt, and mass of $\\sim$1.7~\\msun. These values are very close to (but less dense than) those reported by \\citet{wardthompson99} for a set of very young dense cores and, therefore, typical of even earlier stages of evolution. \\item MAMBO-II maps are in a general good morphological agreement with previous extinction maps \\citep{lombardi06}. By comparing the \\Av\\ peak values of the nine cores from deeper NICER maps \\citep{roman09,roman10}, we derived a proportionality factor \\Av/\\Nhd=(1.27$\\pm$0.12)$\\times$10$^{-21}$~mag~cm$^{2}$, compatible with the standard value (\\tenpow{1.258}{-21}~mag~cm$^{2}$; \\citealp{wagenblast89}). In addition, we found that dust continuum maps underestimate the column density by an \\Av\\ of $\\sim$6.7~mag that may be arising from the diffuse material of the cloud. \\item The orientation of the cores is not correlated with the surrounding diffuse gas magnetic field direction, which suggests that large scale magnetic fields are not important in shaping the cores. On the other hand, the lack of spherical symmetry demands an important anisotropic force, and projection effects might be important. A deeper study of the magnetic field of the dense gas is needed. \\item The analysis of the line widths reports two behaviors depending on the molecular transition: ({\\it i}) a roughly constant value of subsonic turbulent broadening for all the cores (e.g. \\cdo~(1--0) and \\chtoh~(2--1), see also \\citealp{lada08}) and ({\\it ii}) a roughly constant slightly narrower broadening for cores with \\Av\\gsim20~mag and supersonic turbulent broadenings otherwise (e.g. \\cdh~(1--0) and \\ndhp~(1--0)). \\item We observed a set of early- and late-time molecular transitions toward the cores and derived their column densities and abundances. The high spectral resolution molecular normalized line data is in agreement with the lower spectral resolution data presented in \\paperd. The nine starless cores are all very chemically young but show different chemical properties. ``Diffuse'' cores (\\Av\\lsim15~mag: 48 and 74) show emission only in ``ubiquitous'' lines typical of the parental cloud chemistry (e.g. \\co, \\cs, \\chtoh). The denser ``deuterated'' cores (\\Av\\gsim22~mag: 40 and 109) show weaker \\mis\\ for ``ubiquitous'' lines and present emission in nitrogen- (\\ndhp) and deuterium-bearing (\\cthdt) molecules, and in some carbon chain molecules (\\hctn), signposts of a prototypical dense core chemistry. ``Oxo-sulfurated'' cores (\\Av$\\simeq$15--22~mag: 6, 14, 20, and 65) are in a chemical transitional stage between cloud and dense core chemistry. They are characterized by presenting large \\mis\\ of \\chtoh\\ and oxo-sulfurated molecules (e.g. \\so\\ and \\sod) that disappear at higher densities, and they still present significant emission in the ``ubiquitous'' lines. \\chtoh\\ was detected toward the nine cores of the complete sample with abundances of $\\sim$10$^{-9}$, close to the maximum value expected for gas-phase chemistry. \\item Core 47 presents high abundance of \\chtoh\\ and \\ndhp, in spite of being the core with the lowest \\hd\\ column density, and broad line width in some species (\\cdh\\ and \\ndhp). All this is in agreement with the hypothesis given in \\paperd, which suggests that Core~47 could be a failed core. \\item The chemical evolutionary stage is not correlated with the core location in the \\pipe, but it is correlated with the physical properties of the cores (density and size). Thus, the chemically richer cores are the denser ones. The tentative correlation between magnetic field and chemical properties found for the initial subsample of four cores is less clear with the current sample. \\end{enumerate} The \\pipe\\ is confirmed as an excellent laboratory for studying the very early stages of star formation. The nine cores studied show different morphologies and different chemical and magnetic properties. Physical and chemical properties seem to be related, although important differences arise, which evidence the complex interplay among thermal, magnetic, and turbulent energies at core scales. Therefore, a larger statistics is needed to better understand and characterize the \\pipe\\ starless core evolution. In addition, other young clouds with low-mass dense cores, such as the more evolved star-forming Taurus cloud, should be studied in a similar way to prove the presented results as a general trend or, on the contrary, a particular case for a filamentary magnetized cloud." }, "1207/1207.6159_arXiv.txt": { "abstract": "The {\\bf H}$_2${\\bf O} Southern Galactic {\\bf P}lane {\\bf S}urvey (HOPS) has mapped a 100\\,degree strip of the Galactic plane ($-70\\degree>l>30\\degree$, $|b|<0.5^{\\circ}$) using the 22-m Mopra antenna at 12-mm wavelengths. Observations were conducted in on-the-fly mode using the Mopra spectrometer (MOPS), targeting water masers, thermal molecular emission and radio-recombination lines. Foremost among the thermal lines are the 23\\,GHz transitions of \\nhthree~J,K\\,=\\,(1,1) and (2,2), which trace the densest parts of molecular clouds ($n>10^4$\\,cm$^{-3}$). In this paper we present the \\nhthree\\,(1,1) and (2,2) data, which have a resolution of 2~arcmin and cover a velocity range of $\\pm200\\,\\kms$. The median sensitivity of the $\\nhthree$ data-cubes is $\\sigma_{T_{\\rm mb}}=0.20\\pm0.06$\\,K. For the (1,1) transition this sensitivity equates to a 3.2\\,kpc distance limit for detecting a 20\\,K, 400\\,$\\msun$ cloud at the 5$\\sigma$ level. Similar clouds of mass 5,000\\,$\\msun$ would be detected as far as the Galactic centre, while 30,000\\,$\\msun$ clouds would be seen across the Galaxy. We have developed an automatic emission finding procedure based on the ATNF {\\scriptsize DUCHAMP} software and have used it to create a new catalogue of 669 dense molecular clouds. The catalogue is 100 percent complete at the 5$\\sigma$ detection limit ($T_{\\rm mb}=1.0$\\,K). A preliminary analysis of the ensemble cloud properties suggest that the near kinematic distances are favoured. The cloud positions are consistent with current models of the Galaxy containing a long bar. Combined with other Galactic plane surveys this new molecular-line dataset constitutes a key tool for examining Galactic structure and evolution. Data-cubes, spectra and catalogues are available to the community via the HOPS website. ", "introduction": "HOPS ({\\bf H}$_2${\\bf O} Southern Galactic {\\bf P}lane {\\bf S}urvey) is a project utilising the Mopra radio telescope\\footnote{Mopra is a 22-m single dish mm-wave telescope situated near Siding Spring mountain in New South Wales, Australia.} to simultaneously map spectral-line emission along the southern Galactic plane across the full 12-mm band (frequencies of 19.5 to 27.5\\,GHz). Since the survey began in 2007 \\citep{walsh2008} HOPS has mapped 100 square degrees of the Galactic plane, from $l=290^\\circ$, continuing through the Galactic centre to $l=30^\\circ$ and with Galactic latitude $|b|\\leq 0.5^\\circ$. The aim of the survey is to provide an untargeted census of 22.235\\,GHz $\\water$ (6$_{16}-$5$_{23}$) masers and thermal line emission towards the inner Galaxy. Observations were completed in 2010 and the survey properties, observing parameters, data reduction and $\\water$ maser catalogue are described in \\citet{walsh2011} (hereafter Paper~I). The primary thermal lines targeted by HOPS are those of ammonia ($\\nhthree$), whose utility as a molecular thermometer in a broad range of environments is unsurpassed. With an effective critical density of $\\sim10^4$\\,$\\cm3$ \\citep{ho_townes1983}, $\\nhthree$ traces dense molecular gas and it is often associated with the hot molecular core phase of high-mass star formation (e.g., \\citealt{L07A}, \\citealt{morgan2010}), where it exhibits temperatures in excess of 30\\,K. $\\nhthree$ is excited in gas with kinetic temperatures greater than $\\sim5$\\,K \\citep{pickett1998} and is also found associated with cool ($T<10$\\,K) dense clouds. Such regions are too cold for more common gas tracers, like CO, to remain in the gas phase. Instead they are frozen out onto the surfaces of dust grains \\citep{bergin2006}. The J,K\\,=\\,(1,1) inversion transition exhibits prominent hyperfine structure, which can be used to infer the optical depth of the transition. In clouds forming high-mass stars ($\\ge8\\,\\msun$) and under optically thin conditions, the peak brightness of the four groups of satellite lines is approximately half that of the central group \\citep{rydbeck1977}. Comparison of the (1,1) and higher J,K inversion transitions can be used to estimate the rotational temperature of the gas. \\begin{figure*} \\centering \\includegraphics[ angle=90, height=22.6cm, trim=0 0 0 0]{figs/fig1_s2n_map_full.eps} \\caption{Map of the peak \\nhthree\\,(1,1) signal-to-noise (S/N) across the HOPS survey area. The map was made by smoothing the S/N cube spatially (final beam FWHM\\,=\\,2.5~arcmin) and in velocity (hanning-window\\,=\\,5) and measuring the peak brightness in each spectrum. An approximate brightness temperature scale may be found by multiplying by the average root-mean-squared noise temperature $\\langle\\sigma_{T_{\\rm mb}}\\rangle$\\,=\\,0.20\\,K.} \\label{fig:s2n_full} \\end{figure*} \\begin{figure*} \\centering \\includegraphics[ angle=90, height=22.6cm, trim=0 0 0 0]{figs/fig2_NH3_22_sigmap.eps} \\caption{Map of the peak \\nhthree\\,(2,2) signal-to-noise (S/N) across the HOPS survey area. } \\label{fig:s2n_full_22} \\end{figure*} In this second paper we present the HOPS $\\nhone$ and $\\nhtwo$ datasets and the automatic finding procedure used to create catalogues of emission. We illustrate the basic properties of the catalogues, which form the basis for further analysis. A third paper (Longmore et al. {\\it in prep.}) will published the properties of the catalogue derived by fitting the $\\nhone$ and $\\nhtwo$ spectra with model line profiles (i.e. temperature, density, mass and evolutionary state). ", "conclusions": "We have mapped 100 square degrees of the Galactic plane in the J,K\\,=\\,(1,1) and (2,2) transitions of $\\nhthree$ at a resolution of $2$~arcmin using the Mopra radio-telescope. The survey covers the region $-70\\degree 3/2$ is replaced by an initially decelerating wind that reaches a minimum speed and then proceeds to accelerate for $l>r_o$. We showed that the equivalent nozzle function can be used as a means to gauge whether or not the velocity is monotonic without actually finding the transonic solutions. Meliani et al. (2004) also employed equivalent nozzle functions in their relativistic generalization of the polytropic Parker problem using the Schwarzschild metric. There, winds were also found to accelerate for $\\gamma>3/2$, but the velocity profiles were still monotonic. We showed velocity minima to be an effect of adding rotation. Our discussion of the spherically symmetric and rotating Parker wind solutions showed that significant deceleration is associated with a flow regime characterized as having an enthalpy deficit. It is not inconceivable that this type of outflow can exist in an astrophysical setting, so it would be interesting to determine the spectral signatures of a decelerating wind region. We tied the enthalpy deficit regime (i.e. the parameter space giving solutions with $v_\\infty < v_o$) to the appearance of degenerate transonic wind solutions, and we further pointed out that the critical point behavior of the second set of transonic solutions is remarkably similar to that reported by Cur\\'e (2004) in his study of \\textit{isothermal} line-driven stellar wind equations with rotation. Cur\\'e (2004) classified his new solutions as `slow', since they obtain significantly smaller terminal velocities. Might these slow solutions be a different guise of the enthalpy deficit regime? Our main objective was to investigate the dynamical properties of two axially symmetric, thermally driven disc wind models, one with significant adjacent streamline divergence (the Converging model) and another with a complete lack thereof (the CIA model). We emphasize that detailed hydrodynamical simulations show that by taking into account the interactions of neighboring flow tubes, a self-similar streamline geometry emerges (recall \\fig{stefans}). We have neglected to mention elsewhere that CIA-like streamlines have also been found analytically from similarity solutions of idealized models for galactic superwinds. Both the self-similar solutions of Bardeen \\& Berger (1978), which took into account a gravitational potential, and those of Zirakashvili \\& V\\\"olk (2006), which did not involve gravity, are examples. Since the CIA and Converging models only differ by their respective amounts of streamline divergence, we can attribute the differences in the properties of their solutions as being solely due to geometric effects. Our results have implications for kinematic models that adopt a flow geometry similar to the Converging model for the purposes of computing synthetic spectra to compare with observations. Namely, use of the Converging model will significantly overestimate the acceleration of the flow if the true wind configuration more closely resembles the CIA model. The latter model, due to its smaller amount of streamline divergence, features a greatly extended acceleration zone, a more distant sonic surface, a shallower density and temperature falloff, and a smaller mass flux density than the Converging model for similar footprint conditions. Conversely, for a given mass flux density at the wind base, the CIA model will predict a higher density, implying that synthetic line profiles will exhibit stronger line absorption. Ultimately, larger error bars may need to be associated with the inferred mass-loss rate, as spherically diverging winds may tend to over-estimate it. Solving the time-\\textit{dependent} problem is likely to yield insights into the full domain of viable disc wind solutions. Especially considering that we found degenerate transonic solutions, uncovering the effects of time-dependence is a worthwhile task, one that we plan to undertake in a future work. It is likely that one of the degenerate solutions is unstable, and this could be verified using hydrodynamical simulations. The alternative would be more exciting, however, as it is conceivable that the time-dependent solution can settle upon both solutions under various circumstances. \\bigskip \\noindent\\textbf" }, "1207/1207.7038_arXiv.txt": { "abstract": "Outflows from active galactic nuclei may be produced by absorption of continuum radiation by UV resonance lines of abundant metal ions, as observed in broad absorption line quasars (BALQs). The radiation pressure exerted on the metal ions is coupled to the rest of the gas through Coulomb collisions of the metal ions. We calculate the photon density and gas density which allow decoupling of the metal ions from the rest of the gas. These conditions may lead to an outflow composed mostly of the metal ions. We derive a method to constrain the metals/H ratio of observed UV outflows, based on the \\Lya\\ and \\SiIV~$\\lambda\\lambda$1394, 1403 absorption profiles. We apply this method to an SDSS sample of BALQs to derive a handful of candidate outflows with a higher than solar metal/H ratio. This mechanism can produce ultra fast UV outflows, if a shield of the continuum source with a strong absorption edge is present. ", "introduction": "Broad absorption line quasars (BALQs) are a subclass of active galactic nuclei (AGNs) with spectra which exhibit broad, strong, and blue-shifted absorption features (\\citealt{reichardetal03} and references there in). Approximately 15 per cent of quasars are classified as BALQs \\citep{reichardetal03, kniggeetal08, sca09}. Although there are some differences in the continuum and emission lines between BALQs and non-BALQs, they appear to be drawn from the same parent population \\citep{weymannetal91, reichardetal03}. The broad absorption lines (BALs) are the most prominent spectral feature indicating an outflow of material driven from and by the AGN. This outflow is a compelling candidate for the feedback mechanism of the AGN on its host. The outflow is commonly invoked as a growth and activity regulator (\\citealt*{dimatteoetal05} and citations thereafter) and as a mechanism of metal enrichment of the ISM and IGM \\citep{molletal07}. How do the outflows form? Although much theoretical work has been done to model AGN outflows, a work which stems from studies of stellar winds (e.g., \\citealt{lamcas99}), the driving physical mechanism of outflows remains uncertain. The first models involved only gas and radiation pressure as a driving force of the outflow (\\citealt{vitshl88,arali94}; \\citealt*{aravetal94}; \\citealt{murrayetal95}; \\citealt*{progaetal98}; \\citealt{chenet01}). Later work introduced more detailed models of radiation pressure (\\citealt*{progaetal99}; \\citealt{chenet03a}) and magnetic fields \\citep{dekbeg95}. The effect of radiation-driven wind on the emerging spectrum \\citep{arabeg94} and galactic bulge \\citep*{fabianetal06} has also been considered. More recently, with the advance in computational power, more detailed magnetohydrodynamical and magnetocentrifugal models of the outflow have been developed (\\citealt{konkar94,proga00}; \\citealt*{progaetal00}; \\citealt{proga03, everett05, kurpro08}). Currently, the models produce almost no predictions which can help discriminate between them using observations (see a review by \\citealt*{crenshawetal03}). In current models the outflow is treated as a single-component fluid, and the radiation pressure is parametrized by a ``force multiplier'', which encapsulates the effects of radiation pressure due to Thomson scattering and line and dust absorption. What effect does the radiation pressure have on the different constituents of the fluid? As was first pointed out by \\citet{sp92}, for radiation-driven stellar winds, a single-component fluid treatment overlooks the possibility of metal ions decoupling from the mostly-H outflow. \\citet{sp92} have shown that metal ions can run away (i.e., decouple) from the proton-electron fluid, if the coupling between ions and the fluid is small enough, compared to the photon flux absorbed by the ions. Photon absorption by a metal ion results in a net gain of momentum i.e., ion acceleration (assuming illumination by an anisotropic source), while the acceleration of H is inefficient as H is mostly ionized. The ion is decelerated by Coulomb scatterings by the ambient protons and electrons. The deceleration is velocity dependent and is maximal when the ion has a velocity similar to the thermal velocity of each of the constituents of the proton-electron fluid. The ion can run away from the fluid, if the rate of photon absorption (i.e., flux density) is large enough, so that the accumulated gain in velocity of the ion surpasses the fluid thermal velocity before the ion is significantly deflected by Coulomb scatterings i.e., if the acceleration due to the radiation pressure is larger than the deceleration due to the Coulomb force. The metal decoupling prevents the proton-electron fluid from accelerating further and the decoupled metals will produce a fast wind through the proton-electron fluid. Metal ion decoupling from the mostly-H gas may be relevant in AGN outflows due to two reasons. First, gas in close vicinity of the active nucleus is exposed to a large flux of radiation; second, strong ion absorption is observed. Our goal is to explore whether the occurrence of metal decoupling is a plausible process in AGN. Self-consistent models of an outflow and the structure of the ionized gas region, which incorporate the process of metal ion decoupling, are beyond the scope of this study. How can an occurrence of the metal decoupling be observationally verified? The metal decoupling will mostly affect the absolute metallicity of a wind i.e., metal to H abundance ratio. The decoupling effect on the relative i.e., metal to metal abundance ratio, might be negligible. Since the wind can be composed from several types of decoupled metals, the relative abundance ratio of the wind might be similar to that of the wind origin, where the metals are mixed with H. Thus, indirect (i.e., relative) outflow metallicity measurements cannot be used to test the proposed metal enrichment scenario. Direct (i.e., absolute metals to H) metallicity measurements are needed to test the validity of the enrichment scenario in AGN. There is accumulating evidence of metal abundances higher than solar for the broad line region (BLR; \\citealt{hamann97}, \\citealt{hamfer99}). There is also evidence for higher than solar metal abundances in AGN outflows from analysis of narrow absorption lines (\\citealt{aravetal07, wuetal10, hamannetal11}) and broad absorption lines (\\citealt{hamannetal97}). However, for some objects the results are inconclusive (\\citealt{hamann98}) or consistent with solar abundances (\\citealt{aravetal01}). The main difficulty in estimating the element abundance arises from a large uncertainty in the measurement of ionic column densities when the covering factor is velocity dependent (\\citealt{hamann98, aravetal99}). Note however that \\citet{cottisetal10} disfavour radiation-driven outflow altogether, due to the absence of an excess of objects displaying line-locking between \\Lya\\ and \\NV. A metallicity of a few times solar can be explained by stellar enrichment due to extended star formation (e.g., \\citealt{hamannetal02}), as suggested in high-redshift quasars ($z\\geq3.5$, \\citealt{dietrichetal03}). However, if a metal ions runaway occurs, the metal/H column ratio can be larger by orders of magnitude, compared to the solar abundance ratio. Below we derive the conditions for an ion runaway in gas in the vicinity of an active nucleus. We also present a method to set a robust and direct lower limit on the metal/H ratio in outflows based on UV absorption lines. The method is then implemented on a sample of AGN from the Sloan Digital Sky Survey (SDSS; \\citealt{york00}) BALQs catalog \\citep{sca09}. The ion runaway model is outlined in Sec.~2. The direct metallicity estimation method and its implementation are described in Secs.~3 and 4, respectively. We discuss the results in Sec.~5. Our conclusions are summarized in Sec.~6. ", "conclusions": "We analyse the conditions required for a radiation-pressure-driven pure metal outflow, and possible observational evidence for its existence. We find the following: \\begin{enumerate} \\item A metal ion embedded in gas subject to radiation pressure will run away when $\\log f_\\nu(\\lambda_{\\rm ion})-\\log\\nh+ \\log T>-6.2.$ (Eq.~\\ref{eq:runaway}). For an average AGN SED this can be converted to $\\log U\\ga5$. \\item To avoid overionization the gas must be protected by a shield, most likely in the form of a strong He$^+$ edge at 4~Ryd. \\item Photoionization models indicate for solar metallicity an absolute minimum of $N(\\mbox{H$^0$})=2.5N(\\mbox{Si$^{3+}$})$, implying $\\tau(\\Lya)\\geq1.8\\tau(\\mbox{\\SiIV\\ $\\lambda$1394})$, regardless of $U$ and $\\Sigma$. \\item Thus, a comparison of $\\tau(\\SiIV\\,\\lambda\\lambda1394,\\,1403)$ and $\\tau(\\Lya) $, for different absorption geometries, can yield a direct constraint on the metals/H abundance ratio. \\item A search of the SDSS BALQ sample of \\citet{sca09} yields a handful of possible candidates for an outflow with a metallicity above solar. However, no robust evidence (i.e., $Z\\ga10Z_{\\sun}$) is found for a metal runaway. \\item High quality UV observation of lower $z$ quasars, free of intervening \\Lya\\ absorption, can be used to obtain better constraints on the outflow metallicity. \\item An ultra fast outflow of metals is expected to be produced if a runaway takes place. Since the metal ions are not expected to be collisionally coupled, different ions will likely have different terminal velocities. \\end{enumerate} The physical conditions in AGN are definitely more complex than the simplistic scenario outlined in this paper. Clearly, models which simulate more accurately the various microphysics involved, including the caveats mentioned above, are required in order to get some quantitative predictions on metal ions runaway in radiation pressure driven outflows in AGN. Such models can then be integrated into simulations of the large scale structure of AGN outflows." }, "1207/1207.6722_arXiv.txt": { "abstract": "{The radio-quiet quasar BR1202--0725 ($z=4.695$) is a remarkable source with a bright Northwest (NW) companion detected at submm and radio wavelengths but invisible in the optical. In the absence of amplification by gravitational lensing, BR1202--0725 would be the most luminous binary CO and far-IR source in the Universe. In this paper, we report observations with the IRAM Plateau de Bure interferometer of BR1202--0725 in the redshifted emission of the CO(5--4) and (7--6) lines, the [C{\\,\\small I}]($^3$P$_2-^3$P$_1$) line, a high angular resolution ($0.3'' \\times 0.8''$) 1.3\\,mm map of the rest-frame, far-IR dust continuum, and a search for the CO(11--10) line. We compare these results with recent ALMA data in the [C{\\,\\small II}] line. Both the quasar host galaxy and its NW companion are spatially resolved in the molecular line emission and the dust continuum. The CO profile of the NW companion is very broad with a full width at half maximum of $\\sim 1000 \\pm 130$\\,\\kms, compared to $\\sim 360\\pm40$\\,\\kms\\ for the quasar host galaxy to the Southeast (SE). The difference in linewidths and center velocities, and the absence of any lens candidate or arc-like structure in the field, at any wavelength, show that the obscured NW galaxy and the SE quasar host galaxy cannot be lensed images of the same object. Instead, we find morphological and kinematic evidence for sub-structures in both the NW and SE sources. We interpret these results as strong indications that the BR1202--0725 complex is a group of young, interacting, and highly active starburst galaxies. ", "introduction": "Submillimeter observations of galaxies and quasars at high redshift provide invaluable clues about the activity of galaxy mergers leading to the formation of massive galaxies, and about the star-forming environment when the universe was young. These clues come from the redshifted emission of dust heated by newly-formed stars and from molecular and atomic lines from the dense molecular gas in which the stars are born. Since the early 1990s, massive galaxies at high redshift have been observed with ground-based millimeter and submm telescopes that can detect submm-selected starburst galaxies and host galaxies of optically selected quasars. These detections are pushing more and more into the epoch of re-ionization, and are currently out to $z=7.1$; (Venemans et al.\\ 2012). \\begin{figure*} \\includegraphics[width=8.6cm,angle=-0]{f1a.eps} \\includegraphics[width=8.6cm,angle=-0]{f1b.eps} \\caption[CO54 uv-average maps of SE and NW.]{ Integrated CO(5--4) line maps, optimized to show total emission in the SE and NW galaxies separately. Source centroids and sizes are listed in Table~1. Both maps have natural weighting with a beam of $1.6'' \\times 0.7''$ (PA 15$^{\\circ}$; lower-left box in left panel). {\\it Left:} CO(5--4) integrated over a 444\\,\\kms-wide band, centered on the SE galaxy's line center velocity. Contour steps are 0.24\\,mJy (1$\\sigma$). The integrated CO(5--4) flux of the whole SE source, including the extension to the SW, is 1.6\\,Jy\\,\\kms . {\\it Right:} CO(5--4) integrated over a 1067\\,\\kms-wide band, centered on the NW galaxy's line center velocity. Contour steps are 0.15\\,mJy (1$\\sigma$) The integrated CO(5--4) flux of the whole NW source, including the extension to the northwest, over the 1067\\,\\kms\\ band, is 2.6\\,Jy\\,\\kms\\ (see Table~2) } \\label{fig:co54integ} \\end{figure*} One of the first high-$z$ quasars that was detected and studied in detail in both thermal dust emission and in CO lines is the very luminous optically selected quasar BR1202--0725 at $z=4.7$ (McMahon et al.\\ 1994; Ohta et al.\\ 1996; Omont et al.\\ 1996; Benford et al.\\ 1999; Carilli et al. 2002; Iono et al.\\ 2006; Riechers et al.\\ 2006). The source is remarkable because it has two distinct components: a Southeast (SE) source associated with the optically luminous quasar and an obscured Northwest (NW) source with no counterpart in the visible. The two submm dust continuum peaks, separated by $3.7''$ (linear distance 24\\,kpc),\\footnote{ For linear sizes and luminosity distances we took the standard cosmology with $H_{0}$ = 71\\,\\kms\\,Mpc$^{-1}$, $\\Omega_{M}$ = 0.27, $\\Omega_{\\Lambda}$ = 0.73 (Komatsu et al.\\ 2011), and used the cosmological calculator by Wright (2006).} have nearly identical CO redshifts of $z=4.695$ (SE) and $z=4.693$ (NW). The extreme infrared luminosity makes BR1202--0725 one of the brightest high-$z$ sources ($L_{IR}$ = 3.7$\\times 10^{14}$\\,\\Lsun, Leipski et al.\\ 2010, confirming the original estimate by McMahon et al.\\ 1994). The far-IR part of the luminosity is $L_{\\rm FIR}$ = 3.8$\\times 10^{13}$\\,\\Lsun, with 1.2 and 2.6$\\times 10^{13}$\\,\\Lsun\\ for the NW and SE galaxies (Iono et al.\\ 2006). The estimated star-formation rate in each component is $\\geq \\, 1000$\\,M$_{\\odot}$\\,yr$^{-1}$ and the combined molecular gas mass is $\\rm \\sim10^{11} \\, M_{\\odot}$ (Omont et al.\\ 1996; Riechers et al.\\ 2006). The SE optical quasar host galaxy and the optically obscured NW source are both detected in the CO $J= 1-0$ through $7-6$ rotational lines (Ohta et al.\\ 1996; Omont et al.\\ 1996; Carilli et al.\\ 2002; Riechers et al.\\ 2006) as well as in the [C{\\,\\small II}] emission line, the main cooling line in the interstellar medium (Iono et al.\\ 2006; Wagg et al.\\ 2012). At optical wavelengths, the SE component of BR1202--0725 appears as a quasar with at least two faint companion galaxies, seen in the optical (starlight) continuum (see our Fig.~6). The first of these optical continuum sources, found by Fontana et al.\\ (1996) and Hu et al.\\ (1996), is centered 2.6$''$ northwest of the quasar, near to, but not coinciding with, the obscured NW submm source, and it appears to be part of a filament of optical starlight continuum emission extending over 4$''$ away from the quasar. Part of this optical continuum source is prominent in Ly$\\alpha$ emission at $z$ =4.702 (Hu et al.\\ 1996; Petitjean et al.\\ 1996; Ohyama et al.\\ 2004). The second, fainter, optical starlight continuum source is also a Ly$\\alpha$ emission galaxy, located 3$''$ southwest of the quasar (Hu et al.\\ 1997). Up to now, detailed study of the profiles of the molecular and carbon fine structure lines in BR1202--0725 have been hampered by lack of sufficient bandwidth to cover the very broad molecular emission lines of the NW galaxy (Omont et al.\\ 1996; Carilli et al.\\ 2002). To remedy this situation, we present in this paper new wide-bandwidth spectral-line observations of CO and [C{\\,\\small I}] emission in BR1202-0725, along with high-angular resolution observations of the 1.3\\,mm thermal dust continuum. These new data from the Plateau de Bure interferometer (PdBI) are compared with submm data obtained with ALMA that trace the [C{\\,\\small II}] emission line and the 0.9\\,mm thermal dust continuum emission (Wagg et al.\\ 2012). Our observations improve significantly on previous studies and shed new light on the kinematics of the SE and NW sources as well as on their structures. These data indicate that BR1202--0725 is a very complex group of merging starburst galaxies. ", "conclusions": "Our high-resolution millimeter observations of the BR1202-0725 galaxy complex at $z=4.7$, in both line and continuum emission, reveal new aspects of this extreme multiple merger. Morphological and dynamical evidence together with the lack of any lens candidate or arc-like structure in the field, at any wavelength, definitely rule out the hypothesis that the obscured NW galaxy and the SE quasar host galaxy are lensed multiple images of the same background object. Instead, the present data indicate that the images of BR1202--0725 directly show a group of merging and very active starburst galaxies (Fig.~6, right). The SE galaxy merger contains at least two sources. The main SE core source coincides with the optical quasar's position and shows internal kinematics compatible with rotation and a CO line profile (FWZP= 700 km/s) typical of a massive disk galaxy. To the southwest, the CO(5--4) line is narrower, with a linewidth of 200\\,\\kms, and blueshifted by $-$180\\,\\kms\\ relative to the main SE velocity peak. This fainter CO does not coincide in position with the weaker SW continuum source reported in the ALMA data (Wagg et al. 2012), which may be a third, fainter, galaxy in the merger. The NW galaxy contains at least two sources, a main core source, and a fainter companion to the west/northwest at a velocity different by 600\\,\\kms\\ from the main NW core. This second galaxy in the interaction accounts for the unusually broad total CO linewidth (FWZP $>$ 1500 km/s) of the NW galaxy merger. The far-IR luminosities of both sources are high, of order 10$^{13}$\\,\\Lsun, indicating tremendous star-forming activity, with estimated star formation rates greater than 1000 \\,\\Msun\\,yr$^{-1}$. Since BR1202--0725 is radio quiet, its activity must be dominated by the massive starbursts that occur simultaneously in the galaxies interacting in the NW and SE components. We are obviously witnessing an extreme merging event in a group of galaxies, probably the youngest complex merger so far known. Future higher angular resolution observations of BR1202--0725 will allow us to better disentangle the properties of the individual starbursts that are interacting and merging, resulting in the remarkable appearance of this high-redshift object." }, "1207/1207.4727_arXiv.txt": { "abstract": "The recurrent outbursts that characterise low-mass binary systems reflect thermal state changes in their associated accretion discs. The observed outbursts are connected to the strong variation in disc opacity as hydrogen ionises near 5000 K. This physics leads to accretion disc models that exhibit bistability and thermal limit cycles, whereby the disc jumps between a family of cool and low accreting states and a family of hot and efficiently accreting states. Previous models have parametrised the disc turbulence via an alpha (or `eddy') viscosity. In this paper we treat the turbulence more realistically via a suite of numerical simulations of the magnetorotational instability (MRI) in local geometry. Radiative cooling is included via a simple but physically motivated prescription. We show the existence of bistable equilibria and thus the prospect of thermal limit cycles, and in so doing demonstrate that MRI-induced turbulence is compatible with the classical theory. Our simulations also show that the turbulent stress and pressure perturbations are only weakly dependent on each other on orbital times; as a consequence, thermal instability connected to variations in turbulent heating (as opposed to radiative cooling) is unlikely to operate, in agreement with previous numerical results. Our work presents a first step towards unifying simulations of full MHD turbulence with the correct thermal and radiative physics of the outbursting discs associated with dwarf novae, low-mass X-ray binaries, and possibly young stellar objects. ", "introduction": "\\label{intro} Recurrent outbursts in accreting systems are commonly attributed to global instabilities in their associated accretion discs. In particular, it is believed that thermal instability driven by strong variations in the disc's cooling rate causes the observed state transitions in dwarf novae (DNe) and low-mass X-ray binaries (LMXBs) (Lasota 2001), while the rich array of accretion variability associated with quasars could be excited by thermal instability driven by variations in the turbulent heating rate (Shakura and Sunyaev 1976, Abramowicz et al.~1988), by thermal-viscous instability (Lightman \\& Eardley 1974), or assorted dynamical instabilities (e.g.\\ Papaloizou \\& Pringle 1984, Kato 2003, Ferreira \\& Ogilvie 2009). On the other hand, FU Ori outbursts, characteristic of protostellar discs, probably result from the interplay of thermal, gravitational, and magnetorotational instability (MRI) across the intermittently inert dead zone (Gammie 1996, Balbus \\& Hawley 1998, Armitage et al.~2002, Zhu et al.~2009, 2010). Exceptions to this class of model include classical novae whose outbursts can be traced to thermonucleur fusion of accreted material on the white dwarf surface (Gallagher \\& Starrfield 1978, Shara 1989, Starrfield et al. 2000). The most developed, and possibly most successful, model of accretion disc outbursts describes the recurrent eruptions that characterise the light curves of DNe and LMXBs. In this model, the ionisation of hydrogen, occurring at temperatures of around 5000 K, induces a significant opacity change in the disc orbiting the primary star (Faulkner et al.~1983, hereafter FLP83), which then leads to thermal instability and hysteresis in the disc gas (Pringle 1981). The system exhibits a characteristic `S-curve' in the phase plane of surface density $\\Sigma$ and central temperature $T_c$ (cf. Fig.