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JSALT2024 Evaluating LLMs for Research Astronomy 7

2745880

2011A&A...526A.130S

1012

1012.2278_arXiv.txt

The majority of the known exoplanets have been discovered using the radial velocity technique \citep{1995Natur.378..355M}. In recent years, however, an ever increasing number have been discovered as a result of group-based and space-based transit search survey projects. Transiting exoplanets allow parameters such as the mass, radius, and density to be accurately determined, as well as their atmospheric properties to be studied during their transits and occultations \citep{2005ApJ...626..523C,2009MNRAS.394..272S,2009IAUS..253...99W}. The SuperWASP project has robotic observatories in the Canary Islands and South Africa. The wide angle survey is designed to find exoplanets around relatively bright stars in the $V$-magnitude range $9\sim13$. A detailed description of the SuperWASP project is given in \citet{2006PASP..118.1407P}. In this paper we report the discovery of WASP-34b, an exoplanet in orbit around its $V = 10.4$~{mag.} host star 1SWASP\,J110135.89-235138.4 in the constellation Crater. We present the WASP-South discovery photometry, together with Euler Telescope photometry and CORALIE radial velocity measurements.

WASP-34b is a sub-Jupiter mass exoplanet transiting a G5 host star every 4.3177 days. A simultaneous fit to transit photometry and radial-velocity measurements gave a planetary mass of $0.59 \pm 0.01$~$M_{\rm Jup}$ and radius of $1.22 ^{+0.11}_{-0.08}$~$R_{\rm Jup}$. In many respects the WASP-34b system looks like transiting version of 51 Peg \citep{1995Natur.378..355M} with similar periods, separations, masses and host stars. Of the known transiting system, WASP-34b joins an increasing number of exoplanets with masses $\sim0.5M_{\rm Jup}$ and radii $\sim1.0R_{\rm Jup}$, such as WASP-22b \citep{2010AJ....140.2007M}, WASP-25b \citep{2010arXiv1009.5917E}, HAT-P-24b \citep{2010arXiv1008.3389K} and HAT-P-25b \citep{2010arXiv1008.3565Q}. Of these systems, the similarity to WASP-22b is striking considering that it too has a linear trend in radial velocities due to a third-body in the system. \subsection{Properties of the third body} The significant linear trend in the radial velocities of 55 $\pm$4~m\,s$^{-1}$\,y$^{-1}$, noted in Sect.~\ref{CORALIE_Spec}, indicates the presence of a third body in the system with a long period. Using the constant acceleration method of \cite{2009ApJ...703L..99W} the third body satisfies \begin{equation} \frac{M_c \sin i_c}{{a_c}^2} = 0.307 \pm 0.022, \end{equation} where $M_c$ and $a_c$ are the mass (in $M_{\rm Jup}$) and orbital separation (in AU) of the third body. Given that the period of this outer body must be greater than at least twice the RV data baseline, $P_c \ga 490$~days and, using Kepler's third law, $a_c \ga 1.2$~AU. Hence $M_c \ga 0.45$~$M_{\rm Jup}$. Using literature photometry we estimated the bolometric flux at the Earth to be $f_{\rm bol} = 1.97 \pm 0.10 \times 10^{-9}$~erg\,s$^{-1}$\,cm$^{-2}$, which gives $T_{\rm eff}$ = 5740 $\pm$ 140~K using the Infrared Flux Method (IRFM) \citep{1977MNRAS.180..177B}, which is in excellent agreement with the spectroscopic temperature determination. The modified IRFM method of \cite{1993MNRAS.265.1035S} suggests that a main-sequence companion would be cooler than $\sim$4000~K (M-type or later). There is a {\it GALEX} \citep{2007ApJS..173..682M} near-UV flux measurement for WASP-34, and this is in in agreement with that predicted for a star of this temperature. This is no significant UV excess that could be attributed to a hot compact stellar companion, but this does not exclude the possibility of a cool white dwarf companion. Thus, we expect that the companion object is either a low mass star (M-type or cooler), cool white dwarf or another planet in a wide long-period orbit. Further radial velocity measurements are required to constrain the orbit. \subsection{Grazing nature of the system} The transit depth ($(R_{\rm P}/R_{*})^{2}$ = 0.0126) and high impact parameter ($b = 0.90$) suggest that WASP-34b could be the first transiting exoplanet known to undergo grazing transits. For a transit to be truly grazing, we require \begin{equation} b+{R_{\rm P}}/{R_{\star}} > 1 \end{equation} and that this {\it grazing criterion} is significant compared to the observational uncertainties. For WASP-34b we find a value of 1.016 $^{+0.017}_{-0.014}$ for the {\it grazing criterion}. Hence, we might expect grazing transits, unlike the near-grazing transits of TrES-2 \citep{2006ApJ...651L..61O} and HAT-P-14b \citep{2010ApJ...715..458T}, which have {\it grazing criteria} of 0.972$\pm$0.007 and 0.968$\pm$0.022, respectively. Figure~\ref{grazing} shows the 10\,000 MCMC posteriors for ${R_{\rm P}}/{R_{\star}}$ and $b$. A total of 89.6\% of these satisfy the {\it grazing criterion}. If the {\it grazing criterion} were equal to unity, then we would expect 50\% of the points to be higher and 50\% to be lower than one. Using the odds ratio test \citep{2010arXiv1008.3389K}, we find a 82.8\% probability (1.4-$\sigma$) that the system is truly grazing. However, this is lower than a 3-$\sigma$ value which would be a reasonable limit for detection of a truly grazing system. Hence, we conclude that WASP-34b has near-grazing transits. \begin{figure} \includegraphics[height=\columnwidth,angle=-90]{15992fig6.eps} \caption{Plot of the MCMC posteriors $b$ and ${R_{\rm P}}/{R_{\star}}$. The solid-line indicates the position of the stellar limb, i.e. $b+{R_{\rm P}}/{R_{\star}} = 1$. A total of 89.6\% of the points lie above the line and are grazing solutions.} \label{grazing} \end{figure} The eccentricity ($e$) and argument of periastron ($\omega$) are such that the occultation impact parameter will be \begin{equation} b_{\rm occ} = (0.953\pm0.031) \, b_{\rm tra} = 0.862\pm0.035. \end{equation} Hence, $b_{\rm occ}+{R_{\rm P}}/{R_{\star}} = 0.974\pm0.035$ with only 19.2\% of the MCMC posteriors greater than 1. Therefore, the occultations of the planet behind the star are expected to be total. However, the uncertainties in the orbital elements are such that is a possibility that the occultations could be partial. For transits with large $b$, the limb-darkening coefficients are extremely correlated to both $R_{\rm P}/R_{*}$ and $b$ \citep{2010arXiv1006.5680K}. In order to investigate the effects of limb-darkening on the inferred grazing nature of the system, we re-ran our MCMC analysis using limb-darkening coefficients for different filter bands (Table~\ref{limb}). Compared to the $r$-band results, there are small but noticeable, differences. Hence, this heuristic demonstrates that there is a $75\sim85$\% probability that the transit is actually grazing. \begin{table} \caption{Effect of limb-darkening on the derived impact parameter, $b$, and transit depth, $(R_{\rm P}/R_{*})^{2}$, for various filter bands. The grazing criterion, $b+{R_{\rm P}}/{R_{\star}}$, is given along with the percentage of MCMC posteriors that are grazing,\%($>$1), and the confidence values, P($>$1).} \centering \begin{tabular}{llllll}\hline\hline band & $b$ & $(R_{\rm P}/R_{*})^{2}$ & $b+{R_{\rm P}}/{R_{\star}}$ &\%($>$1) & P($>$1)\\ \hline $v$ & 0.898 & 0.01309 & 1.012$^{+0.017}_{-0.013}$ & 83.1 & 74.7 \\[+1mm] $r$ & 0.904 & 0.01261 & 1.016$^{+0.017}_{-0.014}$ & 89.6 & 82.8 \\[+1mm] $i$ & 0.905 & 0.01213 & 1.015$^{+0.016}_{-0.012}$ & 91.1 & 84.9 \\[+1mm] $z$ & 0.907 & 0.01190 & 1.016$^{+0.016}_{-0.013}$ & 90.3 & 83.8 \\[+1mm] \hline \end{tabular} \label{limb} \end{table} The near-grazing nature of the system makes it sensitive to additional planets, via changes to transit duration and shape. Given that there is evidence for another object in the WASP-34 system, further monitoring of the transits is required.

1012.2278

We report the discovery of WASP-34b, a sub-Jupiter-mass exoplanet transiting its 10.4-magnitude solar-type host star (1SWASP J110135.89-235138.4; TYC 6636-540-1) every 4.3177 days in a slightly eccentric orbit (e = 0.038±0.012). We find a planetary mass of 0.59±0.01 M<SUB>Jup</SUB> and radius of 1.22<SUB>-0.08</SUB><SUP>+0.11</SUP> R<SUB>Jup</SUB>. There is a linear trend in the radial velocities of 55±4 m s<SUP>-1</SUP> y<SUP>-1</SUP> indicating the presence of a long-period third body in the system with a mass ⪆0.45 M<SUB>Jup</SUB> at a distance of ⪆1.2 AU from the host star. This third-body is either a low-mass star, a white dwarf, or another planet. The transit depth ((R<SUB>P</SUB>/R<SUB>star</SUB>)<SUP>2</SUP> = 0.0126) and high impact parameter (b = 0.90) suggest that this could be the first known transiting exoplanet expected to undergo grazing transits, but with a confidence of only 80%. <P />Radial velocity and photometric data are only available in electronic form at the CDS via anonymous ftp to cdsarc.u-strasbg.fr (130.79.128.5) or via <A href="http://cdsarc.u-strasbg.fr/viz-bin/qcat?J/A+A/526/A130">http://cdsarc.u-strasbg.fr/viz-bin/qcat?J/A+A/526/A130</A>

12132311

2010arXiv1012.3649Y

1012

1012.3649_arXiv.txt

The fate of a star is determined by its mass, composition, rotation rate and binarity, the mass being the most important parameter. Massive stars with an initial mass between 8 and 140 M$_\odot$ will evolve through all burning stages and form an iron core. The fate of the star follows this sequence as a function of increasing initial mass: SNII--SNIb--SNIc--BH without SN. For initial masses between 140 and 260 M$_\odot$, the stars are thought to die as pair-instability supernovae if they do not lose much mass during their evolution, which is expected for low-metallicity non-rotating stars \cite{Heger02,Langer07}. PISNe are very bright and produce large quantities of iron \cite{Heger02}. Even if they are rare, they contribute significantly to the chemical evolution of galaxies. Their existence is debated because their chemical signature is not observed in extremely metal poor stars \cite{Umeda}. The absence of this chemical signature could be due to the fact that stars more massive than 140 M$_\odot$ do not form. This argument was supported by local observations. Indeed, the initial mass function (IMF) of the most massive young cluster in the Galaxy measured using HST seems to indicate a lack of stars with mass more massive than 150 M$_\odot$ \cite{Figer05}. However, observations of 2007bi \cite{GalYam} indicate that it is consistent with PISN theoretical models. In this work, we present evidence supporting the existence of stars more massive than 150 M$_\odot$. We also investigate the effect of rotation, mass loss and metallicity on the evolution of very massive stars. Main sequence evolutionary models are compared with the observations and we find an excellent agreement for the NGC 3603 and R136 systems \cite{pac10}. Finally, we discuss the fate of these very massive stars.

1012.3649

We recently determined the mass of the most massive star known to the date, R136a1 with a mass at birth 320 times the mass of our sun, as well as the mass of several other stars that are more massive than 150 M. Such massive stars (~150-300 M) may end their life as pair-instability supernovae (PISN) if they retain enough mass until they die. We have calculated a grid of stellar evolution models in order to investigate the impact of mass loss and rotation on the evolution and fate of these very massive stars. As mass loss is very strong at solar metallicity, our models predict that most of the very massive stars will die as type Ic SNe. Only slowly and non-rotating stars at metallicities below that of the LMC might retain enough mass to produce a PISN. This would mean that the first stellar generations might have produced PISN although their chemical signature is not observed in extremely metal poor stars in the halo of our galaxy.

12163443

2011A&A...531A.143D

1012

1012.4858_arXiv.txt

The projected axial rotational velocity of single stars can be measured directly from the broadening of their spectral lines. That is why the rotational velocity is an important observable in statistical studies of stellar astrophysics. Several methods have been developed to measure $v\sin i$, but the problem of systematic differences between different authors and methods has always been present. The first workable model for the stellar rotation was established by \citet {carroll28,carroll33} and \citet {caringra33}, but \citet {stru29} presented a very simple graphical model called the classical model of a rotating star (CMRS) by \citet{ct95}, which was used as a standard for most of the work done in the 20th century \citep[for more details see the excellent review paper by][]{col04}. \citet {stru29} did not consider limb darkening, which was introduced in the CMRS by \citet {carroll33}. In a series of papers Slettebak measured $v\sin i$ for stars in the spectral range O-G \citep{sle54,sle55,sle56,sleho55} by considering a linear limb darkening law. These stars were used in many subsequent papers by different authors to calibrate the full width at half depth of the line used to measure the $v\sin i$ parameter. Therefore, determination of $v\sin i$ was conditioned by the calibration made by the author. In the last quarter of the 20th century \citet{sl75} established a new standard system of stars distributed over both hemispheres and covering a range in spectral type from O9 to F9. The measure of $v\sin i$ was again based on a calibration of $v\sin i$ with the full width at half depth of stellar absorption lines in the star spectrum. \citet {sl75} established this system by comparing their digital data with numerical models constructed by adopting Roche geometry, uniform angular velocity at the surface, von Zeipel gravity darkening, and numerical integration of angle-dependent model atmosphere intensities. The standards for the high-velocity specimens were established by \citet{sle82} by visual comparison of line widths on the spectrograms that were recorded on photographic plates. Systematic effects between the old and new Slettebak systems were studied by \citet{gar84}. The calibrations of the systems described were made for only three lines, \ion{He}{i} 4471 {\AA}, \ion{Mg}{ii} 4481 {\AA}, and \ion{Fe}{i} 4476 {\AA}, and the theoretical profiles were calculated with main sequence models. Besides this, determination of the full width at half depth has shown to be very sensitive to the continuum position, which represents an important error source. A significant improvement was possible thanks to methods based on Fourier transform (FT) of line profiles providing a deeper analysis. In the first place, this tool is used to avoid an external calibration by suppressing the error related to this stage of the process. In second place, with high signal-to-noise ratio (S/N) it is possible to identify another broadening agents present in the line profile, at least in a qualitative way. Finally, it is possible to analyze second-order effects like differential rotation \citep{g77,g82,b81,rs02} and to put a limit on the inclination of the rotational axes \citep {rr04}. Regardless of the method adopted to determine $v\sin i$, it is not an easy task to find the spectral lines that have the minimum conditions required to be used in the measurement of rotational velocity, namely: 1) lines with no blends, 2) lines intense enough to be identified in rapid rotators, and 3) lines mainly broadened by rotation. For instance, in O-type and B-type stars, only the Balmer lines and some helium lines are intense enough, but they all show a significant Stark effect. In stars of spectral type later than A4, almost all lines are blended if $v\sin i$ is grater than 100 \kms, making it impossible to find an isolated line. This problem is evident in the work of \citet{roy02a} where fewer than three lines were measured in A-type stars with $v\sin i > 60$ \kms. Because this problem increases with the spectral type, other methods have been developed. One solution to the blending limitation is the use of a least square deconvolution procedure to derive the broadening function in a selected wavelength region instead of a single line profile. This methodology implies the application of an iterative process to fit the equivalent width of the template's spectral lines, the broadening function, and the continuum position \citep [for details see][]{rs03}. This is a very powerful technique for a detailed study of the rotational profile. However, for extensive applications like the development of a catalog of rotational velocities, a more direct method that does not involve fitting the intrinsic spectrum or including any atmospheric parameter other than limb darkening, might be more suitable. Even though CCF had been originally proposed for determining radial velocities, the projected axial rotational velocity can be inferred from the width of the CCF maximum. This tool has often been applied to determine $v\sin i$ in cool stars (M-L spectral type). The standard procedure is to calculate the CCF between an observed spectrum and a template spectrum of the same temperature with $v\sin i$ $=$ 0 \kms. Then, a fitting profile for the maximum of this function is calculated and the width of the fit can be used to measure $v\sin i$ through an empirical calibration $v\sin i$-width. The absence of single lines is somehow solved by means of the CCF . Nevertheless, as in any other method that depends on empirical calibrations, various strategies, not always equivalent, have been adopted in the literature. In some works the central maximum is fitted with a Gaussian \citep{bj04}, while in others a parabola \citep{tr98}, or even a Gaussian plus a quadratic function are used \citep{wb03}. The adoption of different functions to evaluate the rotational broadening could lead to systematic differences in results from different authors. Templates selection is also very heterogeneous. Even though some authors calculate synthetic template spectra, real stellar spectra have been used in most works based on the CCF \citep[see][]{wb03,mb03,bj04}. Since the CCF contains information from the template and the object spectrum, using an observed template could introduce an external error source from the unknown broadening factors present in the template spectral lines. Nowadays, the best methods dealing with stars that show intense line blending in their spectra require significant computational resources. Motivated by the need for a precise and expeditious technique to be applied extensively for the construction of a catalog of rotational velocities of bright A-type stars (Levato et al. in preparation), we develop here an alternative method based on the CCF and, at the same time, independent of any external calibration. Our procedure uses the FT to measure the parameter $v\sin i$, taking the dependence of the transform with limb darkening into account. In Sect.\,2 a full description of the methodology is presented. Limb darkening consideration and other practical details of the procedure are explained in Sect.\,3. Sect.\,4 describes specific application to A and later-type stars. The precision of the obtained rotational velocities is discussed in Sect.\,5, and our main conclusions are summarized in Sect.\, 6.

To evaluate the performance of our method in determining $v\sin i$, we compared our results with those from \citet{roy02a}. They used FT of line profiles to provide accurate $v\sin i$ of a large sample of A-type stars in the southern hemisphere observed with similar resolving power ($\sim$ 28000). Figure \ref{fig:6} shows Royer et al. values of $v\sin i$ and the ones obtained in this work for 155 objects in common with their sample. Even though a good agreement is found in stars of moderate, projected rotational velocities, a significant deviation is noticed for $v\sin i$ $>$ 150 \kms with an average difference of $5\pm 1$\%, our values larger than theirs. \begin{figure}[t] \centering \includegraphics[width=1\linewidth]{16386fg10.eps} \caption{Comparison between our results and \citet {roy02a} for objects in common. Squares: \textit{normal} stars. Circles: \textit{peculiar} stars. Solid line: \textit{Normal} stars fit $f({\rm \langle v\sin i\rangle})=1.43\;{\rm \langle v\sin i\rangle}^{0.92}$} \label{fig:6} \end{figure} There are several differences between the methodology used by \citet {roy02a} and the present work that can account for the effect present in Figure \ref{fig:6}. First, they measured single lines in the range from 4200 {\AA} to 4500 {\AA}, i.e. a region of 300 {\AA} wide which is the mean width of each region used in this work. Second, the blending effect becomes evident through the dependence of the number of measured lines with $v\sin i$. They used in average less than 3 lines on objects with $v\sin i$ $>$ 60 \kms, and over 90 \kms the average number of measured lines is $ \leq 2$. Even though the dependence on the position of the first zero of the Fourier transform with the limb darkening coefficient was not taken into account by Royer et al., the adopted value $\varepsilon = 0.6$ is typical of the limb darkening coefficient in the range of temperatures of the objects in the sample. However, an error of $\pm 0.01$ in $k_1$ represents an error of 1.5\% in $v\sin i$. Thus, using a wrong $k$ value could account for 1.5\% of the deviation detected in the comparison. Finally, the main source of error is the uncertainty on the continuum position. \citet {roy02a} analyzed the magnitude of this effects by broadening synthetic spectra of $T_\mathrm{eff}$ from 7500 to 10000\,K with $v\sin i$ = 10, 50, and 100 \kms. They found that in the coldest models with the largest $v\sin i$ the error can reach 3\%. This uncertainty affects spectral lines giving as a result a lower value of $v\sin i$ than the real value. To evaluate the intensity of this type of systematic error with our method, we measured $v\sin i$ in a synthetic A5V spectrum broadened by $v\sin i$ = 20, 80, 120, 180, 220, and 280 \kms. Random noise was added to simulate S/N = 100, generating ten spectra with different noise pattern for each value of $v\sin i$. Then, four regions were used to measure $v\sin i$, and the average over the ten measures for each region was calculated. As expected, we found a tendency to measure a lower value of $v\sin i$ than the actual value. This effect is shown in Figure \ref{fig:7}. Nevertheless, with the method proposed in this work the difference between measures and real values stays under 1\% ($\langle \frac{v\sin i_{measured}}{v\sin i_{real}}\rangle_{150-300 \; \mkms} = 0.997 \pm 0.004$) even for $v\sin i$ = 280 \kms. In conclusion, a considerable fraction of the difference between our results and those from \citet {roy02a} can be attributed to the effect of line blending on the continuum position. The influence of the line blending in high rotational velocity stars is significantly less with the proposed method than with Fourier transform of single-line profiles, by virtue of subtracting the secondary maximums from the CCF before calculating the Fourier transform. Similar results would be obtained by removing the blended lines with a similar procedure in the spectrum: modeling the smaller lines by convolving a rotational profile with a synthetic spectrum in which the line of interest has been subtracted. \begin{figure}[t] \centering \includegraphics[width=1\linewidth]{16386fg11.eps} \caption{ Difference between real $v\sin i$ and average over ten consecutive measures with distinct random noise. Dotted line: $f({\rm \langle v\sin i\rangle})=0.0032\;{\rm \langle v\sin i\rangle}+0.0456$.} \label{fig:7} \end{figure} The main originality of the proposed method is the way in which the broadening function is built from the CCF, which, on the one hand, allows measurement of stars with blended spectral lines and, on the other, improves the S/N. Once the broadening function has been obtained, we determine the stellar rotation through the first zero of the FT. However, other alternative analyses are possible, such as direct fitting of the broadening function or Fourier-Bessel transformation \citep{pgp96}. As an illustration, we used our method in combination with the Fourier-Bessel transformation to measure one of the objects in Piters et al. (1996) for which a HARPS spectrum was publicly available. These authors applied the Fourier-Bessel transformation method to individual spectral lines for measuring rotation in F-type stars. Owing to line blending, stars with $v\sin i > 100$ \kms\, could not be measured or had errors of about 10 -- 20 \%. We selected one of their fast rotating objects: HD\,12311, for which they measured 110$\pm$22 \kms. We defined 5 spectral region between 3980 \AA\,and 6530 \AA\, obtaining an average of 157.4 \kms with $\sigma=4.6$ \kms using our method with the FT first zero, and 155.9 \kms\,with $\sigma=1.8$ \kms from our broadening function but using the Fourier-Bessel transform. Therefore, the application of Fourier-Bessel transform method to the broadening function derived from the CCF would be an excellent alternative for late-type fast-rotating stars. In the case of slowly rotating stars, however, the Fourier-Bessel transform is more sensitive to broadening effects that are different from rotation, so the zero of the FT would give better results (see sect. 3.1).

1012.4858

<BR /> Aims: We propose a method for measuring the projected rotational velocity vsini with high precision even in spectra with blended lines. Though not automatic, our method is designed to be applied systematically to large numbers of objects without excessive computational requirement. <BR /> Methods: We calculated the cross correlation function (CCF) of the object spectrum against a zero-rotation template and used the Fourier transform (FT) of the CCF central maximum to measure the parameter vsini taking the limb darkening effect and its wavelength dependence into account. The procedure also improves the definition of the CCF base line, resulting in errors related to the continuum position under 1% even for vsini = 280 km s<SUP>-1</SUP>. Tests with high-resolution spectra of F-type stars indicate that an accuracy well below 1% can be attained even for spectra where most lines are blended. <BR /> Results: We have applied the method to measuring vsini in 251 A-type stars. For stars with vsini over 30 km s<SUP>-1</SUP> (2-3 times our spectra resolution), our measurement errors are below 2.5% with a typical value of 1%. We compare our results with Royer et al. (2002a) using 155 stars in common, finding systematic differences of about 5% for rapidly rotating stars.

12163224

2011A&A...527A..13P

1012

1012.1827_arXiv.txt

\label{sect:intro} Stellar jets from young stellar objects (YSO) play a key role in the star formation process and a wealth of information can be derived analysing their characteristic emission line spectra. The observed lines are collisionally excited in the shock waves generated by the interaction of the jet material with the interstellar medium or previously ejected matter and contain important information on the gas physics/kinematics. Different methods have been proposed to derive the physical conditions of the gas propagating in the jets from emission line ratios, such as the comparison with the predictions of a grid of shock models \citep[e.g., ][]{raga86,hartigan94}, or, alternatively, spectral diagnostics techniques \citep[e.g., ][hereafter referred to as the BE technique]{bacciotti99}. The application of these methods to analyse emission lines allows one to derive the jet structure on different scales, depending on the angular resolution of the observations: from parsec scales to hundreds of AU \citep{hartigan94,bacciotti99,podio06}, and down to $\sim$15 AU from the emitting source with space or Adaptive Optics-assisted observations \citep{lavalleyfouquet00,bacciotti00,woitas02,hartigan07,melnikov09}. In particular, the analysis of emission lines can shed light on the gas conditions at the base of the flow and on the mechanism that generates it. Proposed magneto hydro-dynamical (MHD) models suggest that the wind is accelerated and collimated by magneto-centrifugal forces. It is not clear, however, where the jet originates: from the star itself \citep{sauty94}, from the radius at which the stellar magnetosphere truncates the disk (X wind, \citealt{shu00}), or from an extended region of the disk (Disk wind, \citealt{konigl00}). In addition to the uncertainty on the launch mechanism, other observed properties still lack a proper modeling. For example, velocity resolved observations of the jets from T Tauri stars and younger Class 0/I protostars show that these can be highly asymmetric and can have different velocity components \citep{hamann94,hirth94,hirth97,davis01,garcialopez08,garcialopez10}. Despite some attempts to explain these observational features by means of the existing MHD models \citep[e.g., ][]{pesenti03,ferreira06} their origin is still unclear. A viable possibility to clarify these aspects is to investigate the variation of the jet physical properties in asymmetric jets and in different velocity channels, as it has been attempted in recent works conducted at moderate and high spatial/spectral resolution \citep{lavalleyfouquet00,woitas02,coffey08,garcialopez08,garcialopez10,melnikov09,podio09}. Another interesting issue, which remains not fully clarified, is the dust content of jets. Refractory species such as Ca, Ni, Cr, Fe, are often locked onto dust grains, thus a strong depletion of their gas phase abundance is expected in the interstellar medium (ISM) \citep{savage96}. A few studies have investigated the gas phase abundance of Fe and Ni in HH jets \citep{beck-winchatz94,beck-winchatz96,mouri00,bohm01,nisini02} and only recently the analysis has been extended to other refractory species such as Si, Ca, C, Cr, and Ti \citep{nisini05,nisini07,cabrit07,podio06,podio09,garcialopez08}. These studies showed that the refractory species may be depleted up to 90\% in HH jets and molecular outflows, thus suggesting the presence of dust grains in the ejected material. This in turn can put constraints on the region of the disk from where jets originate. In fact the stellar radiation destroys dust grains in the disk up to the so called dust evaporation radius, R$_{evp}$. This is located between 0.03 and 1.5 AU from the star in T Tauri and Herbig Ae/Be stars depending on the stellar luminosity and the dust properties \citep{isella05,akeson05,eisner07}, but no direct observations and/or modeling of the inner disk structure are available for younger Class 0/I sources, still surrounded by circumstellar matter. In general, we expect that a `dusty' jet can only arise from a region in the disk extended beyond R$_{evp}$. The passage through a shock front, however, can destroy part of the dust grains which may be transported in the jet, complicating the picture \citep[e.g., ][]{may00,guillet09}. Therefore, accurate studies of the gas-phase abundance of refractory species are necessary to determine the dust content in the jet and to constrain the shock efficiency in destroying the dust grains. To improve our understanding of the physics of HH jets from YSO, particularly their launch mechanism, we investigate the physical and kinematical structure of the HH 159 jet emitted by the young source DG Tau B, focusing on the analysis of the velocity components detected at its base, the asymmetry between the two jet lobes, and, finally, the estimate of the dust content in the jet. HH 159 is a bipolar jet, first detected by \citet{mundt83}. The red lobe consists of a chain of bright knots extending to $\sim$55\arcsec\, from the source, while the blue lobe, fainter and less collimated, is detected only up to $\sim$10\arcsec\, from the source \citep{mundt91,eisloffel98}. A large molecular outflow spatially coincident with the redshifted optical jet has been detected in the CO lines by \citet{mitchell94,mitchell97}. The driving star, DG Tau B, has been classified as a Class I source based on its spectral energy distribution \citep{watson04,luhman10} and at optical wavelengths it is obscured by circumstellar optically thick material detected in absorption through broadband imaging with the Hubble Space Telescope (HST) \citep{stapelfeldt97}. The observed dark lane, elongated and perpendicular to the jet, indicates the presence of a circumstellar disk, confirmed by $^{13}$CO millimeter observations \citep{padgett99,padgett00}, and flattened residual envelope. The source position has been determined from 3.5 cm VLA radio continuum observations by \citet{rodriguez95}, who locate it within the observed dark lane. Further HST studies in the near-infrared % indicated the presence of a bipolar reflection nebula, whose axis of symmetry is coincident with the axis of the optical jet \citep{padgett99}. The eastern lobe of the nebula, surrounding the blue jet lobe, is V-shaped, suggesting that the nebula traces the walls of the blueshifted outflow cavity. The western lobe is fainter and more collimated, and encompasses the redshifted optical jet and the CO outflow. Previous spectroscopic \citep{mundt87,eisloffel98} and imaging \citep{mundt91,stapelfeldt97,padgett99} studies highlighted a strong asymmetry between the two lobes in morphology, velocity and degree of excitation. These studies, however, focused only on the \Ha\, and \sii\, emission lines as well as on the continuum emission in the optical and near-infrared bands. On the contrary, in this paper we present for the first time very deep and high-spectral resolution spectroscopic observations of a large number of forbidden and permitted lines in the optical range from 5300~\ang\, to 8500~\ang. The high spectral resolution coupled with long integration times highlight a very complex velocity structure with multiple velocity components evolving along the jet. The large number of detected lines has allowed a detailed study of the gas physical conditions, through the application of spectral diagnostic techniques to selected line ratios. This has provided a rich information on the excitation and dynamics of the flow, as well as on the dust reprocessing. The comparison of our results with those obtained for other jets provides means of investigating the nature of the detected velocity components and the origin of asymmetric jets. This, in turn, can help understanding the generation of outflows and their role in young systems.

\label{sect:conclusions} In this paper we investigate the kinematical and physical structure of the HH 159 jet emitted by the young Class I source DG Tau B, by means of high spectral resolution observations acquired with KECK/HIRES. With respect to previous studies (e.g., \citealt{eisloffel98}) we couple the analysis of the jet kinematics with a detailed study of its physical and dynamical properties in each of its lobes and velocity components. Our method involves the analysis of selected line ratios through the so-called BE diagnostic technique \citep{bacciotti99} to infer the gas physical conditions (\en, \xe, \te, and \nh). Our main conclusions are as follows: \begin{itemize} \item[-] the velocity variations along the jet indicates that the ejection direction is changing due to precession and/or interaction with the surrounding medium. The electron and total density, as well as the ionisation fraction and the temperature, are decreasing along both lobes, due to gas dilution, i.e. to the jet propagation in a conical geometry, and/or to stronger shocks at the jet base, where the ejected material interacts with the dense parental cloud. \\ \item[-] the lines show a complex velocity and excitation structure, with multiple velocity components. The high velocity component, HVC, is extending more than the low velocity component, LVC, which is fading at 500-700 AU from the source, similarly to what previously observed for other jets from Class 0/I sources \citep{garcialopez08} and, on smaller scales, for jets from CTTSs \citep[e.g., ][]{hirth97}. The HVC appears to be denser and more excited suggesting that the LV gas is extracted from an outer region of the disk with respect to the HV gas and is thus accelerated at lower velocities.\\ \item[-] the jet show a strong asymmetry: the blue lobe is faster and more excited than the red lobe but less dense and collimated, suggesting that the interaction of the ejected material with the ambient medium is stronger on the red-lobe side. This asymmetry is similar to the one observed in other jets, and in particular in RW Aur \citep{melnikov09}.\\ \item[-] despite the observed asymmetries, the mass loss rate is similar in the two lobes (\mjet$\sim$6-8 10$^{-9}$ \msolyr), as in other asymmetric jets. This means that the the power transferred by the rotating disk is the same on both sides of the system, as expected if the jet is launched by a magneto-centrifugal engine. The observed asymmetries can be explained in the framework of a MHD disk wind, if the latter propagates in an inhomogeneous ambient medium \citep{ferreira06}.\\ \item[-] the redshifted molecular outflow detected in the CO millimeter lines \citep{mitchell94} supports the idea of an asymmetric ambient medium and a larger mass-load on the red-lobe side. The flux of linear momentum transported by the jet is comparable to that estimated for the molecular outflow, thus the latter can be jet-driven (\pjet$\sim$0.8-2.5 10$^{-6}$ \msolyr \kms).\\ \item[-] the depletion of Ca gas-phase abundance with respect to its solar abundance ([Ca]$_{gas}$/[Ca]$_{solar}$$\sim$0.01-0.15) indicates that most Ca atoms are locked onto dust grains. The minimum depletion is observed close to the source and in the HVC, where the strongest shocks are occurring, in agreement with the predictions from models of dust reprocessing in shocks \citep[e.g., ][]{guillet09}. The presence of dust in the jet may imply that the material is extracted from a region of the disk extending beyond the so-called dust evaporation radius, R$_{evp}$. However, more stringent constraints on the jet launching region can can only be derived by analysing higher angular resolution observations, to infer Ca gas-phase abundance in the first tens of AU from the source where the gas has not been reprocessed by shocks, and for more unembedded sources, for which the location of R$_{evp}$ can be derived from near-infrared interferometric observations \citep[e.g., ][]{akeson05}. \end{itemize}

1012.1827

Context. Jets from young stars can be highly asymmetric and have multiple velocity components. <BR /> Aims: To clarify the origin of jet asymmetries and constrain the launch mechanism, we study as a test case the physical and kinematical structure of the prototypical asymmetric flow emitted by DG Tau B. <BR /> Methods: The analysis of deep, high spectral resolution observations taken with the KECK telescope allows us to infer the properties and the spatial distribution of the velocity components in the two jet lobes. From selected line ratios we derive the gas physical conditions (the electron and total density, n<SUB>e</SUB> and n<SUB>H</SUB>, the ionisation fraction, x<SUB>e</SUB>, and the temperature, T<SUB>e</SUB>), as a function of the distance from the source and the gas velocity. The presence of dust grains in the jet is investigated by estimating the gas-phase abundance of calcium with respect to its solar value. <BR /> Results: The detected lines show broad velocity profiles at the base of the jet (up to ~100 km s<SUP>-1</SUP>), where up to three velocity components are detected. At 5''from the source, however, only the denser and more excited high-velocity components survive and the lines are narrower (~10-30 km s<SUP>-1</SUP>). The jet is strongly asymmetric in the velocity and in its physical structure. The red lobe, which is slower (~140 km s<SUP>-1</SUP>) and more collimated (opening angle: α ~ 3-4°), presents low ionisation fractions (x<SUB>e</SUB> ~ 0.1-0.4) and temperatures (T<SUB>e</SUB> &lt; 5 × 10<SUP>3</SUP> K), while the total density is up to ~2.5 × 10<SUP>4</SUP> cm<SUP>-3</SUP>. The blue lobe, faster (~-320 km s<SUP>-1</SUP>) and less collimated (α ~ 14°), is also less dense (n<SUB>H</SUB> &lt; 10<SUP>4</SUP> cm<SUP>-3</SUP>), but highly excited (T<SUB>e</SUB> up to ~5 × 10<SUP>4</SUP> K and x<SUB>e</SUB> up to 0.9). The estimated mass-loss rate turns out to be similar in the two lobes (~6-8 × 10<SUP>-9</SUP> M<SUB>⊙</SUB> yr<SUP>-1</SUP>), while the flux of the linear momentum is three times higher in the blue one (~2.5 × 10<SUP>-7</SUP> M<SUB>⊙</SUB> yr<SUP>-1</SUP> km s<SUP>-1</SUP>). Calcium is strongly depleted with respect to its solar abundance, indicating that the jet contains dust grains. The depletion is lower for higher velocities, which is consistent with dust destruction by shocks. <BR /> Conclusions: The similar mass-loss rate in the two lobes suggests that the ejection power is comparable on the two sides of the system, as expected from a magneto-centrifugal ejection mechanism, and that the observed asymmetries are caused by a different mass load and propagation properties in an inhomogeneous environment. The presence of dust grains implies that the jet is generated from a region of the disc extending beyond the dust sublimation radius.

12213337

2011MNRAS.413.1845T

1012

1012.5403_arXiv.txt

\label{introduction} Active galactic nuclei (AGNs) are powerful objects in the Universe, which are thought to be powered by gravitational energy through the accretion of surrounding gas onto supermassive black holes (SMBHs) located at the center of the accretion systems. Classically, the accretion paradigm has been applied to extremely powerful objects like quasars and radio galaxies for extragalactic objects. The Eddington ratio $L_{\rm bol}/L_{\rm Edd}$ of these objects, where $L_{\rm bol}$ and $L_{\rm Edd}$ are bolometric luminosity and Eddington luminosity, respectively, is $\sim 10\%$ and thereby geometrically thin, optically thick disk models, so-called standard disks \citep{Shakura1973AA24p337}, have been successfully favored. Recently, the idea that all the galaxies have SMBHs at their centers has been commonplace \citep{Kormendy1995ARAA33p581,1998AJ115p2285}. Optical spectroscopic surveys have revealed that a significant fraction of nearby galaxies has also active nuclei, but much less luminous than the powerful objects ($< 10^{40}$ erg$^{-1}$ in nuclear H$\alpha$ luminosiy), called low luminosity AGNs (LLAGNs) \citep{Ho1997ApJS112p315}. LLAGNs can be furthermore spectroscopically classified into Low Ionization Nuclear Emission Region nuclei (LINERs), Seyfert galaxies, and transition objects. Their Eddington ratios are much lower than the canonical value ($\sim 10\%$) appropriate for standard disks and could reach $\sim 10^{-8}$ in some cases (e.g., \citet{Ho2009ApJ699p626}). In addition to the low radiative efficiency, correlation among radio luminosity, X-ray luminosity, and the mass of nuclear black holes \citep{Merloni2003MNRAS345p1057,Falcke2004AA414p895} and the weak feature of the big blue bump \citep{Ho2008ARAA46p475} indicate optically thin, radiatively inefficient accretion flows (RIAFs) in LLAGNs. Advection dominated accretion flow (ADAF) models \citep{Narayan1994ApJ428L13,Narayan1995ApJ444p231,Abramowicz1995ApJ438p37}, are a kind of RIAFs and have been widely discussed to explain the spectral energy distribution (SED) of LLAGNs. ADAF models successfully explained most parts of the SED of Sagittarius A$^*$ which is a SMBH located at the center of our Galaxy \citep{Narayan1995Natur374p623,Manmoto1997ApJ489p791,Narayan1998ApJ492p554,Mahadevan1998Natur394p651}. ADAF models have been applied to nearby LLAGNs. In order to investigate the origin of radiation from LLAGN core, radio and X-ray bands are often adopted to avoid possible contamination from stellar and dust emission. Follow-up observations of optically selected LLAGNs in radio bands have detected emission dominated with a compact core morphology \citep{Ho2001ApJS133p77,Nagar2002AA392p53}. Despite their low luminosity, a larger fraction of LLAGNs are radio-loud \citep{Ho2001ApJ555p650,Terashima2003ApJ583p145}. \citet{Doi2005MNRAS363p692} detected a spectral bump in high-frequency radio ($\sim$ submillimeter) bands which is a spectral feature of an ADAF component. However, observations have revealed in radio bands (1-10 GHz) that ADAF components are not enough to reproduce the total radio emission, which indicate the contribution of jets \citep{Ulvestad2001ApJ562L133,Anderson2004ApJ603p42,Wu2005ApJ621p130}. These observations imply the coexistence of both disk and jet components in emission from LLAGNs. The respective contribution of disk and jet components from radio to X-ray bands is an open problem. \citet{Merloni2003MNRAS345p1057} found correlation among radio luminosity, X-ray luminosity, and black hole mass and showed that an ADAF model can reproduce observational data better than a radiation model dominated by jets by comparing the data with theoretical models. On the other hand, \citet{Falcke2004AA414p895} suggested that radiation is dominated by non-thermal emission from jets by a similar analysis. The correlation between radio and X-ray luminosities has also implied the common non-thermal origin from jets for radio-loud LLAGNs \citep{Balmaverde2006AA447p97,Panessa2007AA467p519}. Note that \citet{Merloni2003MNRAS345p1057} also mentioned that considering cooling processes of electrons can improve the reproducibility of the data in the jet model and synchrotron radiation from jets can be responsible up to X-rays. Furthermore, \citet{Nemmen2010IAUS267p313} demonstrated the SED fittings of observed LLAGNs and showed that observed X-ray data can be reproduced by both a disk-dominated and jet-dominated radiation models without inconsistency at present. This study suggests that $\gamma$-rays emitted from electrons accelerated in jets can be a direct diagnostic tool for a jet component in radio to X-ray bands. When X-rays are produced by synchrotron radiation of the electrons, the synchrotron photons are also upscattered by the electrons through inverse compton scattering (ICS) into $\gamma$-rays. This is so-called synchrotron self-compton (SSC) scenario \citep{Maraschi1992ApJ397L5,Bloom1996ApJ461p657}. This scenario is well established for the SED modelings of BL Lac objects (e.g., \citet{Anderhub2009ApJ705p1624,Aharonian2009ApJ696L150}). We focus on NGC 4278 as an example of LLAGNs and demonstrate expected $\gamma$-ray flux based on a SSC model on the assumption that non-thermal radiation from jets is dominated. We also survey parameter space in the model and constrain physical parameters required in jets. NGC 4278 is a LLAGN with $z = 0.002$ and the distance of 16.7 Mpc \citep{Tonry2001ApJ546p681} for $H_0 = 70$ km s$^{-1}$ Mpc$^{-1}$, which is Hubble constant. This source is sometimes categorized into LINER or radio-loud LLAGNs. Recent observations in radio bands have revealed a two-sided relativistic parsec scale jet ($\beta \sim 0.75$) closely aligned to the line of sight ($2^{\circ} \leq \theta \leq 4^{\circ}$), where $\beta$ and $\theta$ are the velocity in the unit of speed of light and the viewing angle of the jet, respectively \citep{Giroletti2005ApJ622p178}. These observables lead to the relativistic beaming factor of the parsec scale jet of $2.6$. In addition, a monthly time scale variability by a factor of 3 to 5 has been observed in X-ray bands \citep{Younes2010AA517p33}. In general, the time scale of flux variability limits the size of emission region. This paper is laid out as follows. In Section \ref{sec:model}, we describe a model of SED originating from electrons accelerated in relativistic jets. In Section \ref{sec:gamma}, we fit resultant SED to observed data below X-rays and estimate $\gamma$-ray flux. Moreover, combining observational results in radio and X-ray bands, we constrain the physical parameters of emission region. Finally, we make some discussions and summarize this study in Section \ref{sec:sum}.

\label{sec:sum} We have discussed possible $\gamma$-ray emission from jets in LLAGNs, and have also constrained several physical parameters in our SSC model. The size of an emission region (blob) was limited to $10^{16} \leq R \leq 10^{17.5}$ cm by a monthly time scale variability and synchrotron self-absorption. This range is consistent with the variability on typical time scale of a few days detected in LLAGNs \citep{Anderson2005ApJ627p674}. On the other hand, we have, so far, neglected a small ($\sim 10\%$) hours time scale ($\sim 1.5$ hr) variability in X-ray bands reported by \citet{Younes2010AA517p33}. Here, we discuss the origin of this variability. This indicates the size of an emission region as \begin{equation} R \leq \frac{\delta c \Delta t_{\rm obs}}{1 + z} = 5 \times 10^{14} \delta_3 \Delta t_{\rm obs,1.5h} ~~~{\rm cm}, \end{equation} where $\Delta t_{\rm obs,1.5h} = \Delta t_{\rm obs} / 1.5$ hour is the time scale of the flux variability. This size is comparable with the Schwarzschild radius of the central black hole for $\delta = 2.6$ and is $\sim 2 \times 10^{15}$ cm even for $\delta = 10$. These sizes are not favored due to the limitation of $R$ if the X-rays are emitted from the blob. Thus, it is thought that X-ray radiation from an innermost part of an accretion disk contributes to this variability. Theoretically, ADAF models can predict X-rays via thermal bremsstrahlung of ions in an inner hot part of the disk (e.g., \citet{Manmoto1997ApJ489p791},\citet{Nemmen2010IAUS267p313}). This implies that emission from the disk could contribute to the total X-rays by a small (but significant) fraction, even if a jet component is dominated from radio to X-ray bands suggested by the correlation between radio and X-ray luminosities. In this paper, a linear regime of synchrotron cooling of electrons, i.e., synchrotron cooling under constant magnetic field, was considered. Recently, non-linear radiative cooling of electrons has been discussed \citep{Schlickeiser2007AA476p1,Schlickeiser2008AA485p315}. Non-linear models assume a constant equipartition parameter $\eta^{-1}$ based on the success of spectral modeling of the observed blazars. Since the synchrotron and inverse compton cooling rates of electrons depend on the energy density of relativistic electrons, electron spectrum resultant from a solution of the kinetic equation of electrons is different from that derived from treatment in a linear regime even for a steady electron injection. Although whether equipartition is realized in jets of LLAGNs is not clear observationally at present, it is an intriguing topic to investigate common physical features between blazars (strong radio galaxies) and LLAGNs. To summarize, we have discussed $\gamma$-ray emission from LLAGN jets in the framework of a SSC model on the assumption of radio and X-rays are dominantly produced from jets. The $\gamma$-rays are a direct probe of a jet component in radio to X-ray bands without contamination from the other components, although the predicted flux is not large. Several observational results allowed us to constrain physical parameters in the emission region in jets. In the case of a beaming factor as low as that of parsec scale jets and $R \sim 10^{16}$ cm, CTA may detect the $\gamma$-rays in the near future and test the jet domination of radiation from LLAGNs. The determination of the respective contribution of disk and jet components will gives us a hint of a physical connection between a disk and relativistic jet in LLAGNs.

1012.5403

The respective contribution of disc and jet components to the total emission in low-luminosity active galactic nuclei (LLAGNs) is an open question. This paper suggests that γ-rays emitted from electrons accelerated in jets could be a direct diagnostic tool for a jet component to the total emission. We demonstrate γ-ray flux from jets based on a synchrotron self-Compton model on the assumption that radio and X-rays are dominantly produced from jets in the case of a high state of a nearby LLAGN, NGC 4278. We also survey parameter space in the model. Observational properties of LLAGNs in radio and X-ray bands allow to constrain physical parameters in an emission region. The size of the emission region R is limited to 10<SUP>16</SUP>≤R≤ 10<SUP>17.5</SUP> cm if the observed tight correlation between radio and X-ray emission originates from the same jet component. If the beaming factor of the emission region is close to the observed parsec scale jet of NGC 4278 and R∼ 10<SUP>16</SUP> cm, the γ-rays may be detected by the Cherenkov Telescope Array and the jet domination can be tested in the near future.

12213788

2011MNRAS.415.3462F

1012

1012.5129_arXiv.txt

\label{sec:intro} Extrasolar planetary research has been revitalised in the last decade and so far more than 500 extrasolar planets have been discovered. Improvements in the accuracy of RV measurements have made it possible to detect planets with larger orbital periods and smaller velocity amplitudes. With the flood of new data, more powerful statistical techniques are being developed and applied to extract as much information as possible. Traditionally, planet parameters and their uncertainties were obtained by searching for periodicity in the RV data using the Lomb-Scargle periodogram (\citealt{1976Ap&SS..39..447L, 1982ApJ...263..835S}) to fix the orbital period and then estimating other parameters by using minimisation algorithms. Recent advances in Marko-Chain Monte Carlo (MCMC) techniques (see e.g. \citealt{MacKay}) have made it possible for Bayesian techniques to be applied to extrasolar planetary searches (see e.g. \citealt{2005ApJ...631.1198G, 2005AJ....129.1706F, 2007ASPC..371..189F, 2009MNRAS.394.1936B}). Bayesian methods have several advantages over traditional methods, for example when the data do not cover a complete orbital phase of the planet. Bayesian inference also provides a rigorous way of performing model selection which is required to decide the number of planets favoured by the data. The main problem in applying such Bayesian model selection techniques is the computational cost involved in calculating the Bayesian evidence (see Sec.~\ref{sec:bayesian}). \citet{2007ASPC..371..224C} recently reviewed the state of techniques for model selection from a statistical perspective and \citet{2007ASPC..371..189F} evaluated the performance of a variety of marginal likelihood estimators in the extrasolar planet context. \citet{2007MNRAS.381.1607G} found good agreement (within 28\%) between three estimators: (a) parallel tempering, (b) the ratio estimator, and (c) Restricted Monte Carlo (RMC) for one and two planet models. However, for a 3 planet model the three estimators diverged significantly with the RMC yielding the lowest estimate. \citet{2010MNRAS.403..731G} introduced the Nested Restricted Monte Carlo (NMRC) estimator, an improvement on the RMC estimator. The NRMC estimator is expected to provide a conservative lower bound on the Bayesian evidence in higher dimensions. These Bayesian model selection techniques have already resulted in the discovery of previously unknown planets in existing data-sets, e.g. \cite{2009A&A...496L..13T} discovered a second planet orbiting HD 11506 and \cite{2010MNRAS.403..731G} reported a third planet orbiting 47 Ursae Majoris using Bayesian analysis. Nevertheless, most of the Bayesian model selection techniques employed so far in extrasolar planetary searches have relied on estimates of the Bayesian evidence, with uncertain accuracy. Our aim in this paper is to present a new and efficient method for Bayesian model selection to determine the number of planets favoured by the data, and estimate their parameters, {\em without} having to calculate directly the Bayesian evidence for models containing a large number of planets. The outline of this paper is as follows. We give a brief introduction to Bayesian inference in Sec.~\ref{sec:bayesian} and describe various Bayesian object detection techniques in Sec.~\ref{sec:object_detection}. Our model for calculating radial velocities is described in Sec.~\ref{sec:RV}. In Sec.~\ref{sec:bayes_RV} we describe our Bayesian analysis methodology including the descriptions of likelihood and prior probability functions. We apply our method to simulated data in in Sec.~\ref{sec:mock}, and to real RV data sets on HD 37124, 47 Ursae Majoris and HD 10180 in Sec.~\ref{sec:real}. Finally our conclusions are presented in Sec.~\ref{sec:conclusions}.