~1), and as a consequence the gas falls into one of two stable states: an optically thick hot state, characterized by strong accretion, or an optically thin cooler state, in which accretion is less efficient (e.g. Meyer \\& Meyer-Hofmeister 1981, FLP83). Outbursts can then be modelled via a `limit cycle', whereby the disc jumps from the low accreting state to the high accreting state and then back again as mass builds up and is then evacuated. This mechanism, based on local bistable equilibria, is the foundation for a variety of advanced models, which have enjoyed significant successes in reproducing the observed behaviour of accretion discs in binary systems, even if interesting discrepancies persist (Lasota 2001). The classical model of DNe and LMXBs treats the disc as laminar and assumes that the disc turbulence can be modelled with the $\\alpha$ prescription (Shakura \\& Sunyaev 1973), whereby the action of the turbulent stresses is characterized as a diffusive process with an associated `eddy' viscosity (Balbus \\& Papaloizou 1999). This has been sufficient to sketch out the qualitative features of putative outbursts, but such a crude description has imposed limitations on the formalism that are now impeding further progress (Lasota 2012). On the other hand, fully consistent magnetohydrodynamic (MHD) simulations of disc turbulence generated by the MRI have been performed for almost 20 years (Hawley et al.~1995, Stone et al.~1996, Hawley 2000, etc), though typically with simplified thermodynamics (e.g.\\ isothermality). It is the task of this paper to begin the process of uniting the classic thermal instability models of DNe and LMXBs with full MHD simulations of the MRI, and thus consistently account for both the turbulence and radiative cooling. The first step we take is limited to local models, as even though discs undergo global outbursts, the ability of local annuli to exhibit hysteresis behaviour is key. Local studies will permit us to assess if and how the classic model of thermal instability can operate in the presence of realistic turbulent heating. At the same time, they provide an excellent test of the MRI itself; if the MRI is to remain the chief candidate driving disc accretion it must fulfil its obligations to classical disc theory. We undertake unstratified shearing box simulations of the MRI that include Ohmic and viscous heating and a radiative cooling prescription that is able to mimic the transition between the optically thin and thick states (FLP83). These simulations clearly reproduce S-curves in the $\\Sigma$ and mean temperature plane, and these constitute the main result of our paper. We can thus animate local thermal limit cycles via a sequence of local box runs. In particular, the simulated turbulent heating is found to be `well-behaved' and not so intermittent as to prematurely disrupt the thermal limit cycles required by the classical theory. In addition, simulations with differing net toroidal and vertical fluxes produce S-curves that exhibit variable mean values of $\\alpha$, suggesting that when global simulations are considered, and net-flux is no longer conserved locally, variations in effective $\\alpha$ between the upper and lower branches of the S-curve may be produced, as required by the classical model (e.g.\\ Smak 1984a). Finally, the simulated short-term behaviours of the average viscous stress and the disc pressure reveal only a weak functional dependence, as remarked upon by previous authors (Hirose et al.~2009). This poor correlation means that proposed thermal instabilities driven by a direct heating response to imposed pressure perturbations are unlikely to function in radiation-pressure dominated accretion discs (e.g.\\ Shakura \\& Sunyaev 1976, Abramowicz et al.~1988). This is in marked contrast to the thermal instabilities implicated in DNe and LMXBs, which are driven by the cooling response to imposed temperature perturbations, mediated via opacity variations. Note, however, that over \\emph{long times} we find that the stress and pressure are correlated. The plan of the paper is as follows. In the following section we discuss the thermal instability model for DNe and LMXBs in more detail. In Section 3 we present the governing equations while outlining the radiative cooling prescription the simulations adopt. In section 4 we give the numerical details of the simulations. Our results are presented in Section 5, and we conclude in Section 6. ", "conclusions": "\\label{sec:disc} In this paper we performed a suite of numerical simulations of the MRI in local geometry with the ZEUS and NIRVANA codes. Both viscous and Ohmic heating is included, while the radiative cooling is approximated by a physically motivated cooling function that summarises the strong effect of temperature on the ability of the disc gas to retain heat. Different magnetic configurations (zero flux, net-toroidal flux, net-vertical flux) and parameters were trialed, with little change in the qualitative results. Our simulations unambiguously exhibit the development of thermal instability and hysteresis. In particular, through a sequence of runs we can sketch out characteristic S-curves in the phase space of $(\\Sigma, T_c)$ and $(\\Sigma, \\dot{M})$, which are central to the classical outburst model. It hence appears that MRI turbulence is not so intermittent as to endanger the robustness of the cycle. Temperature fluctuations are well-behaved and relatively small and there is no spontaneous jumping from one stable branch to the other. Only near the `corners' of the S-curve does significant stochasticity enter, as then the existence or not of a local equilibrium is uncertain. This feature will add some degree of low level `noise' to the observed outburst time-series. In addition, the $\\alpha$ we record on the two stable branches differ slightly but systematically. Because of the constraints of our local model, this result is more suggestive than anything else. However, it does indicate that in global disc simulations we may indeed see a systematic difference in the two branches, as required by the classical theory. Finally, on the orbital time the turbulent stresses and pressure only weakly depend on each other in our simulations. Pressure always lags behind the turbulent stress, and thus the causality is from the stress to the pressure, via the turbulent heating (in agreement with Hirose et al.~2009). Moreover, the pressure response is a significantly `smoothed out' echo of the turbulent stress features. Both facts indicate that thermal instability driven by turbulent heating variations, in which the stress is a function of pressure, may not operate in real discs. On long times, however, there is necessarily a feedback of the pressure on the stress, which leads to different accretion rates on the two branches of the S-curve. Our work presents a first step towards unifying simulations of full MHD turbulence with the correct thermal and radiative physics of outbursting DNe and LMXBs, and possibly young stellar objects. We have begun with with the most basic model that yields the correct physics: local simulations of gas inhabiting a single radius in the disc with a simple radiative prescription. This is a `ground zero' test of the compatibility of the MRI with the putative thermal limit cycles of outbursting discs. If MRI-turbulence had failed in this basic setting it would probably fail in more advanced models as well. Now that we are assured of this compatability, a variety of further work may be attempted. For instance, simulations could be performed in vertically stratified shearing boxes with more realistic radiation physics (in the flux-limited diffusion approximation with appropriate opacities). In addition, cylindrical MRI simulations should be performed with the FLP83 cooling prescription, anologous to the global $\\alpha$ disc calculation of Papaloizou et al.~(1983). Such simulations would describe how real turbulence mediates the heating and cooling fronts that propagate through the disc during a transition between branches." }, "1207/1207.1102_arXiv.txt": { "abstract": "Using multiple tracers of large-scale structure allows to evade the limitations imposed by sampling variance for some parameters of interest in cosmology. We demonstrate the optimal way of carrying out a multitracer analysis in a galaxy redshift survey by considering the principal components of the shot noise matrix from two-point clustering statistics. We show how to construct two tracers that maximize the benefits of sampling variance and shot noise cancellation using optimal weights. On the basis of high-resolution $N$-body simulations of dark matter halos we apply this technique to the analysis of redshift-space distortions and demonstrate how constraints on the growth rate of structure formation can be substantially improved. The primary limitations are nonlinear effects, which cause significant biases in the method already at scales of $k<0.1h\\mathrm{Mpc}^{-1}$, suggesting the need to develop nonlinear models of redshift-space distortions in order to extract the maximum information from future redshift surveys. Nonetheless we find gains of a factor of a few in constraints on the growth rate achievable when merely the linear regime of a galaxy survey like EUCLID is considered. ", "introduction": "One of the deepest mysteries of contemporary cosmology is the nature of the observed accelerated expansion of the Universe. So far, its evolution can be described remarkably well by Einstein's theory of gravitation including a nonzero cosmological constant $\\Lambda$. However, in order to fit the astronomical observations (e.g.,~\\cite{Perlmutter1999}), $\\Lambda$ must be many orders of magnitude smaller than what our standard model of particle physics would expect. This \\emph{hierarchy problem} inspired various departures from the cosmological standard model, such as modifications of Einstein's field equations or the introduction of exotic forms of matter. A particularly sensitive probe of cosmology is the large-scale structure (LSS) of the Universe. Galaxy redshift surveys map out large fractions of its observable volume and thereby reconstruct a three-dimensional map of density fluctuations whose statistical properties directly relate to fundamental cosmological parameters \\cite{BigBOSS,EUCLID,SDSS-III}. Unfortunately, this reconstruction is hampered by the fact that galaxies are \\emph{biased} and \\emph{stochastic} tracers of the dominating dark matter density field. Even if the density field could be inferred perfectly well, the finite number $N_k$ of independent Fourier modes in the survey sets a fundamental lower limit on the achievable uncertainty, which is known as \\emph{sampling variance} (or \\emph{cosmic variance}, in case the survey size is the whole observable Universe). For example, in a measurement of the dark matter power spectrum $P$, the sampling variance limit is given by $\\sigma_P/P\\ge\\sqrt{2/N_k}$, an uncertainty floor that propagates into all the parameters one wants to infer from $P$. This limit decreases towards smaller scales as more and more Fourier modes can be sampled, but at the same time linear theory starts to break down and higher-order perturbation theory has to be adopted to model $P$ (see, e.g., \\cite{Bernardeau2002,Crocce2006a,Taruya2008a,Matsubara2008}). Further complication arises in relating the observed galaxy power spectrum to the latter, as galaxy bias becomes nonlinear and nonlocal \\cite{ChuenChan2012,Baldauf2012}. An alternative approach to accurately probe cosmological parameters is to consider multiple tracers of the density field within the well-understood linear regime. The relative clustering amplitude between multiple tracers can be inferred without sampling variance limitation because the underlying density fluctuations cancel out in taking ratios~\\cite{Seljak2009a,McDonald2009}. Therefore, any cosmological information that remains in the relative clustering amplitude between different tracers can potentially be inferred with a much higher accuracy. In this paper we focus on a particular contribution to the clustering amplitude of galaxies coming from redshift-space distortions (RSD). These are caused by peculiar velocities along the line of sight, causing their clustering statistics to become anisotropic. First treated as a contamination, this effect has been realized to be a powerful probe of cosmology, as an understanding of RSD allows to infer the growth rate of structure formation, which is directly tied to the expansion history of the Universe as well as the theory of gravity~(e.g., \\cite{Guzzo2008,Percival2009,Cabre2009,Crocce2011,Blake2011,Ross2011,Nusser2012,Beutler2012,Reid2012,Samushia2012b,Raccanelli2012}). After recapping the fundamentals of RSD and introducing a general formalism for the multitracer analysis in Sec.~\\ref{sec1}, we present our results on the RSD analysis from $N$-body simulations in Sec.~\\ref{sec2}. Finally we draw our conclusions in Sec.~\\ref{sec3}. ", "conclusions": "} In this paper we investigated the benefits of using weights in a multitracer analysis of LSS, with a particular focus on constraining the growth rate of structure formation. On the basis of earlier results on the clustering properties of dark matter halos and their stochasticity \\cite{Hamaus2010}, we argue that the gains from a multitracer analysis in the sense of \\cite{McDonald2009} can be achieved by considering only the two principal components of the clustering signal-to-noise ratio $\\Sigma$ (or, equivalently, the two non-Poisson eigenvectors of the shot noise matrix $\\E$). We present their explicit functional forms in terms of weights, showing that the first one coincides with the weighting function explored in previous work \\cite{Hamaus2010}, giving rise to low stochasticity and high bias. For the second one the weights are also mass dependent, but have a zero crossing, such that the overall bias is low. This yields a high relative galaxy bias $\\alpha$ between the two tracers, maximizing the Fisher information content on the cosmological parameters \\cite{McDonald2009}. All of the other eigenvectors oscillate around zero and add very little information. The advantage of reducing the information to two eigenvectors is that all of the objects in a given catalog can be used to construct the two principal components, while in the conventional multitracer analysis the catalog has to be split into many lower number density subsamples with higher shot noise \\cite{Gil-Marin2010,Bernstein2011}. \\begin{figure}[!t] \\centering \\resizebox{\\hsize}{!}{ \\includegraphics[trim=0 0 0 0,clip]{fig8.eps}} \\caption{Halo model prediction for the improvement on $\\sigma_\\beta$ from the single-tracer analysis to the optimal multitracer analysis (blue) at $k\\simeq0.016h\\mathrm{Mpc}^{-1}$ as a function of the lower halo-mass threshold $M_\\mathrm{min}$ at $z=0$ (solid) and $z=1$ (dashed). The results for the combined analysis of uniform halos (not weighted) and dark matter (green), as well as the combination of optimally weighted halos with the dark matter (red) are also depicted.} \\label{fig8} \\end{figure} On the basis of numerical $N$-body simulations of dark matter halos, we demonstrate that the constraints on $\\beta$ and $f\\sigma_8$ can be improved by up to a factor of 4 relative to a single-tracer method, but most of the improvement comes from large scales (low $k$), while for higher $k$ the gains are smaller and vanish above $k\\sim 0.1h\\mathrm{Mpc}^{-1}$, where nonlinear effects introduce additional stochasticity between the two tracers. This technique is fairly insensitive to the observational uncertainty on the halo masses, as even a $50\\%$ log-normal scatter does not degrade the improvements significantly. Halo model considerations suggest even higher gains of the method with increasing mass resolution. One potential concern for our method is the possibility that galaxies might be bad tracers of their host-halo centers \\cite{Skibba2011} and therefore exhibit less pronounced principal components in the clustering signal-to-noise ratio that are distinct from Poisson sampling. However, there are strong indications that certain types of galaxies do show strong correlations in both position and mass with their host halo, e.g. luminous red galaxies (LRGs) \\cite{Zheng2009}. Techniques to distinguish satellite galaxies from central galaxies have been developed and the satellite fraction can be used as an estimator of the host-halo mass \\cite{Zheng2005}. Mock LRG catalogs obtained from a halo occupation distribution (HOD) model suggest some reduction in stochasticity is possible even without explicit knowledge of halo masses~\\cite{Cai2011}. Therefore, an achievement of the presented gains seems feasible in light of upcoming spectroscopic galaxy surveys such as EUCLID~\\cite{EUCLID}, which will attain galaxy number densities, host-halo mass ranges and a survey volume comparable to the simulations used in this paper~\\cite{Laureijs2011,Majerotto2012}. Whether these gains translate into a useful constraint on the final cosmological parameters depends on our ability to model nonlinear RSD effects. We find nonlinear effects are important for $\\beta$ already at $k>0.03h\\mathrm{Mpc}^{-1}$, although they appear to be important for $f \\sigma_8$ only at $k>0.1h\\mathrm{Mpc}^{-1}$. In the most pessimistic case where the RSD model cannot be trusted for $k>0.03h\\mathrm{Mpc}^{-1}$, the multitracer method provides major gains relative to the single-tracer case, but neither method provides very strong constraints overall because of the limited number of available Fourier modes. In the case where we can use all the modes up to $k\\sim 0.1h\\mathrm{Mpc}^{-1}$, the overall errors are considerably smaller and the multitracer method provides less of an advantage. It is clear that a better modeling of the nonlinear effects in RSD is needed to understand the ultimate reach of RSD in both single-tracer and multitracer methods. In a more idealistic scenario, we also consider the joint analysis of halos and the dark matter density field, which in principle is achievable via a combination of spectroscopic redshift surveys and weak lensing tomography. Here, utilizing optimal weights can yield up to two orders of magnitude improvements in constraining $\\beta$ as compared to a single-tracer analysis, but the method is more prone to uncertainties in the halo mass estimates. It is unlikely that this gain can be achieved in practice, since it is very difficult to measure dark matter clustering in the radial direction directly. A further technique to construct differently biased tracers of the density field makes use of nonlinear transformations \\cite{Seljak2012}. Although it is difficult to describe the effects of a nonlinear transformation on both signal and noise in galaxy clustering data, combined with optimal weights this may provide another tool for the multitracer analysis. In this paper we have focused on the information that can be extracted from RSD, in particular $\\beta$ and $f\\sigma_8$, but our method is not limited to constraints on the growth rate, but may be applied to the analysis of primordial non-Gaussianity \\cite{Hamaus2011}, general relativistic corrections in large-scale clustering \\cite{Yoo2012}, the Alcock-Paczy\\'nski test \\cite{McDonald2009} or any other quantity that influences the effective bias of tracers of the density field. It is possible that a better model of nonlinear RSD may yield a more efficient multitracer method, where the gains relative to the single-tracer analysis described here on large scales can be extended to smaller scales. We leave these directions for the future." }, "1207/1207.3794_arXiv.txt": { "abstract": "We derive constraints on the matter density $\\om$ and the amplitude of matter clustering $\\s8$ from measurements of large scale weak lensing (projected separation $R=5-30\\hmpc$) by clusters in the Sloan Digital Sky Survey MaxBCG catalog. The weak lensing signal is proportional to the product of $\\om$ and the cluster--mass correlation function $\\xicm$. With the relation between optical richness and cluster mass constrained by the observed cluster number counts, the predicted lensing signal increases with increasing $\\om$ or $\\s8$, with mild additional dependence on the assumed scatter between richness and mass. The dependence of the signal on scale and richness partly breaks the degeneracies among these parameters. We incorporate external priors on the richness--mass scatter from comparisons to X-ray data and on the shape of the matter power spectrum from galaxy clustering, and we test our adopted model for $\\xicm$ against N-body simulations. Using a Bayesian approach with minimal restrictive priors, we find $\\s8(\\om/0.325)^{0.501}=0.828\\pm0.049$, with marginalized constraints of $\\om=0.325_{-0.067}^{+0.086}$ and $\\s8=0.828_{-0.097}^{+0.111}$, consistent with constraints from other MaxBCG studies that use weak lensing measurements on small scales~($R \\leq 2\\hmpc$). The $(\\om,\\s8)$ constraint is consistent with and orthogonal to the one inferred from WMAP CMB data, reflecting agreement with the structure growth predicted by General Relativity for a $\\Lambda$CDM cosmological model. A joint constraint assuming $\\Lambda$CDM yields $\\om=0.298_{-0.020}^{+0.019}$ and $\\s8=0.831_{-0.020}^{+0.020}$. For these parameters and our best-fit scatter we obtain a tightly constrained mean richness-mass relation of MaxBCG clusters, $N_{200}=25.4(M/3.61 \\times 10^{14} \\hmsun)^{0.74}$, with a normalization uncertainty of $1.5\\%$ Our cosmological parameter errors are dominated by the statistical uncertainties of the large scale weak lensing measurements, which should shrink sharply with current and future imaging surveys. ", "introduction": "\\label{sec:intro} The most fundamental question about the origin of cosmic acceleration is whether it arises from a new energy component or from a modification of General Relativity~(GR) on cosmological scales. A general strategy to address this question is to compare the growth of cosmic structure --- as measured, e.g., by cosmic shear, redshift--space distortions of galaxy clustering, or the abundance of galaxy clusters as a function of mass --- to the predictions of a GR$+$dark energy model constrained by geometrical probes such as Type Ia supernovae and baryon acoustic oscillations~(BAO). In particular, one can compare measurements of the matter density $\\om$ and the present--day amplitude of matter clustering, characterized by $\\s8$, the rms matter fluctuation in $8\\hmpc$ spheres, to the values expected from extrapolating cosmic microwave background~(CMB) anisotropies forward from recombination to $z=0$.\\footnote{We define $h\\equiv \\mrm{H}_0/(100\\,\\mrm{km}\\,\\mrm{s}^{-1}\\mrm{Mpc}^{-1})$ where $\\mrm{H}_0$ is the Hubble parameter at $z=0$.} Cosmological studies with clusters traditionally use mass proxies derived from X-ray or Sunyaev--Zel'dovich~\\citep[SZ;][]{sunyaev1972} measurements to constrain the cluster mass function $dn/d\\M$~(see \\citeauthor{allen2011} 2011 for a review). In an alternative approach, \\citeauthor{sheldon2009}~(2009; hereafter S09) used stacked weak lensing~(WL) to measure the average mass profiles around clusters in the MaxBCG catalog~\\citep{koester2007} derived from the Sloan Digital Sky Survey~(SDSS; \\citeauthor{york2000} 2000), detecting correlated mass from scales of $0.1\\hmpc$ to $30\\hmpc$. \\citeauthor{rozo2010}~(2010; hereafter R10) used the S09 measurements to constrain the mean relation between optical richness and virial mass for MaxBCG clusters, and they combined this relation with the abundance of clusters as a function of richness to constrain $\\om$ and $\\s8$. (For a general review of this approach in the context of cluster cosmology, see \\S6 of \\citeauthor{weinberg2012} 2012.) In this paper we again target $\\om$ and $\\s8$ with MaxBCG clusters, but we use the {\\it large scale} S09 measurements, from projected separations of 5--30$\\hmpc$. Roughly speaking, stacked weak lensing measures the product of the matter density $\\om$ and the cluster--mass cross--correction function $\\xicm(r)$. More precisely, given knowledge of the distances to lensing clusters and background sources, the mean tangential shear profile of clusters measures the excess surface density profile $\\ds(R)$, which is related to the 3--d $\\xicm(r)$ via \\begin{eqnarray} \\ds(R) & = & \\om\\rho_{\\mrm{c}}\\frac{2}{R^2}\\int_0^{R}\\int_{-\\infty}^{+\\infty}r_p\\xicm\\left(\\sqrt{r_p^2+r_z^2}\\right)dr_zdr_p \\nonumber \\\\ & - & \\om\\rho_{\\mrm{c}}\\int_{-\\infty}^{+\\infty}\\xicm\\left(\\sqrt{R^2+r_z^2}\\right)dr_z \\end{eqnarray} (see~\\S~\\ref{sec:data} for further details). We can understand how large scale $\\ds(R)$ measurements constrain $\\om$ and $\\s8$ by considering the simple case in which optical richness is perfectly correlated with cluster mass, so that a sample of clusters above a richness threshold corresponds to a sample above a mass threshold that has the same comoving space density $\\bar{n}$~(where, for simplicity, we consider a sample at fixed redshift). For a given cosmological model, one can predict the matter correlation function $\\ximm(r)$ and the bias factor $b_c(\\bar{n})$ of halos with space density $\\bar{n}$, and thus the cluster--mass correlation function, which is $\\xicm(r) = b_c(\\bar{n})\\ximm(r)$ on scales in the linear regime. Raising $\\om$ with all other quantities held fixed raises the predicted $\\ds(R)$ proportionally. Raising $\\s8$ increases $\\ximm\\propto\\s8^2$ and thus increases $\\xicm(r)$, but there is a partly compensating decline in $b_c(\\bar{n})$. In the limit of very rare, very highly biased peaks, $b_c(\\bar{n})\\propto\\s8^{-1}$, yielding $\\ds(R)\\propto\\om\\s8$, but for the space densities of typical cluster samples $b_c(\\bar{n})$ drops more slowly than $\\s8^{-1}$. Thus, the combination of cluster abundance measurements, which determine $\\bar{n}$, and large scale weak lensing measurements, which determine $\\ds(R)$, constrains a parameter combination $\\s8\\om^\\gamma$ with $\\gamma<1$. In practice, we will use bins of cluster richness instead of a single sample above a threshold, and the $\\s8$--dependence of $\\xicm(r)$ is different in the linear and mildly non--linear regimes, so there is some leverage to break degeneracy between $\\om$ and $\\s8$. The simplifications of this description point up several complications that must be addressed in our analysis. First, the measurements of $\\ds(R)$ have systematic uncertainties related to the photometric redshifts of the sources and shear calibration. Second, optical richness is a mass indicator with substantial scatter, which makes the bias of clusters in richness bins different from that of mass bins with the same space density. The two principal ``nuisance parameters'' in our statistical analysis are $\\be$, an overall scaling of the $\\ds(R)$ measurements to allow for systematic uncertainty, and $\\sc$, the logarithmic scatter in richness at fixed mass. We discuss these nuisance parameters and the priors we adopt on them in~\\S~\\ref{sec:ana}. We also adopt a prior on the shape of the matter power spectrum, so that a value of $\\s8$ specifies the full shape of $\\ximm(r)$. The inference of comoving space densities itself depends on $\\om$, which affects the volume element transformation between comoving distances and observable angles and redshifts. Incompleteness and contamination of the cluster sample can also affect the inferred space densities and/or bias the estimate of $\\ds(R)$, so they must also be accounted in the analysis. Despite these complications, we find that our constraints are limited by the statistical errors of the weak lensing measurements rather than systematic uncertainties. A complete cosmological analysis of cluster weak lensing would employ $\\ds(R)$ measurements over the full range of observed scales. Here we restrict our analysis to $R\\geqslant5\\hmpc$, in part to avoid the regime where theoretical predictions of $\\xicm(r)$ are uncertain, and in part to keep our results complementary to those of R10, who use the small scale ($R\\lesssim2\\hmpc$) S09 measurements to calibrate their determination of the cluster mass function. One important systematic for interpretation of the small scale measurements is the impact of cluster mis--centering, which must be estimated from simulations of the cluster population and cluster finding technique~\\citep[e.g.][]{johnston2007, george2012}. One advantage of the approach in this paper is that mis--centering has negligible impact at the large scales that we employ. In the following section we briefly review our input data, the MaxBCG cluster catalog and the S09 weak lensing measurements. Section~\\ref{sec:ana} presents our analysis method in detail, including the model parameters and priors and the procedure for computing the likelihood of the data given these parameters. Section~\\ref{sec:simumock} tests our analytic models for $\\ds(R)$~(a modified version of that proposed by~\\citeauthor{hayashi2008} 2008, hereafter HW08) against numerical simulations, and it uses simple mock data sets to test other aspects of our analysis procedures. Section~\\ref{sec:results} presents our cosmological constraints and compares them to those from other cluster analyses and from CMB data. We address systematic uncertainties in~\\S\\ref{sec:systematics}. We close, in~\\S\\ref{sec:dis}, with a summary of our findings and a discussion of future prospects. The reader in a hurry can get an overview of the paper from Fig.~\\ref{fig:model_demo}, which compares our best--fit model to our input data, Fig.~\\ref{fig:pedagogical}, which shows how the $\\ds(R)$ prediction depends on model parameters, and Fig.~\\ref{fig:wmap}, which presents our derived constraints on $\\om$ and $\\s8$. ", "conclusions": "\\label{sec:dis} We have derived cosmological constraints on $\\om$ and $\\s8$ using the combination of large scale cluster--galaxy weak lensing measurements~(S09) and the abundance of MaxBCG clusters as a function of richness. Within the analysis, we have statistically calibrated the cluster masses by requiring consistency between the cosmological model fit and the data, exploiting external priors on the scatter in the richness--mass relation from comparisons to X-ray data~\\citep{rozo2009}, and on the $\\pks$ from galaxy clustering~\\cite{reid2010}. The $68\\%$ confidence ellipse of our cosmological constraints on the $\\om$--$\\s8$ plane can be summarized as \\begin{equation} \\s8(\\om/0.325)^{0.501}=0.828\\pm0.049, \\end{equation} which is consistent with and orthogonal to the WMAP7 constraints on these parameters. This consistency of structure measured in the recombination era and the low redshift universe provides further evidence for the gravitational growth predicted by the $\\lcdm$ model combining GR, a cosmological constant, and cold dark matter. Assuming this model to be correct and combining our analysis with WMAP7, we obtain individual constraints as \\begin{equation} \\om = 0.298\\pm{0.020} \\quad\\mbox{and}\\quad \\s8 = 0.831\\pm{0.020} . \\end{equation} The overall results are consistent with and complementary to two other cosmological constraints from the same underlying clusters but with different input data and systematic uncertainties~(R10 and Tinker12). Collectively these three studies demonstrate consistency of the small scale weak lensing, large scale weak lensing, galaxy content, and abundance of the MaxBCG sample, together with galaxy clustering data. The primary systematic uncertainties in our analysis are the scatter in the cluster richness--mass relation and residual bias in the weak lensing measurements associated with photometric redshifts or shear calibration. However, with the external priors we have adopted, neither of these systematics is a limiting factor in our analysis; the uncertainties in our cosmological constraints are dominated by statistical uncertainties in the large scale $\\ds(R)$ measurements. These statistical errors can be sharply reduced in future surveys with deeper imaging and better seeing. Statistical improvements will require corresponding improvements in the control of systematics. While we have focused in this paper on large scales to complement the R10 analysis, the long term goal should be to derive constraints from the full range of $\\ds(R)$ simultaneously. Achieving this goal will require theoretical and numerical work to construct models that are accurate across the 1-halo and 2-halo transition and to assess uncertainties in the accuracy of the model predictions at all scales. There are opportunities for significant near--term improvements in our analysis using SDSS data. The MaxBCG catalog and weak lensing measurements used here are based on imaging data from DR4. With the recent release of DR8~\\citep{aihara2011}, almost every aspect of the catalog construction and the weak lensing measurements has evolved. The increase in the imaging area will enhance the raw statistical power~(for $7,398$ deg$^2$ VS. $14,555$ deg$^2$), reducing Poisson uncertainties and sample variance in the cluster counts and shape noise in the $\\ds(R)$ measurements. The optical cluster finding algorithm has been improved to produce catalogs with well-controlled selection function and, more importantly, a new richness estimator with reduced intrinsic scatter. \\cite{rykoff2012} considered various modifications of the original richness estimator in MaxBCG and found that the scatter in log--mass at fixed richness could be reduce to $0.2-0.3$ depending on richness, substantially smaller than MaxBCG scatter~\\citep[$0.45\\pm0.10$;][]{rozo2009} that we adopted as a prior in our analysis. When the scatter itself is smaller, then the systematic uncertainty tied to uncertainty in the scatter is also smaller. With improved uncertainty of the selection function, it will be feasible to use higher redshift clusters in the analysis, and while the $\\ds(R)$ measurements will degrade at higher $z$ because of reduced source surface density, the leverage of a wider redshift range may strengthen the cosmological constraints. On the weak lensing side, the main improvement is a better understanding of the photometric redshift distribution of source galaxies. With a much improved spectroscopic training set and better photometric calibration, \\cite{sheldon2011} reconstructed a redshift distribution for DR8 imaging data that is primarily limited by sample variance. Additional improvements come from updates in the photometric pipeline, including better sky subtraction, more refined stellar masks, and better PSF corrections in the shape measurements. Beyond SDSS, our approach can be applied to future, deeper, large--area imaging surveys. In the near term, the Pan--STARRS1~\\cite[PS1;][]{chambers2007} $3\\pi$ survey is expected to have larger area than SDSS, slightly greater depth, and higher image quality that yields a significant increase of the source density for weak lensing. The Dark Energy Survey~\\cite[DES;][]{the_dark_energy_survey_collaboration2005}, expected to start in late 2012, plans to survey $5000$ deg$^2$ to a depth two magnitudes beyond SDSS, with a weak lensing source density a factor of ten higher. It is designed with cluster cosmology and weak lensing as central goals, and our technique is naturally adapted to it. In the longer term, the imaging data sets from LSST, and the Euclid and WFIRST missions will allow radical improvements in the precision of cluster--galaxy lensing analysis, with effective source densities of $20-40$ arcmin$^{-2}$. These imaging surveys can provide their own cluster catalogs identified from the galaxy population, and they can provide stacked weak lensing measurements for clusters identified via X-ray emission or the SZ effect. Comparisons of results from different classes of cluster catalogs allow powerful cross--checks for systematics~\\citep[see, e.g.,][]{rozo2012, rozo2012-1, rozo2012-2} and valuable constraints on mass--observable relations~\\citep{cunha2010}. The SZ effect is a powerful technique for finding massive clusters at very high redshifts~\\citep[e.g.,][]{reichardt2012, williamson2011, marriage2011}, and DES will target the area covered by the SZ survey of the South Pole Telescope. X-ray observables may be more tightly correlated with halo mass than optical observables, and the upcoming eROSITA mission will carry out a sensitive all--sky X-ray survey that will revolutionize cosmological studies with X-ray selected clusters. \\cite{oguri2011} and \\cite{weinberg2012} argue that cluster abundances with masses calibrated by stacked weak lensing can provide constraints on structure growth that are highly competitive with those from cosmic shear analysis of the same WL survey. This conclusion assumes $\\ds(R)$ measurements out to $\\sim 1-2$ cluster virial radii, so the larger scale analysis illustrated here can only strengthen the power of this approach. Cluster-galaxy lensing is analogous to galaxy-galaxy lensing, but the relation between clusters and halos is simpler than the relation between galaxies and halos, and it is less subject to the complexities of baryonic physics. This greater simplicity reduces systematic uncertainty associated with theoretical modeling, which may ultimately compensate for the rarity of clusters relative to galaxies (and consequent lower statistical precision of the WL measurements). The Tinker12 study shows that the constraining power of small scale measurements can be enhanced by bringing in additional information from galaxy clustering and cluster mass-to-number ratios. In future work, we will investigate the generalization of this idea to large scales using the cluster-galaxy cross-correlation function, which can be measured in projection from the same survey used for weak lensing analysis. Nature has provided observable signposts that mark the locations of the most massive halos in the universe, and stacked weak lensing provides a tool to measure the average mass profiles of these halos at high precision over a wide range of scales. Exploiting this combination promises to yield stringent tests of gravity on cosmological scales and of theories for the origin of cosmic acceleration." }, "1207/1207.6980_arXiv.txt": { "abstract": "{ {This paper investigates the hydrodynamic performances of an SPH code incorporating an artificial heat conductivity term in which the adopted signal velocity is applicable when gravity is present. To this end, we analyze results from simulations produced using a suite of standard hydrodynamical test problems.