\label{sec:conclusions} We have presented a new and efficient method to detect extrasolar planets from RV measurements. Our method is not only able to fit for a specific number of planets, but can also infer the number of planets from the data using Bayesian model selection. We have successfully applied our method to simulated data-sets, as well as to the real systems HD 37124, 47 Ursae Majoris and HD 10180. Our method can potentially identify many undiscovered extrasolar planets in existing RV data-sets. One drawback of our method is that it ignores the planet-planet interactions, but these interactions are important only for a very small fraction of planetary systems. Moreover, our basic methodology can be extended to include such interactions. This will be undertaken in further work. Another important avenue of research in extrasolar planet searches is to perform a coherent analysis using different data-sets, e.g. by jointly analysing the RV data and light curves for the same system. This would enable us to place better constraints on the planetary parameters and also to learn about the physical structure of the planets. Once again our basic analysis technique can be easily extended to perform a joint analysis of data sets of different types We plan to extend our approach by incorporating light curve data in a forthcoming paper.

1012.5129

Stellar radial velocity (RV) measurements have proven to be a very successful method for detecting extrasolar planets. Analysing RV data to determine the parameters of the extrasolar planets is a significant statistical challenge owing to the presence of multiple planets and various degeneracies between orbital parameters. Determining the number of planets favoured by the observed data is an even more difficult task. Bayesian model selection provides a mathematically rigorous solution to this problem by calculating marginal posterior probabilities of models with different number of planets, but the use of this method in extrasolar planetary searches has been hampered by the computational cost of the evaluating Bayesian evidence. None the less, Bayesian model selection has the potential to improve the interpretation of existing observational data and possibly detect yet undiscovered planets. We present a new and efficient Bayesian method for determining the number of extrasolar planets, as well as for inferring their orbital parameters, without having to calculate directly the Bayesian evidence for models containing a large number of planets. Instead, we work iteratively and at each iteration obtain a conservative lower limit on the odds ratio for the inclusion of an additional planet into the model. We apply this method to simulated data sets containing one and two planets and successfully recover the correct number of planets and reliable constraints on the orbital parameters. We also apply our method to RV measurements of HD 37124, 47 Ursae Majoris and HD 10180. For HD 37124, we confirm that the current data strongly favour a three-planet system. We find strong evidence for the presence of a fourth planet in 47 Ursae Majoris, but its orbital period is suspiciously close to 1 yr, casting doubt on its validity. For HD 10180 we find strong evidence for a six-planet system.

12167733

2011ApJ...728...43F

1012

1012.0010_arXiv.txt

\label{sec:intro} The protoplanetary nebula stage (PPN) is one of the shortest of a sun-like star's evolution \citep*[e.g.,][]{iben_1983}. Throughout this phase, roughly half of the stellar photosphere is ejected, shocking the gas of the circumstellar envelope (CSE) formed in the AGB stage \citep*[AGB-CSE;][]{kwok_1978}, and unveiling the outermost layers of the nucleus. The stellar UV radiation field is extremely intense in this phase, photodissociating the innermost circumstellar gas and triggering a particularly rich photochemistry \citep*[][hereafter C04]{woods_2002,woods_2003,redman_2003,cernicharo_2004}. CRL618 (Westbrook Nebula, AFGL 618) is a very young PPN \citep*[age $\simeq 200$~yr;][]{kwok_1984}, located at a distance of $\simeq 0.9-1.8$~kpc from the Sun \citep{schmidt_1981,goodrich_1991,knapp_1993,sanchezcontreras_2004a}. It contains a B0 central star embedded in a dusty ultracompact HII region surrounded by a torus and a low-velocity expanding envelope with an external radius $>20$\arcsec, a total mass $\simeq 1$~M$_\odot$, and an expansion velocity $v_\subscript{exp}\simeq 18.0-21.5$~\kms{} \citep{knapp_1985,fuente_1998,pardo_2004}. In addition, it displays gas with velocities as high as 200~\kms{} \citep{burton_1986,cernicharo_1989,gammie_1989}. This high velocity gas (HVG) is the molecular counterpart of the bright optical jets oriented in the E-W direction \citep{trammell_2000} which impact the AGB-CSE and produce the well-known optical lobes. Since its discovery by \citet{westbrook_1975} many molecular species have been detected towards this PPN \citep{cernicharo_1989, cernicharo_2001a, cernicharo_2001b, herpin_2000,remijan_2005,truong-bach_1996}, some of them for the first time in a C-rich CSE \citep*[e.g., formaldehyde, polyacetylenes \diacet{} and \triacet{}, benzene, H$_2$O and OH;][]{cernicharo_1989,cernicharo_2001a, cernicharo_2001b,herpin_2000}. The CSE developed in the AGB phase has been studied in great detail by several authors \citep*[see, e.g.,][and references therein]{cernicharo_2001a, cernicharo_2001b, sanchezcontreras_2004a, sanchezcontreras_2004b,pardo_2004,pardo_2005,pardo_2007a,pardo_2007b}. \citet*[][hereafter W03]{woods_2003} and C04 have modelled the chemistry of CRL618 suggesting that in the innermost envelope the UV photons photodissociate \acet{} producing the radical C$_2$H, which can react with \acet{} to form \diacet{} or with H$_2$ reforming \acet{}. Additionally, C$_2$H can react with \diacet{} forming \triacet{}. These processes lead to a rapid \acet{} polymerization in long carbon chains and clusters. The abundance ratio between consecutive polyacetylenes (C$_{2n}$H$_2$, $n=1,2,3,\ldots$) in CRL618 has been estimated as a factor $\simeq 2-3$ \citep*[][hereafter C01a]{cernicharo_2001a}. Polyacetylenes are symmetric molecules without a permanent dipole moment and, hence, detectable only through their ro-vibrational spectra. The strongest bands of their spectra are expected to fall in the mid-infrared range due to the physical conditions prevailing in the innermost CSE where the polyacetylenes are built up. However, the large telluric opacity in the infrared has largely prevented the exploitation of this frequency range. This issue has been overcome with the launching of the Infrared Space Observatory (ISO) and the Spitzer Space Telescope, and the developing of instruments with high spectral resolving power such as the Texas Echelon-cross-Echelle Spectrograph \citep*[TEXES;][]{lacy_2002}. The analysis of these kinds of observations will allow us to improve our knowledge about the polymerization processes and the formation of complex molecules such as long carbon chains, PAHs, and fullerenes (C$_{60}$ and C$_{70}$), which have been recently observed towards circumstellar and interstellar environments \citep{cami_2010,garciahernandez_2010,sellgren_2010} and whose ions could be the carriers of the Diffuse Interstellar Bands \citep*[DIBs;][]{foing_1994}. In this paper, we present a high resolution mid-infrared spectrum of CRL618. We have identified $\simeq 170$ ro-vibrational lines of bands $\nu_6+\nu_8$, $\nu_6+\nu_8+\nu_9-\nu_9$, $\nu_6+\nu_8+2\nu_9-2\nu_9$, $\nu_6+\nu_8+\nu_7-\nu_7$ of \diacet{}, and $\nu_8+\nu_{11}$ of \triacet{}. In addition, we have observed several lines of \acet{} and \hcn{}. These lines have been analyzed by using a modified version of the model of AGB envelopes developed by \citet*[][hereafter F08]{fonfria_2008}. The observations and the spectroscopic laboratory data of \hcn{} and polyacetylenes required to analyze the observations are presented in \S\S\ref{sec:observations} and \ref{sec:spectroscopy}, respectively. A description of the spectrum and the model, the adopted fitting strategy, and a discussion about the uncertainties of the parameters can be found in \S\ref{sec:modeling}. The results derived from our fits are presented and discussed in \S\ref{sec:results} and, finally, summarized in \S\ref{sec:conclusions}.

\label{sec:conclusions} In this Paper, we present high-resolution mid-infrared observations towards the PPN CRL618. The sampled spectral range ($778-784$ and $1227-1249$~\cm) observed with the high resolving power spectrograph TEXES allow us to resolve bands $\nu_6+\nu_8$ and $\nu_8+\nu_{11}$ of \diacet{} and \triacet, respectively, in addition to several lines of bands $\nu_5$ and $\nu_2$ of \acet{} and \hcn. These rich data have enabled the modeling of the useful ro-vibrational line profiles of these molecular species to estimate their abundances and the physical conditions of the gas and the dust throughout the inner circumstellar envelope. The analysis of the observations has yield the following results among others: \begin{itemize} \item our results support the chemical model suggested by W03 and C04 for the polymerization of \acet; \item most of the \hcn{} and \acet{} are in the inner CSE. The rest come from several dense clumps located in the high velocity gas; \item \diacet{} and \triacet{} are formed in the innermost CSE. Their abundances seem to be negligible in the clumps since the emission from these molecules is undetectable in our spectrum; \item we are not able to detect any trace of the \tetracet{} band $\nu_{10}+\nu_{14}$, expected to fall in the observed range. This implies an even lower abundance for this species compared to previously suggested values. An upper limit to its column density has been estimated; \item there exist large differences between the excitation temperatures (vibrational and rotational) of \hcn, \acet, and \diacet{} which indicate that the inner CSE is out of LTE. \end{itemize} In addition, the results of this work demonstrate the power of IR observations in the determination of the abundances and physical conditions of the gas in complex structured environments such as the innermost envelopes of the evolved stars. Further improvements in the search for \tetracet{} will be made in the future by observing CRL618 and SMP LMC 11 at $\simeq 16~\mu$m with the Echelon-cross-Echelle Spectrograph (EXES) mounted on the Stratospheric Observatory for Infrared Astronomy (SOFIA).

1012.0010

We present a mid-infrared high spectral resolution spectrum of CRL618 in the frequency ranges 778-784 and 1227-1249 cm<SUP>-1</SUP> (8.01-8.15 and 12.75-12.85 μm) taken with the Texas Echelon-cross-Echelle Spectrograph (TEXES) and the Infrared Telescope Facility (IRTF). We have identified more than 170 rovibrational lines arising from C<SUB>2</SUB>H<SUB>2</SUB>, HCN, C<SUB>4</SUB>H<SUB>2</SUB>, and C<SUB>6</SUB>H<SUB>2</SUB>. We have found no unmistakable trace of C<SUB>8</SUB>H<SUB>2</SUB>. The line profiles display a complex structure suggesting the presence of polyacetylenes in several components of the circumstellar envelope (CSE). We derive total column densities of 2.5 × 10<SUP>17</SUP>, 3.1 × 10<SUP>17</SUP>, 2.1 × 10<SUP>17</SUP>, 9.3 × 10<SUP>16</SUP> cm<SUP>-2</SUP>, and lsim5 × 10<SUP>16</SUP> cm<SUP>-2</SUP> for HCN, C<SUB>2</SUB>H<SUB>2</SUB>, C<SUB>4</SUB>H<SUB>2</SUB>, C<SUB>6</SUB>H<SUB>2</SUB>, and C<SUB>8</SUB>H<SUB>2</SUB>, respectively. The observations indicate that both the rotational and vibrational temperatures in the innermost CSE depend on the molecule, varying from 100 to 350 K for the rotational temperatures and 100 to 500 K for the vibrational temperatures. Our results support a chemistry in the innermost CSE based on radical-neutral reactions triggered by the intense UV radiation field.

12132008

2010arXiv1012.1293G

1012

1012.1293_arXiv.txt

\label{intro} The standard model of cosmology emphasizes the role of the gravitational field of the dark matter in structure formation. In this picture, the dark matter determines the overall geometry and rate of growth of structure. While analyzing a cosmological simulation at redshift 5, we found that on scales of a few proper kiloparsecs intergalactic dark matter and baryons form qualitatively different structures \citep{harford08}. In contrast to the dark matter, which tends to occur in many small, quasi-spherical clumps, the baryons occur in a framework of thin, smooth rods that form backbones connecting the large, dark-matter dominated galaxies. These gaseous filaments likely correspond to the filaments described by others that mediate the ``cold mode'' mechanism of gas transport into galaxies \citep{birnboim03,binney04,katz03,keres05,dekel06,ocvirk08, keres09b,dekel09d,brooks09}. \bold{In contrast to the classical picture in which incoming gas is shock-heated to the virial temperature of the dark matter halo, gas accreted by this mechanism is never shock-heated, but rather moves directly to the center of the galaxy along relatively cool filaments. This mechanism predominates for \bolder{haloes} less massive than about $10^{11.5}\dim{M}_\odot$ (dark matter, gas, and stars combined). These \bolder{haloes} appear to be too small to sustain a shock \citep{birnboim03}. If the filaments we have studied are indeed conduits for such accretion, then their structure might have ramifications for understanding how galaxies gain the gas necessary for the varied star formation histories that are observed.} We reasoned that these contrasting distributions of gas and dark matter result from hydrodynamic effects: pressure forces retard the gas as it moves toward the axis of a filament, in contrast to the dark matter which can pass freely through subject only to gravitational forces. The gas would be expected to accumulate along the axis of the filament until the pressure forces are sufficient to counteract the gravitational field. In this paper we explore the hypothesis that the gas in the baryonic core forms a self-gravitating, isothermal cylinder in hydrostatic equilibrium, whose structure is determined primarily by the gravitational and hydrodynamic properties of the gas. Many different versions of an isothermal cylinder are possible with different degrees of concentration of the gas, but they all share the interesting property that the mass per unit length of the gas \bold{is finite and} depends only on the temperature and ionization state \citep{stod63,ostriker64}. We suggest that our findings may place constraints upon models for the movement of gas along filaments. The importance of hydrodynamic forces for the filament structure supports a hydrodynamic mechanism of origin from isothermal sheets as first proposed by \citet{schmidburgk67}. \bold{The plan of the paper is as follows: Section~\ref{ration} summarizes our previous work on intergalactic filaments, which provides the rationale for the present study. The important properties of the infinite, self-gravitating, isothermal cylinder are then presented, followed by the important features of our cosmological simulation. The results of our study are presented in Section~\ref{results}. In Section~\ref{methods} we first review the details of the simulation. Then the details of the procedures for selecting filaments, curve fitting, and the elimination of overlaps are presented. Finally, Section~\ref{discussion} summarizes our conclusions and discusses some broader implications.} \cmdhighway

\label{discussion} \bold{From a comological hydrodynamic simulation at redshift 5, we find that a plausible model for intergalactic filaments is an isothermal gas cylinder whose structure and stability are determined primarily by the gravitational and hydrodynamic properties of the gas. The cylinders have a central gas density of several hundred times the mean total cosmic density, with a peak at about 500. The average temperature of the gas in the cylinders is 1-2 times $10^{4}\dim{K}$. The neutral hydrogen fraction is generally between 0.01 and 0.02.} \bold{The box size of $8h^{-1} \unit{Mpc}$ comoving limits the findings to \bolder{haloes} less than $10^{11}\dim{M}_\odot$ in total mass (gas, dark matter, and stars).} Our findings fit well into the emerging picture of gas transport into galaxies. Except for \bolder{haloes} larger than a few times $10^{11}\dim{M}_\odot$ at low redshifts, gas is believed to enter galaxies primarily through intergalactic filaments at temperatures well below the virial temperature of the galaxy and to never be shock-heated \citep{birnboim03,katz03,keres05,dekel06,ocvirk08, keres09b,dekel09d,brooks09}. The temperatures of our filaments are below a few times $10^{4}\dim{K}$, mostly well below the estimated virial temperatures of the galaxies, which range from about $4\times10^{5}\dim{K}$ down to about $3\times10^{4}\dim{K}$\footnote{The virial temperatures were computed using a spherically symmetric mass profile within the virial radius of the galaxy.}. \bold{The intergalactic medium is thought to be filamentous, and better knowledge of this texture may help to understand the spectra of high redshift quasars. The Lyman alpha forest is the main observational probe of this tenuous medium.} Many of the galaxies in the simulation, particularly small ones, do not have recognizable intergalactic filaments attached to them. For these galaxies the diffuse, hot, ionized gas around them may have few mechanisms to enter. Thus the presence or absence of filaments may effectively divide galaxies into two categories, those that can efficiently accrete gas and form stars and those that cannot. If the baryon-rich cores require a minimum amount of gas for the stability of an isothermal cylinder, it may be that the stellar content of a galaxy is a good indication for overall cosmic density. The importance of the gas environment for accretion has been emphasized by \citet{keres05}. The dearth of satellite galaxies around the Milky Way is sometimes cited as a problem for the currently favored $\Lambda$CDM cosmology, which would predict many more satellite dark haloes. An absence of filaments at an earlier stage in cosmic history may have prevented these haloes from becoming luminous today. Whatever the explanation, it is worth noting that the simulation predicts a clear separation between gas and dark matter at the spatial dimensions of small galaxies. We would argue, then, that an assessment of the number of galaxies that form stars requires a simulation that at a minimum includes some gas hydrodynamics. We invariably see our filaments embedded within thin sheets of gas. This observation suggests that the filaments might emerge from the sheets. \citet{schmidburgk67} has shown that an infinite isothermal cylinder is but one extreme of a series of isothermally balanced structures that range from the infinite isothermal sheet, through intermediate structures containing regularly spaced, parallel, embedded filaments with elliptical cross-sections, to the other extreme of isolated cylinders. A hydrodynamic origin of the filaments may help to explain a general difference between the collapse of dark matter and gas in our simulation. The dark matter can collapse only so far into filaments before fragmenting and collapsing further into spheroidal structures. The gas on the other hand can collapse into denser filaments which are stable to small perturbations.

1012.1293

Using a cosmological simulation at redshift 5, we find that the baryon-rich cores of intergalactic filaments extending outward from galaxies commonly form isothermal gas cylinders in regions favorable to their formation. The central gas density is typically about 500 times the cosmic mean total density, and the temperature is typically 1-2 times 10^4 K, just above the Lyman alpha cooling floor. These findings argue that the hydrodynamic properties of the gas are more important than the dark matter in determining this structure. It is noteworthy that the temperature and ionization state of the gas completely determine a finite total mass per unit length of an isothermal cylinder. Our findings may have implications for understanding the "cold mode" mechanism of gas transport into galaxies.

12231121

2011RAA....11..537X

1012

1012.2918_arXiv.txt

% \label{sect:intro} When a supernova explodes near the molecular clouds (MCs), shock generated by supernova remnant (SNR) can accelerate, compress, heat or even fragment surrounding interstellar MCs. They also can enhance abundances with respect to quiescent cloud conditions of different molecular species. So the SNR-MCs interaction plays a very important role in the evolution of Interstellar Medium (Reynoso et al. 2000). Moreover, the molecular lines observations of MCs adjacent to SNRs can shed light on the SNRs' dynamical evolution and physical properties. So far about half of the Galactic SNRs are expected to be in physical associated with MCs (Reynoso $\&$ Mangum 2001). The association is indicated by morphological correspondence of molecular emission, presence of molecular line broadening within the extent of SNR, and presence of line emission with high high-to-low excitation line intensity ratio, etc (see Jiang et al. 2010). However, only few cases about the SNR-MCs interaction has been well confirmed. Even for some SNRs interacted with the surrounding MCs, the detailed distribution of environmental molecular gas is poorly known. Tycho's SNR is known as a young type Ia SNR located in the Perseus arm (Albinson et al. 1986), with an age of about 400 years. The distance of Tycho's SNR is estimated to be 2.3 kpc (Kamper et al. 1978; Albinson et al. 1986; Strom 1988). The remnant has been widely studied in various wavebands. From the neutral hydrogen HI observation, Schwarz et al. (1995) concluded that the distance should be 4.6$\pm$0.5 kpc; Reynoso et al. (1999) have performed an HI absorption study toward the remnant, suggesting that Tycho's SNR is currently interacting with a dense HI concentration in the northeast ($V_{LSR}$=-51.5 km $s^{-1}$), which locally slows down the expansion of the shock front. In radio continuum (Reynoso et al. 1997) and high-resolution X-ray wavebands (Hwang et al. 2002), they indicated that the radio continuum and X-ray emission along the northeastern edge of Tycho's SNR is strongest. Based on the VLA observation, Dickel et al. (1991) indicates that the northeastern edge of Tycho's SNR is expanding with a very small and decreasing velocity into the dense interstellar medium, and may interact with it. Lee $\&$ Koo (2004) made the CO $J=1-0$ line observation towards Tycho's SNR. They concluded that most of the CO $J=1-0$ emission around Tycho's SNR is between -67 and -60 km $s^{-1}$, but the velocity component in intervals -63.5 $\sim$ -61.5 km $s^{-1}$ appears to be in contact with the Tycho's SNR along its the northeast boundary. Using CO $J=1-0$ line to study the environs of Tycho's SNR, Cai et al. (2009) found that the CO $J=1-0$ emission form a semi-closed molecular shell around the SNR, the emission in velocity range of -69 $\sim$ -58 km $s^{-1}$ is associated with the SNR. Because the calculated virial mass of clumps is greater than their gravitational mass, they suggested that CO molecular clumps are being violently disturbed by Tycho's SNR shock. In order to understand the evolution of SNR interacting with MCs and investigate the detailed distribution of the molecular gas around Tycho's SNR, we have performed CO $J=2-1 $ and CO $J=3-2$ observation toward Tycho' SNR. The observations cover for the first time the whole area of Tycho's SNR in these frequencies. Due to the observed molecular lines at higher frequencies, we can attain higher angular resolution, which is critical to identify relatively compact core. Also, higher $J$ transitions are relatively more sensitive to hot gas. Thus, we can detailedly understand the distribution of the molecular gas in the vicinity of Tycho's SNR.

\label{sect:discussion} Tycho's SNR has complete radio shell on the whole and bright radio knots in southwest. Among the MCs surrounding Tycho's SNR, we have seen that three clouds (cloud A, cloud B and cloud C) are spatially coincident with the SNR. Broad emission lines detected in cloud A and B and bow-shaped morphology suggest that the MCs may be interacting with the SNR. After a careful inspection of the CO component, we find that CO component of cloud A in intervals -67 $\sim$ -59 km $\rm s^{-1}$ (peaked at -63.7 km $\rm s^{-1}$) and CO component of cloud B in intervals -69 $\sim$ -64 km $\rm s^{-1}$ (-66.2 km $\rm s^{-1}$) and -67 $\sim$ -59 km $\rm s^{-1}$ (-63.5 km $\rm s^{-1}$) are well associated with Tycho's SNR. The position-velocity diagram across the peak of cloud B along the north-south and the east-west direction, respectively, are constructed from the CO $J=2-1$ line (see Fig.4). In Figure 4, obviously, the velocity components of cloud A and cloud B are adjacent. For cloud A, CO component in intervals -67 $\sim$ -59 km $\rm s^{-1}$ may belong to cloud B, because cloud A and cloud B are adjacent in space and in velocity, and have the approximately same peak velocity in spectra. We concluded that cloud A and cloud B may not only simply be overlapping along the line of the sight, but also colliding each other. Furthermore, the new cloud C is also associated well with bright knots, the bright knots may be radio emissivity increase due to compression of the shocked material, suggesting that cloud C appears to be swept up by the SNR shock in the southwest area. \begin{figure} \includegraphics[width=48mm,angle=-90]{f4.1.ps} \includegraphics[width=48mm,angle=-90]{f4.2.ps} \vspace{0mm}\caption{\small P-V diagram constructed from the CO $J=2-1$ transition for Cloud B. Left panel: Contour levels are from 0.3 K to 4.4 K by the 0.1 K, with a cut along the north-south direction. Right panel: Contour levels are the same as in left panel, with a cut along the east-west direction. } \label{Fig3} \end{figure} The maximum value of integrated CO line intensity ratio for the three MCs is 0.66 $\pm$ 0.10. The comparison with previously published observations reveals that the $I_{\rm CO}$$_ {J=3-2}$/$I_{\rm CO}$$_ {J=2-1}$ for the MCs associated with Tycho's SNR agrees well with the value measured in in the Milky Way (0.55 $\pm$ 0.08; Sanders et al. 1993), in NGC 253 (0.5 $\pm$ 0.1 in the disk; Wall et al. 1991), and in M33 (0.69 $\pm$ 0.15; Wilson et al. 1997). It also appear near value measured in the starburst galaxies M82 (0.8 $\pm$ 0.2; G\"{u}sten et al. 1993). The high $I_{\rm CO}$$_ {J=3-2}$/$I_{\rm CO}$$_ {J=2-1}$ in starburst galaxies may be due to unusual conditions in these dense and hot regions (Aalto et al. 1997), while for normal MCs the most likely explanation is a significant contribution to the CO emission by low column density material (Wilson $\&$ Walker 1994). Moreover, the high $I_{\rm CO}$$_ {J=3-2}$/$I_{\rm CO}$$_ {J=2-1}$ value (3.4) in the MCs interacting with SNR IC443 (Xu et al. 2010) exceed previous measurement of individual Galactic MCs, implying that the SNR shock has driven into the MCs. For the MCs around Tycho's SNR, the $I_{\rm CO}$$_ {J=3-2}$/$I_{\rm CO}$$_ {J=2-1}$ value (0.66 $\pm$ 0.10 ) may indicate that the Tycho's SNR shock just drove into the surrounding MCs. In addition, the MCs associated with HII regions have higher ratio of $I_{\rm CO}$$_ {J=3-2}$/$I_{\rm CO}$$_ {J=2-1}$ and have higher gas temperature than those sources without HII regions, indicating that the high line ratio may be due to heating of the gas by the massive stars (Wilson et al. 1997). Different masers maybe occur in different astrophysical environment. The $\rm H_{2}O$ masers are located near the MSX sources and within the maximum line intensity ratio $I_{\rm CO}$$_ {J=3-2}$/$I_{\rm CO}$$_ {J=2-1}$ regions, suggesting that $\rm H_{2}O$ masers occur in relatively warm environments (Qin et al. 2008). Thus, we concluded that the line intensity ratios based on the optically thick CO transitions indicate the temperature varies at different positions. The MCs associated with Tycho' SNR have a ratio value gradient increasing from southwest to northeast, may imply that a shock has driven into MCs and compressed MCs. Since the MCs are compressed, it lead to the temperature of molecular gas increases. Hence the line intensity ratio $I_{\rm CO}$$_ {J=3-2}$/$I_{\rm CO}$$_ {J=2-1}$ will increase. We consider that high $I_{\rm CO}$$_ {J=3-2}$/$I_{\rm CO}$$_ {J=2-1}$ is also identified as one good signature of SNR-MCs interaction system. From Table 1, the total mass of MCs is $\sim 2.13\times10^{3}$.

1012.2918

We have carried out CO J = 2 - 1 and CO J = 3 - 2 observations toward Tycho's supernova remnant (SNR) using the KOSMA 3m-telescope. From these observations, we identified three molecular clouds (MCs) around the SNR. The small cloud in the southwest was discovered for the first time. In the north and east, two MCs (Cloud A and Cloud B) adjacent in space display a bow-shaped morphology, and have broad emission lines, which provide some direct evidences of the SNR-MCs interaction. The MCs are revealed at -69∼ -59 km s<SUP>-1</SUP>, coincident with Tycho's SNR. The MCs associated with Tycho's SNR have a mass of ∼ 2.13 × 10<SUP>3</SUP> M<SUB>⊙</SUB>. Position-velocity diagrams show the two clouds to be adjacent in velocity, which means cloud-cloud collision could occur in this region. The maximum value (0.66 ± 0.10) of the integrated CO line intensity ratio (I<SUB>CO J=3-2</SUB>/I<SUB>CO J=2-1</SUB>) for the three MCs agrees well with the previous measurement of individual Galactic MCs, implying that the SNR shock drove into the MCs. The two MCs have a line intensity ratio gradient. The distribution of the ratio appears to indicate that the shock propagates from the southwest to the northeast.

12132133

2010arXiv1012.2773K

1012

1012.2773_arXiv.txt

1012.2773

Is the scaling, \lambda \sim Rm^{1/2}, for the growth rate of small-scale dynamo instability at low magnetic Prandtl numbers and large magnetic Reynolds numbers, Rm, valid in the vicinity of the threshold? Our analysis and even numerical solution (Malyshkin and Boldyrev, 2010) of the dynamo equations for a Gaussian white-noise velocity field (the Kazantsev-Kraichnan model) imply that the answer is negative. Contrary to the claim by Malyshkin and Boldyrev (2010), there are two different asymptotics for the dynamo growth rate: in the vicinity of the threshold and far from the threshold.

950512

2011A&A...526L...1C

1012

1012.0281_arXiv.txt

\label{intro:sec} Most of the stellar mass in low-mass YSOs is assembled within the first 10$^5$\,yr of their evolution~\citep[i.\,e. class\,0, 10$^4$\,yr, and Class\,I, 10$^5$\,yr: see e.\,g.,][]{lada84,andre93}. During this stage, the YSO luminosity is thus expected to be dominated by accretion. However, several studies, including the latest \textit{Spitzer Space Telescope} surveys~\citep[e.\,g.,][]{enoch,evans09}, have found that more than 50\% of the embedded YSOs have $L_{bol}$ and $\dot{M}_{acc}$ values considerably lower than those theoretically predicted~\citep[i.\,e. $\sim$2$\times$10$^{-6}$\,M$_\odot$\,yr$^{-1}$ for solar-mass YSOs;][]{shu,terebey} and roughly of the same order of magnitude as the Classical T Tauri stars~\citep[CTTs; i.\,e. 10$^{-8} \le \dot{M}_{acc} \le$10$^{-6}$\,M$_\odot$\,yr$^{-1}$; e.\,g.,][]{white}. Among several hypotheses, a likely explanation is that the mass accretion is \textit{episodic}, and the protostars with the lowest luminosities are those observed in quiescent accretion states~\citep{enoch,evans09,vorobyov}. Non-steady mass accretion is often observed in CTTs, such as EXors and FUors (lasting from a few months to several decades), in which $\dot{M}_{acc}$ increases by several orders of magnitude up to $\dot{M}_{acc}$$\sim$10$^{-4}$\,M$_\odot$\,yr$^{-1}$~\citep{hart}. It is thus reasonable to believe that a similar mechanism also exists in earlier and more embedded YSOs. Unfortunately, there is little direct observational evidence of outbursts in Class\,I YSOs, and so far, only a few clear cases have been detected~\citep[e.\,g. \object{V\,1647\,Ori} outbursts in 2003 and 2008, or \object{OO Ser} in 1995; see e.\,g.][]{fedele,kospal}. To reconcile theory with observations and improve the quality of the statistics, it is mandatory to detect and study these rare events. With this aim, we started a long-term project to monitor the NIR flux and spectroscopic variability of embedded YSOs (mostly Class\,I and Flat sources) in nearby, young, and active star-forming regions (namely L\,1641, CrA, and the Serpens Molecular Cloud). This letter reports the outburst of an embedded YSO in L\,1641, namely \object{2MASS-J05424848-0816347}, hereafter \object{[CTF93]216-2} ($\alpha_{2000}$=05$^h$42$^m$48.48$^s$, $\delta_{2000}$=-08$\degr$16$\arcmin$34$\farcs$7). This object was identified by our group as a low-mass embedded YSO (spectral type M, circumstellar $A_\mathrm{V}$$\sim$18\,mag) with a flat spectral index ($\alpha$=0.25, derived by fitting all the photometric data points from 2.2 to 24\,$\mu$m) and a bolometric luminosity of $\sim$1.9\,L$_{\sun}$ (Caratti o Garatti et al. in prep., hereafter CoG). It has been named \object{[CTF93]216-2}, because it is relatively close to \object{[CTF93]216}~\citep[][]{chen93,chen94}, located about 38$\farcs$5 SW. Our Spitzer/IRAC images indicate that both YSOs have precessing jets, thus they could be part of a wide binary system ($\sim$17\,300\,AU, assuming a distance $d$=450\,pc). During our recent survey in October 2010 with the robotic telescope REM (see Section\,\ref{observations:sec}), we detected for [CTF93]216-2 a brightness increase of several magnitudes with respect to the 2MASS images. We then compared our new images with the acquisition image in the $K_s$ band and the NIR spectrum of this source acquired in February 2009, discovering that the outburst was already in progress.

\label{discussion:sec} Our discovery of the [CTF93]216-2 outburst gives us a rare opportunity to study boosted accretion in young embedded protostars, probing whether episodic mass accretion can reconcile the low accretion rates observed in young embedded protostars with the theoretical predictions. The Br$\gamma$ luminosity, corrected for the circumstellar extinction, is often used to derive an estimate of the accretion luminosity~\citep[see e.\,g.][]{muzerolle98,natta06}. Assuming that $A_\mathrm{V}$=18$\pm$3\,mag and $d$=450\,pc, we derive $L(Br\gamma)$$\sim$3.5(0.9)$\times$10$^{-3}$\,L$_{\sun}$ and, from the \citet{muzerolle98} relationship, we obtain $L_{acc}$$\sim$22$\pm$8\,L$_{\sun}$, which is close to the derived $\Delta L_{bol}$$\sim$20\,L$_{\sun}$ and indicates that $L_{acc} \sim \Delta L_{bol}$. On the other hand, $\dot{M}_{acc}$ can be derived from $\dot{M}_{acc} = L_{acc} \times 1.25 R_* / G M_*$~\citep[][]{gullbring}, where $M_*$ and $R_*$ are the stellar mass and radius, and the number 1.25 is derived by assuming a value of 5$R_*$ for the inner radius of the accretion disk. Assuming that $T_\mathrm{eff}$=3200$\pm$200\,K and $L_{*}$$\sim$1.9$\pm$0.1\,L$_{\sun}$, from \citet{siess} models, we obtain $M_*$=0.24$\pm$0.04\,M$_{\sun}$ and $R_*$=4.4$\pm$0.4\,R$_{\sun}$, which gives a mass accretion rate of $\sim$1.1-1.3$\times$10$^{-6}$\,M$_{\sun}$\,yr$^{-1}$. While $M_*$ is well constrained by $T_\mathrm{eff}$, the radius $R_*$ depends on $L_{*}$. Thus this last should be considered as an upper limit, since pre-outburst $L_{acc}$ is unknown. However, even a 50\% decrease in $L_{*}$ would affect $R_*$ and thus $\dot{M}_{acc}$ of just 20\%. Compared to typical accretion rates of Class\,I YSOs with similar masses~(i.\,e. 10$^{-8}$\,M$_{\sun}$\,yr$^{-1}$; e.\,g., CoG; \citealt{scholz10,white}), the derived $\dot{M}_{acc}$ is about two orders of magnitude higher. This value is about an order of magnitude lower than what some episodic-accretion models predict for these early YSO bursts~\citep[i\,.e 10$^{-5}$\,M$_{\sun}$\,yr$^{-1}$; e.\,g.][]{vorobyov}, which should resemble the more evolved FUor counterparts. On the other hand, we note that the magnitude of luminosity change during this outburst ($\sim$10) is similar to those of \object{V\,1647\,Ori} and \object{OO Ser} ($\sim$8). Our object would probably differ from the others in its pre-outburst $L_{bol}$, which is about two times lower, because of its lower mass. Moreover, as already noted in previous cases~\citep[e.\,g.][]{hodapp,fedele,kospal}, the [CTF93]216-2 outburst shows similarities with both FUor and EXor events, i.\,e. a NIR featureless FUor spectrum with strong absorption CO band heads, but with Br$\gamma$ line in emission as in EXors. The event duration (if confirmed by additional observations) and amplitude are in-between those of EXors and FUors, thus, as noted by other authors~\citep{gibb,fedele} EXors and FUors might not be distinct categories of eruptive events, but instead part of a continuum of outburst events.

1012.0281

Context. Strong outbursts in very young and embedded protostars are rare and not yet fully understood. They are believed to originate from an increase in the mass accretion rate (dot{M}_acc) onto the source. <BR /> Aims: We report the discovery of a strong outburst in a low-mass embedded young stellar object (YSO), namely 2MASS-J05424848-0816347 or [CTF93]216-2, as well as its photometric and spectroscopic follow-up. <BR /> Methods: Using near- to mid-IR photometry and NIR low-resolution spectroscopy, we monitor the outburst, deriving its magnitude, duration, as well as the enhanced accretion luminosity and mass accretion rate. <BR /> Results: [CTF93]216-2 increased in brightness by 4.6, 4.0, 3.8, and 1.9 mag in the J, H, K<SUB>s</SUB> bands and at 24 μm, respectively, corresponding to an L<SUB>bol</SUB> increase of 20 L<SUB>sun</SUB>. Its early spectrum, probably taken soon after the outburst, displays a steep almost featureless continuum, with strong CO band heads and H<SUB>2</SUB>O broad-band absorption features, and Brγ line in emission. A later spectrum reveals more absorption features, allowing us to estimate T<SUB>eff</SUB> 3200 K, M_* 0.25 M<SUB>sun</SUB>, and dot{M}_acc 1.2 × 10<SUP>-6</SUP> M<SUB>sun</SUB> yr<SUP>-1</SUP>. This makes it one of the lowest mass YSOs with a strong outburst so far discovered. <P />Based on observations collected at the ESO/NTT (082.C-0264), at the REM telescope La Silla, Chile, and at the the Italian Telescopio Nazionale Galileo (TNG), operated on the island of La Palma by the Fundacion Galileo Galilei of the INAF (Istituto Nazionale di Astrofisica).

2529935

2011ApJ...729...19K

1012

1012.0762_arXiv.txt

With evidence of supermassive black holes (SMBH) lurking at the center of nearly all galaxies \citep{Richstone98}, it is pertinent to examine and understand their properties as well as their impacts. Studies have shown a critical relation between SMBHs and their host galaxies in the form of the $M$--$L$ and $M$--$\sigma$ relations. The $M$--$L$ relation between the mass of the SMBH and the luminosity of the bulge suggests an intrinsic link between the SMBH and the amount of mass in the bulge assuming a particular mass-to-light ratio \citep[e.g.][]{Magorrian98,Kormendy95,Gultekin09b} The $M$--$\sigma$ relation between SMBH mass and the velocity dispersion of the host galaxy also implies a physical coupling between formation and growth of the black hole and its surroundings \citep[e.g][]{Ferrarese00,Gebhardt00}. The driving mechanism behind these couplings and the $M$--$\sigma$ relation especially, is thought to be the result of mergers that drive accretion onto the SMBH, which can quench star formation as energy released from the central engine drives the gas out \citep{DiMatteo05}. In particular, the study of an accretion disk around a SMBH begins with observations of the extended blackbody spectrum emitted by the disk; a consequence of the radial dependence of the temperature associated with the accretion disk. In SMBH accretion disks, this spectrum is thought to peak in the UV and is associated with the ``Big Blue Bump" \citep{Elvis94}. Unfortunately, UV flux is extremely susceptible to scattering by dust and modeling to correct for this can induce large uncertainties. In X-rays, emission from inverse-Compton scattering, magnetic flares and magnetic reconnection events associated with the accretion disk are characterized well by a non-thermal power-law \citep[e.g.][]{McHardy04}. Accordingly, X-ray flux can be another proxy for observing accretion disks. In accreting systems, as material migrates toward the center, a fraction is also ejected into outflows that have both radiative and mechanical influences on their environments. The exact physical nature has not yet been observationally determined, but outflows are seen in all types of accreting systems, i.e., proto-stellar objects \citep[e.g.][]{Mundt85}, neutron stars and stellar-mass black holes \citep[e.g.][]{Margon82}, and SMBHs \citep[e.g.][]{Cohen79}. These outflows can reach supersonic speeds when collimated into jets, eventually depositing significant energy into their surroundings \citep[e.g.][]{Cohen79,Fabian02,Allen06}. Material moving outward into their host galaxies also begins to cool via synchrotron radiation \citep[e.g.][]{Jones74}. Observed in the radio frequencies, this non-thermal process emitted in the core of the system is predicted to have a flat spectrum, independent of frequency \citep{Blandford79} making it a great observational tool for characterizing jet emission. Utilizing these two wavelength regimes, \cite{Merloni03}, \cite{Falcke04}, and \cite{Gultekin09} have all suggested a ``fundamental plane" of black hole activity connecting black hole mass, X-ray luminosity and radio luminosity. The plane spans over 9 orders of magnitude in mass, 12 orders of magnitude in radio luminosity and 13 orders of magnitude in X-ray luminosity \citep{Merloni03}. This plane suggests that accretion (traced by X-ray luminosity) and jet production (traced by radio luminosity) are fundamentally linked together. Although the exact coupling is not understood, this relation implies accretion must be driving jet production. A known problem with the relation is the large scatter of the data about the plane, \citep[$\sigma_{\mathrm{radio}}=0.88 {\rm dex}$][]{Gultekin09}. This scatter can be attributed to observational errors or the result of non-simultaneouty between the observations themselves. Measuring the X-ray and radio luminosities at different times may have sampled different fluctuations in the accretion rate in individual sources, driving them away from the relation. The time between X-ray and radio observations of \cite{Merloni03} and \cite{Gultekin09} are known to span a few years to a decade. To address these issues, and in order to examine disk-jet coupling at high mass accretion rates, we undertook a simultaneous X-ray and radio monitoring campaign of the Seyfert-1 AGN NGC 4051. This galaxy is relatively nearby (z=0.002336), and the central black hole mass has been determined through reverberation mapping techniques \citep[$(1.91 \pm 0.78) \times 10^6 M_\odot$][$(1.73 \pm 0.55) \times 10^6 M_\odot$]{Peterson04,Denney09}. NGC 4051 is typically observed to accrete at approximately 5\% of its Eddington luminosity \citep{Peterson04}. The innermost orbital timescale of NGC 4051 is on the order of a few minutes to hours, defined as $t_{dyn} \sim R/v_{\phi}$, where $R$ is the radius assumed to be only a few gravitational radii from the black hole and $v_{\phi}$ is the orbital velocity. While the viscous timescales are on order of days to weeks for typical parameters, defined as $t_{vis} \sim t_{dyn} \alpha^{-1} (H/R)^{-2}$, where $\alpha$ is the viscosity parameter in the standard $\alpha$-disk prescription \citep{Shakura73}, and $H$ is the scale-height of the disk. Variations in the accretion rate of NGC 4051 are thought to occur on these viscous timescales of a day to weeks. This is further supported by its highly variable spectrum, which is most notable in X-rays that vary up to a factor of 10 on weekly timescales \citep{Uttley99}. The variability seen in NGC 4051 was essential for our simultaneous monitoring campaign to probe different accretion rates. In this paper, we present the simultaneous X-ray and radio observations of NGC 4051, in effort to shed light on the implications of the fundamental plane of accretion onto black holes, and to explore jet-production in this Seyfert galaxy.

In this paper, we present eight \emph{Chandra} X-ray observations and six VLA/EVLA radio observations of Seyfert-1 NGC 4051 taken over a seven month period. By simultaneously measuring in X-ray and radio bands every two to four weeks apart, we are able to probe variations in the accretion rate that occur on the viscous timescale of a few days to weeks in NGC 4051. The observations reveal significant variability in both X-ray and radio bands. The first 8.4 GHz observation on MJD 54831.3 also shows evidence of jet production in the form of two distinct radio lobes, contrary to the idea that radio quiet galaxies like NGC 4051 lack jet production. In fact, work by \cite{Falcke01} and \cite{Nagar02} suggest that elongated, non-thermal emission is common in Seyfert and LLAGN. The lobes appear to be resolved into compact impact regions at very high resolution \citep{Giroletti09}. A variability analysis shows that the 2--10 keV X-ray luminosity and 8.4 GHz radio luminosity are inversely correlated according to $L_\mathrm{radio} \propto L_\mathrm{X-ray}^{-0.72 \pm 0.04}$. This differs by $11.6 \sigma$ from the current fundamental plane relations which correlates the radio luminosity to X-ray luminosity as $L_\mathrm{radio} \propto L_\mathrm{X-ray}^{0.67 \pm 0.12}$ for a fixed mass \citep{Gultekin09}. Furthermore, if extended emission is resolved out in the highest resolution image, then excluding the A configuration data point, the correlation is $L_\mathrm{radio} \propto L_\mathrm{X-ray}^{-0.12 \pm 0.05}$. This still differs from the fundamental plane by more than 6$\sigma$ and is not consistent with the relation. The results help to shed light on the fundamental plane of accretion onto black holes \citep{Merloni03, Falcke04, Gultekin09}. This plane demonstrates that accretion is linked to jet production, as evidenced by the positive correlation between X-ray and radio luminosity. The measurements in the fundamental plane are generally not simultaneous, and one goal of this study was to understand if simultaneity between X-ray and radio measurements would reduce the scatter in the plane. NGC 4051 was found to lie near to the fundamental plane, but showed a negative correlation between X-ray and radio luminosities, contrary to the positive correlation of the plane. The inverse correlation suggests that some systems may lie near to the fundamental plane, but vary like NGC 4051, moving across the plane and thus increasing the total scatter. At least a few other case studies have also shown a separate and inverse correlation deviating from the positive correlation suggested by the fundamental plane. The first is 3C 120, a SMBH with a mass of approximately 5.5 $\times 10^7 M_\odot$ that is observed to accrete at approximately 10\% $L_{Edd}$ \citep{Peterson04}. \cite{Chatterjee09} present a five year study of 3C 120 using RXTE and VLBI to obtain X-ray and radio data respectively. They use a discrete cross-correlation function to describe the correlation between the X-ray at 2.4--10 keV and the radio at 37 GHz. They find that the greatest amplitude in the DCCF to be $-0.68 \pm 0.11$, corresponding to an inverse correlation at a 90\% confidence level. This amplitude is consistent with the results of the cross-correlation analysis in NGC 4051. The minimum in of the DCCF of 3C 120 cited in \cite{Chatterjee09} corresponds to a time lag of 120 days $\pm 30$ days, with the X-ray dips leading leading the radio flares. This time delay corresponds to approximately 4 days in NGC 4051 when scaled using their respective masses, which is consistent with our data. At the opposite end of the mass scale, the $14 M_\odot$ black hole GRS 1915+105 \citep{Greiner01} sometimes also shows an inverse relation between simultaneous X-ray and radio observations. \cite{Rau03} present a survey of X-ray and radio observations using RXTE and the Ryle Telescope between 1996 November to 2000 September. The 1--200 keV X-ray flux showed no correlation to the 15 GHz radio observations. However, the 20--200 keV continuum did show an inverse correlation to the 15 GHz flux described by a Spearman's rank-order correlation coefficient of -0.75. The hard X-ray band in GRS 1915+105 excludes direct emission from the disk, which is consistent with using X-ray emission instead UV emission in NGC 4051 and 3C 120. The prototypical stellar-mass black hole, Cygnus X-1, may also show an inverse trend at high luminosity. \cite{Gallo03} study the disk-jet connections in this stellar-mass black hole by analyzing RXTE All Sky Monitoring X-ray data from 2--11 keV and the Ryle Telescope radio data at 15 GHz between 1996 January and January 2003. Cygnus X-1 does follow $L_\mathrm{Radio} \propto L_\mathrm{X-ray} ^{0.7}$, until it reaches approximately 2\% of its Eddington luminosity when the radio flux density turnovers. \cite{Gallo03} describe departures from this relation at high X-ray luminosity as quenching of its jet production as Cygnus X-1 moves into its ``high/soft" state. Observations of NGC 4051, 3C 120, GRS 1915+105, and Cygnus X-1 all show that X-ray and radio flux follow an inverse relation when observing at nearly simultaneous times when the sources are emitting at 1--10\% of Eddington. Given that an inverse correlation is seen in a quasar, a Seyfert, and two stellar-mass black holes at high Eddington fractions, it is possible that a distinct mode of disk-jet coupling holds at high Eddington fractions. This would go beyond a simple quenching of jet production, as discussed by \cite{Maccarone03} and \cite{Gallo03}, since the jet production does not turn off entirely \citep[evidenced by continuous radio emission and jet structures, especially in NGC 4051 and 3C 120;][]{Giroletti09, Chatterjee09}. This is the first study to probe a Seyfert galaxy in X-ray and radio bands on viscous timescales of its inner disk. In the future, we will undertake an extended monitoring campaign of NGC 4051 to further characterize this relation as well as determine any time lags between X-rays and radio fluxes.