} {In accordance with previous findings it is shown that the performances of SPH to describe the development of Kelvin-Helmholtz instabilities depend strongly on the consistency of the initial condition set-up and on the leading error in the momentum equation due to incomplete kernel sampling. On the contrary, the presence of artificial conductivity does not significantly affect the results.} {An error and stability analysis shows that the quartic $B-$spline kernel ($M_5$) possesses very good stability properties and we propose its use with a large neighbor number, between $\\sim50$ (2D) to $\\sim 100$ (3D), to improve convergence in simulation results without being affected by the so-called clumping instability.} {Moreover, the results of the Sod shock tube test demonstrate that in order to obtain simulation profiles in accord with the analytic solution, for simulations employing kernels with a non-zero first derivative at the origin it is necessary the use of a much larger number of neighbors than in the case of the $M_5$ runs.} {SPH simulations of the blob test show that in order to achieve blob disruption it is necessary the presence of an artificial conductivity term. However, it is found that in the regime of strong supersonic flows an appropriate limiting condition, which depends on the Prandtl number, must be imposed on the artificial conductivity SPH coefficients in order to avoid an unphysical amount of heat diffusion.} {Results from hydrodynamic simulations that include self-gravity show profiles of hydrodynamic variables that are in much better agreement with those produced using mesh-based codes. In particular, the final levels of core entropies in cosmological simulations of galaxy clusters are consistent with those found using AMR codes. This demonstrate that the proposed diffusion scheme is capable of mimicking the process of entropy mixing which is produced during structure formation because of diffusion due to turbulence.} {Finally, results of the Rayleigh-Taylor instability test demonstrate that in the regime of very subsonic flows the code has still several difficulties in the treatment of hydrodynamic instabilities. These problems being intrinsically due to the way in which in standard SPH gradients are calculated and not to the implementation of the artificial conductivity term. To overcome these difficulties several numerical schemes have been proposed which, if coupled with the SPH implementation presented in this paper, could solve the issues which recently have been addressed in investigating SPH performances to model subsonic turbulence. } } ", "introduction": "\\label{intro.sec} Smoothed particle hydrodynamics (SPH) is a Lagrangian, mesh-free, particle method which is used to model fluid hydrodynamics in numerical simulations. The technique was originally developed in an astrophysical context \\citep{lu77,gm77}, but since then it has been widely applied in many other areas \\citep{mo05} of computational fluid dynamics. The method has several properties which make its use particularly advantageous in astrophysical problems \\citep{hk89,ro09,sp10}. Because of its Lagrangian nature the development of large matter concentrations in collapse problems is followed naturally, moreover the method is free of advection errors and is Galilean invariant, Finally, the method naturally incorporates self-gravity and possess very good conservation properties, its fluid equations being derived from variational principles. The other computational fluid dynamical method commonly employed in numerical astrophysics is the standard Eulerian grid based approach in which the fluid is evolved on a discretized mesh \\citep{st92,ry93,no99b,fr00,te02}. The spatial resolution of an Eulerian scheme based on a fixed Cartesian grid is often insufficient however to adequately resolve the large dynamic range frequently encountered in many astrophysical problems, such as galaxy formation. This has motivated the development of adaptative mesh refinement (AMR) methods, in which the spatial resolution of the grid is locally refined according to some selection criterion \\citep{be89,kr97,no05}. Additionally, the order of the numerical scheme has been improved by adopting the parabolic piecewise method (PPM) of \\cite{co84}, such as in the AMR Eulerian codes ENZO \\citep{no99b} and FLASH \\citep{fr00}. Application of these different types of hydrodynamical codes to the same test problem with identical initial conditions should in principle lead to similar results. However, there has been growing evidence over recent years that in a variety of hydrodynamical test cases there are significant differences between the results produced the two types of methods \\citep{os05,ag07,wa08,ta08,mi09,read10,vrd10,ju10}. For instance, \\cite{ag07} showed that in the standard SPH formulations the growth of Kelvin-Helmholtz (KH) instabilities in shear flows is artificially suppressed when steep density gradients are present at the contact discontinuities. Moreover, the level of central entropy produced in binary merger non-radiative simulations of galaxy clusters is significantly higher by a factor $\\sim2$ in Eulerian simulations than in those made using SPH codes \\citep{mi09}. The origin of these discrepancies has been recognized \\citep{ag07,read10} as partly due to the intrinsic difficulty for SPH to properly model density gradients across fluid interfaces, which in turn implies the presence of a surface tension effect which inhibits the growth of KH instabilities. A second problem is due to the Lagrangian nature of SPH codes, which prevents mixing of fluid elements at the particle level and leads to entropy generation \\citep{wa08,mi09}. In particular, \\cite{mi09} argue that the main explanation for the different levels of central entropy found in cluster simulations is due to the different degree of entropy mixing present in the two codes. Unlike SPH, in Eulerian codes fluid mixing occurs by definition at the cell level and a certain degree of overmixing is certainly present in those simulations made using mesh-based codes \\citep{sp10b}. Given the advantages of SPH codes highlighted previously, it appears worth pursuing any improvement in the SPH method capable of overcoming its present limitations. Along this line of investigation many efforts have been made by a number of authors \\citep{ab11,pr08,wa08,vrd10,read10,hs10,ch10,mu11}. \\cite{pr08} introduces an artificial heat conduction term in the SPH equations with the purpose of smoothing thermal energy at fluid interfaces. This artificial conductivity (AC) term in turn gives a smooth entropy transition at contact discontinuities with the effect of enforcing pressure continuity and removing the artificial surface tension effect which inhibits the growth of KH instabilities at fluid interfaces. Similarly, \\cite{wa08} suggest that in SPH the lack of mixing at the particle level can be alleviated by adding a heat diffusion term to the equations so as to mimic the effects of subgrid turbulence, thereby improving the amount of mixing. \\cite{read10} present an SPH implementation in which a modified density estimate is adopted \\citep{rt91}, together with the use of a peaked kernel and a much larger number of neighbors. The authors showed that the new scheme is capable of following the development of fluid instabilities in a better way than in standard SPH. \\cite{ab11} presents an alternative derivation of the SPH force equation which avoids the problem encountered by standard SPH in handling fluid instabilities, although the approach is not inherently energy or momentum conserving and is prone to large integration errors when the simulation resolution is low. The method proposed by \\cite{I02} reformulates the SPH equations by introducing a kernel convolution so as to consistently calculate density and hydrodynamic forces. The latter are determined using a Riemann solver (Godunov SPH). The method has recently been revisited by \\cite{ch10} and \\cite{mu11}, who showed that the code correctly follows the development of fluid instabilities in a variety of hydrodynamic tests. A deeper modification than those presented here so far has been introduced by \\cite{hs10}, who replaced the traditional SPH kernel approach with an new density estimate based on Voronoi tesselation. The authors showed that the method is free of surface tension effects and therefore the growth rate of fluid instabilities is not adversely affected as in standard SPH. Finally, a radically new numerical scheme has been introduced by \\cite{sp10b}, with the purpose of retaining the advantages of both SPH and mesh-based codes. In the new code the hydrodynamic equations are solved on a moving unstructured mesh using a Godunov method with an exact Riemann solver. The mesh is defined by the Voronoi tesselation of a set of discrete points and is allowed to move freely with the fluid. The method is therefore adaptative in nature and thus Galilean invariant but, at the same time, the accuracy with which shocks and contact discontinuities are described is that of an Eulerian code. In a recent paper \\cite{ba12} argue that the standard formulation of SPH fails to accurately resolve the development of turbulence in the subsonic regime (but see \\cite{pr12b} for a different viewpoint). The authors draw their conclusions by analyzing results from simulations of driven subsonic turbulence made using the new moving-mesh code, named AREPO, and a standard SPH code. Similar conclusions were reached in a set of companion papers \\citep{sj11,vog11}, in which the new code was used in galaxy formation studies to demonstrate its superiority over standard SPH. However, the code is characterized by considerable complexity which makes the use the SPH scheme still appealing and, more generally, it is desirable that simulation results produced with a specific code should be reproduced with a completely independent numerical scheme when complex non-linear phenomena are involved. It appears worthwhile, therefore, to investigate, along the line of previous authors, the possibility of constructing a numerical scheme based on the traditional SPH formulation which is capable of correctly describing the development of fluid instabilities and at the same time incorporates the effects of self-gravity when present. This is the aim of the present study, in which the SPH scheme is modified by incorporating into the equations an AC diffusion term as described by \\cite{pr08}. However, in \\cite{pr08} the strength of the AC is governed by a signal velocity which is based on pressure discontinuities. For simulations where gravity is present this approach is not applicable because hydrostatic equilibrium requires pressure gradients. An appropriate signal velocity for conductivity when gravity is present is then used to construct an AC-SPH code with the purpose of treating the growth of fluid instabilities self-consistently. The viability of the approach is tested using a suite of test problems in which results obtained using the new code are contrasted with the corresponding ones produced in previous work using different schemes. The code is very similar in form to that presented by \\cite{wa08}, but here the energy diffusion equation is implemented in a different manner. The paper is organized as follows. Sect. \\ref{hymeth.sec} presents the hydrodynamical method and introduces the AC approach. In Sect. \\ref{hydro.sec} we investigate the effectiveness of the method by presenting results from a suite of purely hydrodynamical test problems. These are: the two-dimensional Kelvin-Helmholtz instability, the Sod shock tube, the point explosion or Sedov blast wave test and the blob test. In Sect. \\ref{hygrav.sec} we then discuss results from hydrodynamic tests that include self-gravity. Specifically, we consider the cold gas sphere or Evrard collapse test, the Rayleigh-Taylor instability and the cosmological integration of galaxy clusters. Finally, the main results are summarized in Sect. \\ref{conc.sec} ", "conclusions": "\\label{conc.sec} In this paper, we have presented an SPH numerical scheme which incorporates an artificial conductivity term and uses an appropriate signal velocity for simulations including gravity. The AC formulation has been introduced by \\cite{pr08} as a solution to the problems encountered by standard SPH to correctly follow the development of KH instabilities, due to the inconsistencies of the standard formulation in the description of density at contact discontinuity. Here, a suite of hydrodynamic test problems is investigated with the purpose of validating the new AC-SPH code and to assess its performances when using the specifically adopted signal velocity. The results of the KH instability test are presented in Sect. \\ref{kh2d.sec}, in which the code capabilities have been tested by considering SPH simulations of the KH test performed using a large variety of initial condition set-up and SPH kernels. The set of initial conditions has been chosen as in \\cite{vrd10}, so as to consistently compare the results. These are in accordance with the corresponding ones of \\cite{vrd10} and indicate that, for the version of the KH test analyzed here, the AC implementation is important for the long-term behavior of the simulation. Moreover, the results of the KH test indicate \\citep{vrd10,read10,na11} that the poor performances of standard SPH to properly treat KH instabilities can be explained in terms of two distinct effects. The first is a general problem of consistency, which for the problem under consideration requires that smooth interfaces should be present at contact discontinuities, in order to obtain numerical convergence. The second effect which in standard SPH suppresses the growth of KH instabilities is the leading error in the momentum equation, due to incomplete kernel sampling \\citep{read10} and quantified by the norm $\\vec {E}^{0}$ defined by Eq. (\\ref{en0.eq}). To circumvent this problem several proposals have been made \\citep{vrd10,read10} in which the standard SPH cubic spline $M_4$ kernel is replaced by the new LIQ or CRT kernels with steeper central profiles. These kernels present the advantage of being stable against particle clumping, so that the number of neighbors can be safely increased in order to reduce the error $\\vec {E}^{0}$. In this paper, we consider the possibility of reducing sampling errors by considering higher-order $B-$spline kernels, specifically the $M_5$ or quartic spline kernel. A striking result of the KH runs presented in sect. \\ref{kh2d.sec} is that, for a given value of the ratio $\\eta$ between the smoothing length and the mean interparticle spacing, simulations of the KH test performed using the $M_5$ kernel have amplitudes of the $\\vec {E}^{0}$ error substantially smaller and in line with the behavior of the same quantity for simulations employing the LIQ or CRT kernels. A linear stability analysis reveals that this result follows owing to the very good stability properties of the $M_5$ kernel, in fact the analysis suggests that the clumping instability is absent for values of $\\eta$ up to $\\eta\\simlt2$, which in 3D corresponds to $N_{sph}\\sim520$ neighbors. These findings are consistent with the recent results of \\cite{de12} who proposed to adopt, with the specific purpose of avoiding the clumping instability, the \\cite{we95} functions as a new class of kernels. The authors showed that in terms of stability and accuracy properties, the quartic spline performs extremely well when compared with the proposed Wedland functions. These results are strictly connected to the properties of the kernel Fourier transform, which according to the authors must be non-negative to avoid particle clumping, and are consistent with the findings of sect. \\ref{khstab.sec} since increasing the kernel order both the $B-$spline and the Wedland functions approach the Gaussian. We therefore propose, as a compromise between the need of reducing sampling errors while keeping the computational cost at a minimum, the use of the $M_5$ kernel with neighbor number in the range $N_{sph}\\sim60-120$ as the standard combination which guarantees sufficient accuracy in many SPH simulations of astrophysical problems. Moreover, the results of the Sod shock tube test of sect. \\ref{sod.sec} demonstrate that in order to obtain simulation profiles in accordance with the analytic solution, for simulations employing kernels with a modified shape the use of a much larger number of neighbors than in the case of the $M_5$ runs is necessary. The results of the gravity tests show that the adoption of the AC-SPH scheme significantly reduces, at the level tested in this paper, the differences seen in the hydrodynamics between standard SPH and grid-based simulations of self-gravitating structures. For the cold gas sphere, the entropy profile is in better agreement with the PPM reference solution and for the cosmological cluster simulations, a key result is the final level of core entropies, which are consistent with those of the central entropies produced using AMR codes. Thus, it appears that in hydrodynamic simulations where self-gravity is important the AC term, accompanied by the proposed signal velocity, plays a key role as a mechanism of redistributing thermal energy and hence as a source of entropy mixing. To summarize, results extracted from simulations of hydrodynamic tests where self-gravity dominates are in much better agreement with the corresponding ones obtained using mesh-based codes. The results then demonstrate the capability of the implemented AC-SPH scheme to properly follow the formation of cosmic structures. It is worth noting that the artificial heat conduction term was originally proposed by \\cite{pr08} with the purpose of avoiding the inconsistencies encountered by standard SPH in presence of density steps at contact discontinuities. A complementary view has been proposed by \\cite{wa08}, who introduced the same term, albeit in a different numerical formulation, with the aim of modeling the level of diffusion due to turbulence. The two interpretations are not mutually inconsistent, however the results of the self-gravity tests presented in this paper support the view of a heat diffusion term which in SPH is capable of mimicking the diffusion due to turbulence. In a similar fashion, \\cite{vi07} presented an SPH scheme to model a free-surface incompressible flow which , in analogy with 3D Large Eddy Simulations, assumes a \\cite{sm63} model for the filtered Navier-Stokes equations. The results of the blob test simulations demonstrate that the instabilities leading to the expected cloud disruption can develop only when the SPH energy equation incorporates the AC term. A particularly interesting result is that an appropriate limiting condition must be implemented on the AC coefficients $\\alpha^C$ in order to avoid an unphysical amount of heat diffusion, which in turn leads to a cloud disruption which occurs too early. This limiter has been identified as given by the Prandtl number and, for the AC signal velocity adopted here, it severely limits the amplitude of the AC coefficients in the regime of strong supersonic flows. AC-SPH simulations of the blob test incorporating now the new constraint support this view, since Fig. \\ref{fig:mloss} shows for the new runs a cloud mass-loss rate which is in better agreement with the rates obtained from simulations realized using a completely independent numerical scheme \\citep{sp11}. However, it must be stressed that the condition (\\ref{prac.eq}) has been calculated for a perfect monoatomic gas with $\\gamma=5/3$, so that a physically motivated constraint on the artificial heat diffusion of the simulated medium should be considered problem dependent. However, the code has still several problems which render its use problematic if the development of hydrodynamic instabilities need to be followed in the regime of subsonic flows. The results of the simulations indicate that these shortcomings are not due to the AC implementation, but rather are intrinsic in the standard formulation with which gradients are calculated in SPH and the related errors are subsequently introduced in the momentum equation. In particular, simulations of the RT test show that increasing the kernel order alleviates the problems but does not solve them. These results are in accordance with recent findings \\citep{na11,ga12} and clearly demonstrate that for very subsonic flows, the poor performances of SPH to model hydrodynamic instabilities are strictly connected to the code accuracy in gradient estimates. The formulation of \\cite{ga12}, which has been proposed with the aim of calculating SPH gradients with an as high as possible accuracy while keeping the benefits of a Lagrangian formulation, looks in this aspect very promising and suggests further investigations along the numerical approach proposed in this paper. These would be particularly relevant in the light of the recent results of \\cite{ba12}, who claim that standard SPH fails to properly model the regime of subsonic turbulence. They reach their conclusions by comparing results extracted from simulations of driven subsonic turbulence realized using the moving-mesh code AREPO and GADGET SPH. Their results have been criticized by \\cite{pr12}, for whom the use of a time-dependent AV is critical in SPH simulations of subsonic turbulence. A re-analysis of the \\cite{ba12} simulations using the AC-SPH code presented here, augmented with improved gradient operators, would then be of fundamental importance to achieve a deeper understanding of the capability of different numerical methods to model subsonic turbulence. The latter being expected to have a significant impact in shaping the thermodynamic properties of baryons in cosmological haloes and, subsequently, the process of galaxy formation." }, "1207/1207.2498_arXiv.txt": { "abstract": "We compare the actual WMAP maps with artificial, purely statistical maps of the same harmonic content to argue that there are, with confidence level 99.7\\%, ring-type structures in the observed cosmic microwave background. ", "introduction": " ", "conclusions": "" }, "1207/1207.2451_arXiv.txt": { "abstract": "{ We investigate statistical equilibrium of neutral and singly-ionized strontium in late-type stellar atmospheres. Particular attention is given to the completeness of the model atom, which includes new energy levels, transition probabilities, photoionization and electron-impact excitation cross-sections computed with the R-matrix method. The NLTE model is applied to the analysis of Sr I and Sr II lines in the spectra of the Sun, Procyon, Arcturus, and HD 122563, showing a significant improvement in the ionization balance compared to LTE line formation calculations, which predict abundance discrepancies of up to $0.5$ dex. The solar Sr abundance is $\\log \\epsilon = 2.93 \\pm 0.04$ dex, in agreement with the meteorites. A grid of NLTE abundance corrections for Sr I and Sr II lines covering a large range of stellar parameters is presented.} ", "introduction": "Spectroscopic observations of low-mass stars have shaped our understanding of Galactic evolution and stellar nucleosynthesis. Strontium, as one of the abundant r-process elements, has been extensively investigated in the past few decades. However, its main production site has not yet been identified: the observed abundances of Sr in metal-poor stars are far too large to be explained by conventional rapid neutron-capture nucleosynthesis in SNe II, suggesting some alternative exotic scenarios, such as the light element primary process \\citep{2004A&A...425.1029T}, rp-process in accretion disks around low-mass black holes \\citep{1998PhR...294..167S}, black hole - neutron star mergers \\citep{2008ApJ...679L.117S}, high-entropy winds in SN II \\citep{2010ApJ...712.1359F}, and low-mass electron-capture supernovae \\citep{2011ApJ...726L..15W} to name just a few \\citep[see][and references therein]{2011RPPh...74i6901J}. Until recently, determinations of Sr abundances in metal-poor stars relied almost exclusively on the two near-UV lines of \\ion{Sr}{ii}, which are sufficiently strong to be detected also in the spectra of moderate-to-low resolution and signal-to-noise. The drawback is that in spectra of stars typically used for studies of Galactic chemical evolution, [Fe/H] $ > -2$, these lines saturate and develop pronounced damping wings, overlapping with various atomic and molecular blends. Thus, at higher metallicities a preference is sometimes given to the weak \\ion{Sr}{i} line at $4607.34$ \\AA\\ and/or the \\ion{Sr}{ii} line at $4161$ \\AA. However, there is evidence that the \\ion{Sr}{i} lines may be subject to non-local thermodynamic equilibrium (hereafter, NLTE) effects \\citep{2000MNRAS.311..535B}, which has been supported by \\textit{ab initio} calculations solving for radiative transfer in NLTE for a small grid or red giant model atmospheres \\citep{2006ApJ...641..494S}. In a few studies utilizing near-IR spectra, also the \\ion{Sr}{ii} triplet ($10\\,037$, $10\\,327$, and $10\\,915$ $\\AA$ have been used \\citep{2011A&A...530A.105A}. In the majority of published studies, the preference is given to one ionization stage only \\citep[e.g.,][]{1999A&A...341..241J}, and only a few studies investigated both \\ion{Sr}{i} and \\ion{Sr}{ii} lines \\citep{1994A&A...287..927G, 2002ApJ...572..861C}, finding discrepant results. In this study, we perform for the first time a NLTE analysis of the \\ion{Sr}{i} and \\ion{Sr}{ii} lines in spectra of late-type stars. The new atomic model was constructed from the state-of-art atomic data, computed specifically for this work. The NLTE model atom is tested on a number of reference stars with parameters determined by other independent methods. Furthermore, we present a large grid of NLTE abundance corrections for \\ion{Sr}{i} and \\ion{Sr}{ii} lines. The results presented in this work will be applied to the analysis of a representative sample of metal-poor stars observed at a very-high resolution and signal-to-noise in Hansen et al. (in prep.) Furthermore, we plan to undertake the NLTE Sr abundance analysis of the thick-disk and halo stars with spectroscopic parameters from \\citet{2011ApJ...737....9R}. ", "conclusions": "" }, "1207/1207.5792_arXiv.txt": { "abstract": "{ Recently, high-dispersion spectroscopy has demonstrated conclusively that four of the five globular clusters (GCs) in the Fornax dwarf spheroidal galaxy are very metal-poor with ${\\rm [Fe/H]}<-2$. The remaining cluster, Fornax 4, has ${\\rm [Fe/H]}=-1.4$. This is in stark contrast to the field star metallicity distribution which shows a broad peak around ${\\rm [Fe/H]}\\approx-1$ with only a few percent of the stars having ${\\rm [Fe/H]}<-2$. If we only consider stars and clusters with ${\\rm [Fe/H]}<-2$ we thus find an extremely high GC specific frequency, $S_{\\! N}\\approx400$, implying by far the highest ratio of GCs to field stars known anywhere. We estimate that about 1/5--1/4 of all stars in the Fornax dSph with ${\\rm [Fe/H]}<-2$ belong to the four most metal-poor GCs. These GCs could, therefore, at most have been a factor of 4--5 more massive initially. Yet, the Fornax GCs appear to share the same anomalous chemical abundance patterns known from Milky Way GCs, commonly attributed to the presence of multiple stellar generations within the clusters. The extreme ratio of metal-poor GC- versus field stars in the Fornax dSph is difficult to reconcile with scenarios for self-enrichment and early evolution of GCs in which a large fraction (90\\%--95\\%) of the first-generation stars have been lost. It also suggests that the GCs may not have formed as part of a larger population of now disrupted clusters with an initial power-law mass distribution. The Fornax dSph may be a rosetta stone for constraining theories of the formation, self-enrichment and early dynamical evolution of star clusters. } ", "introduction": "The Fornax dwarf spheroidal galaxy (dSph) is well-known for its high globular cluster (GC) specific frequency, i.e., the number of GCs normalized to a host galaxy $M_V=-15$ \\citep{Harris1981}. Assuming $M_V = -13.2$ \\citep{Mateo1998}, the five GCs \\citep{Hodge1961} correspond to a specific frequency of $S_{\\! N}=26$. Specific frequencies rivalling that of Fornax have only been found in other similarly faint dwarf galaxies \\citep{Peng2008,Georgiev2010}. Spiral galaxies typically have $S_{\\! N}\\approx 1$, while normal elliptical galaxies have $S_{\\! N}\\approx3-5$, slightly higher in clusters than in the field \\citep{Harris1991}. Even the most cluster-rich cD galaxies generally have $S_{\\! N}\\la15$ \\citep{Harris1991,Brodie2006}. These differences constitute the classical ``specific frequency problem''. A second specific frequency problem becomes apparent when considering the numbers of GCs associated with different stellar populations \\emph{within} galaxies. Early-type galaxies are generally redder (more metal-rich) than the average of their GC systems \\citep{Forte1981,Larsen2001,Forbes2001}, and direct comparisons of metallicity distributions for stars and globular clusters have confirmed that the specific frequency tends to increase with decreasing metallicity within galaxies \\citep{Harris2002,Harris2007}. This is also evident in our own Galaxy: of the roughly 150 GCs known in the Milky Way, about 2/3 are associated with the halo \\citep{Zinn1985}, even though the halo only contains 1\\%--2\\% of the stellar mass \\citep{Suntzeff1991,Dehnen1998}. It has, in fact, been suggested that a significant fraction of the Milky Way halo stars originate from disrupted GCs. If the Galactic GC population formed with a power-law mass function similar to that observed in young cluster systems at the present epoch \\citep{Larsen2009}, then the stars lost from dissolving clusters over a Hubble time might be sufficient to account for essentially the entire stellar halo \\citep[e.g.][]{Kruijssen2009}. The above estimate only involves standard dynamical evolution of the clusters. However, a similar conclusion can be reached via a different line of reasoning, starting from the increasing body of evidence that GCs host dual or multiple stellar generations \\citep{Gratton2012}. In the following we will refer to the ``first'' and ``second'' generation of stars, keeping in mind that the actual star formation histories may be more complex. The second generation is characterized, among other things, by anomalous abundances of several light elements that imply $p$-capture nucleosynthesis at high temperatures. The currently favoured sites are massive AGB stars \\citep{DErcole2008}, fast-rotating massive main sequence stars \\citep{Decressin2007}, or massive interacting binaries \\citep{deMink2009}. A fundamental challenge faced by most scenarios for the origin of multiple stellar generations in GCs is the ``mass budget'' problem: at the present epoch, the second generation of stars has a mass similar to, or even greater than the first generation. However, for a standard IMF the ejecta produced by the first-generation stars fall short by large factors compared to the mass required to form the observed second generation, even if a 100\\% star formation efficiency is assumed. This has led to the suggestion that a large fraction of the first generation, up to 90--95\\%, was lost soon after the formation of the second generation \\citep{DErcole2008,Schaerer2011,Bekki2011}. Accounting for subsequent dynamical evolution, one again finds that a significant fraction of the Milky Way stellar halo might come from GCs, even without considering that a significantly larger population of lower-mass clusters may have been present initially \\citep{Gratton2012}. In this letter we discuss the globular cluster system of the Fornax dSph in the context of the second specific frequency problem and self-enrichment scenarios. While these two issues may seem unrelated, the Fornax dSph turns out to be such an extreme case of the second $S_{\\! N}$ problem that it puts tight constraints on the amount of mass that could have been lost from its GCs. ", "conclusions": "We have used the most recent measurements of metallicities for globular clusters and field stars in the Fornax dwarf spheroidal galaxy to estimate the specific frequency of metal-poor GCs. Four of the GCs have ${\\rm [Fe/H]} < -2.0$; associating these GCs with field stars in the same metallicity range we estimate $S_N \\approx 400$, by a wide margin the highest known anywhere. We find that about one quarter of all stars in the Fornax dSph with ${\\rm [Fe/H]}<-2$ belong to these four most metal-poor clusters. Accounting for standard dynamical evolution, an even larger fraction of the metal-poor stars must initially have been born in the clusters. This puts tight constraints on any further amount of stars that could have been lost from metal-poor star clusters due to early ``infant weight loss'' or ``infant mortality''. The extant clusters could at most have been a factor of $\\sim4-5$ more massive initially then they are now, and this requires a rather extreme scenario in which no field stars or other clusters of similar metallicity were formed initially. It would be of significant interest to carry out detailed chemical tagging of the most metal-poor stars in Fornax to determine how many of them can be traced back to the GCs. With sufficient accuracy, it should be possible to detect peaks in the field star metallicity distribution corresponding to stars lost from the GCs. One might also expect to see a significant fraction of metal-poor field stars with peculiar, GC-like abundance patterns. Fornax and other dSph galaxies may be unique laboraties in which the presence or absence of anomalous abundances of light elements in the field stars might shed light on the origin of multiple stellar populations in globular clusters." }, "1207/1207.0809_arXiv.txt": { "abstract": "It is a firm prediction of the concordance Cold Dark Matter (CDM) cosmological model that galaxy clusters live at the intersection of large-scale structure filaments\\cite{1996Nature..380..603B}. The thread-like structure of this ``cosmic web'' has been traced by galaxy redshift surveys for decades\\cite{1978MNRAS.185..357J,1989Sci...246..897G}. More recently the Warm-Hot Intergalactic Medium (WHIM) residing in low redshift filaments has been observed in emission\\cite{2008A&A...482L..29W} and absorption\\cite{2009ApJ...695.1351B,2010ApJ...714.1715F}. However, a reliable direct detection of the underlying Dark Matter skeleton, which should contain more than half of all matter\\cite{2010MNRAS.408.2163A}, remained elusive, as earlier candidates for such detections\\cite{1998astro-ph/9809268K,2002ApJ...568..141G, 2005A&A...440..453D} were either falsified\\cite{2004A&A...422..407G,2008MNRAS.385.1431H} or suffered from low signal-to-noise ratios\\cite{1998astro-ph/9809268K,2005A&A...440..453D} and unphysical misalignements of dark and luminous matter\\cite{2002ApJ...568..141G,2005A&A...440..453D}. Here we report the detection of a dark matter filament connecting the two main components of the Abell 222/223 supercluster system from its weak gravitational lensing signal, both in a non-parametric mass reconstruction and in parametric model fits. This filament is coincident with an overdensity of galaxies\\cite{2002A&A...394..395D,2005A&A...440..453D} and diffuse, soft X-ray emission\\cite{2008A&A...482L..29W} and contributes mass comparable to that of an additional galaxy cluster to the total mass of the supercluster. Combined with X-ray observations\\cite{2008A&A...482L..29W}, we place an upper limit of $\\mathbf{0.09}$ on the hot gas fraction, the mass of X-ray emitting gas divided by the total mass, in the filament. ", "introduction": " ", "conclusions": "" }, "1207/1207.7104_arXiv.txt": { "abstract": "Planets migrate due to the recoil they experience from scattering solid (planetesimal) bodies. To first order, the torques exerted by the interior and exterior disks cancel, analogous to the cancellation of the torques from the gravitational interaction with the gas (type I migration). Assuming the dispersion-dominated regime and power-laws characterized by indices $\\alpha$ and $\\beta$ for the surface density and