1012.0762

We report on the results of a simultaneous monitoring campaign employing eight Chandra X-ray (0.5-10 keV) and six Very Large Array/Extended Very Large Array (8.4 GHz) radio observations of NGC 4051 over seven months. Evidence for compact jets is observed in the 8.4 GHz radio band; this builds on mounting evidence that jet production may be prevalent even in radio-quiet Seyferts. Assuming comparatively negligible local diffuse emission in the nucleus, the results also demonstrate an inverse correlation of L <SUB>radio</SUB> ∝ L <SUP>-0.72±0.04</SUP> <SUB>X-ray </SUB>. If the A configuration is excluded in the case where diffuse emission plays a significant role, the relation is still L_radio ∝ L_{X-ray}^{-0.12 ± 0.05}. Current research linking the mass of supermassive black holes and stellar-mass black holes in the "low/hard" state to X-ray luminosities and radio luminosities suggests a "fundamental plane of accretion onto black holes" that has a positive correlation of L <SUB>radio</SUB> ∝ L <SUP>0.67±0.12</SUP> <SUB>X-ray </SUB>. Our simultaneous results differ from this relation by more than 11σ (6σ excluding the A configuration), indicating that a separate mode of accretion and ejection may operate in this system. A review of the literature shows that the inverse correlation seen in NGC 4051 is seen in three other black hole systems, all of which accrete at near 10% of their Eddington luminosity, perhaps suggesting a distinct mode of disk-jet coupling at high Eddington fractions. We discuss our results in the context of disks and jets in black holes and accretion across the black hole mass scale.

12167225

2011ApJ...726L..16J

1012

1012.2551_arXiv.txt

Recent observations have found ubiquitous plasma jets over the solar atmosphere in various plasma parameters and magnetic field configuration. The solar chromospheric anemone jets show a cusp- or inverted Y-shaped structure which are believed to be a result of magnetic reconnection between a magnetic bipole and a preexisting uniform vertical field~\citep{Shibata2007}. The penumbral microjets show the apparent motion almost along the vertical guide field components in the interlocking-comb structure of magnetic field lines in the sunspot penumbra~\citep{Katsukawa2007}. The umbral light bridge jets are ejected along the vertical field lines emanating from the light bridge in the sunspot umbra~\citep{Shimizu2009}. There are many other chromospheric jets whose footpoints are not well resolved. It has sometimes been proposed that spicules may be produced by magnetic reconnection~\citep{Suematsu2008, Isobe2008}. The three dimensional magnetic field configuration at the footpoints of these jets are still puzzling. Especially, how these jets are accelerated is a fundamental question in solar physics, which would also give a hint to the understanding of the origin of astrophysical jets. The three dimensional (3D) numerical simulation may help us to understand this. For 3D magnetic reconnection, many simulations show the reconnecting process is much more complicated and difficult than two dimensional (2D) case~\citep{Yokoyama1995, Chen1999, Chen2001, Yokoyama2001, Isobe2005, Jiang2010}. Some of the 3D simulations have shown generation of flows parallel to magnetic field lines as a result of non-null reconnection~\citep{Pontin2005, Ugai2010}. However, these parallel flows have not been analyzed in detail. In this paper, we analyzed these flows in detail for the first time, because these flows are important as the origin of chromospheric jets and found that these jets ejected from the diffusion region move along the magnetic guide field and its 3D structure is similar to a fan-shape (hereafter, referred to as fan-shaped jets) which differ from the classical reconnection theory (generally speaking, the reconnection jets move along the ambient magnetic field, hereafter, ordinary reconnection jets)~\citep{Sweet1958, Parker1957, Petschek1964, Priest2000}. In this report, we give a description and analysis for these results.

Let us briefly discuss the application of our results to various jets in the chromosphere as shown in Table~\ref{tab1}. Although the footpoints of these jets are not necessarily well resolved, we assume that the 3D reconnection occurs in the photosphere or in the low chromosphere at the footpoints of these jets, and fan-shaped jets are ejected from the reconnection region along magnetic field lines in these layers. The Alfv\'en speed ($v_A$) in the photosphere and low chromosphere is about 10 km s$^{-1}$ in a typical isolated flux tube outside sunspots, and is about 10-50 km s$^{-1}$ in sunspot umbra and penumbra. Hence the $v_A$ in Table~\ref{tab1} shows such local Alfv\'en speed in the (hypothetical) reconnection region at the footpoint of these jets. Here, the $v_{A\bot}$ is the Alfv\'en speed based on the reconnecting component of magnetic field (we assume that $v_{A\bot} = 0.2 v_A$). The velocity of fan-shaped jets is only half of that of the ordinary reconnection flow ($v_{fan} \sim v_{A\bot}/2$). Once the fan-shaped jets are ejected, the slow mode MHD shock is formed ahead of the jets and propagate along the vertical magnetic field lines. Since the density decreases with height, the velocity amplitude at the slow mode shock increases with height, namely $v_{fan,max} \sim v_{fan}e^{0.5z/H}$ if the slow mode wave energy is conserved or $v_{fan,max} \sim v_{fan}e^{0.23z/H}$ if the shock is strong~\citep{Shibata1982b}, where the value z is the height of the jets measured from the reconnection region and H is the pressure scale height. If z/H = 13.3 (assuming the scale height H is 150 km and the slow mode shock propagates over a height of 2000 km), we get $v_{fan,max} \sim v_{fan}e^{0.23z/H} \sim v_{fan}e^{3.0} \sim 20v_{fan}$. Similar results are obtained also for the case z = 1000 km when the wave energy is conserved ($v_{fan,max} \sim v_{fan}e^{0.5z/H}$). Table~\ref{tab1} shows that the resulting velocity at the shock front ($v_{fan,max}$) when it reaches at the top of the chromosphere is comparable to the actual observed velocity of these chromospheric jets. Hence, our finding of fan shaped jets parallel to field lines are important for understanding the origin of chromospheric jets and seems to be successfully applicable to ubiquitous chromospheric jets. \begin{table}[ht] \caption{Comparison between observation and fan-shaped jets (unit: $km~s^{-1}$).\label{tab1}} \begin{tabular}{cccccc} \hline \hline Jets & Observational velocity & $v_A$ & $v_{A \bot} $ & $v_{fan}$ & $v_{fan,max}$ \\ \hline Chromospheric anemone jets\tablenotemark{1} & $\sim10$ & $10$ & $2$ & $1$ & $20$ \\ Spicules\tablenotemark{2} & $\sim25$ & $10$ & $2$ & $1$ & $20$ \\ Penumbral jets\tablenotemark{3} & $50-150$ & $10-50$ & $2-10$ & $1-5$ & $20-100$ \\ Umbral light bright jets\tablenotemark{4} & $28-180$ & $10-50$ & $2-10$ & $1-5$ & $20-100$ \\ \hline \end{tabular} \tablenotetext{1}{\cite{Shibata2007}} \tablenotetext{2}{\cite{Suematsu2008}} \tablenotetext{3}{\cite{Katsukawa2007}} \tablenotetext{4}{\cite{Shimizu2009}} \end{table} We described fan-shaped jets by simulating the 3D reconnection process using a simple initial shearing magnetic configuration in this report. We found that the fan-shaped jets which are accelerated by both gas pressure gradient and Lorentz force can move along the magnetic guide field lines and the velocity of these jets is about half of the local Alfv\'en speed determined by the reconnecting component of magnetic field. This new finding provides us a new way to understanding the magnetic reconnection in 3D geometry and it is also a new model for explaining the solar ubiquitous chromospheric or more general astrophysical jets. The details of them will be studied in our future papers.

1012.2551

Magnetic reconnection is a fundamental process in space and astrophysical plasmas in which the oppositely directed magnetic field changes its connectivity and eventually converts its energy into kinetic and thermal energy of the plasma. Recently, ubiquitous jets (for example, chromospheric anemone jets, penumbral microjets, umbral light bridge jets) have been observed by the Solar Optical Telescope on board the satellite Hinode. These tiny and frequently occurring jets are considered to be a possible evidence of small-scale ubiquitous reconnection in the solar atmosphere. However, the details of three-dimensional (3D) magnetic configuration are still not very clear. Here, we propose a new model based on 3D simulations of magnetic reconnection using a typical current sheet magnetic configuration with a strong guide field. The most interesting feature is that the jets produced by the reconnection eventually move along the guide field lines. This model provides a fresh understanding of newly discovered ubiquitous jets and moreover a new observational basis for the theory of astrophysical magnetic reconnection.

12306269

2012P&SS...61...15Y

1012

1012.3411_arXiv.txt

\label{intro} Amongst the eight planets in the solar system, Jupiter, in addition to being the largest planet, also has the largest magnetic moment and the largest magnetosphere. The magnetosphere interacts with both the solar wind and the conducting layer or ionosphere, in the planet's upper atmosphere. These interactions can be quite complex and we may use models with some simplifying assumptions (e.g. axial symmetry) to gain insight into the dynamics of the magnetosphere, upper atmosphere and their physical interactions with the solar wind.\\ Several models of Jupiter's magnetosphere and ionosphere have been developed in recent studies \citep{nichols04,cowley05,cowley07,smith09}. These models range from detailed studies of the middle magnetosphere only \citep{nichols04} to global studies of the entire magnetosphere \citep{cowley05, cowley07} and investigations of the coupled magnetosphere, thermosphere and ionosphere systems \citep{smith09} (henceforth SA09). \\ One of the important observations for guiding models is the dominance of Jupiter's magnetosphere by the rapid planetary rotation. Angular momentum is transferred from the planet to the disc-like \Rev{middle} magnetosphere via ion-neutral collisions in the ionosphere. The magnetospheric plasma exhibits a wide range of angular velocities, corresponding \Rev{to} a modest departure from rigid corotation with the planet at distances near Io (\unitSI[6\mbox{--}10]{R_J}) out to regions beyond \unitSI[20]{R_J} which rotate at ${\sim}\unitSI[50]{\%}$ of the planetary rate \citep{McNutt79, hill79, hill1983a, pontius97, vasyliunas1983}. This angular momentum and energy transfer between the ionosphere and magnetosphere is conveyed by two principal current systems. The first of these is related to the rotation of the middle magnetosphere. The second is related to the interaction of the solar wind with the magnetosphere at the high-latitude magnetopause \citep{hill1983bURANUS, isbell1984}. \\ The principal source of plasma for the middle magnetosphere (${\sim}\unitSI[20]{R_J}$ to several tens of \unit{R_J}) is the satellite Io \citep{bagenal1981} which ejects about \unitSI[500\mbox{--}1000]{kg\;s^{-1}} of sulphur dioxide gas which is then ionised \citep{kivelBOOK2004}. Iogenic plasma initially near corotation will lag further behind corotation as it diffuses radially outwards from the Io torus, due to the finitely conducting ionosphere being unable to supply all of the necessary angular momentum via the coupling currents. The electric field in the neutral atmosphere's rest frame depends on the difference in angular velocity between the polar thermosphere and the magnetically conjugate plasma disc, and drives a flow of equatorially directed Pedersen currents. Due to current continuity, \FAC in the steady-state must flow both upwards and downwards along the magnetic field lines which connect the ionosphere and magnetospheric plasma disc. Downward FACs flow from the outermost magnetosphere to the ionosphere. The upward directed FACs are carried by downward precipitating electrons from the magnetosphere \citep{cowbun2001, hill2001, khurana2001, southkiv2001}. These electrons excite emissions in the upper atmosphere and produce the main auroral oval at ${\sim}\unitSI[15]{^{\circ}}$ co-latitude \citep{satoh1996, clarke1998, clarke2004, prange1998, vasavada1999, pallier2001, grodent2003a}. Currents flow radially outward in the equatorial plane of the magnetosphere and, via the \textbf{\emph{J\unit{\times}B}} force, accelerate the plasma towards corotation. The Pedersen, radial and FACs thus represent a complete current `circuit' coupling the magnetosphere and ionosphere. \\ The thermospheric angular velocity at Jupiter partly controls the ionospheric Pedersen currents and thus the dynamics of the magnetosphere. We do not, however, have many measurements of these thermospheric velocities. Studies such as \citet{haungHill89} \Rev{and} \citet{pontius95} have \Rev{attempted} to model these velocities by coupling the magnetosphere, ionosphere and thermosphere, with the assumption that angular momentum was transported through the thermosphere solely by vertical viscous transport. These studies yielded two main

\label{sec:conclusion} In this study, we \Rev{have} expanded on the model of SA09 and described the effects of different solar wind dynamic pressures on the coupled ionosphere-magnetosphere system at Jupiter. We constructed three typical magnetospheric profiles (see \Table{\ref{tb:cases}})\Rev{,} compressed, baseline (average) and expanded. These were then coupled to our global two-dimensional thermospheric model \citep{smith08saturn} and a global conductivity model of the ionosphere (GG). This allowed for a comparison with results from SA09, but also \Rev{provided} a first quantitative investigation of how ionospheric, thermospheric and magnetospheric parameters were affected by differing solar wind conditions.\\ Our results confirm many results from previous studies such as \Rev{those of} \Rev{\citet{southkiv2001,cowbun2003b} and \citet{cowley07}}. We see an increase (resp. decrease) in thermospheric and magnetospheric angular velocities for compressed (resp. expanded) magnetospheres relative to our baseline. The thermosphere super-corotates just equatorward of the middle / outer magnetosphere boundary similarly to SA09. We solve for \unit{\Omega_M} self-consistently in the magnetodisc in all cases using the equations of disc dynamics. The \unit{\Omega_M} value in the outer magnetosphere is a constant, dependent on disc radius i.e solar wind pressure \citep{cowley05}. Magnetospheric angular velocities in the polar cap, are also fixed at a set fraction (${\sim}\unitSI[10]{\%}$) of rigid corotation (\unit{\Omega_J}) \citep{isbell1984}. We also found that the coupling currents showed an increase (${\sim}\unitSI[20]{\%}$) in intensity when going from an average to a more expanded magnetosphere and a decrease (${\sim}\unitSI[40]{\%}$) when going from average to compressed. \\ Our thermospheric model was used to simulate azimuthal and meridional neutral velocities. We see super-corotation in the azimuthal flows equatorward of the edge of the magnetodisc flux shells. There lies a strong sub-corotational jet at mid to upper altitudes in the mapped ionospheric locations of the outer magnetosphere and polar cap. The spatial size of the strong sub-corotation region increases with increased magnetospheric size due to the weaker magnetic field strength in expanded magnetospheres; thus the transfer of angular momentum is less effective at maintaining corotation. \Rev{ Angular momentum is transferred from the thermosphere to the magnetosphere, in order to accelerate the latter towards corotation. If the thermosphere itself is significantly sub-corotating, then there is a lower `reservoir' of available angular momentum that can be transferred. This results in a decreased plasma angular velocity in these outer regions of the magnetosphere.} We see a meridional flow directed polewards at low altitudes and equatorwards at high altitudes. From the poleward edge of the magnetodisc to the centre of the polar cap, \Rev{a region of accelerated poleward flow} exists whose velocity magnitude increases from a compressed to an expanded magnetosphere. This occurs because there is a force imbalance in this region that increases advection of momentum in expanded magnetospheres. Advection restores balance which results in the acceleration discussed above. This accelerated flow produces a `hotspot' in the polar cap, with a maximum temperature \Rev{increase of ${\sim}\unitSI[130]{K}$} from compressed to expanded magnetosphere. The size of the `hotspot' also increases with an expanding magnetosphere. We find that the outer magnetosphere and polar cap are most strongly heated by Joule heating and ion drag. This heat is then distributed across the polar region via advection rather than viscous transport, whilst more equatorial regions are significantly cooled. This aspect of thermospheric flow is consistent with those presented in SA09. These results also suggest that accurate measurements of ionospheric temperature in the polar region could potentially be used to probe magnetospheric conditions.\\ We also showed that the power dissipated in the upper atmosphere (consisting of both Joule heating and ion drag) increases with \Rev{an} expanded magnetospheric configuration. The power used to accelerate the magnetospheric plasma initially increases as we expand the magnetosphere from compressed to average configurations, but then decreases with an expansion from average to expanded. This suggests that power used to accelerate the magnetosphere has a `local' maximum for a magnetosphere size somewhere between a compressed and expanded configuration. The total power extracted from planetary rotation is the net sum of the atmospheric and magnetospheric powers and this is positively correlated with magnetosphere size. Comparing our compressed and average magnetospheres with the `intermediate' and `baseline' cases in \citet{cowley07}, we showed that the use of a two-dimensional thermosphere model results in the transfer of ${\sim}\unitSI[20]{\%}$ more energy from the thermosphere to the magnetosphere in order to accelerate the plasma in the magnetodisc. Using our more realistic model of thermospheric flow also produced increased dissipation of energy in the thermosphere via Joule heating and ion drag than the cases presented in \citet{cowley07}. \\ We have shown that our original compressed case has some unusual current density features due to a relatively high value for \Rev{the} radial current at the outer disc boundary. In order to confirm this we decreased the boundary value \Rev{of} \unit{I_{\rho\infty}} for each case in order to produce alternative models with minimum variance in their \FAC profiles. This led to the selection of \unit{I_{\rho\infty}} of \unitSI[45]{MA}, \unitSI[68]{MA} and \unitSI[80]{MA} for the compressed, average and expanded cases respectively.\\ Decreasing the radial current \unit{I_{\rho\infty}} at the boundary between the middle and outer magnetospheres resulted in all magnetosphere-ionosphere coupling currents being reduced in accordance with the new value of \unit{I_{\rho\infty}}. This is expected under the assumption of current continuity. The main differences between cases with large \Rev{and reduced} radial current\Rev{s} lies mainly within the magnetodisc. For \Rev{the} \FAC density, changes related to \unit{I_{\rho\infty}} were also significant throughout the outer magnetosphere. Thermospheric and magnetospheric angular velocities changed only slightly for the baseline and expanded case but much more substantially for our compressed case. For azimuthal flows we found that decreasing \unit{I_{\rho\infty}} also generally increased the level of sub-corotation throughout high latitudes. For meridional flows we found slight increases in localised regions of accelerated flow, most evident in the alternate compressed case. We also found that the polar region becomes slightly warmer with a decrease in \unit{I_{\rho\infty}}; peak temperatures for the alternative configurations increasing relative to their $\unit{I_{\rho\infty}}{=}\unitSI[100]{MA}$ counterparts. The total integrated powers increased with decreasing \unit{I_{\rho\infty}} for our compressed case, but decreased for our baseline and expanded cases. The integrated magnetospheric power for all cases decreased along with \unit{I_{\rho\infty}}, whilst atmospheric power increased by ${\sim}\unitSI[20]{\%}$ for the alternate compressed case but remained almost equal for our baseline and expanded cases. Thus, it seems that decreasing the boundary radial current \unit{I_{\rho\infty}} effectively decreases the `ability' of the thermosphere to transfer angular momentum to the magnetosphere. This \Rev{behaviour, as expected} decreases the intensity of auroral emissions and produces a slightly warmer polar region.\\ Our calculations suggest that main oval auroral emissions and brightness for an expanded magnetosphere would generally be greater than that of a compressed one. The detailed structure of the \FAC density profile in the magnetodisc is most sensitive to the value of \unit{I_{\rho\infty}} for the compressed case. Compressed magnetospheres in the steady state have larger field strength than expanded ones and are more efficient at maintaining the co-rotating magnetodisc plasma at larger distances. This leads to a smaller shear in angular velocity between the magnetosphere and thermosphere and, consequently, smaller thermospheric temperatures. As a result, auroral emission is brightest for the most expanded magnetospheric systems. We also saw that auroral emissions would increase at the boundary between the outer magnetosphere and the polar region with magnetospheric compression due to the large change in plasma angular velocity at this boundary. Better observational constraints of \unit{\Omega_M} are required to confirm this prediction. \\ This aspect warrants further investigation since we have not attempted to model the change in polar cap angular velocity with solar wind dynamic pressure. Furthermore, the caveat with these predictions is that the system is in a steady-state (where there is no explicit time dependence of the model outputs). We thus view this study as an initial step \Rev{towards} developing a model to study the transient effects of rapid changes in the solar wind dynamic pressure. Results of such studies could provide further insights to the `energy crisis' at Jupiter (SA09), and the physical origin of transient auroral features. \\ Finally, the results presented in this study contribute to a larger set of theoretical investigations which have provided useful quantitative predictions of how the Jovian aurorae would respond to changes in solar wind dynamic pressure. Such results are useful for interpreting auroral observations, and \Rev{for} making more extensive use of such data as remote diagnostics of the physical state of the Jovian magnetosphere.

1012.3411

The coupling of Jupiter's magnetosphere and ionosphere plays a vital role in creating its auroral emissions. The strength of these emissions is dependent on the difference in speed of the rotational flows within Jupiter's high-latitude thermosphere and the planet's magnetodisc. Using an azimuthally symmetric global circulation model, we have simulated how upstream solar wind conditions affect the energy and direction of atmospheric flows. In order to simulate the effect of a varying dynamic pressure in the upstream solar wind, we calculated three magnetic field profiles representing compressed, averaged and expanded ‘middle’ magnetospheres. These profiles were then used to solve for the angular velocity of plasma in the magnetosphere. This angular velocity determines the strength of currents flowing between the ionosphere and magnetosphere. We examine the influence of variability in this current system upon the global winds and energy inputs within the Jovian thermosphere. We find that the power dissipated by Joule heating and ion drag increases by ∼190% and ∼185% from our compressed to expanded model respectively. We investigated the effect of exterior boundary conditions on our models and found that by reducing the radial current at the outer edge of the magnetodisc, we also limit the thermosphere's ability to transmit angular momentum to this region.

12207326

2011JCAP...07..001F

1012

1012.4007_arXiv.txt

The explanation of small neutrino masses \cite{oscil} and dark matter \cite{wmap} seems to be a key ingredient to consider physics beyond the standard model (SM). An interesting possibility of such extensions may be models which can closely relate neutrino masses to dark matter (DM). In this kind of models, a discrete symmetry is often introduced to forbid tree level Dirac neutrino masses. Some of additional particles introduced to commit neutrino mass generation have its charge such that it can forbid the lightest one to decay into the SM particles. This stable particle becomes DM. This DM is a crucial ingredient of the neutrino mass generation in this scenario. The radiative seesaw model proposed by Ma \cite{Ma:2006km} is its simple and interesting example.\footnote{A lot of radiative neutrino mass models exist now. Phenomenology including the DM nature in such models has been studied in a lot of works \cite{cdmmeg,lflavor,ext,scdm,fcdm,ncdm}.} Both the numbers of new particles and free parameters are comparably small. Its supersymmetric extension is also straightforward \cite{sma,fks}.\footnote{A relevant supersymmetric model is also considered in a different context in \cite{e6}.} Moreover, if we introduce an anomalous U(1) symmetry in this extension \cite{anomu}, we could explain the origin of the discrete symmetry, required hierarchical structure of both couplings and masses due to the Frogatt-Nielsen mechanism \cite{fn,yh}. Both neutrino oscillation data and DM relic abundance can also be explained consistently with lepton flavor violating processes such as $\mu\rightarrow e\gamma$. A characteristic feature in such an extension with $R$ parity conservation is that the model has two DM candidates\footnote{Multicomponent DM and its phenomenology are studied in a different model \cite{flnp}.}. One is the lightest superparticle whose stability is guaranteed by the $R$ parity. The other one is a new particle introduced for the neutrino mass generation and its stability is guaranteed by the new $Z_2$ symmetry. As a result, the model shows discriminative differences from the ordinary minimal supersymmetric SM (MSSM) in the DM search. For example, if the recently reported cosmic ray anomalies \cite{pamela,fermi} are considered as the DM signature of the model, they may be explained not by the DM annihilation \cite{mindep,sommerfeld,bwenhance} as in the MSSM but by the DM decay \cite{it,decay,gravitino,rparity,hidden}. In fact, if the $Z_2$ symmetry is violated by the anomaly effect, the DM guaranteed its stability by the $Z_2$ symmetry can decay into the lightest neutralino \cite{fks,anomu}. Direct search of the DM could also show the difference from the MSSM. In this paper, we study signals of the DM in the supersymmetric extension of the Ma model. The model is considered as an effective model due to spontaneously broken anomalous U(1) gauge symmetry. It naturally brings the weakly broken $Z_2$ symmetry to the model in addition to the conserved $R$ parity. We discuss signatures due to the decay of the unstable DM and also the direct detection of the DM through the elastic scattering with nuclei. The paper is organized as follows. In section 2 we address the model and explain the nature of the DM sector which is imposed by various experimental results. In section 3 several signals expected in the DM sector are analyzed. In particular, the decaying DM is studied to explain the cosmic ray anomalies reported recently. A feature of the monochromatic gamma yielded through the DM radiative decay is also studied. Finally, we discuss the direct search of the DM. Section 4 is devoted to the summary.