eccentricity profiles, we calculate the net torque on the planet. We consider both distant encounters and close (orbit-crossing) encounters. We find that the close and distant encounter torques have opposite signs with respect to their $\\alpha$ and $\\beta$ dependences; and that the torque is especially sensitive to the eccentricity gradient ($\\beta$). Compared to type-I migration due to excitation of density waves, the planetesimal-driven migration rate is generally lower due to the lower surface density of solids in gas-rich disk, although this may be partially or fully offset when their eccentricity and inclination are small. Allowing for the feedback of the planet on the planetesimal disk through viscous stirring, we find that under certain conditions a \\textit{self-regulated} migration scenario emerges, in which the planet migrates at a steady pace that approaches the rate corresponding to the one-sided torque. If the local planetesimal disk mass to planet mass ratio is low, however, migration stalls. We quantify the boundaries separating the three migration regimes. ", "introduction": "Bodies immersed in gaseous or particle disks migrate radially. Very small particles, strongly coupled to the gas, are carried by the gas. Thus, they follow the accretion flow or are dispersed by turbulent motions \\citep{Ciesla2009,Ciesla2010}. Larger particles tend to move on Keplerian orbits. However, in protoplanetary disks the gas is partially pressure-supported, which causes solids to drift inwards due to the headwind they experiences \\citep{AdachiEtal1976,Weidenschilling1977}. This effect peaks for $\\sim$m-size bodies (or their aerodynamic equivalents) at which they spiral in in as little as $\\sim$100 orbital periods. Larger, km-size bodies (planetesimals) are more resistant against drag-induced orbital decay due to their large inertia. The motions of these bodies will be predominantly determined by gravitational encounters, rather than gas drag. The gravitational interaction with the gas also causes a drag force on the planet. The picture here is that of a massive body gravitationally perturbing the disks, which causes an excess density structure that backreacts on the planet. One can regard the force that the planets experiences a manifestation of \\textit{dynamical friction} -- a concept that is perhaps more familiar with collisionless systems, but which can also be applied to gaseous disks \\citep{Ostriker1999,KimKim2007,KimKim2009,MutoEtal2011,LeeStahler2011}. When the planet is small, the resulting migration from the gravitational interaction with the gas is known as type I \\citep{GoldreichTremaine1980,Ward1986}. Like dynamical friction, the type-I migration rate increases linearly with mass. Although rather insignificant for planetesimals, it becomes very efficient for Earth-mass planets resulting in migration timescales as short as $10^5$ yr at 1 AU \\citep{TanakaEtal2002}. For these reasons (gas-driven) migration is often invoked to explain the existence of close-in, Neptune- and Jupiter-mass planets (`hot Jupiters'), since conditions very close to the star are thought to be ill-suited to form giant planets \\textit{in situ} \\citep{IkomaHori2012}. However, a clear understanding of type-I migration is somewhat complicated by the fact that it is a higher order effect; that is, the net torque on the planet results from a near cancellation of two large but opposite torques, corresponding to the respective contributions from the inner and outer disks. In addition, the net co-orbital and Lindblad torque may have different signs. As a result the sign of type-I migration is very sensitive to the local distribution of matter, which in turn is determined by the thermodynamic properties of the disk \\citep[\\eg][]{PaardekooperMellema2006}. Similar to `gas-driven' migration, scattering of solid bodies also causes a planet to migrate. This effect of planetesimal-driven migration (PDM) has been mostly explored through $N$-body studies \\citep{HahnMalhotra1999,KirshEtal2009,BromleyKenyon2011,CapobiancoEtal2011}. In some cases, these authors found an migration instability, at which the planet migrates at a rate determined by the one-sided torque \\citep{IdaEtal2000}. Under these conditions, PDM is fast. Other studies have investigated the embryo-planetesimal interaction analytically \\citep[\\eg][]{Ida1990,IdaMakino1993,TanakaIda1996,Rafikov2003iii}. Mostly, these studies consider the effect of the embryo on the planetesimal disk, \\eg\\ the rate at which the protoplanet excites the planetesimal's eccentricity or how it opens a gap by scattering. In this paper, on the other hand, we will study the \\textit{recoil} of the scattering on the planet for given planetesimal properties. These calculations provide, for the first time, an analytic expression for the two-sided torque for planetesimal scattering -- the analogue to the type-I migration torque. We assume the following: \\sumi\\ a smooth disk where the spatial distribution of surface density and eccentricity are power-laws; \\sumii\\ the dispersion-dominated regime (relative velocities are given by the eccentricity of the planetesimals at close encounter); \\sumiii\\ Keplerian orbits for the planetesimals; \\sumiv\\ a circular orbit for the planet. We account for both distant and close encounters, corresponding to orbits that do or do not cross the planet (see \\fg{disk}). We then compute the recoil of the planet due to scatterings with planetesimals on orbits both interior and exterior to the planet, which results in the PDM rate. PDM can be divided into three regimes, depending on the ratio of the planet mass compared to the mass of the solids with which it interacts: \\begin{enumerate} \\item Low mass planets. They do not exert a (strong) feedback on the disks. Correspondingly, the gradients in planetesimal's eccentricity and surface density are those of the background disk ($\\alpha$ and $\\beta$ in \\fg{disk}) and can be assumed fixed during the migration; \\item Massive planets. They have difficulty to migrate over large distances due to their inertia. Instead, the planet scatters away the planetesimals, leaving a gap \\citep{Rafikov2003iii,Rafikov2003}. \\item Intermediate-mass planets. They exert some feedback on the disk but not enough to halt their migration. \\end{enumerate} In \\sesto{model}{torques} the first regime is assumed. In \\se{model} the calculation for the migration rate due to distant and close encounters are presented. In \\se{torques} the net torque and the corresponding migration timescale are computed and compared to the type-I migration timescale. Furthermore, the approach is sketched how a distribution in eccentricity must be incorporated. In \\se{diff} the importance of diffusive motions (`noise') is investigated. Then, in \\se{sus-migr} we focus on the third regime and find that the migrating planet regulates the local eccentricity profile. Furthermore, we will outline the boundaries dividing the regimes and find that the intermediate regime covers a large region of the parameter space. We summarize our results in \\se{diss}. \\begin{figure}[t] \\plotone{fig1.eps} \\caption{Sketch of the disk profile. A planet on a circular orbit ($e=0)$ at semi-major axis $a_0$ interacts with planetesimals either through close or distant encounters. Planetesimals that are able to cross the planet's semi-major axis interact via close encounters; otherwise the encounters are distant. We allow for a power-law profile of surface density and eccentricity with indices $\\alpha$, $\\beta$ (see \\eq{sige0}) and compute the net migration rate $da_0/dt$ that the planet experiences due to scattering of the planetesimals in the dispersion-dominated regime. A nonzero $\\beta$ causes the transition between the regimes (dotted and dashed vertical lines) to shift by an amount $\\approx$$\\beta e_0^2 a_0$ (see text).} \\label{fig:disk} \\end{figure} \\begin{deluxetable}{lp{7cm}} \\tablecaption{\\label{tab:constants}List of frequently-used symbols.} \\tablehead{ Symbol & Description } \\startdata $\\Delta v_i$ & Change in the $i$-th component of the relative velocity \\\\ $\\Gamma$ & Dimensional torque on planet\\\\ $\\Lambda$ & Coulomb factor \\\\ $\\Sigma(a),\\Sigma_0$ & Surface density of planetesimals at disk radius $a$ or at reference radius $a_0$ \\\\ $\\Omega_0,\\Omega_a$ & Orbital frequency at corresponding to the reference position or to semi-major axis $a$ \\\\ $\\alpha$ & Exponent in surface density power-law (\\eqp{sige0})\\\\ $\\beta$ & Exponent in eccentricity power-law (\\eqp{sige0}) \\\\ $\\gamma_X$ & Dimensionless torque for close or distant encounters (\\eqp{gam-X}) \\\\ $\\gamma_\\mathrm{cv}$ & Curvature component of dimensionless torque (\\eqp{gam-X}) \\\\ $\\gamma_\\nabla$ & Gradient component of dimensionless torque (\\eqp{gam-X}) \\\\ $\\delta$ & Exponent in inclination power-law \\\\ $\\phi$ & Azimuthal coordinate in cylindrical coordinate system \\\\ $\\nu$ & Viscosity or diffusion rate in semi-major axis (\\eqp{nu-scat}) \\\\ $\\theta$ & True anomaly ($\\theta=0$ indicates periapsis; Appendix) \\\\ $\\Diff{i}$ & Diffusion coefficient: rate of change in relative velocity (\\eqp{Diffi})\\\\ $G_N$ & Newton's gravitational constant \\\\ $M_\\star$ & Stellar mass \\\\ $M_p$ & Planet mass \\\\ $Q_\\mathrm{pd}$ & Combination of $q_d$ and $q_p$, \\eq{Qpd} \\\\ $R$ & Radial coordinate in cylindrical units \\\\ $R_h$ & Hill radius \\eqp{Rhill} \\\\ $\\tilde{P}_{Rz}$ & Probability density of finding a planetesimal near $R=a_0$ and $z=0$ (\\eqp{PRz-tilde})\\\\ $T_\\mathrm{migr}$ & Timescale to migrate globally over distance $a_0$ \\\\ $T_\\mathrm{migr}^\\ast$ & Timescale to migrate locally over distance $e_0 a_0$ \\\\ $T_\\mathrm{syn}$ & Synodical period \\\\ $T_\\mathrm{type-I}$ & Type-I migration timescale \\\\ $T_\\mathrm{vs}$ & Viscous stirring timescale (\\eqp{Tvs-0})\\\\ $V_{k0}$ & Kepler (orbital) velocity corresponding to $a_0$ \\\\ $a_\\mathrm{[in,ou]}$ & Inner or outer-most semi-major axis from where planetesimals cross the planet's orbit \\\\ $a_0$ & Semi-major axis of the planet; reference radius \\\\ $b$ & Distance between semimajor axis planet and planetesimal \\\\ $[e,e_0]$ & Eccentricity of planetesimals (at $a=a_0$) \\\\ $e_h$ & Hill eccentricity (\\eqp{eh}) \\\\ $e_h^\\star$ & Lower range of the Hill eccentricity for which self-regulated migration applies (\\eqp{ehast}) \\\\ $i$ & Inclination of planetesimals \\\\ $f_\\Lambda$ & Coulomb term, $f_\\Lambda = \\log (1+\\Lambda^2)$ \\\\ $g_\\Lambda$ & Coulomb term, $g_\\Lambda = \\Lambda^2/(1+\\Lambda^2)$\\\\ $m$ & Mass of individual planetesimal \\\\ $q_d$ & Dimensionless `disk mass' (\\eqp{qdisk}) \\\\ $q_p$ & Dimensionless mass of the planet (=$M_p/M_\\star$) \\\\ $r$ & Radial coordinate in polar coordinate system \\\\ $v$ & Relative velocity between planet and planetesimal at $a=a_0$ \\\\ $v_i$ & Relative velocity of $i$-th component \\\\ $z$ & Vertical coordinate in cylindrical units \\enddata \\end{deluxetable} ", "conclusions": "In this paper, we have employed detailed analytical and numerical calculations to obtain the net torque acting on a planet due to the recoil from gravitational interactions with planetesimals in the dispersion-dominated regime. We have included both distant and close encounters and obtained the net migration rate by summing the torques from the interior and the exterior disks. We list our conclusions: \\begin{enumerate} \\item While the magnitude of the migration rate is primarily determined by the local values of the surface density ($\\Sigma_0$) and eccentricity ($e_0$), the direction is given by the local gradient in these quantities ($\\alpha$ and $\\beta$) and by the Coulomb factor $f_\\Lambda$ (\\fg{Itot}). Usually, the contribution from close encounters will determine the migration direction, unless $f_\\Lambda$ (and by implication $e$) are low. \\item The expressions for the migration timescale \\eq{Tmigr} due to planetesimal scattering display similarities to type-I migration (\\eq{Ttype-I}), if one replaces the disk mass in planetesimals by that of the gas and the eccentricity by $c_s/V_k$. Since the disk mass in solids is lower, the planetesimal-driven migration timescale is generally longer. \\item Under certain conditions a much faster migration mode (rivaling that of type-I) is obtained when the feedback of the planet on the disk is accounted for. The planet then self-regulates the value of the eccentricity gradient to $\\beta_\\mathrm{sr}$, which is a function of the local physical parameters (\\eq{beta-2}). Generally, $|\\beta_\\mathrm{sr}| \\gg 1$ and migration is quite rapid (\\eq{Tmigr}). \\item As function of the dimensionless disk mass $q_d$ (\\eq{qdisk}), planet mass $q_p=M_p/M_\\star$, and planetesimal eccentricity $e$, we have identified three migration regimes (\\fg{Qpd}) representing: (I) low mass planets, for which disk excitation is negligible; (II) high-mass planets, too massive to migrate significantly; and (III) intermediate-mass planets, which exert a mild feedback on the disk and migrate in the self-regulated mode. \\end{enumerate}" }, "1207/1207.0512_arXiv.txt": { "abstract": "The merger of two white dwarfs (WDs) creates a differentially rotating remnant which is unstable to magnetohydrodynamic instabilities. These instabilities can lead to viscous evolution on a time-scale short compared to the thermal evolution of the remnant. We present multi-dimensional hydrodynamic simulations of the evolution of WD merger remnants under the action of an $\\alpha$-viscosity. We initialize our calculations using the output of eight WD merger simulations from \\citet{Dan11}, which span a range of mass ratios and total masses. We generically find that the merger remnants evolve towards spherical states on time-scales of hours, even though a significant fraction of the mass is initially rotationally supported. The viscous evolution unbinds only a very small amount of mass $(\\la 10^{-5} \\Msun)$. Viscous heating causes some of the systems we study with He WD secondaries to reach conditions of nearly dynamical burning. It is thus possible that the post-merger viscous phase triggers detonation of the He envelope in some WD mergers, potentially producing a Type Ia supernova via a double detonation scenario. Our calculations provide the proper initial conditions for studying the long-term thermal evolution of WD merger remnants. This is important for understanding WD mergers as progenitors of Type Ia supernovae, neutron stars, R Coronae Borealis stars and other phenomena. ", "introduction": "\\label{sec:intro} Systems consisting of two white dwarfs (WDs) are natural outcomes of binary stellar evolution. These binaries are not static; absent any other torques the loss of angular momentum via gravitational wave (GW) emission will drive the binary together. Programs such as the SWARMS survey \\citep{Mullally09} and the ELM survey \\citep{Brown10} have dramatically increased the number of known WD binaries, including some systems that will merge within a Hubble time \\citep{Kilic12}. The Galactic population of WD binaries is expected to be a source of unresolved GW foregrounds at mHz frequencies, though only a handful of presently known systems would be individually detectable by a space-based GW interferometer mission \\citep{Nelemans09}. Details of the inspiral, in particular whether tidal torques cause the binary to be synchronized and the location of the tidal heating, are active areas of inquiry that can have a significant impact on the dynamics of the binary and the thermal state of the WDs \\citep{Fuller12}. As the orbital separation shrinks, the less massive (and hence larger) WD will eventually overflow its Roche lobe and begin transferring mass to the companion. The stability of this mass transfer depends on e.g., whether the material forms a disc or flows directly onto the companion, which in turn depends on the mass ratio ($q$) and total mass ($M_{\\mathrm{tot}}$) of the binary \\citep[e.g.][]{Marsh04}. Those systems that do undergo unstable mass transfer and subsequently merge have been of substantial theoretical interest. In particular, such systems have received attention as the possible progenitors of Type Ia supernovae \\citep{Iben84,Webbink84}. Considerable work exists exploring this ``double degenerate'' scenario and recent observational results have begun to favor it \\citep[e.g.][]{Bloom12, Schaefer12}. Another possibility is that double white dwarf binaries with total masses exceeding the Chandrasekhar mass undergo accretion induced collapse to form a neutron star \\citep[e.g.][]{Saio85}. Less massive double degenerate systems are likely to have non-explosive outcomes and have been invoked to explain objects like the R Coronae Borealis stars and extreme helium stars \\citep{Webbink84,Saio00,Clayton07}. An accurate simulation of the merger process requires a 3D code without prescribed geometry and with good numerical conservation properties. For these reasons, the pioneering study of \\citet{Benz90} used smoothed particle hydrodynamics (SPH). More recent studies \\citep[e.g.][]{Dan11,Raskin12,Pakmor12} have improved on these first results by contributing additional physics, more accurate initial conditions, higher resolution and more sophisticated numerical techniques. These simulations follow the evolution of the binary through the tidal disruption of one of the components. In some cases the merger is sufficiently violent that an explosion may result \\citep{Pakmor10, Dan12}. When the merger itself does not trigger an explosion, some material from the disrupted lower mass WD forms a shock-heated layer at the surface of the primary WD while the rest of the material forms a thick disc at larger radii. The evolution of such systems has frequently been treated in the literature as a long-lived ($\\sim 10^5$ yr) phase of accretion from a disc at the Eddington limit \\citep[e.g.][]{Nomoto85}. This picture was improved by \\citet{Yoon07}, who considered accretion at a similar rate but onto a hot envelope, and by \\citet{van-Kerkwijk10}, who made simple $\\alpha$-disc estimates of the accretion time-scale and found it to be far more rapid ($\\sim$ hours) than the time-scale for accretion at the Eddington limit. Recently, \\citet{Shen12} provided a new model of the different evolutionary phases of WD merger remnants. They argued that the evolution is much more ``star-like'' than the accretion disc oriented models that have dominated the literature. More concretely, \\citet{Shen12} showed that the rapid dynamical evolution of the merger ($\\sim 10^{2}$ s) gives way to a longer lived viscous phase driven by magnetohydrodynamic instabilities ($\\sim 10^{4} - 10^{8}$ s) before the onset of a long ($\\sim 10^4$ yr) thermal phase. In contrast with previous work, this implies that the long term evolution of a white dwarf merger remnant is not determined by accretion, but rather by the internal redistribution of heat/momentum and the external cooling rate of the viscously heated, nearly shear-free remnant. In \\citet{Shen12}, the viscous evolution was calculated in 1D using a $\\gamma$-law equation of state. The goal of this work is to refine the understanding of the outcome of the viscous evolution of WD merger remnants using higher dimensional numerical simulations. In addition, we consider a wider variety of WD+WD systems than \\citet{Shen12}, who focused on roughly Chandrasekhar mass CO+CO mergers. In \\S2 we outline the numerical methods we use, including how we construct our initial conditions from simulations by \\citet{Dan11}. In \\S3 we present the results of each of our calculations. \\S4 provides a discussion of the end states of the calculations. In \\S5 we state our conclusions and propose avenues for future work. In an Appendix, we show various test calculations that confirm the results we focus on in the main text. ", "conclusions": "The merger remnants of binary white dwarfs are differentially rotating and unstable to MHD instabilities like the MRI. As outlined by \\citet{Shen12}, MHD stresses give rise to a viscous phase of evolution which occurs on a time-scale much less than the thermal time. To investigate the outcome of this viscous evolution, we perform multi-dimensional hydrodynamic calculations of the evolution of WD binary remnants under the action of an $\\alpha$-viscosity. The initial conditions for these calculations are the SPH simulations by \\citet{Dan11}. We find that these remnants evolve towards spherical states on time-scales of hours. This confirms the arguments in \\citet{Shen12} that the post-merger evolution of WD merger remnants is via viscous redistribution of angular momentum that leads to nearly solid body rotation. The transport of angular momentum outwards removes rotational support from the majority of the mass leading to a nearly spherical remnant. {\\em The dynamics during this phase is not consistent with accretion at the Eddington limit,} as in previous models of WD merger remnants \\citep[e.g.][]{Nomoto85,Saio98,Saio04,Piersanti03a,Piersanti03b}. Instead, the viscous evolution of WD merger remnants is much more analogous to that of a differentially rotating star. Viscous heating associated with the approach to solid body rotation unbinds only a very small amount of mass ($\\la 10^{-5} \\Msun$ in our fiducial calculation). This is in contrast to some of the intuition developed in the context of radiatively inefficient accretion flows, which predict outflows. To understand this, we perform simple accretion tori calculations which indicate that the relatively small radius difference between the disc and the surface of the WD can explain why only a small amount of mass becomes unbound (see \\S\\ref{sec:COCO}). Viscous heating causes one of the systems we simulate to reach conditions of nearly dynamical He burning, so it is possible that the post-merger viscous evolution triggers a detonation in some cases. Recently \\citet{Dan12} presented a suite of more than 200 WD merger simulations which more thoroughly populate the $q$-$M_{\\mathrm{tot}}$ plane. They found that many of these systems reached the conditions for detonation during the merger (see for example their figures 6 \\& 8). In our calculations, $\\min(t_{\\mathrm{burn}} / t_{\\mathrm{dyn}})$ decreases by a factor of $\\sim 10$ during the viscous phase (see \\S 4.2). We speculate that systems that have $t_{\\mathrm{burn}} / t_{\\mathrm{dyn}} \\la 10$ at the merger may reach conditions for detonation during the subsequent viscous phase. However, we estimate that the number of systems which would satisfy this condition but have not reached $\\min(t_{\\mathrm{burn}} / t_{\\mathrm{dyn}}) <1$ during the dynamical phases of the merger is likely to be small. If other earlier detonation mechanisms do not prove to be robust, viscous heating could potentially trigger a surface detonation after the merger, causing either a .Ia supernovae \\citep{Bildsten07} or a Type Ia supernova via a double detonation scenario. Our purely hydrodynamic simulations cannot address the effects of magnetic fields. MHD simulations resolving the action of the MRI would allow a more realistic treatment of the viscous stresses than an $\\alpha$-viscosity\\footnote{It is worth noting that MHD simulations which capture the evolution of the entire remnant promise to be quite challenging. The instabilities in regions where $d\\Omega/dr > 0$ are likely to be short wavelength non-axisymmetric modes that have a different time-scale and spatial scale than the MRI modes that operate where $d\\Omega/dr < 0$. Correctly capturing the physics both inside and outside the rotation peak will be extremely difficult.}, though the quantitative insensitivity of our results to the value of $\\alpha$ leads us to think that our conclusions are robust. Converting our fiducial value of $\\alpha$ to a magnetic field strength gives $|B| \\sim \\sqrt{4 \\pi \\alpha \\rho c_S^2} \\sim 10^{10}$G. The implications of this estimate for the subsequent evolution of the merger remnant depend on the structure of the field. The generation of a large-scale field could lead to the formation of a strongly magnetized WD, which would be rapidly rotating and would quickly spin down via a magnetized wind. The presence of a strong magnetic field would also affect the conduction of heat in the interior of the WD. Alternatively, it is possible that the strong field is relatively small scale and so efficiently redistributes angular momentum in the interior of the remnant but does not significantly affect its global properties. The end states of our calculations provide a starting point for investigations of the long-term thermal evolution of WD merger remnants. In our fiducial case, we expect that the luminosity from the nuclear burning will drive convection, establishing an extended convective envelope with its base at slightly larger radii than the temperature peak. The object will likely grow to have a radius comparable to that of a giant star and correspondingly a relatively cool effective temperature like the models presented in \\citet{Shen12}. There are clear opportunities for future work in the self-consistent thermal evolution of these objects and their consequences for Type Ia supernovae, neutron stars, R Coronae Borealis stars and other phenomena." }, "1207/1207.1722_arXiv.txt": { "abstract": "We present a novel proposal strategy for the Metropolis-Hastings algorithm designed to efficiently sample general convex polytopes in 100 or more dimensions. This improves upon previous sampling strategies used for free-form reconstruction of gravitational lenses, but is general enough to be applied to other fields. We have written a parallel implementation within the lens modeling framework GLASS. Testing shows that we are able to produce uniform uncorrelated random samples which are necessary for exploring the degeneracies inherent in lens reconstruction. ", "introduction": "\\label{introduction} Some inversion problems in astrophysics make it desirable to search or sample a high dimensional solution domain $S \\subset \\mathbb{R}^n$ such that $S$ is bounded by the linear constraints \\begin{equation} \\label{inequalities} \\mat Ax \\le b \\end{equation} where $\\mat A \\in \\mathbb{R}^{m\\times n}$ and $b$ is a constant vector. A classic application is Schwarzschild's construction of triaxial stellar systems in equilibrium \\citep{1979ApJ...232..236S}. Given a three-dimensional discretized target density function $\\rho_j$, the number of stars $c_i$ on a given orbit $i$ is found by solving \\begin{equation} \\rho_j = \\sum_i c_i \\sigma_{i,j} \\end{equation} where $\\sigma_{i,j}$ is the orbit density. The orbit density is calculated \\emph{a priori} using test particles in a fixed potential corresponding to $\\rho$. However, searching the model space was not feasible at the time and only some particular models were considered. More recent work has further developed this technique \\citep[and references therein]{1982ApJ...263..599S,1999PASP..111..129M,2006MNRAS.366.1126C}. In this paper we consider applications to gravitational lensing. Lensing has had quite a long history, beginning with the first direct evidence of general relativity, but until 1979 with the discovery of the extra-solar lens Q0957+561 \\citep{1981ApJ...244..736Y} the field was largely of only theoretical interest \\citep{1964MNRAS.128..295R,1964MNRAS.128..307R}. Today more than one hundred strong lensing objects are known with many studied in great detail \\citep[e.g.,][]{1999AIPC..470..163K,2008ApJS..176...19F,2009ApJ...705.1099A}. Future surveys promise to deliver thousands more. Of utmost interest is the mass distribution of the lensing object. Characterizing this distribution is important for understanding the properties of galaxies and clusters \\citep{2007ApJ...667..645R,2010A&A...518A..55S}, galaxy formation and evolution \\citep{2010ApJ...721L...1T,2011A&A...529A..72F}, the nature of dark matter \\citep{2006ApJ...648L.109C}, as well as estimating cosmological parameters \\citep{2001PhR...340..291B} and the age of the universe \\citep{2006ApJ...650L..17S,2007ApJ...660....1O,2008ApJ...679...17C}. Crucially, the equations governing gravitational lensing are linear in the projected mass density $\\kappa$. As detailed in \\secref{framework}, one can discretize $\\kappa$ onto a grid of pixels and solve for physically motivated solutions by imposing constraints in the form of \\eqnref{inequalities}. Several versions of this idea have been developed by \\citet{2004AJ....127.2604S}, \\citet{2008ApJ...681..814C}, and \\citet{2005MNRAS.363.1136K}. This free-form approach is more flexible than simple analytic models, which assume a functional form of the mass profile and may unintentionally break degeneracies. However, this creates a large system of linear equations that is highly underconstrained. To understand the range of degeneracies we therefore require a technique that can explore the space of solutions $S$. One possible technique is to choose a random point $x$ and accept it if $x$ lies in $S$. This might be a reasonable method in low dimensions $n$, as is done in Monte Carlo integration, but the probability of acceptance rapidly approaches zero as $n$ increases. Each of the pixels in the discretization of $\\kappa$ represents one dimension and typically $n$ is greater than 100. Complex systems, where multiple lenses are used, can easily have more than 1000 dimensions. The priors can also be arbitrary, so the simplex will have a very complex shape, although by construction it will always be convex. General sampling of probability distributions has been a topic of statistics research for many years \\citep[e.g.,][]{1995...Chib...Greenberg,Robert:2005:MCS:1051451}. In the case of lensing, the PixeLens algorithm \\citep{2004AJ....127.2604S} is frequently used. We show, however, that the sampling of this algorithm is not uncorrelated. We address these details and related issues in \\secref{PixeLens} and suggest an alternative based on the Metropolis-Hastings algorithm in \\secref{new algorithm}. In \\secref{implementation} we discuss the implementation and demonstrate in \\secref{algorithm eval} that even for high dimensions we are able to sample our solution space to achieve a uniform uncorrelated random sample. We also achieve significant speed improvements over PixeLens. In \\secref{outlook} we discuss future work and applications. ", "conclusions": "" }, "1207/1207.6102_arXiv.txt": { "abstract": "{We analyse high-resolution near-UV and optical spectra of the afterglow of \\grb{}, obtained with the Very Large Telescope Ultraviolet and Visual Echelle Spectrograph (VLT/UVES), to investigate the circumburst environment and the interstellar medium of the gamma-ray burst (GRB) host galaxy. The VLT rapid-response mode (RRM) enabled the observations to start only 13 minutes after the \\textit{Swift} trigger and a series of four exposures to be collected before dawn. A low neutral-hydrogen column-density (log\\,$N$(\\hi{}) $=18.7$) is measured at the host-galaxy redshift of $z=2.42743$. At this redshift, we also detect a large number of resonance ground-state absorption lines (e.g., \\cii{}, \\mgii{}, \\alii{}, \\siii{}, \\crii{}, \\civ{}, \\siiv{}), as well as time-varying absorption from the fine-structure levels of \\feii{}. Time-varying absorption from a highly excited \\feiii{} energy level ($^7$S$_{3}$), giving rise to the so-called UV\\,34 line triplet, is also detected, for the first time in a GRB afterglow. The \\crii{} ground-state and all observed \\feii{} energy levels are found to depopulate with time, whilst the \\feiii{} $^7$S$_{3}$ level is increasingly populated. This absorption-line variability is clear evidence of ionization by the GRB, which is for the first time conclusively observed in a GRB afterglow spectrum. We derive ionic column densities at each epoch of observations by fitting absorption lines with a four-component Voigt-profile model. We perform CLOUDY photo-ionization modelling of the expected pre-burst ionic column densities, to estimate that, before the onset of the burst, [C/H] $=-1.3\\pm0.2$, [O/H] $<-0.8$, [Si/H] $=-1.2\\pm0.2$, [Cr/H] $=+0.7\\pm0.2$, and [Fe/H] $=+0.2\\pm0.2$ for the integrated line profile, indicating strong overabundances of iron and chromium. For one of the components, we observe even more extreme ratios of [Si/Fe] $\\leq-1.47$ and [C/Fe] $\\leq-1.74$. These peculiar chemical abundances cannot easily be explained by current models of supernova yields. They are indicative of a low dust content, whilst dust destruction may also contribute to the marked Fe and Cr overabundances. The overall high iron enhancement along the line-of-sight suggests that there has been negligible recent star formation in the host galaxy. Thus, the occurrence of a GRB indicates that there has been episodes of massive star formation in the GRB region.} ", "introduction": "\\label{sec intro} Gamma-ray burst (GRB) afterglows, which are produced by the shock between the burst jet and either the surrounding interstellar medium (ISM) or stellar wind, radiate their extremely powerful and featureless continuum through their host galaxies \\citep[see][for reviews]{Piran04,Meszaros06}. This offers a unique opportunity to study the properties of such distant and faint galaxies. In particular, optical and near-ultraviolet (near-UV) spectroscopy of long-duration ($t_{\\rm obs}>2$~s) GRBs reveals the chemical composition of the gas along the line-of-sight through absorption lines imprinted in the afterglow spectrum \\citep[e.g.,][]{Prochaska07}. However, the vast majority of spectra have to date been taken at low resolution \\citep[e.g.,][]{Fynbo09,Christensen11}, hampering metallicity studies. At the high-resolution end ($R\\sim45,000$), the Very Large Telescope (VLT) Ultraviolet and Visual Echelle Spectrograph \\citep[UVES,][]{Dekker00} and the Keck High-Resolution Echelle Spectrometer \\citep[HIRES,][]{Vogt94} allow detailed studies of the metal abundances, physical states, and gas kinematics within the host galaxies and possibly lower-redshift intervening systems \\citep{Fiore05,Prochaska07,Vreeswijk07,Fox08,Ledoux09,Vergani09}, albeit for a small sample of bright afterglows. Among the six \\textit{Swift}-era GRB afterglows with metallicity determinations from VLT/UVES spectroscopy, four absorbers have low metallicities, [X/H$]<-1$, where X is taken to be either Zn or S. These generally have low dust depletion factors, [X/Fe], indicating a low dust content, while in one case a higher [S/Fe] ratio is suggestive of an $\\alpha$-element enhancement \\citep{Ledoux09}. On the other hand, \\citet{Prochaska07} found that [$\\alpha$/Fe] and [Zn/Fe] tend to be higher in GRB absorbers than along quasi-stellar object (QSO) lines-of-sight, based on high- and low-resolution spectroscopy of a sample of 16 absorbers, indicating significant contributions from massive stars and/or differential dust depletion. However, the current GRB absorber samples are still too limited in size and possibly too biased to determine what the typical abundances of GRB host galaxies are. A larger collection of metal column densities in 29 GRB absorbers, drawn from both high- and low-resolution spectroscopy \\citep{Schady11}, reveals that the ranges of relative chemical abundances measured in GRB absorbers are generally similar to those determined along QSO lines-of-sight. In this paper, we present peculiar metal abundances in the host galaxy of \\grb{} and discuss their possible origin. Thanks to time-resolved high-resolution spectroscopy, the variability of fine-structure lines has been detected in GRB absorbers and explained by photo-excitation induced by the incident GRB afterglow flux \\citep[i.e., UV pumping,][]{Vreeswijk07,Prochaska06a}. On the other hand, the variability of resonance metal lines (i.e., those associated with the ground-state level) is expected in the case of photo-ionization \\citep[e.g.,][]{Perna98} but has never been conclusively detected until now \\citep[the possible variability of \\lya{} towards GRB\\,090426 was reported by][]{Thone11}. Here we present the detection of significant variability in both fine-structure and resonance metal absorption lines in the afterglow spectrum of \\grb{}. In addition, a triplet of lines arising from a highly excited level of \\feiii{} is observed in this absorber, which is a primer in a GRB afterglow spectrum. As we discuss in this paper, these observations clearly show that photo-ionization by the GRB has taken place. Self-consistent photo-excitation and ionization modelling to determine the GRB-absorbing cloud distance will be presented in a companion paper \\citep[][hereafter referred to as paper II]{Vreeswijk12}. This paper is organized as follows. In Sect.~2, we report on the observations and the data reduction. In Sect.~3, we derive ionic column densities from Voigt-profile fitting of identified absorption lines and in Sect.~4 we present CLOUDY photo-ionization modelling. We discuss our results in Sect.~5 and conclude in Sect.~6. Throughout the paper, we adopt ions cm$^{-2}$ as the linear unit of column density $N$. The relative abundance of two chemical elements, $X$ and $Y$, is defined as $\\left[X/Y\\right] \\equiv \\log{\\frac{N(X)}{N(Y)}} - \\log{\\frac{N(X)_\\odot}{N(Y)_\\odot}}$. We estimate its uncertainty by adding in quadrature the errors in the observed column densities and in the solar abundances involved. For the reference solar abundances appearing in the second term of this formula, we follow the recommendations of \\citet{Lodders09}, adopting either their meteoritic estimates, the photospheric values of \\citet{Asplund09}, or the average of these both. The solar abundances used in this paper are summarized in Table~\\ref{tab solar}. \\begin{table} \\caption{Solar abundances used in this paper} \\centering \\begin{tabular}{ @{}c c c | c c c @{}} \\hline \\hline \\rule[-0.2cm]{0mm}{0.8cm} Element & $A(\\rm El)\\pm\\sigma_A$ $^a$ & S$^b$ & Element & $A(\\rm El)\\pm\\sigma_A$ $^a$ & S$^b$\\\\ \\hline \\rule[-0.0cm]{0mm}{0.4cm} H & $\\equiv12$ & s & Cr & $5.64\\pm0.04$ & a \\\\ C & $8.43\\pm0.05$ & s & Fe & $7.47\\pm0.04$ & a\\\\ O & $8.69\\pm0.05$ & s & Ni & $6.21\\pm0.04$ & a\\\\ \\rule[-0.2cm]{0mm}{0.4cm} Si & $7.51\\pm0.01$ & m & Zn & $4.63\\pm0.04$ & m\\\\ \\hline\\hline \\end{tabular} \\tablefoot{$^a$ Abundances $A(\\rm El) \\equiv \\log N(\\rm El)/N(\\rm H) + 12$ are taken from \\citet{Asplund09} following the recommendations of \\citet{Lodders09}. $^b$ Source of the estimate: solar photosphere (s), meteorites (m), or the average between the two (a).} \\label{tab solar} \\end{table} \\begin{table*} \\caption{Log of VLT/UVES observations} \\centering \\begin{tabular}{ @{}l | c c @{\\hspace{3mm}} c @{\\hspace{3mm}} c @{\\hspace{3mm}} c c @{\\hspace{2mm}} c @{\\hspace{2mm}} c @{\\hspace{3mm}} c c @{}} \\hline \\hline \\rule[-0.2cm]{0mm}{0.6cm} Ep. & $t_{\\rm start}^a$ & $\\Delta\\,t^b$& $t_{\\rm exp}$ & Dic. & Setting & Coverage & FWHM$^c$ & Mean & S/N$^d$ & $R=\\frac{\\lambda}{\\Delta\\lambda}$\\\\ \\rule[-0.2cm]{0mm}{0.6cm} & UT & (min) & (min) & \\# & $\\lambda_{\\rm c}$ (nm)&$\\lambda_{\\rm obs}$ (nm)& blue/red ($\\arcsec$) &airmass & blue/red & blue/red\\\\ \\hline & & & & & & & & & & \\\\ I& 08:51:05 & 14.55 & 3 & 1 & 346+580 & 303--388; 476-- 684 & 1.1/1.0 & 1.12 & 1.9--2.0/4.0--4.5 & 48,500/45,600\\\\ II& 08:56:46 & 21.23 & 5 & 2 & 437+860 & 373--499; 660--1060 & 1.2/0.9 & 1.12 & 3.4--4.7/2.7--8.4 & 47,900/45,300\\\\ III& 09:04:20 & 31.25 & 10 & 1 & 346+580 & 303--388; 476-- 684 & 1.5/1.3 & 1.14 & 3.1--3.7/6.8--7.6 & 48,500/45,600\\\\ IV& 09:16:36 & 50.06 & 23 & 2 & 437+860 & 373--499; 660--1060 & 1.6/1.2 & 1.17 & 6.0--8.7/4.8-14.3 & 47,900/45,300\\\\ & & & & & & & & & & \\\\ \\hline \\hline \\end{tabular} \\tablefoot{$^a$ March 10 2008. $^b$ Mid-exposure time after the burst event in the observer's frame. $^c$ Full width at half maximum of the spatial profile in the two-dimensional spectrum. $^d$ Signal-to-noise per pixel; min--max ranges over the spectral regions where absorption lines are observed.} \\label{tab_log} \\end{table*} ", "conclusions": "Our time-series of high-resolution VLT/UVES spectra of the \\grb{} afterglow has revealed the unique features of the ISM of the GRB host galaxy. We have reported the detection of several resonance absorption lines commonly observed in the ISM, as well as \\feii{} fine-structure lines sometimes associated with GRB afterglows. In the case of \\grb{}, both the \\feii{} ground-state and the fine-structure lines vary with time. We decomposed the complex spectral-line profiles of several ions and modelled them altogether using a four-component Voigt profile, resulting in column-density determinations for each component. These components might be associated with different clouds along the line-of-sight within the GRB host galaxy. Interestingly, a low $\\log N$(\\hi{}) $=18.7\\pm0.1$ is derived from the \\lya{} absorption. We also detected the \\feiii{} UV\\,34 $\\lambda\\lambda\\lambda$1895,1914,1926 line triplet in absorption. These transitions arise from the $^7$S$_3$ energy level of \\feiii{}, which is observed for the first time in a GRB afterglow spectrum. The \\feii{} and \\feiii{} time variability, of both the ground-state and excited levels, is clear evidence that we are witnessing the ISM being gradually photo-ionized by the GRB afterglow radiation. Through the analysis of chemical abundances measured in the \\grb{} absorber and CLOUDY photo-ionization modelling, we inferred pre-burst ionization-corrected ratios of [C/Fe] $=-1.5\\pm0.2$, [O/Fe] $<-1.0$, and [Si/Fe] $=-1.4\\pm0.2$ (where [Fe/H] $=+0.2\\pm0.2$ and [Cr/H] $=+0.7\\pm0.2$), while, typically, $0\\lesssim$ [Si/Fe] $\\lesssim+0.7$ is observed in QSO and/or GRB absorbers. Furthermore, we observed an even more extreme iron overabundance ([Si/Fe] $\\leq-1.47$ and [C/Fe] $\\leq-1.74$) in the gas associated with component ``b'' of the absorption profile. Such a low [Si/Fe] ratio has never been observed before in either QSO- or GRB-DLAs and cannot easily be explained by current models of SN Ia and Pop~III SN chemical yields. A potential explanation might be provided by the destruction of iron-rich dust grains, thereby recycling heavy elements into the gas phase. The dust could be destroyed by the GRB itself or GRB-unrelated processes such as SNe shock waves. The high iron column-density measured in the gas phase generally suggests a low dust content in the absorber. The strong overabundance of iron compared to silicon and carbon also suggests that there has been negligible recent star formation along the line-of-sight. The occurrence of the GRB then indicates that there has been episodic massive-star formation in the GRB region." }, "1207/1207.1664_arXiv.txt": { "abstract": "A general formalism to include experimental reaction cross sections into calculations of stellar rates is presented. It also allows to assess the maximally possible reduction of uncertainties in the stellar rates by experiments. As an example for the application of the procedure, stellar neutron capture reactivities from \\cite{kadonis} are revised and the remaining uncertainties shown. Many of the uncertainties in the stellar rates are larger than those obtained experimentally. This has important consequences for s-process models and the interpretation of meteoritic data because it allows the rates of some reactions to vary within a larger range than previously assumed. ", "introduction": "\\label{sec:intro} An increasing number of reaction cross section measurements are performed with the aim to improve reaction rates for nucleosynthesis studies. The questions of how to assess the impact of the experimentally obtained cross sections on astrophysical reaction rates and how to convert laboratory cross sections to stellar rates inevitably arise. Closely related is the question concerning the estimate of remaining uncertainties in the stellar rates. These questions are addressed in the following. The stellar enhancement factor (SEF) -- the (theoretically predicted) ratio of stellar and laboratory rate -- was used in the past to derive the stellar rate from a measurement. It was shown in \\citet{xfactor} that the SEF is not adequate for this purpose and the ground state contribution $X_0$ was introduced to replace the SEF. Here, a further generalization of this concept is presented and a formalism for including data and their uncertainties into an improved stellar rate with revised uncertainties is laid out. The generalized relative level contribution to the stellar rate is introduced in \\S~\\ref{sec:The-ground-state}. The proper inclusion of data into stellar rates is discussed in \\S~\\ref{sec:Renormalization-of-theory}. The following \\S~\\ref{sec:Determining-uncertainty-factors} then explains how the attached error changes when using experimental information to improve a stellar rate. As an important application, neutron capture rates for the s-process are revised in \\S~\\ref{sec:application}. They are commonly believed to be strongly constrained by precise data but it will be shown that the remaining uncertainties are not dominated by the experimental errors in many cases. ", "conclusions": "A general formalism to include experimental reaction cross sections into calculations of stellar rates was developed, which also allows to assess the reduction of uncertainties in the stellar rates by experiments. As an important example for the application of the procedure, stellar neutron capture reactivities from \\cite{kadonis} were revised and the remaining uncertainties have been shown. Although the uncertainties in the stellar rates are close to the experimental values for a number of nuclei, many of the uncertainties remain considerably larger. This has important consequences for s-process models \\citep[see, e.g.,][]{arl99} and the interpretation of meteoritic data \\citep[see, e.g.,][]{qin,burk12} because it allows the rates of some reactions to vary within a larger range than previously assumed. The revised reactivities will be included in the upcoming new version of KADoNiS." }, "1207/1207.3382_arXiv.txt": { "abstract": "We present spectroscopic observations of ultra compact dwarf (UCD) galaxies in the Fornax and Virgo Clusters made to measure and compare their stellar populations. The spectra were obtained on the Gemini-North (Virgo) and Gemini-South (Fornax) Telescopes using the respective Gemini Multi-Object Spectrographs. We estimated the ages, metallicities and abundances of the objects from measurements of Lick line-strength indices in the spectra; we also estimated the ages and metallicities independently using a direct spectral fitting technique. Both methods revealed that the UCDs are old (mean age $10.8\\pm 0.7$ Gyr) and (generally) metal-rich (mean [Fe/H] = $-0.8\\pm 0.1$). The alpha-element abundances of the objects measured from the Lick indices are super-Solar. We used these measurements to test the hypothesis that UCDs are formed by the tidal disruption of present-day nucleated dwarf elliptical galaxies. The data are not consistent with this hypothesis because both the ages and abundances are significantly higher than those of observed dwarf galaxy nuclei (this does not exclude disruption of an earlier generation of dwarf galaxies). They are more consistent with the properties of globular star clusters, although at higher mean metallicity. The UCDs display a very wide range of metallicity ($-1.7<$[Fe/H]$<0.0$), spanning the full range of both globular clusters and dwarf galaxy nuclei. We confirm previous reports that most UCDs have high metalliticities for their luminosities, lying significantly above the canonical metallicitiy-luminosity relation followed by early-type galaxies. In contrast to previous work we find that there is no significant difference in either the mean ages or the mean metallicities of the Virgo and Fornax UCD populations. ", "introduction": "Ultracompact dwarf (UCD) galaxies were originally discovered as very compact stellar systems (CSS) in the nearby Fornax galaxy cluster \\citep{Hilker1999,Drinkwater2000}. Although as luminous as dwarf galaxies ($-11-1$), although the larger sample measured by \\citet{Chilingarian2011} included three objects with lower metalicity ($-1.39<$[Fe/H]$<-1.2$). We are now finding that UCDs exhibit the full range of (low) metalicities shown by globular cluster populations. Our data also allow us to compare the UCD populations between the Virgo and Fornax clusters as we used very similar Gemini (North and South) data for both clusters and identical analysis. Furthermore the objects observed had the same luminosity distribution in both clusters. Using the model fits we obtain mean ages of $(10.4 \\pm 1.1)$ Gyr for Fornax and $(11.2 \\pm 0.9)$ Gyr for Virgo with no significant differences in the distributions\\footnote{The Kolmogorov-Smirnov, t-test and F-test all gave $p>0.5$ that the samples were from the same distributions.}. Similarly there is no significant difference in the mean metalicities of [Fe/H] $=-0.79 \\pm 0.14$ for Fornax and [Fe/H] $=-0.90 \\pm 0.14$ for Virgo. (In each case, for age and metallicity we quote the standard error of the respective mean value.) This contradicts the earlier suggestions made \\citep{Mieske2006_UCD, Firth2009_UCD} that the Fornax UCD population was younger than that of Virgo due to differences in the cluster environment. The cluster environments are different (Virgo is larger and much less relaxed), but UCDs are mostly only found close to the central galaxies of each cluster, so we would argue that they experience similar local environments in the two clusters, resulting in similar internal properties. For two of the Virgo cluster objects with HST imaging, VUCD~3 and VUCD~5, \\citet{Evstigneeva2007_VUCD} published dynamical masses and $V$-band dynamical mass-to-light ratios: $(M/L)_{\\rm{dyn}}=5.4 \\pm 0.9$ and $3.9 \\pm 0.6$ in Solar units. Using our precise age and metallicity measurements, we can derive their stellar masses and hence estimate their dark matter fractions defined as $((M/L)_{\\rm{dyn}}-(M/L)_{*})/(M/L)_{\\rm{dyn}}$. The $V$-band stellar mass-to-light ratios for VUCD~3 and VUCD~5 are $(M/L)_{*}=4.5 \\pm 0.2$ ($8.1 \\pm 0.4$) and $2.85 \\pm 0.13$ ($5.0 \\pm 0.2$) $(M/L)_{\\odot, V}$ respectively, where values in parentheses correspond to the \\citet{Salpeter1955} IMF. The corresponding dark mass fractions are: $15 \\pm 25$ ($-50 \\pm 25$) and $25 \\pm 20$ ($-28 \\pm 20$) per cent. Negative values indicate that the stellar mass estimate exceeds the dynamical one, suggesting that the Salpeter IMF is not compatible with the observations \\citep[see discussion by][]{Chilingarian2011}. The fact that we do not detect dark matter in these two Virgo UCDs is in contrast to previous work suggesting that the Virgo UCDs had higher mass-to-light ratios than the Fornax UCDs. \\citet{Hasegan2005} reported mass-to-light ratios between 6 and 9 for their ``probable'' UCDs, significantly higher than comparison measurements for UCDs in the Fornax Cluster (1-3; see their Figure 11). Although we do not have data for a large sample of Virgo UCDs, VUCD~3 and VUCD~5 are the most massive UCDs in Virgo and for these at least we find no evidence of dark matter. \\begin{figure*} \\centering \\epsfig{file=LZrelation.ps,width=0.8\\linewidth,angle=0} \\caption{ The metallicity-luminosity relation for UCDs and other stellar systems. The model-fitted metalicities of each UCD are plotted against their absolute $B$ magnitudes. The dashed line traces the canonical metallicity-luminosity relation for early-type galaxies \\citep[e.g.\\ Fig.~8 of][]{Chilingarian2011}. Comparison values for different types of galaxy and globular clusters are taken from the literature according to the following codes: A496 - \\citet{Chilingarian2008abell} C08 - \\citet{Chilingarian2008fornax}, C09c - \\citet{Chilingarian2009cE}, C09v - \\citet{Chilingarian2009virgo}, C11 - \\citet{Chilingarian2011}, M98 - \\citet{Mateo1998}, MV05 - \\citet{McLaughlin2005}, P09 - \\citet{Price2009}. Data for the nearby cE galaxies are from \\citet{Chilingarian2008m59} (M59cO), \\citet{Chilingarian2010n5846} (N5846cE), \\citet{SanchezBlazquez2006} (N4486B), \\citet{Graham2002} (M32 luminosity) and \\citet{Worthey2004} (M32 metallicity). } \\label{fig-lmetal} \\end{figure*} In Fig.~\\ref{fig-lmetal} we plot the model-fitted UCD metallicities against their luminosities compared to previous measurements of UCDs and dwarf galaxies, as well as reference data for globular clusters and elliptical galaxies. (We use the model metallicities rather than the those derived from the Lick indices as they use more of the spectra dalta and have smaller statistical uncertainties.) It is immediately obvious that our data are much more precise than most of the previous work. With this precision, the two UCDs with lower metallicity stand out as being significantly different to the rest of the UCD population. We should note, however, that the nature of our sample selection means that we cannot exclude future measurements of intermediate objects. In fact, with increasing measurements we now see there is a continuous range of objects with intermediate properties (both luminosity or metallicity) filling the space between UCDs and compact elliptical (``cE'' or ``M32-like'') galaxies. The UCDs in Fig.~\\ref{fig-lmetal} with new Gemini measurements show no evidence of a metallicity-luminosity relation. We also measured the metallicity of the objects as a function of projected distance from the nearest large galaxy in both clusters and found no correlation in those parameters. In particular the two objects from the intra-cluster Virgo field (VUCD1232+0944 and VUCD1233+0952 at 550 kpc from M87) have metallicities entirely consistent with the other Virgo objects. Fig.~\\ref{fig-lmetal} does, however, show that (with two exceptions) the UCDs lie well above the canonical metallicity-luminosity trend for early-type galaxies (including nucleated dwarf elliptical galaxies), as was demonstrated previously for Fornax Cluster UCDs \\citep{Chilingarian2011}. This is exactly what we might expect from tidal disruption of current-day dwarf galaxies: the luminosity would decrease while the metallicity would remain high. On the other hand, the figure also shows that many globular clusters (albeit at lower metalicities and luminosities) also lie above the canonical metallicity-luminosity trend. Thus the UCDs could represent either the high-luminosity end of the globular cluster metallicity-luminosity relation, or the natural end point of a constant-metallicity stripping process of dwarf galaxies. There are now sufficient intermediate objects on both sides that the data shown in this figure cannot rule out either hypothesis. A promising approach to help separate the two populations would be to measure the physical sizes of the objects. This has recently been done by \\citet{Brodie2011_new} who found evidence that UCDs in the Virgo Cluster formed a distinct population with significantly larger sizes than globular clusters. For each of our Fornax cluster and Virgo cluster UCDs, we measured both the Lick/IDS line strength indices and performed the {\\sc nbursts} full spectral fitting technique, and compared these results to SSP models. Our findings are summarised below: \\begin{itemize} \\itemsep12pt \\parskip0pt \\parsep0pt \\item While much of the evidence, such as super-Solar $\\alpha$-abundance ratios, old age measurements, and a large spread in the metallicity measurements, seems to imply that the formation of UCDs is consistent with that of GCs, it does not rule out the threshing formation model. The data are only inconsistent with stripping of {\\em present-day} nucleated dwarf galaxies, but not with much older disruptions before the nuclei evolved to their current composition. \\item Many of our measurements contrast those from previous work. The differences are outlined below: \\begin{enumerate} \\itemsep6pt \\parskip0pt \\parsep0pt \\item Our metallicity measurements exhibit a larger range of metallicity than indicated in previous work, consistent with the full range of metallicities of GCs. \\item Our age and metallicity measurements, along with luminosity distributions, indicate similar UCD populations in both the Virgo and Fornax clusters. Previous work suggests that the different cluster environments cause the Fornax UCD population to be younger than that of Virgo. \\item We do not detect dark matter in the two Virgo UCDs with published dynamical masses and dynamical mass-to-light ratios, which happen to be the most massive of the Virgo UCDs measured, despite previous work indicating that the Virgo UCDs had higher mass-to-light ratios than the Fornax UCDs. \\end{enumerate} \\item The UCDs did not conform to a metallicity-luminosity relation, but did mostly lie above the metallicity-luminosity trend for early-type galaxies, which would be consistent with formation from present-day tidal stripping of dE,Ns. Intermediate objects on either side of the data indicate that neither hypothesis can be ruled out in this way. \\item The current data provide evidence for both formation hypotheses, and are insufficient to rule out either one. \\end{itemize}" }, "1207/1207.4870_arXiv.txt": { "abstract": "Compact stars can have either hadronic matter or can have exotic states of matter like strange quark matter or color superconducting matter. Stars also can have a quark core surrounded by hadronic matter, known as hybrid stars (HS). The HS is likely to have a mixed phase in between the hadron and quark phase. Observational results suggest huge surface magnetic field in certain neutron stars (NS) called magnetars. Here we study the effect of strong magnetic field on the respective EOS of matter under extreme conditions. We further study the hadron-quark phase transition in the interiors of NS giving rise to hybrid stars (HS) in presence of strong magnetic field. The hadronic matter EOS is described based on relativistic mean field theory and we include the effect of strong magnetic fields leading to Landau quantization of the charged particles. For the quark phase we use the simple MIT bag model. We assume density dependent bag pressure and magnetic field. The magnetic field strength increases going from the surface to the center of the star. We construct the intermediate mixed phase using Glendenning conjecture. The magnetic field softens the EOS of both the matter phases. The effect of magnetic field is insignificant unless the field strength is above $10^{14}$G. A varying magnetic field, with surface field strength of $10^{14}$G and the central field strength of the order of $10^{17}$G has significant effect on both the stiffness and the mixed phase regime of the EOS. We finally study the mass-radius relationship for such type of mixed HS, calculating their maximum mass, and compare them with the recent observation of pulsar PSR J1614-2230, which is about 2 solar mass. The observations puts a severe constraint on the EOS of matter at extreme conditions. The maximum mass with our EOS can reach the limit set by the observation. ", "introduction": " ", "conclusions": "" }, "1207/1207.1178_arXiv.txt": { "abstract": "We present combined VLA and Green Bank Telescope images of \\ammonia\\ inversion transitions (1,1) and (2,2) toward OMC2 and OMC3. We focus on the relatively quiescent Orion cores, which are away from the Trapezium cluster and have no sign of massive protostars nor evolved star formation, such as IRAS source, water maser, and methanol maser. The 5\\arcsec\\ angular resolution and $0.6~\\rm{}km\\,s^{-1}$ velocity resolution of these data enable us to study the thermal and dynamic state of these cores at $\\sim{}0.02~\\rm{}pc$ scales, comparable to or smaller than those of the current dust continuum surveys. We measure temperatures for a total of 30 cores, with average masses of $11\\,\\Ms$, radii of $0.039~\\rm{}pc$, virial mass ratio $\\overline{R_{vir}}$ = 3.9, and critical mass ratio $\\overline{R_{C}}$ = 1.5. Twelve sources contain \\textit{Spitzer} protostars. The thus defined starless and protostellar subsamples have similar temperature, line width, but different masses, with an average of $7.3\\,\\Ms$ for the former and $16\\,\\Ms$ for the latter. Compared to others Gould Belt dense cores, mores Orion cores have a high gravitational-to kinetic energy ratio and more cores have a larger thant unity critical mass ratio. Orion dense cores have velocity dispersion similar to those of cores in low-mass star-forming regions but larger masses for fiven size. Some cores appear to have truly supercritical gravitational-to-kinetic energy ratios, even when considering significant observational uncertainties: thermal and non-thermal gas mothins alone cannot prevent collapse. ", "introduction": "Stars form in molecular clouds. Within molecular clouds, discrete structures with observable column density contrast, particularly in high density tracers, with its surroundings are referred to as cores (e.g.\\ Benson \\& Myers 1989; Ward-Thompson et al.\\ 2007). The relatively high extinction and density of cores make them the likely site for the onset of collapse. Recent core surveys (e.g.\\ Ikeda et al.~ 2007; K{\\\"o}nyves et al.\\ 2010; Sadavoy et al.\\ 2010) are facilitated by focal plane imaging arrays of growing sizes. Although many surveys have large sample size in the hundreds, the majority of these studies are of dust continuum, which tend to focus on core mass function. Direct measurement of cores' thermal and dynamic structures require spectra maps of high density tracers and preferably cover the same spatial dynamic ranges as those of the dust emission. There seems to be an observational dichotomy between low mass star formation and high mass star formation. Massive stars are formed exclusively in giant molecular clouds and with higher efficiency. Due to the large distances of most massive star forming regions, many massive cores are under-resolved, showing signs of star formation well underway, such as H$_2$O masers (Mookerjea et al.~2004) and/or compact HII regions (Reid \\& Wilson 2005). At about 437 pc (Hirota et al.~2007; Menten et al.\\ 2007), Orion molecular clouds are the closest massive star forming region with an OB cluster, thus particularly suited for studying the early stages of star formation in massive cores and/or under the influence of young massive stars. In a series of papers, we identified 'quiescent' Orion clouds/cores (containing no HII region, no IRAS point sources, and are at last 1 pc away from the OB cluster) with \\ammonia\\ and \\n2h\\ with a beam size of about 1\\arcmin\\ (Li, Goldsmith \\& Menten 2003; hereafter paper I), presented high resolution 350 \\micron\\ images with a beam size of 9\\arcsec\\ (Li et al.\\ 2007; hereafter paper II), and revealed that two thirds of Orion cores have signatures of ongoing dynamic evolution, both outflows and inflows (Velusamy et al.\\ 2008 hereafter paper III). Based on dust mass and dust core size (paper II), the majority of the cores are seemingly supercritical, i.e., no adequate support from thermal pressure, turbulence(c.f. McKee \\& Tan 2003), or static magnetic field. This is consistent with the majority of the cores being hydrostatically unstable (paper III), which, however, could not be directly tested due to a lack of high resolution spectroscopic data. Spectroscopic survey of CO 1-0 (Williams, Plambeck \\& Heyer 2003) and CO 3-2 (Takahashi et al.~2008) reveal that part of the Orion molecular cloud, namely, OMC2 and OMC3, contains many molecular outflows. OMC2 and OMC3 have also been mapped in high density tracers. Using the Nobeyama 45m telescope, Tatematsu et al.\\ (2008) identifed 34 cloud cores in \\n2h\\ 1-0 with a beam width of around 18\\arcsec. Ikeda et al.\\ (2007) studied dense gas in the same region using H$^{13}$CO$^+$ also with Nobeyama. These studies find higher column density for respective tracers compared with low mass star forming regions, such as Taurus. \\ammonia\\ inversion transitions are particularly suited for studying dense cores due to their lack of depletion and their sensitivity to kinetic temperature (Ho \\& Townes 1983). Wiseman \\& Ho (1998) obtained \\ammonia\\ maps of a 8\\arcmin\\ $\\times$ 6\\arcmin\\ region around Orion-KL using VLA. The gas morphology there is dominated by quasi-parallel filaments severely influenced by the energy output of young massive stars. In this letter, we present the combined VLA and GBT ammonia survey of the cores in OMC2 and OMC3. The spatial resolution of $\\sim$ 5\\arcsec\\ and the spectral resolution of 0.6 \\kms\\ facilitate a detailed examination of the thermal and dynamic states of massive quiescent Orion cores. ", "conclusions": "We have mapped OMC2 and OMC3 in \\ammonia\\ (1,1) and (2,2) with both VLA and GBT. The combined single dish plus interferometric data provide a rare detailed look into the thermal and dynamic properties of a collection of massive quiescent cores. Our main results are: \\begin{enumerate} \\item We obtain good temperature measurement for 30 dust cores. The median core temperature is $17~\\rm{}K$. The typical uncertainty for derived temperature of each pixel is about $2~\\rm{}K$. \\item 12 cores are associated with a protostar. The average temperatures of the protostellar and the starless cores are similar suggesting that the heating in OMC2-3 region is primarily external. \\item A total of 24 cores (80\\%) are gravitationally bound ($R_{vir}>1$), and 11 cores (37\\%) achieve $R_{vir}>3$. Compared to other Gould Belt clouds, a much higher fraction of cores are tightly bounded. \\item 12 out of 30 cores (40\\%) are more massive than the critical mass defined as the combination of Jeans mass and mass supported by a steady magnetic field of $100~\\rm{}\\mu{}G$. \\end{enumerate} In summary, this sample of Orion cores, identified in dust emission with temperature and turbulence measured in \\ammonia\\ inversion lines, are proven to be well bounded by gravity and contains a substantial fraction of supercritical cores. They will evolve rapidly, either collapsing to form a star or fragmenting." }, "1207/1207.4041_arXiv.txt": { "abstract": "We investigate the dwarf ($M_B>\\bmagcutoff$) galaxies in the Virgo cluster in the radio, optical, and ultraviolet regimes. Of the 365 galaxies in this sample, 80 have been detected in H \\textsc{i} by the Arecibo Legacy Fast ALFA survey. These detections include 12 early-type dwarfs which have H \\textsc{i} and stellar masses similar to the cluster dwarf irregulars and BCDs. In this sample of 12, half have star-formation properties similar to late type dwarfs, while the other half are quiescent like typical early-type dwarfs. We also discuss three possible mechanisms for their evolution: that they are infalling field galaxies that have been or are currently being evolved by the cluster, that they are stripped objects whose gas is recycled, and that the observed H \\textsc{i} has been recently reaccreted. Evolution by the cluster adequately explains the star-forming half of the sample, but the quiescent class of early-type dwarfs is most consistent with having recently reaccreted their gas. ", "introduction": "Dwarf ellipticals (dEs) are rare in the field and in galaxy groups, but are the most common galaxies in clusters (\\citealt{Binggeli1985}; \\citealt{Caldwell1987}; \\citealt{Trentham2009}). An ongoing question is whether cluster dEs form in situ, or whether the cluster environment strips infalling late-type galaxies of their gas, transforming them into dEs (\\citealt{TullyShaya1984}; \\citealt{Boselli2006}; \\citealt{Boselli2008}). Driving such evolution are such processes as ram-pressure stripping \\citep{Gunn1972}, starvation \\citep{Larson1980}, and galaxy harassment \\citep{Lin1983}. There is significant evidence that at least some subset of the dwarf elliptical population were once late-type galaxies. In the most luminous of early-type dwarfs (ETDs), late-type features such as stellar disks and faint spiral arms can be observed \\citep{LiskerDisk}. Furthermore, rotational support or even rotational flattening is observed in a significant fraction of ETDs (\\citealt{vanZee2004}; \\citealt{Beasley2009}; \\citealt{Toloba2011a}), and such galaxies share a Tully-Fisher plane with the dwarf irregulars. Rotationally supported ETDs are also observed to have younger stellar populations and are most likely to be found in the cluster outskirts \\citep{Toloba2009}. The Virgo cluster is a rich ($>$1000 members), relatively nearby ($d\\sim$17 Mpc) cluster. There is no clear single core, but rather two: a compact subcluster around M49, and a more massive and extended subcluster around M87. There are also several near-background `clouds' which are falling onto the cluster. Both its three dimensional structure (\\citealt{Mei2007}) and velocity dispersion (\\citealt{Binggeli1993}; \\citealt{Drinkwater2001}; \\citealt{Conselice2001}) suggest that the Virgo cluster is dynamically young. Due to its proximity, Virgo has been the target of many surveys, and all of its members have been targeted observationally. The most recent comprehensive compilation of Virgo member data can be found in GOLDMine (\\citealt{GOLDMine}; \\citealt{Gavazzi2005}), a multiwavelength aggregate of many data sources. Because it is so close, Virgo is also optimal for sensitive H \\textsc{i} observations of both the dwarf galaxies and more massive members. It is the first cluster in which H \\textsc{i} deficiency, a measure of how much less H \\textsc{i} a galaxy contains compared to a field sample of similar optical properties, has been measured \\citep{Davies1973}. Targeted H \\textsc{i} observations exist for all massive late-type galaxies (\\citealt{Gavazzi2005}), bright dwarf irregulars (\\citealt{Hoffman1987}), and a significant number of dwarf ellipticals (\\citealt{MHRG1986}; \\citealt{Huchtmeier1986}; \\citealt{vanDriel2000}; \\citealt{Conselice2003}). Moreover, blind H \\textsc{i} surveys have covered large fractions of Virgo, such as the H \\textsc{i} Jodrell All Sky Survey (HIJASS; \\citealt{Davies2004}) and the Arecibo Legacy Fast ALFA survey (ALFALFA; \\citealt{ALFALFA1}). Recently, the VLA Imaging of Virgo in Atomic Gas (VIVA; \\citealt{Chung2009}) survey has produced high-resolution H \\textsc{i} maps of late-type galaxies in Virgo; this has allowed extensive study of stripping as it takes place (\\citealt{Chung2007}; \\citealt{Vollmer2008}). Direct H \\textsc{i} evidence of galaxy harassment has been found by ALFALFA, which observed a $\\sim$250 kpc tail off of NGC 4254 \\citep{Haynes2007}. The focus of this article is the dwarf galaxy population in the Virgo cluster, and specifically how dwarf galaxies form and evolve in the cluster environment, primarily through the study of their H \\textsc{i} content. For this purpose, we will make extensive use of ALFALFA\\footnote{The Arecibo Observatory is operated by SRI International under a cooperative agreement with the National Science Foundation (AST-1100968), and in alliance with Ana G. M\u00e9ndez-Universidad Metropolitana, and the Universities Space Research Association.