We have studied the nature of the DM sector in a supersymmetric extension of the radiative neutrino mass model. An anomalous U(1) symmetry is introduced to explain the hierarchical structure of the coupling constants and mass scales in the model. The spontaneous breaking of this symmetry can induce a new $Z_2$ symmetry which guarantees the stability of the lightest odd parity particle. As a result, the model has two DM components as long as $R$ parity is assumed to be conserved. However, since one of these discrete symmetries which guarantee the stability of DM is not exact due to the anomaly, one DM component is unstable to decay through a hugely suppressed term which is nonperturbatively induced via the anomaly effect. These DM components could be detected through the indirect search of the yields of the decaying DM and the direct search of the elastic scattering from nuclei by taking account that the DM relic abundance is composed of these. Positrons generated by the decaying DM can explain the cosmic ray anomaly reported recently. Parameter regions predicted by the direct detection can be different from the MSSM case since two DM components may contribute the relic abundance in the same order. If the line shape gamma is observed in the cosmic ray, we might confirm the model by combining it with the direct search of the DM. Forth coming experiments for DM can give fruitful information to the model. \vspace*{5mm} We would like to thank Martin Holthausen for careful reading of the manuscript. This work is partially supported by a Grant-in-Aid for Scientific Research (C) from Japan Society for Promotion of Science (No.21540262) and also a Grant-in-Aid for Scientific Research on Priority Areas from The Ministry of Education, Culture, Sports, Science and Technology (No.22011003). The numerical calculations were carried out on SR16000 at YITP in Kyoto University. \vspace*{7mm} \newpage

1012.4007

Supersymmetric radiative neutrino mass models have often two dark matter candidates. One is the usual lightest neutralino with odd R parity and the other is a new neutral particle whose stability is guaranteed by a discrete symmetry that forbids tree-level neutrino Yukawa couplings. If their relic abundance is comparable, dark matter phenomenology can be largely different from the minimal supersymmetric standard model (MSSM). We study this in a supersymmetric radiative neutrino mass model with the conserved R parity and a Z<SUB>2</SUB> symmetry weakly broken by the anomaly effect. The second dark matter with odd parity of this new Z<SUB>2</SUB> is metastable and decays to the neutralino dark matter. Charged particles and photons associated to this decay can cause the deviation from the expected background of the cosmic rays. Direct search of the neutralino dark matter is also expected to show different features from the MSSM since the relic abundance is not composed of the neutralino dark matter only. We discuss the nature of dark matter in this model by analyzing these signals quantitatively.

12214427

2011MNRAS.415.2580K

1012

1012.1880_arXiv.txt

Observations show that the cosmic star formation rate (SFR) has declined by more than an order of magnitude since $z \simeq 1$ \citep{2004ApJ...615..209H}. However, a combined census of the cold gas, the fuel for star formation, and stellar components is still largely missing in observations. The cold gas fraction of a halo is a crucial ingredient in models of galaxy formation and constitutes the link to how galaxies obtain gas and subsequently convert it to stars. Hence, measurements of HI in the post-reionization era can place tight constraints on different models of galaxy formation \citep{2009astro2010S.241P}. After the epoch of reionization, the neutral hydrogen (HI) survives in dense clouds, e.g.~damped Lyman-$\alpha$ systems (DLAs) and Lyman-limit systems (LLS), that are high-redshift equivalents of the HI-rich galaxies that we see at the present epoch. The baryon fraction locked up in HI, $\Omega_{\rm HI}$, in star-forming galaxies in the post-reionization epoch can be determined from the study of damped Lyman-$\alpha$ systems in absorption for $0.5 \le z \le 5$ \citep{2005ApJ...635..123P,2006ApJ...636..610R,2009A&A...505.1087N}. Even though these observations give clues about aggregate behaviour of star formation as a function of redshift, they cannot be used to infer the total HI mass of these systems because they are seen in absorption. At $z\simeq 1$, even the detection of HI in damped Lyman-$\alpha$ has not been easy as the Lyman-$\alpha$ frequency is not accessible to ground-based telescopes. At this redshift, constraints on the global HI fraction come from associated $\rm MgII$ systems, HST observations \citep[for details see][and references therein]{2006ApJ...636..610R}, and the absorption of 21~cm radiation from bright background radio sources \citep{2009MNRAS.396..385K}, but with significant uncertainties on the estimated HI fraction. Direct observation of HI in emission and its detailed modelling has only been possible at $z \simeq 0$ thus far \citep{2005MNRAS.359L..30Z}. Direct observation in 21cm emission of ten massive galaxies have been reported for $0.17 < z < 0.25 $, with the Arecibo telescope \citep{2008ApJ...685L..13C}. At higher redshifts, the HI emission from individual clouds is too weak to be detectable with present radio instruments \citep{2010MNRAS.407..567B}. A long integration time is required for detecting even the brightest objects since the peak signal is a few tens of micro Jansky whereas the system noise is of the order of hundreds of micro Jansky. A possible approach to circumvent the difficulty of detecting individual clouds lies in stacking the HI emission of galaxies with known redshifts. This approach has been attempted for both cluster galaxies and the field galaxies in the recent past \citep{2007MNRAS.376.1357L,2009MNRAS.399.1447L}. In particular, a similar line of study has resulted in the recent detection of HI at $z \simeq 0.8$ \citep{2010Natur.466..463C}. An alternative approach rests on the possible detection of the fluctuation in the redshifted HI emission from high redshifts \citep{2001JApA...22..293B,2008PhRvL.100i1303C,2009PhRvD..79h3538B}. On the theoretical side, semi-analytical models of galaxy formation have looked at the evolution of cold gas (both in atomic and molecular form) in galaxies and their results match with observations at $z =0$ \citep{2009ApJ...696L.129O,2009ApJ...698.1467O,2009ApJ...703.1890O,2009MNRAS.400..665O, 2010MNRAS.406...43P,2010MNRAS.409..515F,2010arXiv1003.0008K}. However observations at higher redshifts are needed to better constrain the evolution of cold gas predicted by these models. Given the importance of connecting cold gas and stars at $z \simeq 1$ over a wide range of galaxy environments, it is crucial to make predictions for various detection strategies for current and upcoming telescopes. In this work we focus on the stacking method of individual galaxies with known redshifts to predict how well the HI mass function at $z = 1$ can be constrained with existing surveys and telescopes (in particular the DEEP2\footnote{http://deep.berkeley.edu} survey and the GMRT\footnote{http://gmrt.ncra.tifr.res.in}; but note that our method is generic and can be extended to future surveys and instruments). By stacking we can also study the contribution of small satellite galaxies, which are undetected in an optical survey but (as we shall show) contain non-negligible amounts of HI, to the total $21\,{\rm cm}$ signal in emission and also examine the constraints that one can put on the HI mass function. We model the HI in dark matter halos in a large $N-$body simulation by refining the model of \cite{2010MNRAS.407..567B}. Given the paucity of observations at the redshifts under consideration, and our limited understanding of how HI populates dark matter halos at these redshifts, we consider a variety of models. These are constrained by observations of HI at low redshift, simulations of DLAs in small-volumes at high redshift, as well as by some results of semi-analytical models of galaxy formation at intermediate redshifts. In particular, the models that we consider are consistent with recent observations of HI in emission at $z \simeq 0.8$ \citep{2010Natur.466..463C}. Our paper is organised as follows. We present our large dark matter simulation in Section~\ref{sec_nbody}, and describe our model for the HI distribution in the simulation along-with specifications of the DEEP2 survey as well as the GMRT in Section~\ref{sec_h1model}. We discuss our stacking procedure of individual galaxies and the contribution of undetected satellites to the stacked HI spectra in Section~\ref{sec_h1signal}. In Section~\ref{sec_results}, we present our results and discuss the prospects of detection with the GMRT, and the constraints that one can put on the HI mass function. We revisit the issue of undetected satellites and its effect on the HI mass function and discuss whether their presence can be detected. Finally, we present our conclusions in Section~\ref{sec_conclusions}. \begin{table} \begin{center} \begin{tabular}{c|c|c|c} \hline $L_{\mathrm{box}}$ & $N_{\mathrm{part}}$ & $m_{_{\mathrm{DM}}}$ & $\epsilon$ \\ $\left(h^{-1}\mathrm{Mpc}\right)$& & $\left(10^8 h^{-1}\mathrm{M}_{\odot}\right)$ & $\left(h^{-1}\mathrm{kpc}\right)$ \\ \hline 400 & $2448^3$ & 3.1 & 6.5 \\ \hline \end{tabular} \end{center} \caption{Basic simulation parameters for our dark matter run. The columns list the size of the simulation box, $L_{\mathrm{box}}$, the number of dark matter particles used in the simulation, $N_{\mathrm{part}}$, the mass of a single dark matter particle, $m_{_{\mathrm{DM}}}$, and the gravitational softening length, $\epsilon$. All length scales are in comoving units.} \label{table_simparam} \end{table}

\label{sec_conclusions} In this paper, we have studied the prospects for detecting HI in emission at $z \simeq 1$. This is a crucial epoch in the study of galaxy formation, since the cosmic star formation rate starts to decline around this time and the missing link in observations is an accurate census of cold gas, which fuels star formation, at these redshifts and beyond \cite[for more discussion of the importance of this issue see e.g.][]{2009astro2010S.241P}. We make a case that an existing instrument like the GMRT can put strong constraints on the amount of cold gas contained in galaxies, when it is combined with a survey like DEEP2. In this work, we have only focused on the overlapping volume of DEEP2 and GMRT, which represents a quarter of the total DEEP2 volume. Our study is representative of what might be achievable by combining the already existing optical data and the presently operational radio interferometers. The HI signal is too weak in emission for the detection of individual objects at $z\simeq1$. However, this can be circumvented by a stacking strategy, similar to \cite{2007MNRAS.376.1357L,2009MNRAS.399.1447L}, which we here apply to look at the prospects of detection. Our conclusions are: \begin{itemize} \item We find that a detection of HI in emission at redshifts of $z \simeq 1$ is possible even with existing instruments like the GMRT when combined with the DEEP2 survey. Such an observation will be able to constrain the HI mass function in the halo mass range $10^{11.4} h^{-1}\msun \leq M \leq 10^{12.5} h^{-1}\msun$. The detection significance is in the range of $5\hbox{--}12\sigma$ for an optimistic noise level of $71\mu$Jy with 24 hours of integration. \item The models that we consider are consistent with recent observations of \cite{2010Natur.466..463C}, who computed the cross-correlation of the density field of DEEP2 galaxies and the 21cm intensity field with the GBT. However these observations allow for all the three models that we consider. On the other hand, we find that using the stacking technique it will be possible to discriminate between the different scenarios with an instrument like the GMRT, at least for the optimistic level of noise. \item Combining our estimates of HI bias with the observations of \cite{2010Natur.466..463C}, we find that the most conservative constraint on the cosmic HI fraction at $z \simeq 0.8$ to be $\Omega_{\mathrm{HI}} = (1.16 \pm 0.30)\times 10^{-3}$. \item We find that undetected satellites in the optical produce a non-negligible contribution to the stacked HI spectra. Their signature is better seen in the stacked spectra rather than in the mass function, since we integrate over one parameter, i.e.~the width of the spectra, to obtain the mass function. For a noise of $71 \mu \mathrm{Jy}$, features of satellites can be seen at the $4\sigma$ level in the stacked spectra for a mass threshold of $M \geq 10^{11.4}h^{-1}\msun$. This detection significance for satellites increases by more than $7\sigma$ for $M \geq 10^{12.0} h^{-1}\msun$ (see e.g. Fig~\ref{fig_chisq} and Fig.~\ref{fig_cummassfn}). In comparison, the mass function discriminates satellites at the $\sim 1\sigma$ level. \item We have also considered a much higher level of noise, i.e.~$420 \mu \mathrm{Jy}$, which should represent an upper bound on noise in the GMRT. With this amount of noise, a detection of the mass function is possible at the $1.7\hbox{--}3\sigma$ level. We expect that the real detection significance is bracketed by our the optimistic and conservative noise levels. \item For the higher noise, the effect of satellites on the stacked spectra can be seen only at the $1\hbox{--}3\sigma$ level across the ranges of mass that we consider. The best-fit parameters of the spectra however are incorrect for the larger mass cuts when compared to the theoretical numbers. \item Cosmic variance affects the mass function more strongly than the average stacked spectra. For this reason, if HI is populated in halos according to model 2 one cannot quantify the effect of subhalos on the mass function due to the effect of cosmic variance. For models 1 and 3, cosmic variance does not swamp the errors due to noise. \end{itemize} One can use the stacking strategy to independently probe $\Omega_{\mathrm{HI}}$ \citep{2009MNRAS.399.1447L}. This is not the case in the cross-correlation approach which constrains $br\Omega_{\mathrm{HI}}$. As in \cite{2007MNRAS.376.1357L}, it would be useful to also target a subset of galaxies in DEEP2 whose SFR has been measured. This would provide the link between the SFR and the amount of cold gas in galaxies and provide insight into models of galaxy formation. Since spectroscopic surveys are accurate but expensive it would be worthwhile to first try this stacking strategy on future surveys like the LSST, which are designed to give photometric redshifts of $\simeq 10^{10}$ galaxies. Photometric redshifts are more prone to errors, but it has to be seen if the larger sample of a photo-$z$ survey like LSST could beat down the noise by its sheer number of objects. In this study, we have modelled the HI in all the halos, centrals and satellites, and we have seen how the satellite population affects the HI mass function as well as the stacked HI profile. The possibility to see the effect of satellites missing in an optical survey in the corresponding $21\,{\rm cm}$ survey is an exciting prospect. On the one hand, we find that stacking can distinguish between models, but the effect of satellites on the stacked profile is model dependant, and to see their effect one may need to combine it with the cross-correlation approach. In the cross-correlation method the optical density field does not contain all the satellites, whereas the HI intensity field does. If we use the same mass threshold when constructing the HI intensity field, one naively expects a stronger cross-correlation between the two fields. A preliminary investigation shows that this is indeed the case. We also expect that the HI bias and stochasticity will be sensitive to subhalos. A combination of both approaches would shed light on both the model and the contribution of satellites. The other approach is to observe the auto-correlation function or the power spectrum of HI, and to constrain the HOD of HI galaxies \citep{2010MNRAS.404..876W} from it. Such an inferred model of HOD when combined with a direct detection as is done here could reveal the contribution of the satellite population on the total signal. We will look into these aspects of the analysis in a forthcoming paper. Currently operational radio instruments -- both single dish and interferometers -- have the capability to detect HI in emission at $z \simeq 1$, as already demonstrated by \cite{2010Natur.466..463C}. We explored the potential of these complementary strategies. In particular, we studied in detail the efficiency of stacking, possible only with interferometers. In the near future, we expect larger optical galaxy samples at $z\simeq 1$ and radio observations with wider field-of-views and spectral coverage using upcoming radio instruments \cite[e.g.][]{2007PASA...24..174J}. This observational progress will enable a better determination of the HI signal using either of the strategies, thereby substantially improving our estimate of the HI content of galaxies at $z \simeq 1$.

1012.1880

We use a large N-body simulation to examine the detectability of H I in emission at redshift z≃ 1, and the constraints imposed by current observations on the neutral hydrogen mass function of galaxies at this epoch. We consider three different models for populating dark matter haloes with H I, designed to encompass uncertainties at this redshift. These models are consistent with recent observations of the detection of H I in emission at z≃ 0.8. Whilst detection of 21-cm emission from individual haloes requires extremely long integrations with existing radio interferometers, such as the Giant Meter Radio Telescope (GMRT), we show that the stacked 21-cm signal from a large number of haloes can be easily detected. However, the stacking procedure requires accurate redshifts of galaxies. We show that radio observations of the field of the Deep Extragalactic Evolutionary Probe 2 (DEEP2) spectroscopic galaxy redshift survey should allow detection of the H I mass function at the 5-12σ level in the mass range 10<SUP>11.4</SUP> h<SUP>-1</SUP> M<SUB>⊙</SUB>≤M<SUB>halo</SUB>≤ 10<SUP>12.5</SUP> h<SUP>-1</SUP> M<SUB>⊙</SUB>, with a moderate amount of observation time. Assuming a larger noise level that corresponds to an upper bound for the expected noise for the GMRT, the detection significance for the H I mass function is still at the 1.7-3σ level. We find that optically undetected satellite galaxies enhance the H I emission profile of the parent halo, leading to broader wings as well as a higher peak signal in the stacked profile of a large number of haloes. We show that it is in principle possible to discern the contribution of undetected satellites to the total H I signal, even though cosmic variance limitation make this challenging for some of our models.

12168488

2011ApJ...737...62B

1012

1012.5284_arXiv.txt

\label{sec:intro} The large-scale dynamics of astrophysical plasmas is theoretically described within the framework of magnetohydrodynamics (MHD). In many cases of interest, these astrophysical flows are characterized by extremely large Reynolds numbers, which in turn implies that a wide range of spatial scales are relevant to properly describe their dynamical behavior. At sufficiently small spatial scales, kinetic plasma processes might become non-negligible under certain circumstances. For instance, in a fully ionized plasma, whenever one reaches spatial scales as small as the ion skin depth $c/\omega_{pi}$ (where $c$ is the speed of light and $\omega_{pi}$ is the ion plasma frequency), the Hall effect should not be neglected. This regime correspond to fully ionized plasmas with sufficiently low ion densities. However, it can also arise in cold plasmas with a low ionization fraction $\chi$ in wich case the relevant Hall scale is given by $(c/ \omega_{pi})\chi ^{-1/2}$ (see \citet{pandey} for details on Hall-MHD of partially ionized plasmas). In cold plasmas such as those present in protoplanetary disks, another kinetic effect known as ambipolar diffusion might become relevant \citep{brazwe}, especially toward the disk surface \citep{pandey}. The relative importance of non-ideal effects such as Hall, ohmic dissipation and ambipolar diffusion has been extensively discussed by \citet{balter} as well as by \citet{pandey}. To highlight the relevance of the Hall effect in astrophysical plasmas, it is useful to compare the orders of magnitude of the ohmic ($O$), inductive ($I$), and Hall ($H$) terms in the generalized Ohms's law. For example, in a typical protostellar disk, the relevant ratios are $H/O \sim 10^{2}$ and $H/I \sim 10^{4}$ \citep{balter}; while, for dwarf nova disk $H/I \sim 1$ \citep{sanostone1} and for the crust of a neutron star $H/O \sim 10^{3}$ \citep{holrud}. There are many examples of astrophysical plasma flows for which the role of the Hall effect has been studied: regions of star formation \citep{norhey}, dense molecular clouds \citep{warng}, the interstellar medium \citep{spa, kin} or even the early universe \citep{taj}. Due to their intense magnetic fields, the Hall currents can be relevant also in white dwarfs and neutron stars \citep{urpyak, shaurp, pot}. Also, the role of Hall currents in the generation of magnetic fields by turbulent dynamo activity has been studied by \citet{min1}; see also \citet{min2,min3} and references therein. Even though the dynamics of small-scale structures is often unobservable in astrophysical flows, they may play an important role through nonlinear interactions with the large-scale part of the flow. In many cases, the small-scale dynamics of the fluids are instrumental in changing the transport properties of the large-scale dynamics of the fluids, and therefore it is relevant to identify potential instabilities in the microscale. At these small spatial scales, the large scale velocity field can be reasonably approximated by a linear shear flow. The so-called shear-driven instabilities are those that originate as a result of the presence of a large-scale velocity shear. In this paper we study the potential relevance of the Hall effect in the presence of an external magnetic field as well as a linear shear flow. In particular, we focus our attention on the following two types of flows with Hall effect: (a) non-rotating shear flows leading to what we call Hall magneto-shear instability (Hall-MSI); and (b) differententially rotating flows leading to Hall magneto-rotational instability (or Hall-MRI). In the absence of both rotation and shear, the linear modes in Hall-MHD correspond to right-hand polarized {\it whistlers} and left-hand polarized {\it ion-cyclotron} waves (see for instance \citet{mahajan05} and references therein). When these modes propagate embedded in a shear flow, the ion-cyclotron mode might become unstable. This instability takes place when the shear is steep enough to be larger than the ion-cyclotron frequency. A linear analysis of this instability has been recently reported by \citet{kunz}, for the case of weakly ionized plasmas. Another study on the influence of the Hall effect on weakly ionized plasmas subjected to differential rotation, was reported by \citet{rudkit} for finite magnetic Reynolds numbers up to $300$ (see also \citet{rudsha}). In the present study we adopt a one-dimensional configuration, perform a linear analysis to identify potential instabilities and then compare with numerical simulations. The set of equations as well as the simplifying assumptions that we adopt are listed in \S\ref{sec:eqs}. The dispersion relation for the linear regime is shown in \S\ref{sec:disprel}. We briefly describe the numerical code employed in \S\ref{sec:simul}. The role of the Hall effect on non-rotating shear flows is presented in \S\ref{sec:hmsi}, while the action of Hall currents on the well known magneto-rotational instability is discussed in \S\ref{sec:hmri}. The non-linear behavior is tackled through a qualitative approach in \S\ref{sec:nonlinear}. Finally, we summarize our conclusions in \S\ref{sec:conclu}.

\label{sec:conclu} The present work is a comprehensive study of the instabilities arising from the interplay between the Hall effect and a linear shear flow for a one-dimensional model. In other words, we analyze the role of the Hall effect in shear-driven instabilities. We find that an instability develops when the Hall effect is present, which we term Hall magneto-shear instability. Also, we recover the magneto-rotational instability, as a particular case, and we quantitatively evaluate the influence of the Hall currents on it. More specifically, we investigate the stability of the system in the parameter space set by the wavenumber and the Hall parameter. In non-rotating plasmas, we determine the region in the $(k^{2}, \varepsilon)$ diagram where the Hall magneto-shear instability takes place. In rotating plasmas (such as accretion disks), we examine three cases: sub-keplerian, keplerian, and super-keplerian. For each rotation profile, we establish the region in the $(k^{2}, \varepsilon)$ diagram where the Hall magneto-rotational instability occurs. The standard MRI is recovered in the particular case for zero Hall parameter ($\varepsilon=0$). In both unstable modes, we find a very good agreement between the theoretical model and the numerical simulations. In addition, we explore the influence of the plasma parameter in two asymptotic cases: $\beta \cong 1$ and $\beta \gg 1$. We find that the linear behavior is independent of this parameter. In the non-linear stage of the large beta regimes, the flow dynamics seems to evolve like in an MHD system. Whitin the framework of astrophysics, the Hall magneto-shear instability could be relevant in the interface between a jet and the surrounding environment where a strong shear is present. Astrophysical jets have a very high degree of collimation, probably as a consequence of the presence of magnetic fields. The one-dimensional model adopted in this work, even though quite simple, might properly describe this circumstance: the azimuthal component of the helicoidal velocity field arround the jet can be represented as a function of the radial direction and fulfills the periodicity condition in a straightforward fashion. The relevant instability in these strongly sheared flows is Kelvin-Helmholtz, which is a purely hydrodynamic instability. The presence of external magnetic fields modify the corresponding growth rate, depending on their strength and spatial orientation, but typical numbers quoted in the literature remain a small fraction of the imposed velocity shear (\cite{ferrari, bodo} and \citep{huba} for results from Hall-MHD simulations). Therefore, the Hall magneto-shear instability, with maximum groth rate of $0.5$ times the externally imposed shear, is definitely relevant in these strongly sheared flows. On the other hand, it seems clear that a three dimensional extension of the present study is necessary for a more realistic description of this instability, especially when it comes to its nonlinear stage.

1012.5284

The large-scale dynamics of plasmas is well described within the framework of magnetohydrodynamics (MHD). However, whenever the ion density of the plasma becomes sufficiently low, the Hall effect is likely to become important. The role of the Hall effect has been studied in several astrophysical plasma processes, such as magnetic reconnection, magnetic dynamo, MHD turbulence, or MHD instabilities. In particular, the development of small-scale instabilities is essential to understand the transport properties in a number of astrophysical plasmas. The magneto-rotational instability (MRI), which takes place in differentially rotating accretion disks embedded in relatively weak magnetic fields, is just one example. The influence of the large-scale velocity flows on small-scale instabilities is often approximated by a linear shear flow. In this paper, we quantitatively study the role of the Hall effect on plasmas embedded in large-scale shear flows. More precisely, we show that an instability develops when the Hall effect is present, which we therefore term as the Hall magneto-shear instability. As a particular case, we recover the so-called MRI and quantitatively assess the role of the Hall effect on its development and evolution.

12168302

2011ApJ...734...47P

1012

1012.2371_arXiv.txt

The solar infrared absorption lines from the fundamental band of the CO molecule at 4666~nm are well-known diagnostics of the temperature minimum and chromosphere of the solar atmosphere. Initial discovery of off-limb emission in the quiet Sun from these lines \citep{noyes1972} instigated the development of thermally bifurcated models of the solar atmosphere, \citep{wiedemann1994, ayres1996} and the latest models draw particular attention to the dynamical events and chemical equilibria in the atmosphere \citep{ascensio2003, wedemeyer2005}. The CO lines are known to show 5 minute period Doppler oscillations \citep{noyes1972, uiten1994, solan1996, uiten2000a} which are presumably driven from the solar global p-modes, and images of the solar surface in the lines reveal that magnetic regions of the solar atmosphere show a reduced molecular line strength \citep{uiten2000a}. Models of the line formation \citep{uiten2000a} suggest the formation height of the strong CO lines at 4666~nm, particularly the 3-2 R14 transition at 4665.8~nm, varies from roughly 150 to 580~km above the $\tau_{500} = 1$ formation height (z=0). Studies of the solar p-modes at various heights in the solar atmosphere have been done using a variety of spectral lines. It is thought that the p-modes are reflected back into the Sun by the steep gradient of density at the photosphere, and the wave solution requires an evanescent component (characterized by an exponential amplitude drop and a constant velocity phase with height) to be present above the reflecting layer. Simultaneous measurements in various spectral lines have shown that for the solar p-mode frequencies the velocity oscillations are in phase at different heights \citep{lites1979} but that the amplitude of the velocity oscillation increases with height \citep{cram1977,ruizcobo1997} In addition to velocity oscillations, intensity oscillations are also observed in the solar atmosphere; since the early work of \cite{schmeider1976} the relationship between the two has become an important diagnostic of the physical conditions of the solar atmosphere. The phase between the temperature and velocity oscillations at various heights in the solar atmosphere has been probed by many authors through the use of line depth and Doppler shift measurements of a different solar spectral lines \citep{deubner1992, al1998, oliviero1999, strauss1999}. A motivation here is to understand if the solar atmosphere reacts adiabatically to the p-mode oscillations or not, and if not, then to determine what other physical processes may be important. Since the velocity oscillation amplitudes increase with height, the solar atmosphere undergoes compression when the velocity oscillations move the gas downwards towards the photosphere, as the higher layers catch up to lower layers. If energy is conserved by the atmosphere then the gas temperature will reach a maximum value when the compression reaches a maximum; this corresponds to the time when the gas velocity is zero just before it begins to move away from the solar photosphere. If the gas temperature is plotted versus the atmospheric velocity (defined as positive when the gas moves away from the photosphere, which is the opposite direction as the observed Doppler shift) in an adiabatic atmosphere there will be a 90 degree phase shift between the two quantities, with the temperature peak leading the atmospheric velocity peak. The phase between the temperature and the atmospheric velocity is referred to as the "I-V phase". In summary, observations from many different spectral lines show that the I-V phase is less than 90 degrees in the photosphere, around 90 degrees from the temperature minimum to the middle chromosphere, and above 90 degrees in the high chromosphere \citep{deubner1990, massiello1998}. Complicating the spectral line measurements of I-V phase is the fact that chromospheric spectral lines are often formed in a non-LTE environment, and the spectral line depth can change non-linearly as the temperature of the solar atmosphere changes. Investigations from Kopp and co-workers \citep{kopp1990,kopp1992} addressed this problem by observing temperature oscillations using thermal infrared continuum channels at wavelengths from 50 to 800 microns from the airborne NASA Kuiper Astronomical Observatory (KAO). The continuum intensity depended only on the continuum opacity, and so the observed intensities were linearly related to the temperature of the solar atmosphere and provided excellent thermal probes. These wavelengths sampled the solar atmosphere from heights of roughly 340 to 820~km above z=0 and thus the temperature minimum through low chromosphere could be explored. Since the continuum observations lacked a way to determine the gas velocity, simultaneous ground-based spectral observations were made using both photospheric and chromospheric spectral lines. The KAO observations showed an I-V phase shift consistent with results from spectral lines which sampled the same atmospheric heights. In particular, the I-V phase was observed to be 122 degrees at 340~km, and decreased with height to approach 95 degrees at 600~km; the solar atmosphere became more adiabatic as one moved from the temperature minimum up to the lower chromosphere. In order to explain the non-adiabatic quality, \cite{kopp1992} introduced a simple thermal relaxation term into the hydrostatic gas thermal balance equation and solved for the phase in a simple model assuming an isothermal plane-parallel non-magnetic atmosphere. The resulting analytic expressions were used to determine the values of the thermal relaxation frequency $\omega_r$. In the following paper we discuss our observations of temperature and Doppler oscillations using the 4666~nm absorption lines of CO. Since the CO spectral lines are formed in LTE conditions, the line depth provides a more direct probe of the gas temperature than other chromospheric lines; and as spectral lines their Doppler shift gives a direct measurement of the atmospheric velocity and thus they are superior to observations using infrared continuum channels. In Section~2 we discuss the data collection, Section~3 we outline the data reduction and analysis procedures, and in Section~4 we discuss the results obtained from the intensity, velocity, and I-V phase measurements from these lines as they vary across the solar surface and with height in the solar atmosphere. We apply the simple analytic model of Kopp to our I-V phase spectra to estimate the gas properties in the solar atmosphere.

The shifts of the CO line centers across the solar disk are consistent with the expected solar photospheric rotation velocity. After removing solar rotation, a low pass temporal filter was applied to the time series in order to extract the supergranulation component of the disk velocity. In these images the supergranulation cells are clearly visible, and a fit to the outflow velocities (both positive and negative lobes) suggests a horizontal outflow velocity of 186~m~s$^{-1}$. This is in agreement with the distribution of outflow velocities seen from Hinode feature tracking measurements in the UV continuum channels (Tian et al 2010) which originate from the same approximate heights. The behavior of the supergranulation outflow Doppler signals with $\mu$ suggests a height variation in the outflow speed, which will be investigated in future work. An analysis of the line bisectors of the CO lines showed no variations, which was likely due to the rather low spectral resolution achieved with these observations. \subsection{Spatial Variations} Using the low pass temporal filter, maps of the line depth across the solar disk were made. The data show that on a global scale the CO line absorption appears weakened in magnetic features associated with the solar magnetic network; confirming results from previous higher resolution studies \citep{uiten1994, ayres1996}. In Figure~2 we show an image produced from a magnetogram taken on the same day from the NSO SOLIS instrument (i.e. \cite{jones2002}); here the absolute value of the magnetogram is taken, and then the map is smeared with box-car averaging to a resolution of about 10 arcseconds (a little sharper than the expected CO map resolution). In Figure~2 we plot the mean CO line depth (averaged over the three spectral lines) across the solar disk, excluding pixels near the limb where seeing and fitting noise dominate the map. There is a very clear correspondence between the position of stronger magnetic fields in the SOLIS map with the regions of weaker CO line depth. The beating of solar p-modes with slightly different frequencies complicate the simple cross-correlation of the line depth and velocity time series. This is especially true if the power distribution among the oscillation modes varies in the line depth compared to the velocity measurements, and such a different power distribution is a natural consequence of non-adiabatic atmospheric conditions. Furthermore, if the I-V phase varies with oscillation frequency we expect even more difficulty with a simple cross-correlation technique. Nevertheless computing the cross-correlation was very robust. As previously noted, the cross-correlation lag was measured in units of seconds of time, but then was converted to an angular phase shift assuming a dominant oscillation period of five minutes. In most cases, the peak of the cross-correlation function occurred with the line depth leading the velocity by about 90 degrees; in some cases however the peak in the cross-correlation function at -270 degrees was higher. To unwrap the measurements, we forced our analysis software to fit the peak near 90 degrees. Figure~2 shows a map of the phase shift across much of the solar disk. As the Sun had very low activity, only small variations are noted; however these variations are significant, and are not tightly correlated with the line depth map of Figure~2b. The phase map shows something different than the line depth map, and presumably the variations represent regions of more or less adiabatic response in the solar atmosphere. As we will discuss in detail later, one interpretation of this is that the phase map show regions with faster and slower thermal relaxation. The Fourier analysis of the line depth and velocity time series was done to more fully investigate the properties of the oscillations. Both line depth and velocity showed power spectrum peaks at the roughly 3.3~mHz, equivalent to a period of five minutes. Unlike other studies of oscillations with these spectral lines, these data show no power for oscillations near 5~mHz (3 minute period). It must be stated that the signal-to-noise ratio of these data benefit significantly from the ability to average across many pixels from the solar surface; in some of the following discussions we average over tens of thousands of pixels, and so our power spectra has a factor of 100 times better signal-to-noise than previous works which report spectra from single positions on the solar disk. The oscillation power integrated from $ 3.0 < \nu < 4.2$~mHz using the CO velocity power spectrum measured pixel-by-pixel across the solar disk is also shown in Figure~2. Regions of high power and lower power are clearly seen. By examining the 500 weakest and strongest line depth pixels on the CO line depth map, (also corresponding to regions of stronger and weaker line-of-sight magnetic fields) a power deficit in magnetic regions is seen as the p-mode power (after background subtraction). In the regions of weaker line depth pixels (stronger magnetic fields) the power is about a factor of 0.5 times the power in the strongest line depth pixels. There is also a weak correspondence between oscillation power and phase, with regions of high (low) power showing non-adiabatic (adiabatic) phases. The interest in observing the full solar disk with these observations was to look for both large and small spatial scale variations of the oscillations parameters. The variation from center-to-limb can be used to characterize the height variation within the solar atmosphere by using the Eddington-Barbier relationship. Here the approximation is made that the source function arises solely from a thin layer at $\tau=1$, and that the height of this layer changes across the solar surface with $\mu$. Using our measured I-V phases, it becomes possible to convert from $\tau$ to height $z$ in the solar atmosphere. As discussed by Kopp et al. (1992), the formation of the infrared continuum in the solar atmosphere is more easily understood than the formation of spectral lines, as the continuum is formed in LTE and the dominant source of opacity is H$^{-}$ free-free absorption. Given a solar atmospheric model, (they use model~C from \cite{vernazza1981}) they determine the formation heights of their 50, 100, 200, and 400~$\mu$ IR continuum channels as 340, 420, 480, and 600~km above $z=0$. By subtracting their observed phase shifts from 180 degrees (since they use Doppler velocity rather than atmospheric velocity), we can compute I-V phase lags for these heights as 122, 122, 109 and 95 degrees respectively (with errors of about $\pm$10 degrees). Figure~3 shows the mean I-V phase shift as a function of $\mu$ from our CO observations. Here we have averaged the phase spectra from 3.0 to 4.0~mHz, and averaged over annuli on the solar disk with a width of $\Delta\mu = 0.01$. The I-V phase drops from about 116 degrees at disk center to about 100 degrees at $\mu=0.5$. These values are clearly within the range of phases observed by Kopp et al. (1992) and we can now determine the atmospheric heights from which they arise, assuming that equal phases in our data correspond to equal heights in the IR continuum data. In detail, we use the Eddington-Barbier relation to replace $ln(\tau)$ with $ln(\mu)$. We determine the relationship between $z$ and $\phi_{I-V}$ from the data from \cite{kopp1990}, and then use the relationship between $\mu$ and $\phi_{I-V}$ in our observations to determine a $ln(\mu)$ vs $z$ function for the formation height of our CO line depth. Our CO I-V phases between $0.5 < \mu < 1.0$ correspond to heights between $560 > z > 425$~km. and the scale height which we derive for the variation of the CO I-V phase gives $H_1=195$~km. We can now examine the power in the CO oscillations as a function of disk position, and therefore height. Figure~4 shows the CO velocity power spectra averaged over the spatial positions from $0.99 < \mu < 1.00$. It shows a significant background power which varies with frequency, and a broad peak near 3.3~mHz which represents the well-known p-mode oscillations. We fit a polynomial function to the log of this background using the frequency ranges of $1.5 < \nu < 2.0$~mHz and $5.0 < \nu < 8.0$~mHz and interpolate the background power through the other frequencies. After subtracting this background, we fit a Gaussian function to the power spectrum. The fits to the power spectrum are shown as the solid line in Figure~4. The fits to the oscillation power across the solar disk were examined. The peak of the p-mode power given by the Gaussian fit was corrected for two projection effects: first, the power was divided by $\mu$ to compute the oscillation power assuming purely radial motions, and second the power was corrected for the smearing in spatial wavenumber caused by projection effects as one moves from the center to the limb. (This correction, basically a distribution of power with spatial wavenumber, was computed by averaging the velocity signal in pixels near disk center and then computing the power spectra of these averaged velocities.) After making these two correction, we transform the $\mu$ coordinate into a height coordinate $z$ and examine the power changes with height. Figure~4 shows that the total power (p-mode and background) increases with height, whereas the p-mode power alone decreases with height. The increase of the total power is consistent with previous studies \citep{ruizcobo1997, simoniello2008} and supports the idea that the velocity oscillation power must increase with height as the density drops. An exponential fit to this background power gives a scale height of $H_2=230$~km, consistent with values used for chromospheric models. By examining the p-mode power alone, after subtracting the background power, we find that it decreases with height, as would be expected from an evanescent solution to the wave equation. This decrease is very rapid; the scale height is much smaller and consistent with photospheric scale heights, $H_3=90$~km. Another parameter derived from the power spectra fitting is the peak central frequency for the p-mode oscillations. While many studies find 3-minute period oscillations at chromospheric heights and lower (using these CO spectral lines) we find no significant power peaks at that frequency. Moreover, the central frequency position of the fit to the p-mode power does not change with height through the range of $560 > z > 425$~km. No 3-minute period power is seen in an analysis of the line depth oscillation power spectra, and the central frequency of the p-mode oscillations as measured with the line depth also remains constant with disk position. \subsection{Diagnostic Diagrams} In Figure~5 we present diagnostic diagrams, or l-$\nu$ diagrams of several parameters from the CO data; these diagrams are produced using the three dimensional Fourier transform of the region near the center of the solar disk as discussed in Section~3. A simple correction for solar rotation was made by aligning the ridge structure for positive and negative frequencies before coadding; but no correction is made for projection effects. The CO velocity l-$\nu$ diagram clearly shows several ridges of power, and these six ridges are at the position of the well-known n=2 through n=7 ridges for the global oscillation modes. The CO line depth diagram is noisier, but shows a couple of ridges at the positions of the global oscillation modes. In both parameters, no distinct ridge structure is seen at high frequencies. At $4< \nu < 5$mHz the ridges fade into the background power; and no other structure is seen at the $\nu=5$mHz frequency where some studies have reported to have seen chromospheric oscillation power. The coherence between the velocity and line depth in Figure~5 also shows strong ridge structure. While the resolutions $\Delta$l and $\Delta\nu$ not especially high, there is some evidence that the coherence between the global p-mode ridges does not go to zero. This is consistent with \cite{oliviero1999} who show significant background power at the inter-ridge positions in measurements of global p-modes near these spatial frequencies. The coherence also shows significant non-zero values at high temporal frequencies, up to about $\nu=6.5$mHz, especially at low spatial degree. Although there are no ridges or peaks at these frequencies suggesting distinct oscillation modes, this large coherence value suggests that waves at this chromospheric oscillation frequency may be seen with these CO lines. The I-V phase diagnostic diagram in Figure~5 shows a correlation with the coherence diagram: the phase values are noisy where the coherence is small, and the phase values change smoothly where the coherence is large. There is very little evidence for a ridge structure in this data, and the phase is rather constant at all spatial wavelengths. The phase does show a gradual decrease as one moves to lower temporal frequencies, from about $\phi=130$~degrees at $\nu=2.5$~mHz to roughly $\phi=90$~degrees at $\nu=4.5$~mHz. (It should be noted that trend this agrees with $180 -\phi$ from \cite{oliviero1999}, assuming that the phase in that work was measured using Doppler velocities.) There is no evidence for a sharp change in the phase or a plateau of low phase below $\nu=2$~mHz as seen in other work \citep{deubner1990}; as suggested by the low coherence at these frequencies, this may simply reflect that the phase in these observations is not well measured here. \subsection{Phase Spectra Fits with Analytic Model} The Kopp model \citep{kopp1990, kopp1992} is based upon several simplifying assumptions. First it uses hydrostatics, thus ignoring the gas dynamics and the magnetic fields which are present in the solar temperature minimum and chromosphere (effectively the magnetic field is assumed to be less than about 200 Gauss). The model is based upon an isothermal atmosphere which simplifies the temperature profile in this region of the atmosphere. And finally, the model assumes a plane-parallel atmosphere which may have implications with these full-disk measurements. Nevertheless, the Kopp model is one of the few analytic models (in addition to \cite{worrall2002}) to date which describes the I-V phase for solar oscillations in physical parameters; because of this it is still very useful. \cite{kopp1990} develop expressions for the I-V phase in terms of the atmospheric acoustic cutoff frequency $\omega_c$, the scale height $H$ the ratio of specific heats $\gamma$ and the thermal relaxation frequency $\omega_r$. The work makes no assumption regarding the actual physical mechanism for the thermal relaxation process, it just assumes that the process follows a Newton cooling law. What is explicitly lacking from the analysis of \cite{kopp1990} is a direct expression for the I-V phase spectrum, $\phi(\nu)$. In the Appendix below, we continue the development from \cite{kopp1990} one extra step to derive and expression for $\phi(\nu)$ which depends solely on the physical variables listed above. Our expression is listed in Equation 5 of the Appendix. We use our expression for $\phi(\nu)$ to fit our I-V phase spectra as determined at different disk positions from our data. The points Figure~6 shows a sample I-V spectrum measured near disk center. An adiabatic model with no thermal relaxation would produce a flat spectrum with a 90 degree phase shift at all frequencies; this is clearly contradicted by the data which show a decrease from about 120 degrees to about 105 degrees from frequencies of about 2.8 to 4.0 mHz (these frequencies are where the I-V coherence is high). Following \cite{kopp1990} we can assume that the acoustic cutoff frequency is $\omega_c / (2\pi) =4.5$mHz and do a one parameter fit to the data to determine the best value of the thermal relaxation frequency $\omega_r$. This 1d fit is shown as the dashed line in Figure~6; it has a high $\chi^2$ value and does not produce a satisfying fit to the data points. If we let the acoustic cutoff frequency be a free parameter and do a 2d fit to the spectrum, we significantly reduce the value of $\chi^2$ and derive a more satisfying fit to the data points. An examination of the $\chi^2$ surface which results from a 2d fit using $\omega_c$ and $\omega_r$ reveals that in some cases the analytic function only weakly constrains $\omega_c$. This combined with the inherent non-linearity of the $atan$ function have foiled our attempts to fit this data using standard non-linear least squares software. Instead, for each I-V spectra at a given disk position ($\mu$) we simply examine the $\chi^2$ over the range of $0 < \omega_r /(2\pi) <10$mHz and $0 < \omega_c /(2\pi) <10$mHz and find the point of minimum $\chi^2$. We estimate the 68\% confidence ranges by using specific $\Delta\chi^2$ intervals on this surface according to standard numerical analysis techniques \citep{press1994}. Figure~6 shows the values we derive from fitting the I-V phase spectra over the range of $0.25 < \mu <1.0$mHz. As in the case of Figure~6, we plot the results for the 1d fit using the dashed line, and the 2d fit results using solid lines. In the 1d case we see that the thermal relaxation frequency $\omega_r$ drops from the center of the disk to the limb; this suggests that as the data probe higher layers of the solar atmosphere the thermal relaxation process becomes less important and the atmosphere becomes more adiabatic. The same behavior for $\omega_r$ is shown with the 2d fits where the relaxation frequency decreases at $\mu$ decreases, although the values are consistently 0.5mHz higher than in the 1d fits. And finally, the value for the acoustic cutoff frequency is around 4.5mHz, although the data suggest the value may significantly drop near the solar limb (and thus at higher heights in the solar atmosphere). The values of $0.5 < \omega_r /(2\pi) < 2.8$mHz correspond to a cooling time scales between 2000 and 360 seconds which are longer than the values of 290 to 76 seconds for this range of heights in the solar atmosphere measured with different fitting methods in \cite{kopp1992}. While the relaxation rate (defined by \cite{kopp1990}) of ${\omega_r} \over {\omega_c}$ which we obtain is similar to the values determined by \cite{kopp1990}, our values of $\omega_r$ are somewhat lower since our fitted values of $\omega_c$ are somewhat lower than the 4.5mHz assumed in their work.

1012.2371

Oscillations were observed across the whole solar disk using the Doppler shift and line center intensity of spectral lines from the CO molecule near 4666 nm with the National Solar Observatory's McMath/Pierce solar telescope. Power, coherence, and phase spectra were examined, and diagnostic diagrams reveal power ridges at the solar global mode frequencies to show that these oscillations are solar p-modes. The phase was used to determine the height of formation of the CO lines by comparison with the IR continuum intensity phase shifts as measured in Kopp et al. we find that the CO line formation height varies from 425 km &lt; z &lt; 560 km as we move from disk center toward the solar limb 1.0 &gt; μ &gt; 0.5. The velocity power spectra show that while the sum of the background and p-mode power increases with height in the solar atmosphere as seen in previous work, the power in the p-modes only (background subtracted) decreases with height. The CO line center intensity weakens in regions of stronger magnetic fields, as does the p-mode oscillation power. Across most of the solar surface the phase shift is larger than the expected value of 90° for an adiabatic atmosphere. We fit the phase spectra at different disk positions with a simple atmospheric model to determine that the acoustic cutoff frequency is about 4.5 mHz with only small variations, but that the thermal relaxation frequency drops significantly from 2.7 to 0 mHz at these heights in the solar atmosphere.

12213456

2011MNRAS.412.2445P

1012

1012.0004_arXiv.txt

\label{sec:intro} In recent years, stellar population synthesis (SPS) models have rapidly become a fundamental tool in the study of both Galactic and extragalactic stellar populations (see \citealt{basti4,galev09,coelho07,starb99,maraston05,bc03} for some recent examples). Despite the increasing sophistication of the underlying stellar models used in SPS and ever-growing libraries of spectra (both empirical and synthetic) available to modellers, there are still areas in which models, and methods for fitting them to observational data, can be significantly improved and expanded. Some of the current shortcomings are due to the need to adopt simplifying assumptions about the stellar population in question in order to make some problems tractable, e.g. the fitting of simple stellar population (SSP) models to galaxies, assuming that they can be represented by a single age, single metallicity model. A key area in which simplifying assumptions are implicitly made is the morphology of the horizontal branch (HB) for individual SSPs -- this is largely due to the limitations of theoretical isochrones, which are used as the basis of all SPS models (except those of Maraston and co-workers, who use a fuel-consumption based method -- see \citealt{maraston05} and references therein). However it is well known that the morphology of the HB can impact strongly on the integrated light of stellar populations especially if an extended blue component is present, potentially affecting both colours and line indices, and hence impacting on inferred ages for these systems (see e.g. \citealt{ocvirk,conroy09,schi04,lee2002,lee2000}). Hence it is vital to assess the ways which these `simplified' HBs inherent in theoretical isochrones impact on SPS models, and whether a more realistic and detailed treatment of the HB is necessary, or indeed practical. This is the purpose of the work presented here. For unresolved stellar populations our knowledge of their ages and elemental abundances is largely derived through the fitting of diagnostic spectral indices, hence the focus of our work is to investigate the effect of a detailed treatment of the HB on key spectral indices, such as H$\beta$. However, it is important to realise that our results will also impact on analyses which use full-spectrum fitting methods, as will be discussed in Section~\ref{sec:disc}. In real stellar systems the horizontal extension of the HB is governed by stochastic mass loss in stars approaching the tip of the red giant branch (TRGB). All the stars in a particular SSP leave the TRGB with the same core mass but individual stars have different masses remaining in the outer layers, which determine the position of the star along the zero-age horizontal branch (ZAHB). In stellar models this mass loss is parameterised by the mass loss parameter, $\eta$, which is assigned some fixed value calculated according to the Reimers mass loss relation \citep{reimers}. Hence in theoretical isochrones the HB comprises a single mass point, with no extension. It is worth remembering here that an isochrone consists of a series of evolutionary points (EPs) which define the locus of points for an SSP. Whilst each EP defines the appropriate stellar parameters for any star located at that point, in terms of mass, effective temperature and luminosity, they do not represent individual stars themselves. In order to create an SSP, an isochrone is ``populated'' according to some initial mass function (IMF) which effectively gives the appropriate weighting to each EP. In most population synthesis work the isochrone is treated as an analytical function for this purpose, so that all EPs along the isochrone are smoothly populated and the ZAHB remains as a single mass point. In order to create a model with an extended HB, effectively incorporating a spread in mass loss, the isochrone (or, at least, the HB portion) must be populated with a discrete number of stars, so that each star in the HB phase can be assigned a specific mass from within some range of masses -- this requires interpolation between individual core-helium burning stellar tracks. Creating an integrated spectrum for these extended HB models is a computationally expensive procedure since the time taken is proportional to the number of points in the simulation (for details on how integrated spectra for SSPs are produced the interested reader is referred to \citealt{basti4}, hereafter P09). For the analytical case, the number of points is just the number of EPs along the isochrone, which varies between a few hundred and around two thousand depending on which isochrone set is used. To populate an isochrone with individual stars, one has to consider the problem of statistical fluctuations which are likely to arise as a result of low numbers of stars (and hence poor sampling) in the later stages of evolution, including the HB phase. This is of particular relevance in the case of an SSP with an extended blue HB, since these hot blue HB stars give rise to strong Balmer lines, which are generally used as the primary age indicators for stellar populations. This problem can only be overcome by using large numbers of points in the simulation, potentially up to $\sim10^6$. In fact, part of the work presented here is to explore how large these statistical fluctuations can be in real stellar clusters, and how many input stars are needed in the models to avoid significant uncertainties in ages and metallicities derived from the final integrated SSP spectra. Another problem in creating models with extended HBs is that there is still no theory which predicts mass loss rates from stellar parameters and so the value of $\eta$ is chosen arbitrarily, usually to reflect the typical HBs seen in Galactic globular clusters. This is complicated by the fact that clusters apparently with the same age and chemical composition can have different HB morphologies. However, in principle $\eta$ can take any value and so producing a database of SPS models to cover all possibilities is not feasible. The pioneering work of \citet{lee2000} (hereafter L2000) highlighted the importance of including realistic HBs in stellar population work. L2000 created models for 15 SSPs using a single value for the mass loss parameter for all their models, which was chosen to replicate the range of HBs seen in Galactic globular clusters. They included a Gaussian spread in the HB mass distribution of $\sigma_{M}$=0.02$M_\odot$ to simulate the observed extension of the HB. This work graphically demonstrated the systematic variation of HB morphology with age and metallicity (see their Figure 5), showing that, at fixed $\eta$, HBs generally become bluer as metallicity decreases and as age increases. L2000 note that the strength of the H$\beta$ line in the integrated population does not increase monotonically as metallicity decreases at a fixed age, but peaks at some intermediate metallicity and then falls again as metallicity continues to decrease. This is because H$\beta$ reaches a maximum strength in stellar spectra when the effective temperature, $T_{eff}$, is around 9500K and so H$\beta$ is maximised if the distribution of stars along the HB centres around this value. In these cases the contribution to H$\beta$ from the HB completely dominates over the contribution from the main-sequence turnoff and makes the population look spuriously young. It is apparent from the L2000 figures, but not specifically noted by them, that the actual extension (i.e. from red to blue) of the HB in the colour-magnitude diagrams is very different for the various SSPs, even though the value of $\eta$ and $\sigma_{M}$ are the same. This is because the stellar effective temperature in this phase is extremely sensitive to very small changes in envelope mass and so the resulting HB morphologies can be very different, even when the same mass loss prescription is used. In this paper we use similar techniques to L2000 to model extended HBs, populating them with discrete numbers of stars by interpolating between core-helium burning stellar tracks and using a Gaussian spread in the mass loss distribution, and we extend that work in 3 key ways. Firstly we explore the effects of varying $\eta$ at fixed age and metallicity, which enables us to model high mass loss in metal rich systems to see whether we can mimic a system such as M32, which has approximately solar metallicity, but an extended blue HB component. We are also able to create SSPs at different metallicities with bimodal HBs, similar to those observed in several Galactic globular clusters. Secondly, we create several models which we populate with numbers of stars representative of typical stellar clusters, and use Monte Carlo techniques to explore the impact of statistical fluctuations on measured diagnostic line indices, such as H$\beta$. Thirdly, we create full integrated high resolution spectra for all our models and simulations which enables us to assess the behaviour of all diagnostic line indices (within a wavelength range 2500\AA\ to $\sim$6000\AA), as well as the continuum flux. We also investigate whether there are any diagnostic indices capable of breaking the degeneracy between an old SSP with extended blue HB (hence strong H$\beta$) and a truly young or intermediate age SSP.