}, a blind H \\textsc{i} survey which, in its current data release covering 40\\% of the final survey area ($\\alpha.40$; \\citealp{ALFA40}), covers the Virgo cluster at declinations between 4\\degrees\\ and 16\\degrees. Examining the H \\textsc{i} content of Virgo has always been a goal of ALFALFA (\\citealt{ALFALFANVirgo}; \\citealt{Kent2008}), with the first global study performed by \\citet{Gavazzi2008}, who investigated the scaling relations between the H \\textsc{i} and luminous properties of ALFALFA-detected galaxies in the cluster. More recently, colors and star-formation properties using $\\alpha.40$ and H$\\alpha$ observations were described by \\citet{Gavazzi2011a} and \\citet{Gavazzi2011b}. If dwarf ellipticals are the products of evolved dwarf irregulars and low luminosity spirals, then the efficient stripping of H \\textsc{i} precedes quenched star formation and morphological transformation. In fact, dwarf irregulars in Virgo have detectable reservoirs of H \\textsc{i}, and the bulk of dwarf ellipticals are relatively gas-free, having low gas fraction (M$_\\text{HI}$/M$_*$). Nonetheless, a small fraction of ALFALFA detections are elliptical galaxies. Early-type H \\textsc{i}-detected galaxies included in previous ALFALFA data releases have been studied both in Virgo \\citep{Alighieri2008} and in the field \\citep{Grossi2009}. They found that very few Virgo ETGs had H \\textsc{i}, and those that did were either massive galaxies which may be accreting from a companion, or dwarfs with peculiar morphologies. In the field, they found that a surprising 44\\% of massive early-type galaxies were detected, possibly the results of major mergers, but only 13\\% of dwarf ellipticals were detected---both much higher than the 2\\% detection rate in Virgo. With the near complete coverage of Virgo made available by $\\alpha.40$, we compile a sample of 12 low-luminosity (M$_\\text{B}>-16$) ETDs which are detected in H \\textsc{i} by ALFALFA; this sample represents an almost two-fold increase from samples based on earlier ALFALFA catalogs. We also probe beneath the detection limit of ALFALFA using spectral stacking methods. In addition to their H \\textsc{i} properties, we investigate their stellar populations and star formation activity using optical and ultraviolet photometry from the seventh data release of the Sloan Digital Sky Survey (SDSS DR7; \\citealt{Abazajian2009}) and the GALEX satellite. An overview of the ALFALFA dwarf population, selected by H \\textsc{i} mass, is presented in \\citet{Huang2011}. An important result is that the within the ALFALFA dwarf sample, Virgo cluster members have lower gas fractions at a given M$_\\text{HI}$ with a wide spread in the distributions of specific star formation rate (SFR/M$_*$). Lastly, we consider possible evolutionary paths that would lead to the existence of H \\textsc{i}-bearing dwarf galaxies with early-type morphologies in a cluster environment. The scenarios considered include the possibility of intrinsically gas-rich objects that have recently infallen onto the Virgo cluster, stripped objects with newly recycled H \\textsc{i} from stars near the end of their lifetimes, and recent accreation of gas by formerly gas-free galaxies. The paper is organized as follows: In section \\ref{sec:datasample} we present the dwarf galaxy sample, and describe the selection process, as well as the H \\textsc{i}, optical and UV data extraction process. Section \\ref{sec:results} describes the H \\textsc{i} content, colors and star formation of the sample and the reference samples, and we discuss differences among dwarf galaxies with different morphologies. In section \\ref{sec:discussion} we present possible evolutionary paths for our sample of H \\textsc{i}-detected ETDs, and argue about their consistency with our data. ", "conclusions": "We examined the sample of all 365 VCC dwarf ($M_B > -16$) galaxies with known redshifts. While not as sensitive as observations targeting individual galaxies, ALFALFA's blind coverage gives us a statistically complete view of the atomic gas content of the Virgo cluster population, both for detected and non-detected galaxies. We use ALFALFA's current $\\alpha.40$ \\citep{ALFA40} data release, which expands the previous Virgo coverage (\\citealt{ALFALFANVirgo}; \\citealt{Kent2008}) to declinations $4^\\circ<\\delta<16^\\circ$. Of the 365 dwarf galaxies, 80 are detected in H \\textsc{i} in the $\\alpha.40$ catalog. As is to be expected, a large fraction of these galaxies (68 out of 80) are late-types and peculiar galaxies, but 12 are classified as early-types by the VCC. The gas-bearing population of dwarf galaxies show a clear tendency to avoid the regions of Virgo closest to M87 and M49. Despite being early-type galaxies, the H \\textsc{i}-detected ETDs have H \\textsc{i} gas fractions similar to the H \\textsc{i}-detected LTDs of the same stellar mass. However, the upper limit on the gas fraction for the non-detected ETD population is approximately 2.5 dex lower than the H \\textsc{i} detected dwarfs, and there does not appear to be a significant population bridging the gap between the gas bearing and the gas poor populations. % In agreement with \\citet{LiskerColors} and \\citet{Kim2010}, we find that the dwarf ellipticals follow a tight color-magnitude relation with bluer $g-r$, NUV$-r$, and FUV$-r$ colors at fainter M$_r$. The H \\textsc{i}-detected late-type dwarfs trace a very broad distribution which is well-separated from the red sequence in $g-r$ (0.4 magnitudes); the separation is larger in NUV$-r$ and FUV$-r$ (2 and 4 magnitudes, respectively). The H \\textsc{i} detected ETDs are unusual in that they do not as a class sit wholly in either the red sequence or blue cloud: half are clearly red while the other half are blue. The blue gas-bearing ETDs have SSFRs and SFEs similar to the LTDs and the `other' dwarfs, while the red ETDs typically have rates a few dex lower than the blue subclass. Lastly, we consider the various possible evolutionary histories of the bas-bearing ETDs. It is quite possible that the blue ETDs are on their first pass through the cluster, and are unstripped or being stripped. VCC 281 stands out from the other blue ETDs in that it has a `blue-center' suggestive of a disk galaxy evolved by galaxy harassment but is also consistent with ram pressure stripping. However, stripping cannot explain the red subclass of ETDs, which are gas-rich but not forming stars. Models of gas returned into the ISM by stars at the end of their lives \\citep{Boselli2008} cannot account for the mass of H \\textsc{i} observed in either the blue or red subclass. The gas-bearing red ETD subset is most consistent with having already traveled through the cluster core, been stripped, and now at lying the edge of the cluster where H \\textsc{i} is being accreted from the cosmic web. All of the red ETDs, with the exception of VCC 956 ($\\sim1^\\circ$ from M87), are in a region of the cluster where the dominant hard X-ray flux arises from background sources, rather than the ICM. The dynamical timescales for these galaxies to re-accrete their gas are significantly less than one orbital time and typically larger than the time it would take the ICM to evaporate the newly accreted gas. \\\\ \\subsubsection*{Acknowledgements} The authors would like to acknowledge the work of the entire ALFALFA collaboration team in observing, flagging, and extracting the catalog of galaxies used in this work. The ALFALFA team at Cornell is supported by NSF grants AST-0607007 and AST -1107390 and by a grant from the Brinson Foundation.\\\\ \\\\ The research leading to these results has received funding from the European Community's Seventh Framework Programme (/FP7/2007-2013/) under grant agreement No 229517.\\\\ \\\\ GALEX is a NASA Small Explorer, launched in 2003 April operated for NASA by the California Institute of Technology under NASA contract NAS5-98034. We gratefully acknowledge NASA's support for construction, operation and science analysis for the GALEX mission, developed in cooperation with the Centre National d'Etudes Spatiales of France and the Korean Ministry of Science and Technology. EP, MPH and RG acknowledge support for this work from the GALEX Guest Investigator program under NASA grant NNX10AI01G.\\\\ \\\\ Funding for the SDSS and SDSS-II has been provided by the Alfred P. Sloan Foundation, the participating institutions, the National Science Foundation, the US Department of Energy, the NASA, the Japanese Monbukagakusho, the Max Planck Society and the Higher Education Funding Council for England. The SDSS Web Site is http://www.sdss.org/. The SDSS is managed by the Astrophysical Research Consortium for the participating institutions. The participating institutions are the American Museum of Natural History, Astrophysical Institute Potsdam, University of Basel, University of Cambridge, Case Western Reserve University, University of Chicago, Drexel University, Fermilab, the Institute for Advanced Study, the Japan Participation Group, Johns Hopkins University, the Joint Institute for Nuclear Astrophysics, the Kavli Institute for Particle Astrophysics and Cosmology, the Korean Scientist Group, the Chinese Academy of Sciences (LAMOST), Los Alamos National Laboratory, the Max Planck Institute for Astronomy, the MPA, New Mexico State University, Ohio State University, University of Pittsburgh, University of Portsmouth, Princeton University, the United States Naval Observatory and the University of Washington." }, "1207/1207.4055_arXiv.txt": { "abstract": "Inside neutron stars, the hadron-quark mixed phase is expected during the first order phase transition from the hadron phase to the quark phase. The geometrical structure of the mixed phase strongly depends on the surface tension at the hadron-quark interface. We evaluate the shear modulus which is one of the specific properties of the hadron-quark mixed phase. As an application, we study shear oscillations due to the hadron-quark mixed phase in neutron stars. We find that the frequencies of shear oscillations depend strongly on the surface tension; with a fixed stellar mass, the fundamental frequencies are almost proportional to the surface tension. Thus, one can estimate the value of surface tension via the observation of stellar oscillations with the help of the information on the stellar mass. ", "introduction": "\\label{sec:I} Neutron stars are formed in supernova explosions, which arise at the final stage of the evolution of massive stars. The density can exceed the standard nuclear density of $\\rho_0\\approx 0.17$ fm$^{-3}$ inside neutron stars. Since such a high density is almost impossible to be realized on the Earth, neutron stars may be a unique ``laboratory\" to investigate the properties of matter around and beyond the nuclear density. One of the possible ways to see the properties of neutron star matter could be the observations of gravitational waves emitted from neutron stars. Since the gravitational waves with high permeability will bring us raw information on the wave sources, one can directly see the properties of neutron star matter. % Thus observations of gravitational waves will provide us with the astrophysical data to reveal basic properties of dense matter (e.g., \\cite{AK1996,AK1998,KS1999,AC2001,Sotani2001,Sotani2003,BFG2005,Erich2011}), and to examine the theory of strong gravity (e.g., \\cite{Sotani2004,Sotani2009a,Sotani2010,BPPL2012}). Now, the worldwide projects are going on to detect gravitational waves associated with the astrophysical phenomena involving compact objects \\cite{Barish2005,LIGO1,LIGO2}. Another way to see the properties of neutron-star matter could be the direct observations of global oscillations of neutron stars. One would know then the stellar mass, radius, and equation of state (EOS) from such observations. This method is often referred to asteroseismology, which is quite similar to the helioseismology. % Unlike gravitational waves, fortunately the observational evidences of neutron-star oscillations have been detected, i.e., the quasi-periodic oscillations (QPOs) in the X-ray afterglow of the giant flares in soft gamma repeaters (SGRs). Up to now, at least three giant flares have been detected in SGR 0526-66, SGR 1900+14, and SGR 1806-20. Furthermore, through the timing analysis of the X-ray afterglow in those giant flares, the specific QPO frequencies have been also extracted in the range from tens Hz up to a few kHz \\cite{WS2006}. Since the central objects in SGRs are considered to be magnetars which are strongly magnetized neutron stars, the discovered QPOs in giant flares could be due to neutron star oscillations. In order to explain these QPO frequencies theoretically, a lot of numerical attempts have been done not only by the shear oscillations in the crust of neutron stars but also by the magnetic oscillations \\cite{Levin2006,Lee2007,SA2007,Sotani2007,Sotani2008a,Sotani2008b,Sotani2009,CBK2009,CSF2009,GCFMS2011,CK2011}. In addition, ascribing the observed QPOs to shear oscillations in the crust of neutron stars, the possibilities are also pointed out to reveal the properties of inhomogeneous nuclear matter in the crust \\cite{SW2009,Sotani2011,GNHL2011,Sotani2012,Sotani2012b}. The structure of a neutron star is considered as follows: ocean of liquid iron exists in the vicinity of stellar surface up to the density $\\sim 10^6-10^8$ g cm$^{-3}$, subsequently the crust region exists from the bottom of iron ocean up to the density of the order of $\\rho_0$. Then a fluid core exists in higher density region. In most part of the crust, nuclei form a bcc lattice due to the Coulomb interaction, while the existence of exotic nuclear shapes at the bottom of crust is also suggested in the recent studies \\cite{Ravenhall83,Hashimoto84,Lorenz1993,Oyamatsu1993,SOT1995, maru2005,new,oka}. According to such studies, with increasing density, the shape of nuclear matter changes from sphere (bcc lattice\\footnote{ Very recently, we and our collaborators have found that fcc lattice of spherical nuclei can be the ground state at some densities, choosing the optimum sizes of the cell and nuclei as well as the inhomogeneous electron distribution \\cite{oka}. }), to cylinder, slab, cylindrical hole, spherical bubble, and uniform matter (inner fluid core), which is collectively called ``nuclear pasta.\" On the other hand, there are still many uncertainties in the core region. For example, hyperons appear in beta-equilibrium nuclear matter when the density becomes higher than $\\sim 2-3\\rho_0$, or non-hadronic quark matter might exist in the innermost stellar core \\cite{rev,rev1}. Depending on the presence of these exotic components, structure of neutron stars dramatically changes \\cite{BBSS2002,Maruyama2007}. As for the hadron-quark (HQ) phase transition, there are further uncertainties such as EOS of quark matter and/or the deconfinement mechanism. Anyhow the HQ mixed phase may emerge as a consequence of the Gibbs conditions, supposing that the HQ phase transition is of the first order \\cite{gle}. The HQ mixed phase looks like the nuclear pasta and strongly depends on the Coulomb interaction and the surface tension at the hadron-quark interface, which are called ``fine-size effects\" \\cite{hei,vos}. Namely, whether non-uniform (pasta) structures appear in the HQ mixed phase is subject to a balance between the surface tension and the Coulomb repulsion \\cite{Maruyama2007,Yasutake2009}, which is similar to the situation in the crust \\cite{maru2005} and kaon condensed matter \\cite{Maruyama2006}. However, to determine the surface tension with the experiments on the Earth is quite difficult because temperature becomes too high for the HQ mixed phase to appear during the relativistic heavy-ion collisions. One of the possibilities to distinguish the finite-size effects on the HQ mixed phase in neutron stars might be the observations of astronomical phenomena. Actually, we suggested such a possibility via direct observations of gravitational waves emitted from neutron stars with the HQ mixed phase \\cite{SYMT2011}. Such attempts are very challenging, but there are very few literatures. On the other hand, shear modulus is one of the specific properties of the HQ mixed phase due to the existence of non-uniform structures. In fact, the shear modulus becomes zero without non-uniform structures. % As mentioned before, however, the properties of such a phase are quite uncertain since the deconfinement mechanism is not clarified. Hence, this article aims mainly at giving the shear modulus and the shear speed in the HQ mixed phase in order to get an insight of the HQ mixed phase. Additionally, as an application, we also explore shear oscillations to see the dependence on the surface tension. We examine the frequencies of shear oscillations within the relativistic Cowling approximation, modeling neutron stars with the HQ mixed phase. This article is organized as follows: In Sec.\\ \\ref{sec:II}, we describe the equilibrium of neutron stars and the adopted EOS, where we will discuss the shear modulus in the HQ mixed phase. In Sec.\\ \\ref{sec:III}, we show equations governing shear oscillations and the boundary conditions to determine the eigenfrequencies. Additionally, the obtained spectra of such oscillations will be shown. At the end, we make a conclusion in Sec.\\ \\ref{sec:IV}. We adopt the unit of $c=G=1$ in this article, where $c$ and $G$ denote the speed of light and the gravitational constant, respectively, and the metric signature is $(-,+,+,+)$. ", "conclusions": "\\label{sec:IV} In this article, we have focused on the hadron-quark (HQ) mixed phase, which can appear inside neutron stars if hadron matter makes a phase transition into quark matter. Depending on the surface tension at the HQ interface, the non-uniform structures can appear in the HQ mixed phase, which produces the shear properties. With the EOS involving the non-uniform structure in the HQ mixed phase, we have estimated the shear modulus in this article. Then, we have found that the shear modulus depends strongly on the surface tension, which becomes $\\sim 10^3$ times larger than that in the crust region. Probably, this is caused by the difference of the charge number including in the droplet. Meanwhile, as an application, we have calculated the shear oscillations in the HQ mixed phase. % As a result, we have found that the frequencies of shear oscillation in the HQ mixed phase could be around ten times larger than that in the crust region and those frequencies depend strongly on the value of the surface tension for the HQ interface. % We also found that the frequencies of fundamental oscillations with the fixed stellar mass are almost proportional to the surface tension. Thus, with the help of the observation about the stellar mass, one might be able to determine the value of surface tension with using the observations of the frequencies of shear oscillations in the HQ mixed phase. Finally, the resulting frequencies of fundamental oscillations in the HQ mixed phase are order of 100 Hz. This means that some of the QPO frequencies observed in giant flares, for example 150 Hz and even 626.5 Hz in SGR 1806-20 or 155 Hz in SGR 1900+14 \\cite{WS2006}, might be associated with the shear oscillations in the HQ mixed phase. % Although it has not been successful to get the collective view about the observed QPOs yet, the consideration of the shear oscillations in the HQ mixed phase may be able to solve the puzzle for the theoretical explanation of the QPO frequencies observed in giant flares. H.S. is grateful to N.~Yasutake for his warm hospitality and fruitful discussions. This work was supported in part by Grants-in-Aid for Scientific Research on Innovative Areas through No.\\ 23105711, No.\\ 24105001, and No.\\ 24105008 from MEXT, by Grant-in-Aid for Young Scientists (B) through No.\\ 24740177 from JSPS, by the Yukawa International Program for Quark-hadron Sciences, and by the Grant-in-Aid for the global COE program ``The Next Generation of Physics, Spun from Universality and Emergence\" from MEXT. \\appendix" }, "1207/1207.3396_arXiv.txt": { "abstract": "We determine the radio and optical luminosity evolutions and the true distribution of the radio loudness parameter $R$, defined as the ratio of the radio to optical luminosity, for a set of more than 5000 quasars combining SDSS optical and FIRST radio data. We apply the method of Efron and Petrosian to access the intrinsic distribution parameters, taking into account the truncations and correlations inherent in the data. We find that the population exhibits strong positive evolution with redshift in both wavebands, with somewhat greater radio evolution than optical. With the luminosity evolutions accounted for, we determine the density evolutions and local radio and optical luminosity functions. The intrinsic distribution of the radio loudness parameter $R$ is found to be quite different than the observed one, and is smooth with no evidence of a bi-modality in radio loudness for $\\log R \\geq -1$. The results we find are in general agreement with the previous analysis of Singal et al. 2011 which used POSS-I optical and FIRST radio data. ", "introduction": "\\label{intro} Quasars are distant active galactic nuclei (AGN) with emission seen across the electromagnetic spectrum, from radio to X-rays. The optical emission from quasars is thought to be dominated by the radiation from the accretion disk around supermassive black holes, while the radio emission is dominated by the plasma outflowing from the black hole/accretion disk systems. Because of this, different but complementary information can be gathered from both photon energy ranges regarding the cosmological evolution of AGN. It is therefore important to determine in detail the redshift evolutions of quasars in both radio and optical regimes. In a previous paper \\citep[hereafter QP1]{QP1} we explored the luminosity evolutions and radio loudness distribution of quasars with a dataset consisting of the overlap of the FIRST (Faint Images of the Radio Sky at Twenty Centimeters) radio survey with the Automatic Plate Measuring Facility catalog of the Palomar Observatory Sky Survey (POSS-I), as presented by \\citet{White00}. In this paper we present the results from a much larger dataset consisting of the overlap of FIRST with the Sloan Digital Sky Survey (SDSS) Data Release 7 \\citep{SDSSDR7} quasar catalog, which contains a factor of $\\sim 10$ more objects and spans redshifts from 0.065 to 5.46. In the literature, the evolution of the quasar luminosity function (LF) has been described not only for optical and radio luminosities but also for X-ray, infrared, and bolometric luminosities \\citep[e.g.][]{Ueda03,R06,Matute06,Hopkins07,Croom09}. The shape of the LF and its evolution are usually obtained from a flux limited sample in the $a$ waveband $f_a > f_{m,a}$ with $f_{m,a}$ denoting the flux limit and $L_a = 4 \\, \\pi \\, d_L^2 f_a / K_a$, where $d_L(z)$ is the luminosity distance and $K_a(z)$ stands for the K-correction. For a pure power law emission spectrum of index $\\varepsilon_a$ defined as $f_a \\propto \\nu^{-\\varepsilon_a}$, one has $K_a(z) = (1+z)^{1-\\varepsilon_a}$. This simple form may be modified for optical data by the presence of emission lines, as in this work. In general, the determination of the LF and its evolution requires analysis of the bi-variate distribution $\\Psi_a(L_a,z)$. The first step of the process should be the determination of whether the variables of the distributions are correlated or are statistically independent. A correlation between $L_a$ and $z$ is a consequence of luminosity evolution. In the case of quasars with the optical and some other band luminosity, we have at least a tri-variate function. One must determine not only the correlations between the redshift and individual luminosities (i.e. the two luminosity evolutions) but also the possible intrinsic correlation between the two luminosities, before individual distributions can be determined. Knowledge of these correlations and distributions are essential for not only constraining robustly the cosmological evolution of active galaxies, but also for interpretation of related observations, such as the extragalactic background radiation \\citep[e.g.][]{Singal10,Hopkins10}. A related question is the distribution of the `radio-loudness parameter', $R=L_{\\rm rad}/L_{\\rm opt}$, for the quasar population, defined as the ratio of the 5 GHz radio to 2500 \\AA \\, optical luminosity spectral densities, and the distinction between so-called `radio loud' ($R>10$; RL for short) and `radio quiet' ($R<10$; RQ for short) quasars. Weak hints of a bi-modality in the distribution of the radio loundess parameter described by \\citet{Kellerman89} suggested that $\\log R = 1$ could be chosen as the radio loud/quiet demarcation value. Using this value for the division between RL and RQ quasars, the differences between the two classes have been investigated, including the possibility of distinct cosmological evolution of the RL and RL populations \\citep[e.g.][]{Miller90,Goldschmidt99,Jiang07}. Still, the more recent analyses of different samples of objects reported in the literature so far gave rather inconclusive results on whether any bi-modality in the distribution of the radio loudness parameter for quasars is inherent in the population \\citep[see][]{Ivezic02,Cira03,Ivezic04,Zamfir08,Kimball11,Mahony12}. In QP1 we found no evidence for a bi-modality in $R$ in the range $-1<\\log R<4$. In addition to the above cited works, there have been many papers dealing with this ratio and RL vs RQ issue, as well as luminosity ratios at other wavelengths, e.g. IR/radio, Optical/X-ray etc. However, to the best of our knowledge, except for QP1, none of these works have address the correlations between the radio and optical luminosities, which is necessary for such and analysis. Additionally, in general they have not concentrated on obtaining the {\\it intrinsic} --- as opposed to the raw observed --- distribution (and/or evolution) of the ratio, which is related to the tri-variate LF $\\Psi (L_{\\rm opt}, L_{\\rm rad}, z)$ by\\footnote{Equation \\ref{Gtrue} arises because by definition $\\int {G_R(R,z) \\, dR} = \\int \\int {\\Psi(L_{\\rm opt}, \\, L_{\\rm rad} \\, z) \\, dL_{\\rm opt} \\, dL_{\\rm rad}}$, and, following from the definition of $R$, $dL_{\\rm rad} = L_{\\rm opt} \\, dR$ and $dL_{\\rm opt} = -(L_{\\rm rad} / R^2) \\, dR$.} \\begin{eqnarray} G_R(R,z) = \\int_0^{\\infty} { \\Psi (L_{\\rm opt}, R \\,\\, L_{\\rm opt}, z) \\, L_{\\rm opt} \\, dL_{\\rm opt} } \\nonumber \\\\ = \\int_0^{\\infty} { \\Psi \\left( {L_{\\rm rad} \\over R}, L_{\\rm rad}, z \\right) \\, L_{\\rm rad} \\, {{ dL_{\\rm rad} } \\over {R^2}} } . \\label{Gtrue} \\end{eqnarray} In Appendix A of QP1 we showed how, even in the simplest cases, the observed distributions can be very different from the intrinsic ones. Thus, for determination of the true distributions the data truncations must be determined and the correlations between all variables must be properly evaluated. \\citet[EP for short]{EP92,EP99} developed new methods for determination of the existence of correlation or independence of variables from a flux limited and more generally truncated dataset, which were further expanded in QP1. Our aim in this paper is to take all the selection and correlation effects into account in determination of the true evolution of optical and radio luminosities and their ratio and to find their distributions, using the larger SDSS DR7 QSO $\\times$ FIRST dataset. In \\S \\ref{datasec} we describe the radio and optical data used. \\S \\ref{simlumf} contains a general discussion of luminosity evolution and the sequence of the analysis. In \\S \\ref{evsec} we apply the EP method to achieve the luminosity-redshift evolutions and the correlation between the luminosities. We determine the density evolution in \\S \\ref{dev}, and the local luminosity functions in \\S \\ref{local}. A discussion of the radio loudness distribution is presented in \\S \\ref{Rdist}. \\S \\ref{testass} we investigate some of the assumptions used, and \\S \\ref{disc} contains a discussion of the results. This work assumes the standard $\\Lambda$CDM cosmology throughout, with $H_0=71$\\,km\\,s$^{-1}$\\,Mpc$^{-1}$, $\\Omega_{\\Lambda}=0.7$ and $\\Omega_{m}=0.3$. ", "conclusions": "\\label{disc} We have used a general and robust method to determine the radio and optical luminosity evolutions simultaneously for the quasars in the SDSS DR7 QSO $\\times$ FIRST dataset, which combines 1.4 GHz radio and $i$-band optical data for thousands of quasars ranging in redshifts from 0.64 to 4.82 and over five orders of magnitude in radio loudness. As can be seen, the results for the SDSS DR7 QSO $\\times$ FIRST dataset employed here are similar in bulk to the results for the White et al. dataset that we presented in QP1. In this work we show that, importantly, the major conclusions are not sensitive to whether the correlation between the radio and optical luminosities is considered to be induced by similar redshift evolution or intrinsic. In the previous work, we assumed the correlation to be intrinsic. Additionally, the luminosity evolution found for the optical only dataset matches that found for the combined radio-optical dataset, providing a test that the technique has properly handled the truncations and correlations inherent in the data. We have seen that the results of this analysis are nearly identical for the two different radio-optical matching radius criteria considered, 1) that adopted by J07 with a $5''$ radius for single matches and a $30''$ radius for multiple matches, and 2) a universal $5''$ matching radius which we feel is more appropriate. The reason for this preference is that $30''$ radius corresponds to a physical scale $> 100$\\,kpc \\emph{projected}, at any redshift $z>0.2$. Extended lobes in radio galaxies are known to reach similar or even larger sizes. These are however objects viewed at large inclinations, unlike quasars observed, by definition, at much smaller viewing angles \\citep{Barthel89}. With such small viewing angles, projection effects are expected to limit the projected sizes of the extended structures in radio-loud quasars. In addition, evolutionary effects may play a role as well, with distant quasars being typically younger, and therefore more compact in radio, than nearby evolved radio galaxies. This is supported by a detailed investigation of radio morphologies of SDSS (DR3) $\\times$ FIRST quasars by \\citet{dev06}, who noted that overwhelming majority of the selected sources are characterized by compact radio morphologies. For all these reasons, we believe that applying a universal $5''$ matching radius criterion in assigning FIRST counterparts to the SDSS bona fide \\emph{quasar} sources is more appropriate, while $30''$ criterion for multiple matches has a high probability of including radio flux from objects unrelated to the optical quasar that they are supposedly linked to. \\subsection{Differential evolutions toward radio loudness with increasing redshift}\\label{rldisc} Quasars are seen to exhibit positive luminosity evolution in both wavebands, with stronger positive evolution in radio than optical. The differential evolution can be seen in the evolution of $G_{R'}\\!(R')$ as shown in Figure \\ref{psir}. This is in agreement with the results presented for a combined FIRSTxPOSS-I dataset in QP1, but in likely disagreement with J07 which claimed that the fraction of RL quasars decreases with increasing redshift. However, \\citet{Miller90} have noted that the fraction of RL quasars may increase with redshift and \\citet{Donoso09} reached the same conclusion by computing radio and optical LFs at different redshifts. \\citet{Cira06} also found that the RL fraction of quasars may modestly increase at high redshift, and the conclusions of \\citet{Balokovic12} favor $R$ increasing with redshift. We note that a trend toward increasing radio loudness with redshift is visible to the eye even in the raw observed data as in Figure \\ref{LvsL}, although this data suffers from biases so that the true evolution of $R$ can only be recovered with an analysis method such as employed here. We attribute our apparent disagreement with the J07 result to three causes: 1) They present one particular moment of the $R$ distribution --- the``radio loud fraction'' --- in bins, while we are presenting the distribution itself, so direct comparison is rather difficult, 2) they are calculating this fraction by including optical sources that have no radio detection, which is a potentially correctable bias, and 3) they are not including the effect of neglecting sources that are radio bright enough to appear in FIRST but don't appear in the SDSS quasar sample that they use. There are potentially many AGN sources in the FIRST catalog that are not present in the SDSS quasar sample, due to their low optical fluxes. The sample used in J07 to calculate the radio loud fraction in bins would be missing this population of sources with detectable radio emission but lower optical fluxes than would allow an SDSS detection or classification as a quasar source, while we have used an analysis which takes both flux limits properly into account. While for the joint dataset the present work has used only quasars detected in both radio and optical and including both optical-color-selected and radio-match-selected sources, J07 used as their sample quasars detected in optical regardless of a radio detection and their quasars were selected based solely on optical color. Stepping back, there are four basic options when constructing a combined optical-radio dataset of quasars for an analysis of the radio loudness distribution and its evolution, given that many more confirmed quasars have an optical detection than a radio one: I. quasars detected in both radio and optical II. quasars detected in optical regardless of a radio detection A. optical quasar candidates selected for followup based on either optical colors or presence of a radio match B. optical quasar candidates selected for followup based only on optical colors Under the above rubric, we have used IA for the joint radio-optical dataset, while J07 used IIB. We believe that while using option II is appropriate to investigate the optical luminosity function and its evolution, such a dataset is not appropriate for investigating the radio loudness distribution and its evolution. This is because with such a dataset one will always be introducing a bias by not including objects bright in radio but dim in optical, and at the same time, not being able to say if such a population exist at all by means of some censored data method. The situation could be different if one were able to select two quasar samples based separately on the radio and optical data, down to given radio and optical flux limits of the surveys, and then to match the two samples forming a master-list with both radio and optical fluxes provided for each object included. Only then non-detections in either radio or optical bands could be considered, and the censored data method could be applied \\citep[see in this context, e.g.,][]{Feigelson83,Feigelson85,Isobe86}. However since an optical identification is needed to claim a quasar nature of a source the first place, this is not possible, so in using an only optically detected set one will be always biased against a radio-loud population. On the other hand, if, as in the present work, one uses option I, and then analyzes the data in such a way as to take the flux limits into account, an unbiased reconstruction of the intrinsic properties of the population can be achieved. The question of whether option A or B is most appropriate is less straightforward. However, quasar optical color may be correlated with radio loudness \\citep{Kimball11}. If this is the case, then by selecting quasars based on the optical colors solely, one may introduce some bias when dealing with the radio loudness distribution and evolution. Option IA means that the color criterion is not important: by including only those quasars which are detected in radio, and by allowing for radio-selected optical quasar candidates, we basically end up with a radio-selected (and not color-selected) and radio and optical flux limited sample of quasars. \\subsection{Lack of bi-modality in radio loudness}\\label{lbmdisc} Our results favor a single population of quasars with no bi-modality in the radio loudness parameter for the range of $R$ considered here, in agreement with the conclusions of QP1. Some physical implications of a single population in this range of $R$ are discussed in that work, including that there is thus no evidence for the existence of two physically different population of quasars with respect to the radio/jet production efficiency. Although it is a somewhat different issue, there have also been studies of radio populations investigating the distribution of quasars, Seyferts, and other types of AGN in the plane of radio loudess versus accretion rate (note that by definition all the quasars accrete at high and very high rates, unlike radio galaxies or Seyferts). These results indicate that there may be two `tracks' in the space, characterized by a similar non-monotonic dependence of the radio loundess parameter on the Eddington ratio, but differing in normalization, and corresponding to what might be considered as a RL population hosted exclusively by elliptical galaxies, and a RQ population of AGN hosted by both types of galaxies \\citep{Sikora07}. Such studies were however done --- and in fact had to be done --- on incomplete and inhomogeneous samples of AGN, and were claimed to depend on whether core or total radio luminosities are used \\citep[e.g][]{BF11}. The host-related bimodal distribution of radio loudness in the entire AGN population is however distinct from the main problem addressed in this work, dealing strictly with a well-defined population of quasar-type AGN (for which, notably, we use the FIRST --- i.e. the total --- radio fluxes). And in particular, the conclusion that there is no radio bimodiality for the quasar population does not contradict the findings presented in \\citet{Sikora07}, as these authors emphasized that while a complete sample of elliptical-hosted AGN (including quasars) is expected to show continuous distribution of the radio loudness down to the ``radio quiet'' regime, the point is that the spiral-hosted AGN do never reach high values of the radio-loudness parameters (i.e., values comparable to those of radio galaxies), even though at very low accretion rates they are often characterized by $\\log R >1$ (see section 3 of that work and the discussion therein). In \\S \\ref{Rdist} we explored additional arguments against the presence of a bi-modality in $R$ at lower values of $R$, below those considered in this analysis. The caution to the above conclusion is that the sample analyzed here is dominated by RL objects, as we do not, for the reasons discussed above, consider SDSS quasars with no detected FIRST counterparts. Figure \\ref{LvsL} indicates that with this SDSS $\\times$ FIRST dataset we are sensitive to regions where additional RQ objects would lie were they to exist (e.g. below the $R=10$ line and above the $L_{rad,min}$ lines in that figure) but there does not appear to be an overabundance of those objects. However, we cannot formally exclude the possibility that there may be a population of very radio quiet quasars which would form a separate branch in Figure \\ref{psir} were the abscissa extended to the left, resulting in a bi-modal distribution of radio loudness for the entire quasar population. Such a numerous population of very radio-quiet quasars, even if absent in the FIRST database, would not be missed in deeper radio surveys probing sub-mJy flux levels. In fact, \\citet{Kimball11} have used EVLA observations to detect 179 SDSS DR7Q quasars in the range $0.220\\sigma$ detection significance and time-averaged energy flux $F_E > 3 \\times 10^{-11}{\\rm erg\\,cm^2\\,s^{-1}}$ -- there are 249 such sources. All but six presently have lower energy identifications: blazars, pulsars and a few binaries. The handful of sources {\\it not} yet associated with one of these source classes provide the best prospect for new types of $\\gamma$-ray emitter. We have initiated a campaign to characterize these unidentified sources. The first result from this effort is the discovery of dramatic optical and X-ray variability for the high latitude ($|b|=62^\\circ$) source J2339$-$0533, implying that it is a short period `black-widow'-type binary millisecond pulsar (MSP) \\citep{rs11,ket12}. Here we report progress on the next high latitude unidentified source in this set, 2FGL J1311.7$-$3429 (hereafter J1311) at $|b|=28^\\circ$. With a $43\\sigma$ detection significance (the highest among the 2FGL unidentified sources) and an energy flux of $F_{0.1-100GeV} = 6.2 \\times 10^{-11}{\\rm erg\\,cm^{-2}\\,s^{-1}}$ (the second brightest UnIDed 2FGL), this is a top candidate for follow-up. The `Variability index' value 19.5 and `Curvature Significance'=6.3 mark this as a steady source with a substantial spectral cut-off: a prime pulsar candidate. It has been searched for $< 30$\\,Hz $\\gamma$-ray pulsations \\citep[eg.][]{blind} and for radio pulses down to $\\sim$\\,ms periods \\citep{ray12}, with no detection. Thus it is unlikely to be an isolated young pulsar or a radio-loud MSP. We describe here an optical campaign to identify and characterize a counterpart. ", "conclusions": "" }, "1207/1207.3219_arXiv.txt": { "abstract": "\\noindent Canonization of $F(R)$ theory of gravity to explore Noether symmetry is performed treating $R - 6(\\frac{\\ddot a}{a} + \\frac{\\dot a^2}{a^2} + \\frac{k}{a^2}) = 0$ as a constraint of the theory in Robertson-Walker space-time, which implies that $R$ is taken as an auxiliary variable. Although it yields correct field equations, Noether symmetry does not allow linear term in the action, and as such does not produce a viable cosmological model. Here, we show that this technique of exploring Noether symmetry does not allow even a non-linear form of $F(R)$, if the configuration space is enlarged by including a scalar field in addition, or taking anisotropic models into account. Surprisingly enough, it does not reproduce the symmetry that already exists in the literature \\cite{18} for scalar tensor theory of gravity in the presence of $R^2$ term. Thus, $R$ can not be treated as an auxiliary variable and hence Noether symmetry of arbitrary form of $F(R)$ theory of gravity remains obscure. However, there exists in general, a conserved current for $F(R)$ theory of gravity in the presence of a non-minimally coupled scalar-tensor theory \\cite{27,28}. Here, we briefly expatiate the non-Noether conserved current and cite an example to reveal its importance in finding cosmological solution for such an action, taking $F(R) \\propto R^{\\frac{3}{2}}$. ", "introduction": "Increasing interest in $F(R)$ theory of gravity (see \\cite{1} for a recent review) initiates to probe deeply into it and to explore its merits-demerits from different angles. Cosmological consequences of $F(R)$ theory of gravity have been explored \\cite{2} taking into account different sets of field equations which appear through the metric variation and Palatini variation formalisms. Although, both set of field equations lead to late time cosmic acceleration without the requirement of dark energy in the form of scalar fields, different results have emerged following these two different techniques. For example, unification of early inflation and late time acceleration \\cite{3} on one hand, and violent instability \\cite{7}, cosmologically non-viability \\cite{8}, big bang nucleosynthesis and fifth-force constraints altogether \\cite{9}, on the other, are the outcome of metric formalism. On the contrary, no instability appears as such in Palatini formalism \\cite{10} and it leads to correct Newtonian limit \\cite{11}, while curvature corrections are found to induce effective pressure gradients which create problem in the formation of large-scale structure \\cite{12}. Further, even density perturbation in the two techniques, produce different results \\cite{13}. Nevertheless, it has been shown that under conformal transformation the two formalisms are the same dynamically \\cite{14} and also are equivalent for a large class of theories \\cite{15}. Despite such contrasting results on one hand and equivalence of the two formalisms on the other, at the end it is required to choose some particular form of $F(R)$ out of indefinitely many. This is only possible by imposing Noether symmetry. In the present work we discuss this issue in the metric variation formalism.\\\\ \\noindent Noether symmetry when applied for the first time in scalar-tensor theory of gravity \\cite{16}, the idea was to find a form of the potential that might give rise to a cyclic co-ordinate and hence a conserved current. That is, out of indefinitely many choice of the potential, Noether symmetry selects one, corresponding to which there exists a cyclic coordinate that can't be found a-priori through inspection. In the Robertson-Walker background metric the exponential form of the potential thus found was rather encouraging, because it could trigger inflation in the early Universe. Since then, it has been applied in nonminimally coupled scalar-tensor theory of gravity both in the background of isotropic and several anisotropic models \\cite{17}, together with nonminimally coupled scalar-tensor theory in the presence of $R^2$ term in the action \\cite{18} and recently, to classify different dark energy models \\cite{19}.\\\\ \\noindent To explore Noether symmetry, it is first required to express the action in canonical form, which is possible by introducing auxiliary variable in higher order theory of gravity. This has been attempted by several authors \\cite{21} for $F(R)$ theory of gravity, treating $R - 6(\\frac{\\ddot a}{a} + \\frac{\\dot a^2}{a^2} + \\frac{k}{a^2}) = 0$ as a constraint in the Robertson-Walker minisuperspace model. In the process, the Lagrangian is spanned by a set of configuration space variables $(a, \\dot a, R, \\dot R)$, i.e., $R$ is treated as an auxiliary variable and the result is $F(R) \\propto R^{\\frac{3}{2}}$ in the vacuum and matter dominated era. Such a form of $F(R)$ produces $a \\propto t^{\\frac{3}{4}}$ instead of the standard $\\sqrt{t}$ in the radiation dominated era \\cite{22}. This creates problem in explaining BBN, since Universe expands at a much faster rate than the standard Friedmann model. Further, it has been shown that \\cite{22} under appropriate choice of variable $h_{ij} = a^2 = z$, required for canonical quantization of higher order theory of gravity in the Robertson-Walker metric \\cite{22a}, $z$ becomes cyclic for $F(R) = R^{\\frac{3}{2}}$. It means $R^{\\frac{3}{2}}$ has inbuilt Noether symmetry for $F(R)$ theory in vacuum or in matter dominated era. Now, scalar field appears naturally along with higher order curvature invariant terms in the weak energy limit of string theories \\cite{23}, $4$-dimensional Brane world effective action \\cite{24}, supergravity theories \\cite{25} and in the Chern-Simons gravitational theories \\cite{26}. Therefore, adding more degrees of freedom in the form of a scalar field, $z$ is no longer cyclic for $F(R) = R^{\\frac{3}{2}}$ and situation might improve, yielding a better form of $F(R)$ to explain presently available cosmological data. Additionally, the symmetry already obtained for $F(R) = f(\\phi)R + \\beta R^2$, in the presence of a scalar field \\cite{18} must be reproduced, if canonization treating $R$ as an auxiliary variable is correct. \\\\ \\noindent Higher order theory of gravity may be expressed in canonical form choosing other auxiliary variables too, different from $R$. All lead to correct classical field equations but different quantum dynamics. Horowitz \\cite{h} suggested to choose auxiliary variable as the derivative of the action with respect to the highest derivative present in the action, i.e., as $Q = \\frac{\\partial A}{\\partial \\ddot a}$. This suggestion is flawed, since in that case auxiliary variable may be introduced even in linear theory as shown by Pollock \\cite{p} leading to completely wrong quantum dynamics. Later, in a series of works \\cite{22a} it was suggested that auxiliary variable should be introduced following the prescription of Horowitz \\cite{h} only after removing appropriate total derivative terms from the action. Following this technique, we found Noether symmetry for $F(R) = f(\\phi) R + \\beta R^2$, $R$ being non-minimally coupled to a scalar field \\cite{18}. Thus our intention is to see if same result \\cite{18} is reproduced taking $R$ as an auxiliary variable, and if Noether symmetry at all exists for other form of action.\\\\ \\noindent Thus in the following section 2, we take up a theory of gravity consisting of scalar field being nonminimally coupled with $F(R)$, follow the usual technique of finding Noether symmetry \\cite{21} to end up with the result that Noether symmetry for $F_{,RR} \\ne 0$ is obscure. In section 3, following the same technique \\cite{21}, for a nonminimally coupled scalar-tensor theory of gravity along with $F(R)$ term we obtain the same result that $F_{,RR} \\ne 0$ is impossible. This is surprising, since, as already mentioned, we have explored Noether symmetry for the same action, with $F(R) = f(\\phi)R + \\beta R^2$ and the Brans-Dicke coupling parameter $\\omega(\\phi) = 1$, in the conformally flat Robertson-Walker metric \\cite{18}. In doing so, we have used an auxiliary variable $Q$, different from $R$, and the Lagrangian was spanned by the set of configuration space variables ($a, Q, \\phi, \\dot a, \\dot Q, \\dot\\phi$), instead. Thus, the absence of Noether symmetry in both the cases automatically raises doubt about the technique of treating $R - 6(\\frac{\\ddot a}{a} + \\frac{\\dot a^2}{a^2} + \\frac{k}{a^2}) = 0$, as a constraint of the theory. To check if indeed the technique is flawed, in section 4, we take up a set of anisotropic models, which also increases the configuration space variables ($a, b, R, \\dot a, \\dot b, \\dot R$), to explore Noether symmetry of $F(R)$ theory in vacuum. Here again we find that $F(R)$ does not admit any nonlinear form. Thus we conclude that symmetry in isotropic space-time obtained for $F(R) = R^{\\frac{3}{2}}$ in vacuum and matter dominated era, was an accident, since it makes $z = a^2$ cyclic, but the technique \\cite{21} is flawed, and it is not possible to explore Noether symmetry of $F(R)$ theory of gravity, in general.\\\\ \\noindent The next question that automatically arises, ``is it possible to find an integral of motion in general for $F(R)$ theory of gravity being coupled with a scalar field, including a linear term in addition?\". The answer is yes and already appears in the literature \\cite{27}, where a general conserved current in connection with a general higher order theory of gravity coupled with a dilatonic scalar has been explored in a metric independent way. In section 5, we briefly discuss the technique of finding the same and as an example, show how to use such conserved current to find solutions in the context of cosmology. We have given a simple solution of the scale factor directly for the action containing $F(R) = R^{\\frac{3}{2}}$ being added to a linear term ($f(\\phi)R$) along with a scalar field $\\phi$, in the Robertson-Walker metric, both in radiation and matter dominated era. It produces Friedmann type solution in the radiation era and a coasting solution in the matter era, which fits SnIa data perfectly although early deceleration remains absent creating problem in structure formation. Finally, we end up with conclusions.\\\\ ", "conclusions": "Both the metric and Palatini variation methods have vices and virtues and it is a matter of taste to adopt one. However at the end one requires a form of $F(R)$, any nonlinear form of which makes field equations extremely difficult to solve. Noether symmetry not only fixes a form out of indefinitely many, but also yields an integral of motion which makes field equations somewhat tractable. For such purpose one requires a canonical point Lagrangian. The only technique known so far to express $F(R)$ in such a form is to introduce $R - 6(\\frac{\\ddot a}{a} + \\frac{\\dot a^2}{a^2} + \\frac{k}{a^2}) = 0$ as a constraint in Robertson-Walker metric, which gives correct field equations, in view of the auxiliary variable $R$. Following this technique, Noether symmetry yields $F(R) = R^{\\frac{3}{2}}$ both in vacuum and matter dominated era. Although such a form is reasonably good in the early Universe, suffers from the disease that it does not produce standard Friedmann radiation era \\cite{22}. In anticipation that a linear term in addition could have cured the disease, here we have enlarged the configuration space by adding a scalar field. Rigorous and detailed calculation shows that Noether symmetry remains obscure for $F_{,RR} \\ne 0$. We have also attempted to find Noether symmetry for pure $F(R)$ gravity in the background of some anisotropic models and ended up with the same result. Nature might not allow additional symmetry, but the problem is that Noether symmetry for a scalar-tensor theory of gravity being coupled with $F(R) = R^2$, already exists in the literature \\cite{18}, while the present technique fails to reproduce the same. This raises doubt about the technique \\cite{21} under consideration.\\\\ \\noindent Obviously, one should ask ``how come the technique then yields $F(R) = R^{\\frac{3}{2}}$ in vacuum and matter dominated era?\". The answer has been given in a recent work \\cite{22}, where it has been shown that under appropriate choice of co-ordinate $h_{ij} = a^2 =z$, an action containing $R^{\\frac{3}{2}}$ in vacuum and in the matter dominated era, makes the variable $z$ cyclic. Thus, the symmetry obtained in Robertson-Walker space-time yielding $F(R) = R^{\\frac{3}{2}}$, is due to the very speciality of the isotropic and homogeneous space-time and so is a mere accident. Since there is no way to find a canonical point Lagrangian for $F(R)$ theory of gravity, other than the one under consideration, so we conclude that such a symmetry is obscure for a general $F(R)$ gravity.\\\\ \\noindent It therefore appears that there is practically no way to handle $F(R)$ theory of gravity. However, the situation is not as bad, since a scalar-tensor theory of gravity in the presence of $F(R)$, indeed carries a general conserved current \\cite{27,28}, which is independent of the metric and also on the forms of parameters involved, viz., $F(R)$, $f(\\phi)$, $\\omega(\\phi)$ and $V(\\phi)$. Here in section (5), we have briefly presented the conserved current and cited an example taking $F(R) = R^{\\frac{3}{2}}$, to understand its use in finding solutions for higher order curvature invariant terms. The presence of a linear term in the action modifies the solution in the radiation era to the standard Friedmann type ($a \\propto \\sqrt t$). Thus BBN remains unaltered, which is of-course a brilliant result. On the other hand matter dominated era admits only a coasting solution in the form $a \\propto t$, which although fits SNIa data perfectly, does not allow a transition from early deceleration and so $R^{\\frac{3}{2}}$ is not a good choice to explain late time cosmological observations. \\\\" }, "1207/1207.7326_arXiv.txt": { "abstract": "We describe the automated spectral classification, redshift determination, and parameter measurement pipeline in use for the Baryon Oscillation Spectroscopic Survey (BOSS) of the Sloan Digital Sky Survey III (SDSS-III) as of the survey's Ninth Data Release (DR9), encompassing 831,000 moderate-resolution optical spectra. We give a review of the algorithms employed, and describe the changes to the pipeline that have been implemented for BOSS relative to previous SDSS-I/II versions, including new sets of stellar, galaxy, and quasar redshift templates. For the color-selected ``CMASS'' sample of massive galaxies at redshift $0.4 \\la z \\la 0.8$ targeted by BOSS for the purposes of large-scale cosmological measurements, the pipeline achieves an automated classification success rate of 98.7\\% and confirms 95.4\\% of unique CMASS targets as galaxies (with the balance being mostly M stars). Based on visual inspections of a subset of BOSS galaxies, we find that approximately 0.2\\% of confidently reported CMASS sample classifications and redshifts are incorrect, and about 0.4\\% of all CMASS spectra are objects unclassified by the current algorithm which are potentially recoverable. The BOSS pipeline confirms that $\\sim$51.5\\% of the quasar targets have quasar spectra, with the balance mainly consisting of stars and low signal-to-noise spectra. Statistical (as opposed to systematic) redshift errors propagated from photon noise are typically a few tens of km\\,s$^{-1}$ for both galaxies and quasars, with a significant tail to a few hundreds of km\\,s$^{-1}$ for quasars. We test the accuracy of these statistical redshift error estimates using repeat observations, finding them underestimated by a factor of 1.19 to 1.34 for galaxies, and by a factor of 2 for quasars. We assess the impact of sky-subtraction quality, signal-to-noise ratio, and other factors on galaxy redshift success. Finally, we document known issues with the BOSS DR9 spectroscopic data set, and describe directions of ongoing development. ", "introduction": "The Sloan Digital Sky Survey III (SDSS-III, \\citealt{Eisenstein11}) is the third phase of the SDSS \\citep{York00}.\\footnote{Throughout this paper, we will refer to the earlier SDSS phases collectively as SDSS-I/II\\@.} Within the SDSS-III, the Baryon Oscillation Spectroscopic Survey (BOSS, \\citealt{Dawson12}) is currently mapping a larger volume of the universe than any previous spectroscopic survey. The Ninth Data Release of the SDSS-III (DR9, \\citealt{Nine12}, released publicly on 2012 July 31) is the first SDSS-III data release to include BOSS spectroscopic data, and comprises good observations of 831 unique plate-pluggings of 813 unique tilings (plates worth of targets) on the sky. Each plate delivers simultaneous spectroscopic observations of 1000 lines of sight with optical fibers that feed a pair of two-arm spectrographs, giving a total of 831,000 BOSS DR9 spectra. The main science goal of BOSS is to trace the large-scale mass structure of the universe using massive galaxies and quasar Ly$\\alpha$ absorption systems, in order to measure the length scale of the ``baryon acoustic oscillation'' feature in the spatial correlation function of these objects \\citep[e.g.,][]{Eisenstein05}, and thereby to constrain the nature of the dark energy that drives the accelerated expansion of the present-day universe. To meet this goal, the BOSS project has specified a series of scientific requirements, including: (1) an RMS galaxy redshift precision better than 300 km\\,s$^{-1}$; (2) a galaxy redshift success rate of at least 94\\%, including both targeting inefficiency and spectroscopic redshift failure; (3) a catastrophic galaxy redshift error rate of less than 1\\%; and (4) spectroscopic confirmation of at least 15 quasars at $2.2 < z < 3.5$ per degree$^2$ from among no more than 40 targets per degree$^2$. To satisfy these requirements within such a large survey, automated spectroscopic calibration, extraction, classification, and redshift measurement methods are essential. This paper, one of a series of technical papers describing SDSS-III DR9 in general and the BOSS data set in particular, presents the automated classification and redshift measurement software for the main galaxy and quasar target samples implemented for the BOSS project. This software is written in the IDL language, and is titled \\texttt{idlspec2d}. Earlier versions of this code were used to analyze SDSS-I/II data (see \\citealt{Aihara11}), alongside the complementary and independently developed pipeline software \\texttt{spectro1d} (see \\citealt{Subbarao02} and \\citealt{Adelman06}); for the BOSS project, the \\texttt{idlspec2d} software has been adopted as the primary code, due to its robust error estimation methods and its tight integration of redshift measurement and classification with the lower-level operations of raw data calibration and extraction. The code has also been upgraded with new redshift-measurement templates and several new algorithms in order to meet the scientific requirements of the BOSS project. The tagged software version \\texttt{v5\\_4\\_45} was used to process all BOSS spectroscopic data for DR9\\footnote{The DR9 tagged version of \\texttt{idlspec2d} can be obtained at \\texttt{www.sdss3.org/svn/repo/idlspec2d/tags/v5\\_4\\_45/}.}, and the classification and redshift results delivered by this code have been used for recently published BOSS DR9-sample cosmological analyses \\citep{Anderson12,Manera12,Nuza12,Reid12,RossA12,Sanchez12,Tojeiro12}. An overview of the BOSS project, including experimental design, scientific goals, observational operations, and ancillary programs, is given in \\citet{Dawson12}. A description of the \\texttt{idlspec2d} calibration and extraction methods which transform raw CCD pixel data into one-dimensional object spectra will be presented in \\citet{Schlegel12}. The organization of this paper is as follows. Section~\\ref{sec:data} presents an overview of the spectroscopic data sample of BOSS DR9. Section~\\ref{sec:pipeline} describes the classification and redshift pipeline algorithms and procedures, including the core redshifting algorithm (\\S\\ref{subsec:zmeasure}), special classification handling for the galaxy target samples (\\S\\ref{subsec:z_noqso}), measured spectroscopic parameters (\\S\\ref{subsec:params}), and output files (\\S\\ref{subsec:outfiles}). Section~\\ref{sec:templates} describes the templates constructed for the automated spectroscopic identification and redshift analysis of BOSS galaxies (\\S\\ref{subsec:galtemp}), quasars (\\S\\ref{subsec:qsotemp}), and stars (\\S\\ref{subsec:startemp}). Section~\\ref{sec:performance} analyzes the completeness, purity, accuracy, and precision of the samples classified and measured by the \\texttt{idlspec2d} pipeline. Section~\\ref{sec:issues} documents known issues in the DR9 release of BOSS data, and \\S\\ref{sec:summary} provides a summary and conclusions. ", "conclusions": "\\label{sec:summary} We have described the ``1D'' component of the \\texttt{idlspec2d} pipeline that provides automated redshift measurement and and classification for the SDSS-III BOSS DR9 data set, which comprises 831,000 optical spectra. This software is substantially similar to the \\texttt{idlspec2d} redshift analysis code used for SDSS-I/II data, but has been upgraded with new templates and several new algorithms for application to the BOSS project, and has been presented in great detail for the first time in this work. The pipeline also provides additional parameter measurements, including emission-line fits for all objects, and velocity-dispersion likelihood curves for objects classified as galaxies. The redshift success rate of the \\texttt{idlspec2d} pipeline is well in excess of the scientific requirements of the BOSS project. The software provides first-principles estimates of statistical redshift errors that are Gaussian distributed and accurate to within small correction factors. The ``2D'' component of the \\texttt{idlspec2d} pipeline that extracts spectra from raw CCD pixels is the subject of \\citet{Schlegel12}. Full data-model information for both the 2D and 1D BOSS pipeline outputs can be found at the SDSS-III DR9 website (\\url{http://www.sdss3.org/dr9/}). Development work continues on data-reduction software for BOSS, both in the calibration and extraction of spectra, and in the classification and redshift analysis procedures. Subsequent BOSS data releases will be accompanied by similar documentation of the implemented results of this ongoing development." }, "1207/1207.0506_arXiv.txt": { "abstract": "The search for diffuse non-thermal, inverse Compton (IC) emission from galaxy clusters at hard X-ray energies has been underway for many years, with most detections being either of low significance or controversial. In this work, we investigate 14--195 keV spectra from the \\swifts BAT all sky survey for evidence of non-thermal excess emission above the exponentially decreasing tail of thermal emission in the flux-limited \\his sample. To account for the thermal flux contribution at BAT energies, \\xmms EPIC spectra are extracted from coincident spatial regions so that both the thermal and non-thermal spectral components can be determined simultaneously. We find marginally significant IC components in 6 clusters, though after closer inspection and consideration of systematic errors we are unable to claim a clear detection in any of them. The spectra of all clusters are also summed to enhance a cumulative non-thermal signal not quite detectable in individual clusters. After constructing a model based on single temperature fits to the \\xmms data alone, we see no significant excess emission above that predicted by the thermal model determined at soft energies. This result also holds for the summed spectra of various subgroups, except for the subsample of clusters with diffuse radio emission. For clusters hosting a diffuse radio halo, relic, or mini-halo, non-thermal emission is initially detected at the $\\sim5\\sigma$ confidence level, but modeling and systematic uncertainties ultimately degrade this significance. This marginal detection is driven by the mini-halo subgroup, suggesting low average magnetic field strengths ($B \\sim 0.1 \\mu$G) in the cores of these clusters. ", "introduction": "\\label{sec:bathi:intro} A number of observations, mainly at radio frequencies, have established that relativistic particles and magnetic fields are part of the intracluster medium (ICM) of galaxy clusters \\citep[e.g.,][]{GF04}. The large ($\\sim$Mpc) scale, diffuse structures known as radio halos and relics are produced by relativistic electrons spiraling around $\\sim$$\\mu$G magnetic fields. Because halos and relics are not detected in every cluster, but are only found in clusters with ongoing major merger activity \\citep{Buo01, SBR+01}, mergers probably temporarily reaccelerate underlying relativistic populations \\citep[e.g.,][]{Sar99, BB05}. It is important to fully characterize the non-thermal phase if the dynamics and general state of the ICM is to be understood; the proportion of energy tied up in these relativistic components, if significant, may bias inferred mass estimates necessary to use clusters as cosmological probes \\citep[e.g.,][]{MAE+08, Vik+09, Van+10}. Unfortunately, synchrotron emission alone cannot separately determine particle and magnetic field energy densities, and so the total energy in the non-thermal phase remains relatively unconstrained. However, the electron population can be independently observed through inverse Compton (IC) emission due to scattering of the ubiquitous Cosmic Microwave Background (CMB) photons, which are up-scattered to X-ray energies and may be observable if the electron population is sufficiently large \\citep{Rep79}. Detections of IC emission, therefore, have the potential to determine whether the non-thermal phase is energetically negligible or, particularly if the average magnetic field is large, it is sizable enough to affect the dynamics and structure of the thermal gas. Thermal emission clearly dominates at $\\sim$keV energies, so searches for excess emission due to an IC spectral component are more easily undertaken at very soft or hard ($>$ 10 keV) energies. The latter range is particularly promising, given the exponential decline in the thermal spectrum and the lack of Galactic and solar wind charge exchange foregrounds that can hamper searches at soft energies \\citep{KLK+09, THF+07, BLB09}. In particular, the \\swifts BAT all sky survey \\citep{Tue+10} provides a deep map of hard energy (14--195 keV) emission from which non-thermal excesses can be identified. Its uniform coverage and impressive sensitivity makes it the most complete dataset from which to study the brightest objects in a given class \\citep[e.g.,][]{WMR+09}. Whereas previous searches have concentrated on long pointed observations of individual clusters, this survey allows a larger, more uniform sample to be searched, as similarly done by \\citet{Aje+09, Aje+10} for detected BAT clusters. To take full advantage of this capability, we have chosen the flux-limited \\his sample \\citep{RB02}, which contains the brightest clusters in the sky outside the Galactic plane. The selection of the brightest clusters may provide the greatest opportunity to detect IC emission, as in most models the nearest and most luminous clusters are expected to have the strongest IC signal. Also, because these clusters are bright and contained within a well-defined survey, there already exist good observations at lower X-ray energies, which can be used to strongly constrain the thermal properties of the ICM -- an important prerequisite for the robust detection of an IC excess. Finally, the fact that \\his is a complete flux-limited survey allows one to discuss the statistical properties of their hard excesses by stacking the individual cluster observations. Because they are nearby and bright, many of the clusters in \\his have been targets of IC searches with other telescopes, including A3667 \\citep{FSN+10}, A3112 \\citep{BNL07}, A3376 \\citep{Kaw+09}, A2256 \\citep{FLO05}, A1367 \\citep{HM01}, A2199 \\citep{KS00}, and A2163 \\citep{RGA06}. Most often clusters are targeted because they host a radio halo or relic, as the IC flux then leads to a direct measure of the average magnetic field strength. A large fraction of \\his clusters were also included in an analysis of all long exposure \\saxs observations \\citep{NOB+04}, which found marginal evidence for non-thermal excesses in individual clusters but a substantial excess in a stacked spectrum. In general, an IC component distinct from thermal emission in the hard band has been difficult to clearly identify, with perhaps the only counter example being an exceptionally deep observation of the Ophiuchus cluster \\citep{EPP+08}. The cluster most thoroughly searched for non-thermal emission, also in \\hi, is the Coma cluster. Controversial \\citep{RM04} detections with \\rxtes \\citep{RG02} and \\saxs \\citep{FOB+04} have recently been challenged with comparable \\suzakus \\citep{WSF+09} observations and a detailed analysis of the \\swifts BAT survey data \\citep{WSF+11}. To perform the deepest hard X-ray survey of non-thermal emission in clusters to date, we jointly fit high quality \\xmms EPIC and \\swifts BAT spectra, extracted from identical regions and cross-calibrated to make their absolute spectral responses as consistent as possible. We describe the data and its calibration in Section~\\ref{sec:bathi:obs}. In Section~\\ref{sec:bathi:separate}, the thermal and non-thermal character of the spectra are separately analyzed, and in Section~\\ref{sec:bathi:joint} they are jointly fit for each individual cluster. We also search for a statistical hard excess in sets of stacked spectra for the entire sample and for several subsamples in Section~\\ref{sec:bathi:stack}. Lastly, the implications of our results are discussed in Section~\\ref{sec:bathi:disc}. We assume a flat cosmology with $\\Omega_M = 0.23$ and $H_0 = 70$ km s$^{-1}$ Mpc$^{-1}$. Unless otherwise stated, all uncertainties are given at the 90\\% confidence level. ", "conclusions": "" }, "1207/1207.2673_arXiv.txt": { "abstract": "{The attachment of free electrons to polycondensed aromatic ring molecules (PAHs) is studied for the variety of these molecules with different numbers of condensed rings and over a broad range of electron temperatures, using a multichannel quantum scattering approach. The calculations of the relevant cross sections are used in turn to model the corresponding attachment rates for each of the systems under study, and these rates are parametrized as a function of temperature using a commonly employed expression for two-body processes in the interstellar medium (ISM).} {The scope of this work is to use first principles to establish the influence of chemical properties on the efficiency of the electron-attachment process for PAHs.} {Quantum multichannel scattering methods are employed to generate the relevant cross sections, hence the attachment rates, using integral elastic cross sections computed over a broad range of relevant energies, from threshold up to 1000 K and linking the attachment to low-energy resonant collisions.} {The rates obtained for the present molecules are found to markedly vary within the test ensemble of the present work and to be lower than the earlier values used for the entire class of PAHs anions, when modelling their evolutions in ISM environments. The effects of such differences on the evolutions of chemical networks that include both PAH and PAH- species are analysed in some detail and related to previous calculations.}{} ", "introduction": "\\label{intro} \\noindent The widespread presence in the interstellar medium (ISM) of polycyclic aromatic hydrocarbons (PAHs), i.e. organic molecules made up of several aromatic rings fused together, as well as the existence in the same environments of their dehydrogenated, ionized, and/or protonated counterparts, has been largely inferred from the extensive presence of unidentified infrared emission bands that have been observed in the range of 3 to 14 micrometres \\citep{Hendecourt1997, Rhee2007,Parker2012,Ricks2009}. Current astrochemical models for the PAH formation have largely been derived from the combustion chemistry community that suggest how these formation processes should occur in warm and dense circumstellar envelopes of carbon-rich stars through sequential reactions of smaller radicals with acetylene at temperatures around 1000 K (\\citet{Cherchneff1992} and erratum). Even more recently, their formation at temperatures down to 10 K has been surmised after an interesting investigation that combines cross molecular beam studies and quantum chemistry calculations \\citep{Parker2012}. Another low-energy path to possible formations of PAHs has come from investigating the role of the free electrons that are present in the circumstellar and interstellar media and in protoplanetary atmospheres \\citep{Herbst2008} in order to assess from calculations and possible experimental data the likely attachment efficiency of these free projectiles to the PAHs and the ensuing formation of very reactive anionic species, which can in turn react with radicals present in those environments \\citep{DeMarais2012}. That even temporary anions of smaller members of this family of aromatic molecules could be involved in subsequent reactions with cationic radicals has also been suggested in our earlier computational studies of benzyne anions \\citep{Carelli2010, Carelli2011}. Recent experimental work has shown that indeed anionic PAHs like phenide (C$_6$H$_5^-$), naphthalenide (C$_{10}$H$_7^-$), and anthracenide (C$_{14}$H$_9^-$) can efficiently react with H, H$_2$, and D$_2$ as observed in flowing afterglow-selected experiments \\citep{DeMarais2012,Yang2011}. It therefore follows from the above that to establish the possible effects of the presence of anionic PAHs on the chemistry of dark circumstellar regions and to model both the possible efficiency of formations and the consequences of such additional partners on the evolutionary history of that chemistry becomes directly interesting for improving our understanding of the PAH chemistry in the ISM in general. In particular, the present work wishes to analyse whether the chemical properties of the individual component partners in the large series of postulated PAHs play any significant role in the evolutionary studies of the chemical networks and if it is still a realistic choice to view their efficiency as given by a single rate value with a unique temperature dependence. In the following we therefore intend to approach the problem of identifying, if at all possible, differences in behaviour for a subset of PAHs chosen as initial examples for that entire class of molecules, which are related to their differences in chemical properties. In the next section we therefore summarize first the methods employed in earlier models for obtaining the attachment rates for the PAHs and then present our quantum scattering calculations in some detail to obtain integral cross sections from threshold up to several eV of energy and further link such observables with the model we use to yield the relevant attachment rates. Section 3 reports our results for a small group of PAH molecules and discusses their differences in relative efficiencies, while section 4 describes the final effects of such differences on the evolutionary history of the chemical networks currently employed for simulations. Our conclusions are given in section 5. ", "conclusions": "\\label{conclusions} In the present work we have endeavoured to analyse in some detail how one can realistically generate electron-attachment cross sections over a broad range of energies for gas-phase polycondensed aromatic molecules, which is a specific subset of those molecules chosen to be representative of the more general behaviour of PAHs in the interstellar and circumstellar environments. In particular, we selected a dynamical model that employs the above cross sections to obtain attachment rates over a range of temperatures that are of direct interest in the ISM environments. The corresponding rates are essentially considered to be upper bounds to the true attachment rates, but are nevertheless found to already be less than those employed in other evolutionary models, e.g. see \\citet{Wakelam2008}, where no direct scattering calculations were used to generate attachment rates. Additionally, the parametric representation of computed rates is chosen to be the one usually employed in the literature to describe 2B processes in the ISM \\citep{Woodall2007} and is found by our present work to depend on at least two chemical properties of the aromatic molecules: (i) the presence of a permanent dipole moment, a features that greatly increases the size of the attachment rates, especially at the threshold temperatures; and (ii) a minor effect of the number of C atoms or of the number of condensed rings in the molecular system. The capture rates were further employed to model the time evolution within a pseudo time-dependent dense cloud description that follows the one given earlier by \\citet{Wakelam2008}. In particular, we tried to analyse the effects of the attachment rate values produced by the present work on the evolutions of negatively charged species with time (e.g. see Fig.\\ref{EA2_negative}) and the evolutions of PAH/PAH$^-$ abundances throughout the same model, as well as the electron number density changes over the same time period (see Fig.\\ref{EA2_PAH}). The corresponding calculations suggest rather clearly that: \\begin{itemize} \\item[(i)] the attachment rates are lower, at least over the relevant range of temperatures, than the values previously employed; \\item[(ii)] they depend on the chemical properties of the specific PAH being considered, especially when comparing anthracene and coronene; \\item[(iii)] the electron densities remain higher over a much broader range of time and reach a steady state much later that thought before \\citep{Wakelam2008}; \\item[(iv)] the ``soaking-up'' power of the PAHs we have considered here turns out to be less than expected, hence reducing the efficiency of the neutralization processes. \\end{itemize} Although we feel that further numerical experiments are needed in order to more extensively include a broader variety of polar PAHs within the evolutionary model, the set of exemplary molecules presented here are already making it clear to us that chemical variety plays a role that is bigger than previously expected and that a molecule-specific modelling of the role of electrons in evolutionary studies should be sought as much as possible. In conclusion we postulate here that the formation of a variety of TNIs and of zero-energy virtual states are all features included in our scattering calculations and are instrumental mechanisms in yielding the final electron attachment rates that we in turn have employed for the evolutionary studies reported in the present work." }, "1207/1207.5053_arXiv.txt": { "abstract": "We present the current results from the development of a wide integral field infrared spectrograph (WIFIS). WIFIS offers an unprecedented combination of etendue and spectral resolving power for seeing-limited, integral field observations in the $0.9-1.8$ $\\mu$m range and is most sensitive in the $0.9-1.35$ $\\mu$m range. Its optical design consists of front-end re-imaging optics, an all-reflective image slicer-type, integral field unit (IFU) called FISICA, and a long-slit grating spectrograph back-end that is coupled with a HAWAII 2RG focal plane array. The full wavelength range is achieved by selecting between two different gratings. By virtue of its re-imaging optics, the spectrograph is quite versatile and can be used at multiple telescopes. The size of its field-of-view is unrivalled by other similar spectrographs, offering a 4.5\\arcsec$\\times$ 12\\arcsecl integral field at a 10-meter class telescope (or 20\\arcsec$\\times$ 50\\arcsecl at a 2.3-meter telescope). The use of WIFIS will be crucial in astronomical problems which require wide-field, two-dimensional spectroscopy such as the study of merging galaxies at moderate redshift and nearby star/planet-forming regions and supernova remnants. We discuss the final optical design of WIFIS, and its predicted on-sky performance on two reference telescope platforms: the 2.3-m Steward Bok telescope and the 10.4-m Gran Telescopio Canarias. We also present the results from our laboratory characterization of FISICA. IFU properties such as magnification, field-mapping, and slit width along the entire slit length were measured by our tests. The construction and testing of WIFIS is expected to be completed by early 2013. We plan to commission the instrument at the 2.3-m Steward Bok telescope at Kitt Peak, USA in Spring 2013. ", "introduction": "\\label{sec:intro} % With the advent of megapixel, CMOS-based detectors, integral field spectroscopy (IFS) has become practicable in the near-infrared (NIR). IFS allows one to obtain spectra over both spatial dimensions, providing spectral information at each spatial co-ordinate on the sky, which is crucial for the study of various types of interesting astronomical objects (e.g., kinematics, chemistry, and stellar properties of high-redshift galaxies). However, most NIR integral field spectrographs used in large telescopes are coupled with adaptive optics (AO) systems to produce high spatial resolution, integral field spectra at moderate spectral resolution ($R \\sim 3000$). As a result, the etendue ($A\\times\\Omega$), a figure of merit that reflects how efficiently a telescope is used per unit time, of these spectrographs is relatively small. In Figure \\ref{fig:ifuetendue} we compare the etendues and spectral resolving powers of current/upcoming NIR integral field spectrographs such as KMOS, along with the values for our own Wide Integral Field Infrared Spectrograph (WIFIS) instrument that we describe in this paper. Most of these spectrographs predominantly focus on scientific problems that require high angular resolution and small fields at moderate spectral resolution. However, there are a myriad of scientific problems that require IFS over a larger field. These include the characterization of the ionization properties of gas in nearby extended star forming regions or supernova remnants, the determination of the kinematics and stellar properties of nearby galaxies, and the measurement of kinematics and star-forming properties of moderate redshift ($z \\sim 0.5-1.0$) merging galaxies. In all of these cases, the etendue of the instrument becomes critical because one needs to obtain spectroscopic information of fairly extended and faint emission. For this reason, we set out to construct WIFIS, an integral field spectrograph with an unprecedented etendue (see Figure \\ref{fig:ifuetendue}). \\begin{figure} \\begin{center} \\begin{tabular}{c} \\includegraphics[height=7cm]{ifu_plot_v2.pdf} \\end{tabular} \\end{center} \\caption[ifuetendue] { \\label{fig:ifuetendue} The etendue values of several near-IR integral field spectrographs currently in use or ones that will be commissioned in the near future such as KMOS. We also show the expected etendue of our previous, WIFIS1, and current, WIFIS2, designs. SINFONI can be used in both AO and non-AO mode, where its non-AO mode etendue is shown with an asterix. KMOS has 24 IFUs that can be deployed across a large field. We only show the etendue of a single KMOS IFU because there are mechanical restrictions on how close the individual IFUs can be deployed. WIFIS clearly has an unprecedented combination of etendue and spectral resolving power over a single field. } \\end{figure} \\par WIFIS is an image slicer-based imaging spectrograph designed to have the maximum size of the integral field affordable with the Hawaii 2RG (H2RG) NIR array of $2048\\times2048$ pixels with a broad-band spectral coverage at a medium spectral resolution (R $\\sim 3,000-5,000$) in a single exposure. The desire to maximize both spatial and spectral coverage while maintaining satisfactory optical performance is quite challenging, resulting in a very fast optical system for the spectrograph camera (see below). Two separate optical designs were developed for WIFIS: (1) WIFIS1 was designed to be a cryogenic instrument that operated over the $0.9-2.5\\mu$m range with a spectral resolving power R$\\sim 5000$ and a field-of-view of 4.5\\arcsec$\\times$ 12\\arcsec at a 10-meter telescope. The full spectral range was not accessible simultaneously and different gratings had to be used to obtain complete spectral coverage. The optical design of this instrument is discussed in our previous work\\cite{chou2010}. (2) WIFIS2 does not require fully cryogenic optics and maintains the same field-of-view (and etendue), but has a reduced resolving power of $R\\sim3000$ and operating wavelength range. The WIFIS2 design operates over the $0.9-1.8$ $\\mu$m spectral range with the $0.9-1.35$ $\\mu$m bandpass offering the best spectral resolving power and sensitivity. The full spectral coverage is also achieved using two different gratings. In the interest of reducing the complexity, cost of the spectrograph, and deployment time, we chose to construct our WIFIS2 design. Hereafter, we will refer to the WIFIS2 design as WIFIS. \\begin{figure} \\begin{center} \\begin{tabular}{c} \\includegraphics[height=7cm]{ngc2623v2.pdf} \\end{tabular} \\end{center} \\caption[ngc2623] { \\label{fig:ngc2623} An optical Hubble Space Telescope colour composite of NGC 2623, a nearby major merger of galaxies. The intensities are shown in a logarithmic scale to reveal the low surface brightness features. We overlay the expected 4.5\\arcsec$\\times$12\\arcsecl field-of-view of WIFIS if this galaxy was at a redshift of 0.5, the distance at which H$\\alpha$ emission is visible in the WIFIS bandpass. Each red rectangular region is a 0.25\\arcsec$\\times$ 12\\arcsecl slice of WIFIS's integral field when coupled with a 10-meter telescope. One can see the benefits of the large etendue afforded by WIFIS when observing merging galaxies at moderate redshift. } \\end{figure} \\par WIFIS has been designed to be versatile and can operate at different telescopes by simply replacing its reimaging optics. Its immediate intended destination is the Steward Bok 2.3-meter telescope. While the aperture size is relatively small, the size of its integral field-of-view is notably large (20\\arcsec$\\times$ 50\\arcsec) and is perfectly suited for mapping line emission (e.g. Pa$\\beta$, [Fe II]) from extended galactic star forming regions, supernova remnants, and measuring the kinematics and stellar properties of nearby galaxies. For a 10-m class telescope, such as the 10.4-meter Gran Telescopio Canarias (GTC), WIFIS has a smaller field (4.5\\arcsec$\\times$12\\arcsec) with the increased sensitivity optimized to study the star forming properties of merging galaxies at moderate redshift $z \\sim 0.5-1.0$ where the H$\\alpha$ line is redshifted into WIFIS's most sensitive bandpass. An example observation is shown in Figure \\ref{fig:ngc2623} where a local major merger, NGC 2623, has been placed at a redshift of $z \\sim 0.5.$ The large field-of-view of WIFIS is well-matched to the physical size of this merger, which clearly illustrates the necessity of large field IFS in the NIR wavebands. \\par In this paper, we discuss the final optical design and system layout of WIFIS along with its predicted performance, including its on-sky characteristics, sensitivity, and spectral resolving power at both the Bok telescope and the GTC. We also present the results from the laboratorial characterization of its integral field unit. Finally, we discuss the current status of the project and our future plans. ", "conclusions": "In this work, we present the final optical design and system layout of WIFIS, a wide integral field IR spectrograph, designed to operate in the $0.9-1.8$ $\\mu$m wavelength range. WIFIS has an unrivalled etendue when it comes to observing a single field in the NIR. This is achieved through the use of image-slicer type IFU, complex optics, and a large format H2RG detector array. The wide-field capability of WIFIS benefits a unique class of scientific problems, previously inaccessible by small-field NIR integral field spectrographs, that range from ionization and chemical properties of nearby star forming regions to kinematics and star forming properties of moderate redshift galaxies. The versatility of WIFIS also allows one to use the instrument in different classes of telescopes ranging from a 2.3-meter telescope to a 10-meter one. Each telescope destination provides a unique set of capabilities because the field-of-view of the instrument scales inversely with telescope size even though the sensitivity is reduced at a smaller telescope. We also carry out simulations to determine the predicted on-sky sensitivity and spectral resolving power of the instrument at the Bok and GTC. We find that our sensitivities are comparable to other similar instruments that will be commissioned in the near future. Finally, we present results from our laboratorial characterization of the integral field unit. We plan on commissioning this instrument at the Steward Bok telescope during Spring 2013." }, "1207/1207.1050_arXiv.txt": { "abstract": "Recent N-Body simulations are in favor of the presence of a co-rotating Dark Disk that might contribute significantly (10\\%-50\\%) to the local Dark Matter density. Such substructure could have dramatic effect on directional detection. Indeed, in the case of a null lag velocity, one expects an isotropic WIMP velocity distribution arising from the Dark Disk contribution, which might weaken the strong angular signature expected in directional detection. For a wide range of Dark Disk parameters, we evaluate in this Letter the effect of such dark component on the discovery potential of upcoming directional detectors. As a conclusion of our study, using only the angular distribution of nuclear recoils, we show that Dark Disk models as suggested by recent N-Body simulations will not affect significantly the Dark Matter reach of directional detection, even in extreme configurations. ", "introduction": "Within the standard Dark Matter halo paradigm, the local Dark Matter distribution is assumed to be smoothly spatially distributed and to be well-described by a Maxwellian velocity distribution. However, the hierarchical structure formation model indicates that the Galactic Dark Matter halo results from successive small halo accretions, thus directly linking its structure to its merging history. The presence of substructures in the Milky Way halo is inferred from recent results of N-body simulation \\cite{nbody.vl2,nezri,Kuhlen:2012fz,lisanti1,vogel,moore,klypin,kravtos}. Such substructures may be classified as follows : Dark Matter tidal streams (spatially localized), debris flows (spatially homogenized but with velocity substructures) and a Dark Disk. The latter has received much interest since late sub-halo merging is expected to lead to the formation of a co-rotating Dark Disk \\cite{nezri,Read:2008fh,Read:2009,Purcell:2009yp} that may affect the expected WIMP signal both in direct and directional detection. While the influence of the dark disk on Dark Matter signals has been exhaustively investigated for direct \\cite{nezri,Ling:2009cn,Bruch:2009rp,Green:2010gw} and indirect \\cite{Bruch:2008rx} detection, it is still unclear how it may affect directional detection. Following a previous work from A.~M.~Green \\cite{Green:2010gw}, we aim at evaluating the influence of the presence of a co-rotating Dark Disk on the discovery potential of a forthcoming directional detector. In particular, we are interested in determining the values of the Dark Disk parameters for it to significantly affect the Dark Matter reach of directional detection. In order to be model independent from the background energy modelling, the study has been done by considering only the angular distribution of nuclear recoils $dR/d\\Omega_r$.\\\\ Since the pioneering paper of D.~N.~Spergel~\\cite{spergel}, the contribution of directional detection to the field of Dark Matter has been addressed through a wealth of studies~\\cite{Kuhlen:2012fz,billard.exclusion,henderson,morgan1,morgan2,copi1,copi2,copi3,green1,green2,billard.disco,billard.profile,green.disco,billard.ident,Alves:2012ay,Lee:2012pf,albornoz,Bozorgnia:2011vc,Bozorgnia:2012,Creswick:2010dm,Lisanti:2009vy,Alenazi:2007sy,Gondolo:2002np}. Depending on the unknown WIMP-nucleon cross section, directional detection may be used to : exclude Dark Matter \\cite{billard.exclusion,henderson}, reject the isotropy hypothesis \\cite{morgan1,morgan2,copi1,copi2,copi3,green1,green2}, discover galactic Dark Matter with a high significance \\cite{billard.disco,billard.profile,green.disco} or constrain WIMP and halo properties \\cite{billard.ident,Alves:2012ay,Lee:2012pf}. In particular, for neutralino Dark Matter, a large fraction of MSSM configurations with a neutralino lighter than 200 $\\rm GeV/c^2$ would lead to a significance greater than 3$\\sigma$ (90\\% CL) in a 30 kg.year $CF_4$ directional detector \\cite{albornoz}.\\\\ In the following, we focus on the effect of a co-rotating Dark Disk on the potential of forthcoming directional detectors to discover Dark Matter \\cite{billard.profile}. The paper is organized as follows. Section \\ref{sec:dd} presents the current knowledge on the Dark Disk. In particular, we define its parameterization used throughout. The directional framework is recalled in sec.~\\ref{sec:directional}, with emphasize on the directional statistic methods used to exploit forthcoming data. Then, for a wide range of Dark Disk parameters, we evaluate in \\ref{sec:disco} the effect of such substructure on the discovery potential of upcoming directional detectors. ", "conclusions": "A co-rotating Dark Disk, as predicted by recent N-Body simulations, might contribute (10\\%-50\\%) to the local Dark Matter density, with a potentially dramatic effect on directional detection. In this letter, we have evaluated the effect of Dark Disk model on the discovery potential of upcoming directional detectors. We conclude that, if a co-rotating Dark Disk is present in our Galaxy and has the properties predicted by N-Body simulations \\cite{nezri}, the discovery potential of directional detection would be strictly unchanged. Only an extreme and unrealistic Dark Disk model (high co-rotational velocity and high velocity dispersion) might affect significantly the Dark Matter reach of upcoming directional detectors, by increasing the discovery limit by a factor of three at high WIMP mass ($m_{\\chi} \\sim 1000$ GeV/c$^2$). Additionally, we also have shown that anisotropic features in the Dark Matter velocity distribution of the Dark Disk will only have a small effect on the expected directional signal. Hence, according to our results we believe that the possibility of the existence of a co-rotational Dark Disk in our galaxy shouldn't be a threat for upcoming directional detection experiments.\\\\ Interestingly, note that even if the impact of Dark Disk contribution to the local Dark Matter distribution only mildly affects the discovery potential of directional detection, it may significantly affect the mass and cross section determination \\cite{billard.ident}. Indeed, as explained in \\cite{Green:2010gw}, WIMP events arising from the Dark Disk contribution will induce an excess at low recoil energies which can lower the estimation of the WIMP mass when considering a standard halo model. As outlined in \\cite{Lee:2012pf}, the presence of a Dark Disk restricts the ability to constrain the Dark Matter parameters (both from the halo and particle physics). Of course, a measurement of the parameters of the Dark Disk itself remains challenging with the exposure of the next generation of directional detectors (30 kg.year). This highlights the fact that even if a co-rotating Dark Disk is not a threat to the discovery potential of directional detection, it has to be characterized in order to consistently constrain the Dark Matter properties." }, "1207/1207.3113_arXiv.txt": { "abstract": "Two significant progresses have been made in the past years on our understanding of hot accretion flows. One is that only a small fraction of accretion flow available at the outer boundary can finally falls onto the black hole while most of them is lost in outflow. Another one is that electrons may directly receive a large fraction of the viscously dissipated energy in the accretion flow, i.e, $\\delta\\sim 0.1-0.5$. The radiative efficiency of hot accretion flow when these two progresses are taken into account has not been systematically studied and is the subject of the present paper. We consider two regimes of hot accretion model. One is the advection dominated accretion flows (ADAFs) which lie on low accretion rate regime, $\\la 10\\alpha^2\\ledd/c^2$; another being the luminous hot accretion flows (LHAFs) which lie above this accretion rate. For the latter, we assume that the accretion flow will has a two-phase structure above a certain accretion rate, and a simplification is adopted in our calculation of the dynamics. Our results indicate that the radiative efficiency of hot accretion flow increases with the accretion rate and is highly enhanced by the direct viscous heating to electrons compared to the previous case of $\\delta\\ll 1$. When the accretion rate is high, the radiative efficiency of hot accretion flow is comparable to that of the standard thin disk. Fitting formulae of radiative efficiency as a function of accretion rate for various $\\delta$ values are presented. ", "introduction": "\\label{s:intro} One of the most important parameters in accretion theory is the radiative efficiency. This parameter describes the significance of converting rest-mass energy into radiative energy, \\begin{equation} \\epsilon\\equiv {L\\over\\mdot c^2},\\label{eq:epsilon} \\end{equation} where $L$ is the total luminosity emitted from the accretion flow and $\\mdot$ is the corresponding mass accretion rate of the system. Take the standard thin disc model \\citep[][hereafter SSD]{ss73} as an example, its radiative efficiency lies in the range $0.057 - 0.43$, depending on the spin of the black hole \\citep{nt73}. According to the differences of temperature and mass accretion rate of accretion flows, we now have four accretion models which belong to two series, namely cold and hot ones. In the cold series, when the accretion rate is lower than the Eddington rate, $\\dot{M}\\la \\medd (\\equiv 10 \\ledd/c^2)$, we have the standard thin disc model \\citep{ss73}. When $\\dot{M}\\ga \\medd$, the model is the slim disk \\citep{a88}. The temperature of the gas in these two models is roughly within the range of $10^5-10^7{\\rm K}$. The accretion flows are optically thick, emitting multi-temperature blackbody spectrum \\citep{fkr02}. The radiative efficiency of the former is high, and is independent of the accretion rate \\citep{nt73}. In a slim disk, the optical depth is so large that photons are trapped in the accretion flow and advected into the black hole; therefore the radiative efficiency is lower \\citep{a88,m00,s09}. In the hot series of model, the temperature of the accretion flow is almost virial. When the mass accretion rate is below a critical value, $\\dot{M}_{\\rm cr,ADAF}\\approx 5\\theta_e^{3/2}\\alpha^2\\medd$ with $\\theta_e\\equiv kT_e/m_ec^2$, we have the advection-dominated accretion flows (ADAFs; \\citealt{i77,r82}, \\citealt{ny94,a95}; see \\citealt{nmq98,nm08} for reviews)\\footnote{As we will state below, the mass accretion rate of hot accretion flows is a function of radius. In addition, the value of $\\dot{M}_{\\rm cr,ADAF}$ is a function of parameter $\\delta$. The value of $\\dot{M}_{\\rm cr,ADAF}$ cited here was obtained when $\\dot{M}(R)=constant$ and $\\delta=10^{-3}$. The result of other cases will be presented in the present paper.}. In an ADAF the gas is tenuous. The Coulomb coupling between electrons and ions is not strong enough thus the flow is two-temperature with the ions being much hotter than the electrons \\citep{ny95}. The critical accretion rate $\\dot{M}_{\\rm cr,ADAF}$ is determined by the balance between the Coulomb collision and viscous heating in the ions energy equation. When $\\dot{M}\\ll \\dot{M}_{\\rm cr,ADAF}$, most of the viscously liberated energy is stored as the gas internal energy and advected into the black hole rather than being transferred from ions to electrons and radiated away; therefore, the radiative efficiency of an ADAF is very low. With the increase of $\\dot{M}$, more and more viscously dissipated energy will be transferred into electrons and radiated away until $\\dot{M}_{\\rm cr,ADAF}$ is reached at which advection is no longer dominated. ADAFs have been widely applied to the low-luminosity black hole sources including the supermassive black hole in our Galactic center, Sgr A*, low-luminosity AGNs, and the quiescent and hard states of black hole X-ray binaries \\citep{yqn03,n05,y07,nm08}. When $\\dot{M}\\ga \\dot{M}_{\\rm cr,ADAF}$, Coulomb collision cooling becomes stronger than the viscous heating. \\citet{y01} found that in this case, up to another critical accretion rate, there exist another hot accretion solution in which the sum of the compression work ($PdV$ work) and viscous heating balances the cooling. Compared to ADAFs, this model corresponds to higher accretion rates and radiative efficiency; thus it is called luminous hot accretion flow (LHAF; \\citealt{y01}). LHAFs have been invoked to explain the origin of hard X-ray emissions detected in luminous X-ray sources such as Seyfert galaxies and luminous hard state of black hole X-ray binaries \\citep{yz04,yz07}. In the present paper we focus on the radiative efficiency of hot accretion flows. Using the data from \\citet{e97}, \\citet{nmq98} presented in their Fig. 7 the relationship between bolometric luminosity and the accretion rate of ADAFs (see also \\citealt{ny95}). \\citet{y01} investigated the radiative efficiency of LHAFs and found that $\\epsilon_{\\rm LHAF}$ is higher than the typical ADAF value. Both calculations are, however, based on the ``old'' version of hot accretion flow models in the sense that it is assumed that the mass accretion rate is a constant of radius and the value of parameter $\\delta$, which describes the fraction of turbulent dissipation that heat the electrons directly, is very small, $\\delta\\sim 10^{-3}$. Both assumptions are now known to be no longer correct after the development of the accretion flow theory in the recent years, as we will illustrate in detail in \\S2. In this paper we systematically revisit the efficiency of ADAFs and LHAFs after taking into account the new progresses of hot accretion flow theory. This paper is organized as follows. We first give a brief introduction to the recent progresses on hot accretion flows in Section\\ \\ref{s:out_vis}. We then describe our model in Section\\ \\ref{s:model}. Our calculation results are presented in Section\\ \\ref{s:result}. The last section is devoted to a discussion. ", "conclusions": "\\label{s:dis} In our calculations, we only consider the local Compton scattering, namely the scattering between photons and electrons occurred at the same region where the photons are produced. However, since a hot accretion flow is usually optically thin in the radial direction, the photons produced at one certain radius can in principle travel for a long distance and collide with electrons at another radius. Such a ``global'' Compton scattering effect has been systematically investigated in previous works \\citep{po01,po07,yxo09,x10,nxz12}. It was found that it plays a significant cooling and heating roles in the region of $50 R_{\\rm s}\\la R \\la 100 R_{\\rm s}$ and $R\\ga 5\\times 10^3 R_{\\rm s}$, respectively, when the accretion rate is high enough so that the total luminosity emitted from the accretion flow $L_{\\rm bol}\\ga 2\\%\\ledd$ \\citep{yxo09,x10}. One consequence is that the radiative efficiency will be lower by a factor of 2 compared to the case that this effect is not taken into account \\citep{x10}. In addition, the highest luminosity a hot accretion flow can emit will be constrained to be $L_{\\rm bol}\\la 1\\%\\ledd$. Above this limit, the global Compton cooling and heating will be so strong that no steady hot solution can be found and the system will ``oscillate'' \\citep{yxo09}. Note that if $L_{\\rm bol} \\la 2\\%\\ledd$, or outer boundary radius of hot flow is small, i.e. $R_{\\rm out}\\la 50-100R_{\\rm s}$, the global Compton scattering effects will be unimportant. The latter is the case of luminous hard state of black hole X-ray binaries. In our model, we assume that the magnetic filed is tangled and weak, thus it does not play any dynamical role. Numerical simulations have shown that a large-scale toroidal magnetic field is likely to exist in the inner region of the accretion flow, imposed on the stochastic component \\citep[e.g.][]{hkdh04}. The effect of such a field has been studied by self-similar approaches \\citep{af06,agn08,b09} or global calculations \\citep{o07,o12}. Especially, the global solution with strong large-scale magnetic fields indicates an increase in the highest luminosity a hot accretion flow can achieve \\citep{o12}. Throughout this paper, we fix the outer boundary condition ($T_{\\rm i}, T_{\\rm e}, v$) in our calculations. The effect of outer boundary condition on the dynamics of accretion flow has been studied in \\citet{y00}. Obviously, it will also influence the radiative efficiency. A correlation between the radio and the 2-10 keV luminosity ($L_{\\rm X}$) has been found among the hard state of black hole X-ray binaries and low-luminosity active galactic nuclei \\citep{c00,c03,g03,mhd03,gmf12,c12}, which is well described by a power-law, $L_{\\rm Radio} \\propto L_{\\rm X}^p$ with index $p \\sim 0.6$. Recently \\citet{z11} found that this correlation extends to intermediate and soft states for Cyg X-1, if only the luminosity from hot disc is used. This correlation has been quantitatively explained by the coupled jet-ADAF model in \\citet{yc05}, in which the radio and X-ray emissions are dominated by the radiation from the jet and ADAF, respectively. It is interesting to note that there are now growing number of sources which show that when $L_{\\rm X}\\ga 4\\times 10^{36}\\ergs$ the radio/X-ray correlation follows a steeper power-law, with index $p \\sim 0.98$ or $1.4$ \\citep{c11,gmf12,c12}. Below this critical luminosity, the sources return to the $\\sim 0.6$ correlation at $\\sim 10^{35}\\ergs$. Between $4\\times 10^{36}\\ergs$ and $10^{35}\\ergs$, the radio luminosity remains almost unchanged. \\citet{c11} proposed that one way to explain the steep ($p \\sim 1.4$) correlation is that the radiative efficiency of the hot accretion flow is independent of the accretion rate, if the ratio of the mass loss rate in the jet and the accretion rate in the accretion flow is a constant. As shown by Fig. 1, this is the case for our two-phase accretion flow (Type II LHAF). Moreover, as also shown by this figure, the efficiency curve of Type I LHAF is very steep, which means that a small change of accretion rate will result in a large change of $L_{\\rm X}$. This feature is obviously attractive to explain the ``flat transition'' between $4\\times 10^{36}\\ergs$ and $10^{35}\\ergs$. The reason why some sources follow a single $\\sim 0.6$ correlation while others follow three branches is simply because of different values of $\\alpha$ among these sources. If $\\alpha$ is large, $\\dot{M}_{\\rm cr,ADAF}$ will be large, so no transition to LHAFs will occur throughout the evolution of $\\dot{M}$ during the outburst. This is why we can only observe one single $\\sim 0.6$ correlation. If on the other hand $\\alpha$ in a source is small, $\\dot{M}_{\\rm cr,ADAF}$ will be small thus the sources will enter the two-phase LHAFs regime during the outburst. In this case, three branches of correlation should be expected. In a future work we plan to investigate the correlation in detail." }, "1207/1207.3922_arXiv.txt": { "abstract": "We consider static cosmological solutions along with their stability properties in the framework of a recently proposed theory of massive gravity. We show that the modification introduced in the cosmological equations leads to several new solutions, only sourced by a perfect fluid, generalizing the Einstein Static Universe found in General Relativity. Using dynamical system techniques and numerical analysis, we show that the found solutions can be either neutrally stable or unstable against spatially homogeneous and isotropic perturbations. ", "introduction": "The exact solution of Einstein's equations known as the Einstein Static (ES) Universe is a static closed Friedmann-Robertson-Walker model sourced by a perfect fluid and a cosmological constant (see \\cite{Hawking:1973uf}). Its stability properties have been widely investigated. The ES Universe is unstable to homogeneous perturbations \\cite{Edd:1930}, moreover it is always neutrally stable against small inhomogeneous vector and tensor perturbations and neutrally stable against adiabatic scalar inhomogeneities with high enough sound speed \\cite{Barrow:2003ni}. Furthermore, the ES Universe was recently shown to be unstable to Bianchi type-IX spatially homogeneous perturbations in the presence of nontilted and tilted perfect fluids with $\\rho + 3P>0$ \\cite{arXiv:1108.3962} and for several kinds of matter fields sources (see \\cite{Barrow:2009sj} and references therein). The renewed interest in the ES Universe, besides its historical importance, comes from the Emergent Universe scenario \\cite{Ellis:2002we}, an inflationary cosmological model in which it plays a crucial role as initial state. This model, in turn, suffers from a fine-tuning problem which is ameliorated when modifications to the cosmological equations of GR are present. For this reason, analogous solutions have been considered in several different modified gravity models \\cite{Gergely:2001tn} and quantum gravity models \\cite{Mulryne:2005ef,Parisi:2007kv,Boehmer:2009yz,Park:2009zra,Wu:2009ah,Canonico:2010fd}. Indeed, when dealing with modified cosmological equations, many new static solutions are present, whose stability properties, depending on the details of the single theory or family of theories taken into account, are substantially different from those of the classical ES solution of GR. In particular, neutrally stable solutions are present thus the fine-tuning problem is ameliorated but then a mechanism is needed to get out of the phase of infinite expansions and recollapses and to trigger the expanding phase of the Universe\\cite{graceful}. Here we study static cosmological solutions in the framework of a covariant Massive Gravity (MG) model recently proposed in \\cite{derham10,derahm11}. In order to construct a consistent theory, nonlinear terms should be tuned to remove order by order the negative energy state in the spectrum \\cite{boulware72}. The model under investigation follows from a procedure originally outlined in \\cite{arkani03,creminelli05} and has been found not to show ghosts at the complete nonlinear level with an arbitrary reference metric \\cite{Hassan:2011tf,Hassan:2011ea}. The considered theory exploits several remarkable features. Indeed the graviton mass typically manifests itself on cosmological scales at late times thus providing a natural explanation of the presently observed accelerating phase \\cite{car2012}. Moreover, the theory allows for exotic solutions in which the contribution of the graviton mass affects the dynamics at early times. Indeed, in contrast with GR where, in order to have static solutions, a cosmological constant term and a positive curvature term are needed in addition to a suitable perfect fluid source term, we find that in the considered MG theory it is possible to have static cosmological solutions only sourced by a perfect fluid. These solutions can be either unstable or neutrally stable and they exist even for spatially flat (i.e. $\\mathcal{K}=0$) cosmological models. This paper is structured as follows. In Sec. \\ref{massive}, the nonlinear MG framework considered in this work is shortly described and the modified Friedmann equations under investigation are introduced. In Sec. \\ref{stab} static cosmological solutions are found, a linearized analysis is performed and the stability properties are discussed in details. In Sec. \\ref{num} the dynamics near the fixed points is described using numerical integrations. The phase diagrams of the system are drawn both in the $(a, \\dot{a})-$plane and $(H,\\rho)-$plane. In Sec. \\ref{concl}, some conclusions are eventually drawn. ", "conclusions": "We have considered static solutions in the cosmological sector \\cite{chamseddine11} of a recently proposed theory of massive gravity \\cite{derham10}. We have shown that the effect of a massive graviton is to enrich the phase space of the cosmological equations, enlarging the ranges of existence of static solutions and affecting their stability properties. The solutions found exhibit stability properties rather different from those of the standard ES solution of GR, which requires a positive cosmological constant and a positive spatial curvature in addition to a perfect fluid with an equation of state parameter $w>-1/3$. Two kinds of solutions are present: neutrally stable solutions and unstable solutions (of the saddle type). Summing up, for spatially closed ($\\mathcal{K}=1$) models three cases are possible: $i)$ both the unstable and the neutrally stable solutions are admitted, $ii)$ either the unstable solution or the neutrally stable solution is admitted, or $iii)$ no static solutions are admitted. For spatially flat ($\\mathcal{K}=0$) and open ($\\mathcal{K}=-1$) models, three cases are possible: $i)$ both the unstable and the neutrally stable solutions are admitted, $ii)$ only the unstable solution is admitted, or $iii)$ no static solutions are admitted. Notice that, in the considered region of the parameter space, the neutrally stable solution requires a negative equation of state parameter $w$; in particular it must be $-1