\label{sec:disc} To summarise, we have created integrated spectra for 16 SSPs, 4 each at 4 different metallicities, all with an underlying age of 14~Gyr, with a range of extended HB morphologies. This was done by varying the mass loss prescription for each individual SSP in 2 ways, firstly setting a mean mass, $<M>$, for stars coming on to the ZAHB, and then adding a spread in mass loss, $\sigma_{M}$. We find that the H$\beta$ strength for each SSP depends on the exact temperature distribution of stars along the HB, which in turn depends on the exact details of the mass loss prescription coupled with the metallicity of the population in question. For any of the modelled SSPs with any amount of blue HB extension, H$\beta$ is increased relative to the fixed $\eta$=0.2 case (i.e. the `standard' models), implying younger ages than the actual SSP age of 14~Gyr. In the worst-case scenario modelled here, a solar metallicity 14~Gyr population with a blue HB which has a peak in its distribution of stars around 9000K, has an implied age of around 2~Gyr from the strength of the H$\beta$ line. Our preliminary investigation to identify spectral features which might be capable of breaking the degeneracy between an old SSP with extended blue HB, and a truly intermediate age or young SSP, indicates that the Ca{\sc ii} index defined by \cite{rose84} is a very promising candidate, at least for populations with low velocity dispersions, such as extra-galactic globular clusters (but see caveat in the following paragraph). There is also tentative evidence that the Mg{\sc ii} index, in combination with Mg$b$, could also be a very useful tracer of hot blue HBs. There are also indications that the UV continuum shortward of the Mg{\sc ii} feature could also potentially be a useful additional tool, although more work is needed to model this part of the spectrum in the required detail. However, it is interesting to note that the extreme blue HB models (i.e. models 3 and 4 presented here) have very strong UV continuum flux, as demonstrated by Figure~\ref{fig:diffspec}, even at solar metallicity. This supports the idea that extreme blue HBs can be a significant contributor to the UV upturn identified in old elliptical galaxies (see review by \citealt{oconnell}). An important caveat to all the tests and results presented here is that all our models have been created using SSPs, i.e. single age, single metallicity systems. However another scenario which could explain strong H$\beta$ lines in a predominantly old population would be the presence of a small fraction of very young stars. Although this is not likely to be an issue for globular clusters (Galactic or extra-galactic) it is a possibility which is hard to completely rule out for elliptical galaxies. Preliminary tests with our population synthesis code indicate that adding even a very small percentage ($<$1\%) of a young SSP (300~Myr or less) to a 14~Gyr SSP would significantly strengthen H$\beta$ and imply an intermediate age, around 5-6~Gyr (see also \citealt{serra}). In fact, this scenario of a `frosting' of young stars in an otherwise old stellar population has been explored by \citet{smith} for (apparently) quiescent galaxies in the Shapley supercluster. However, \citet{smith} use the Ca{\sc ii} index as a tracer of hot stars and, as we have shown here, this index also traces hot HB stars in an SSP. So far we have not identified any completely unambiguous tracers that can distinguish between the 3 scenarios; 1) a small fraction of hot young stars in an otherwise old population, 2) an old SSP with a hot HB, 3) an intermediate age SSP. It seems likely that no single tracer will be able to disentangle these 3 cases and a combination of well understood spectral indices and colours may well be needed -- this is the focus of our ongoing investigation. Another factor potentially linked to extended blue HBs is the presence of a stellar subpopulation with a high helium fraction. \citet{lee05} showed that the same level of helium enrichment required to reproduce the bluer main sequence in the massive GC $\omega$ Centauri would also naturally produce the extreme blue HB stars seen in that cluster. Very recent observational results have provided some evidence for varying helium fractions within individual GCs and there is at least circumstantial evidence linking helium enriched subpopulations to extended blue HBs (see \citealt{brag2010} and references therein). The theoretical support for the HB morphology--enhanced-helium connection is that stars born with a higher helium abundance will display a lower mass at the turn off for a given cluster age. A lower mass at the TO will favour lower mass -- hence hotter and bluer -- HB stars. We were able to test the potential impact of a higher helium abundance on our present work by utilising the helium-enhanced isochrones in the BaSTI database in combination with a grid of helium-enhanced spectra, which have been produced for a subsequent paper (Coelho, Salaris \& Percival, 2011, in preparation). We chose helium-enhancement at the level $Y=0.3$, for the [Fe/H]=$-0.7$ isochrone, which is a reasonable enhancement over the standard cosmological value, given current observational constraints (e.g. \citealt{brag2010}) and our new spectra incorporate the same Fe and He abundance ratios as the isochrones. As a preliminary test, we measured the strength of spectral features in individual spectra for both the standard helium and enhanced helium versions, matching the spectra in $T_{eff}$ and log$g$. We found that for all the diagnostic lines considered in this paper, the line strengths are practically indistinguishable for the two cases, with any differences being less than 1 percent (i.e. within the measurement errors). We also created an integrated spectrum for the [Fe/H]=$-0.7$, 14~Gyr SSP using the helium-enhanced isochrone in combination with the helium-enhanced spectra. Again we found that the strengths of all the indices under scrutiny here are negligibly different from the standard helium case. We note here that this result is consistent with that of \citet{girardi2007} who find that a similar level of helium enhancement has a negligible effect on broadband colours. The important point to stress here is that, whilst enhanced helium is a likely mechanism for producing extended blue HBs, it in no way impacts on the work presented in this study. Our diagnostic models and tests are not intended to predict, $a~priori$, the existence of hot HB stars, but rather to find diagnostics which can distinguish their presence in a stellar population from other hot components, such as a young sub-population. Whether a blue HB morphology arises from a large mass loss along the RGB in stars with `normal' helium, or from a more moderate mass loss from stars with higher helium, the resulting integrated spectrum is largely unaffected -- the only parameter that matters here is the temperature range of stars on the HB, irrespective of how they have been produced. In fact, the results of our enhanced helium tests, described above, show that enhanced helium has absolutely no impact on the diagnostic indices discussed here, such as H$\beta$, Mg{\sc ii} and Ca{\sc ii}, nor on the results of the statistical fluctuations tests discussed below. Finally, the results of the tests presented here in Section~\ref{sec:statfluc} demonstrate that several key diagnostic line indices are significantly affected by statistical fluctuations, even in average to large mass GCs. Equally problematic is the fact that not all indices fluctuate at the same level -- some are strongly affected, including all the Balmer lines, whilst others are only negligibly affected, such as Ca{\sc ii} (see Table 3). This is potentially a significant source of uncertainty if full-SED fitting methods are used to derive ages, metallicities and/or star formation histories from integrated spectra, since different features in a spectrum can be giving conflicting best fit parameters. The problem is illustrated in \citet{dias} who present integrated spectra for 14 Magellanic Cloud stellar clusters for which they derive best fit ages and metallicities using 2 different fitting codes, {\sc starlight} \citep{cidf} and $ULySS$ \citep{koleva}, and 3 different sets of SSP models, from \citet{bc03}, \citet{pegase} (PEGASE-HR) and \citet{vaz2010}. Half of the clusters studied have masses in the range $1-2.5\times10^4M_\odot$, which corresponds to the low mass cluster modelled here, in Section~\ref{sec:statfluc}. For several clusters, \citet{dias} find that the results from the different fitting routines and models are completely discrepant. As an example, cluster HW1 yields best-fit ages of 3.2, 5.8, 7.9, 9.0, 9.4 and 10~Gyr from the 6 different combinations of SSPs models and SED-fitting routines, whilst its actual age is known to be around 6~Gyr from isochrone fitting to the CMD. Other clusters are even more unconstrained, yielding ages ranging from $<$1~Gyr to 10~Gyr for the same cluster. Some of these variations in best fit parameters are likely to be due to systematic uncertainties in stellar parameters inherent within the SSP models themselves (see \citealt{smp}), but statistical fluctuations in the observed spectra are almost certainly contributing to the problem. Hence we urge caution if these types of stellar clusters are to be used as empirical calibrating objects for various aspects of SPS models.

1012.0004

The presence of an extended blue horizontal branch (HB) in a stellar population is known to affect the age inferred from spectral fitting to stellar population synthesis models. This is due to the hot blue component which increases the strength of the Balmer lines and can make an old population look spuriously young. However, most population synthesis models still rely on theoretical isochrones, which do not include realistic modelling of extended HBs. In this work, we create detailed models for a range of old simple stellar populations (SSPs), with metallicities ranging from [Fe/H]=-1.3 to solar, to create a variety of realistic HB morphologies, from extended red clumps, to extreme blue HBs. We achieve this by utilizing stellar tracks from the BaSTI data base and implementing a different mass-loss prescription for each SSP created. This includes setting an average mass and a Gaussian spread in masses of individual stars coming on to the zero-age HB for each model, and hence resulting in different HB morphologies. We find that, for each metallicity, there is some HB morphology which maximizes Hβ, making an underlying 14-Gyr population look ∼5-6 Gyr old for the low- and intermediate-metallicity cases, and as young as 2 Gyr in the case of the solar metallicity SSP. We explore whether there are any spectral indices capable of breaking the degeneracy between an old SSP with extended blue HB and a truly young or intermediate-age SSP, and find that the Ca II index of Rose and the strength of the Mg II doublet at 2800 Å are promising candidates, in combination with Hβ and other metallicity indicators, such as Mgb and Fe5406. We also run Monte Carlo simulations to investigate the level of statistical fluctuations in the spectra of typical stellar clusters. We find that fluctuations in spectral indices are significant even for average to large globular clusters and that various spectral indices are affected in different ways, which has implications for full-spectrum fitting methods. Hence, we urge caution if these types of stellar clusters are to be used as empirical calibrating objects for various aspects of stellar population synthesis models.

4018667

2011A&A...526A..99S

1012

1012.1528_arXiv.txt

The discovery of exoplanets continues at a very high rate and has recently passed the 450th detection. The radial velocity technique gives strong input for this number, supported by the several dedicated observing programs that almost continually observe stellar spectra in different high-resolution instruments spread across the world. The HARPS spectrograph is one of the leading instruments that over the past five years has alone spotted more than 85 of the exoplanets now known to orbit stars other than the Sun. All of these new discoveries are providing new clues for the formation and evolution of stars and planets. One of these clues is the very well known and established correlation between the metallicity of the stars and the presence of an orbiting giant planet \citep[][]{Gonzalez-1997, Gonzalez-2001, Santos-2001, Santos-2004b, Fischer_Valenti-2005, Udry-2006, Udry-2007b}. This observational correlation suggests that giant planets are more easily formed around metal-rich stars, supporting the core accretion idea as the main mechanism in the formation of giant planets \citep[][]{IdaLin-2004, Benz-2006} instead of the alternative model focused on the idea of the disk instability\citep[][]{Boss-2002}. Although this correlation seems to be true for giant planets, it might not be so for lower mass planets. With the new discoveries reaching lower and lower masses, the new ``lighter systems`` are starting to reveal that these new planet-host stars present a different and wider metallicity distribution \citep[][]{Sousa-2008}. That can also be explained by recent models of the core accretion idea \citep[][]{Mordasini-2009}. \begin{table*}[!t] \centering \caption[]{Details of the spectral data for each star.} \begin{scriptsize} \begin{tabular}{lccccc|lccccc} \hline \hline \noalign{\smallskip} Name & totsp & sumsp & maxsn & minsn & sumsn & Name & totsp & sumsp & maxsn & minsn & sumsn \\ \hline \object{HD\,102200} & 7 & 3 & 164.50 & 128.10 & 251.75 & \object{HD\,197890} & 5 & 0 & 0.00 & 0.00 & 0.00 \\ \object{HD\,104800} & 6 & 6 & 123.70 & 70.20 & 235.46 & \object{HD\,199288} & 15 & 11 & 411.10 & 272.80 & 1137.58 \\ \object{HD\,105004} & 5 & 5 & 81.60 & 45.10 & 138.91 & \object{HD\,199289} & 5 & 3 & 140.80 & 98.60 & 201.48 \\ \object{HD\,107094} & 12 & 7 & 123.40 & 95.60 & 287.16 & \object{HD\,199604} & 6 & 4 & 157.10 & 144.00 & 301.09 \\ \object{HD\,108564} & 6 & 6 & 130.70 & 75.70 & 246.54 & \object{HD\,199847} & 7 & 6 & 118.40 & 67.60 & 210.07 \\ \object{HD\,109310} & 15 & 9 & 180.80 & 127.40 & 434.18 & \object{HD\,206998} & 6 & 4 & 132.00 & 92.80 & 224.53 \\ \object{HD\,109684} & 6 & 4 & 142.20 & 78.00 & 239.40 & \object{HD\,207190} & 5 & 2 & 232.50 & 134.30 & 268.50 \\ \object{HD\,111515} & 5 & 2 & 193.30 & 188.60 & 270.06 & \object{HD\,207869} & 17 & 17 & 135.30 & 41.20 & 425.99 \\ \object{HD\,111777} & 6 & 4 & 178.20 & 126.30 & 309.55 & \object{HD\,210752} & 17 & 12 & 229.60 & 159.50 & 697.15 \\ \object{HD\,113679} & 6 & 6 & 92.60 & 37.80 & 192.95 & \object{HD\,215257} & 37 & 31 & 288.30 & 108.50 & 976.52 \\ \object{HD\,11397 } & 33 & 33 & 151.40 & 33.50 & 649.63 & \object{HD\,218504} & 15 & 15 & 192.00 & 73.60 & 582.42 \\ \object{HD\,119949} & 5 & 2 & 192.60 & 192.50 & 272.31 & \object{HD\,221580} & 54 & 54 & 127.20 & 42.40 & 649.58 \\ \object{HD\,121004} & 5 & 5 & 115.70 & 50.00 & 178.49 & \object{HD\,223854} & 4 & 3 & 154.50 & 134.30 & 253.96 \\ \object{HD\,123517} & 9 & 6 & 93.60 & 70.30 & 209.56 & \object{HD\,224347} & 8 & 5 & 159.90 & 33.50 & 261.12 \\ \object{HD\,124785} & 17 & 17 & 131.20 & 46.10 & 397.18 & \object{HD\,224817} & 30 & 23 & 190.70 & 119.50 & 726.06 \\ \object{HD\,126681} & 14 & 13 & 103.70 & 49.50 & 285.61 & \object{HD\,22879} & 36 & 19 & 362.70 & 224.30 & 1280.67 \\ \object{HD\,126793} & 7 & 4 & 191.80 & 94.30 & 292.44 & \object{HD\,25704} & 20 & 8 & 190.70 & 124.60 & 458.90 \\ \object{HD\,126803} & 7 & 6 & 98.00 & 43.20 & 188.84 & \object{HD\,31128} & 37 & 37 & 127.40 & 38.70 & 600.57 \\ \object{HD\,128340} & 5 & 3 & 140.30 & 70.30 & 177.67 & \object{HD\,38510} & 5 & 2 & 151.90 & 116.20 & 191.25 \\ \object{HD\,128575} & 2 & 0 & 0.00 & 0.00 & 0.00 & \object{HD\,40865} & 30 & 20 & 169.70 & 112.20 & 599.71 \\ \object{HD\,129229} & 5 & 5 & 165.50 & 47.10 & 225.19 & \object{HD\,51754} & 21 & 21 & 143.00 & 50.00 & 472.26 \\ \object{HD\,131653} & 4 & 4 & 111.20 & 59.70 & 186.37 & \object{HD\,56274} & 14 & 10 & 222.60 & 161.80 & 611.40 \\ \object{HD\,134088} & 4 & 3 & 219.00 & 154.00 & 311.59 & \object{HD\,59984} & 45 & 33 & 498.40 & 241.90 & 1774.18 \\ \object{HD\,134113} & 28 & 23 & 184.30 & 73.20 & 611.32 & \object{HD\,61902} & 7 & 4 & 193.50 & 106.10 & 270.13 \\ \object{HD\,134440} & 10 & 10 & 124.50 & 52.60 & 292.34 & \object{HD\,62849} & 17 & 17 & 101.80 & 49.50 & 300.91 \\ \object{HD\,144589} & 11 & 11 & 91.10 & 45.80 & 249.85 & \object{HD\,68089} & 7 & 7 & 84.60 & 44.80 & 177.31 \\ \object{HD\,145344} & 5 & 4 & 139.20 & 65.50 & 187.81 & \object{HD\,68284} & 10 & 5 & 220.40 & 114.30 & 371.61 \\ \object{HD\,145417} & 5 & 2 & 244.00 & 127.90 & 275.49 & \object{HD\,69611} & 6 & 3 & 188.10 & 123.40 & 260.23 \\ \object{HD\,147518} & 4 & 3 & 104.60 & 60.20 & 158.86 & \object{HD\,75745} & 14 & 9 & 115.50 & 76.60 & 288.92 \\ \object{HD\,148211} & 34 & 30 & 248.10 & 72.80 & 826.44 & \object{HD\,77110} & 16 & 16 & 143.20 & 55.80 & 458.88 \\ \object{HD\,148816} & 7 & 4 & 260.90 & 161.90 & 413.57 & \object{HD\,78747} & 26 & 17 & 237.80 & 165.80 & 832.53 \\ \object{HD\,149747} & 9 & 9 & 93.00 & 26.60 & 166.58 & \object{HD\,79601} & 16 & 11 & 215.10 & 134.70 & 588.14 \\ \object{HD\,150177} & 30 & 17 & 475.60 & 177.20 & 1135.11 & \object{HD\,88474 } & 6 & 3 & 136.00 & 63.70 & 188.49 \\ \object{HD\,161265} & 2 & 1 & 48.50 & 48.50 & 48.50 & \object{HD\,88725} & 22 & 16 & 233.40 & 185.50 & 857.18 \\ \object{HD\,164500} & 2 & 2 & 89.40 & 81.60 & 121.04 & \object{HD\,90422} & 7 & 4 & 208.30 & 85.60 & 291.89 \\ \object{HD\,167300} & 9 & 9 & 109.70 & 47.10 & 255.80 & \object{HD\,91345} & 8 & 8 & 94.80 & 30.70 & 160.61 \\ \object{HD\,16784 } & 3 & 1 & 160.60 & 160.60 & 160.60 & \object{HD\,94444} & 7 & 4 & 166.30 & 90.00 & 264.05 \\ \object{HD\,171028} & 48 & 39 & 184.80 & 85.70 & 835.80 & \object{HD\,95860} & 7 & 7 & 94.40 & 48.20 & 180.14 \\ \object{HD\,171587} & 14 & 14 & 169.60 & 75.60 & 492.28 & \object{HD\,967} & 34 & 28 & 175.90 & 108.30 & 770.54 \\ \object{HD\,175179} & 3 & 3 & 113.70 & 91.00 & 175.98 & \object{HD\,97320 } & 6 & 4 & 183.50 & 95.00 & 263.43 \\ \object{HD\,17548 } & 10 & 9 & 185.90 & 53.60 & 252.61 & \object{HD\,97783} & 6 & 4 & 132.00 & 105.40 & 227.50 \\ \object{HD\,175607} & 7 & 5 & 134.10 & 85.90 & 246.40 & BD+062932 & 4 & 4 & 68.10 & 51.00 & 119.78 \\ \object{HD\,17865 } & 21 & 17 & 184.50 & 98.60 & 591.48 & BD+063077 & 1 & 1 & 34.10 & 34.10 & 34.10 \\ \object{HD\,181720} & 29 & 21 & 228.20 & 114.10 & 704.80 & BD+083095 & 3 & 3 & 80.10 & 44.00 & 103.79 \\ \object{HD\,187151} & 1 & 1 & 105.00 & 105.00 & 105.00 & BD-004234 & 1 & 1 & 15.10 & 15.10 & 15.10 \\ \object{HD\,190984} & 46 & 44 & 158.40 & 51.90 & 718.18 & BD-032525 & 2 & 1 & 52.10 & 52.10 & 52.10 \\ \object{HD\,193901} & 3 & 3 & 107.60 & 57.20 & 137.78 & BD-084501 & 4 & 4 & 62.30 & 37.10 & 101.90 \\ \object{HD\,195633} & 4 & 2 & 160.40 & 93.60 & 185.71 & CD-231087 & 35 & 35 & 84.50 & 41.00 & 421.16 \\ \object{HD\,196892} & 3 & 3 & 98.80 & 66.80 & 150.62 & CD-436810 & 9 & 8 & 88.20 & 43.40 & 171.55 \\ \object{HD\,197083} & 12 & 12 & 128.40 & 52.30 & 353.37 & CD-4512460 & 4 & 3 & 52.10 & 40.50 & 81.31 \\ \object{HD\,197197} & 21 & 18 & 207.60 & 83.20 & 486.02 & CD-452997 & 6 & 1 & 24.20 & 24.20 & 24.20 \\ \object{HD\,197536} & 3 & 2 & 157.80 & 129.00 & 203.82 & CD-571633 & 7 & 5 & 107.00 & 66.00 & 202.12 \\ \hline \end{tabular} \tablefoot{ \textit{totsp} is the total number of spectra observed; \textit{sumsp} is the number of spectra used for the combination of the final spectrum for each star; \textit{maxsn} is highest S/N value from the combined spectra for each star; \textit{minsn} is the lowest value from the combined spectra for each star ; and \textit{sumsn} is the final S/N for the combined spectrum.} \end{scriptsize} \label{tabspec} \end{table*} Several programs have been compiled to try to understand the distribution of the planets and the metallicity correlation. In particular, some are focused on metal-poor stars with the goal of not only checking the frequency of giants and low-mass planets in these stars, but also of measuring the lower limit in metallicity where it is possible to form and find giant planets. One of these programs is part of the HARPS GTO planet search program \citep[][]{Mayor-2003}. In this paper we present the precise derivation of fundamental spectroscopic stellar parameters and make an estimate of the masses for the stars in this sample. We present a catalog of spectroscopic stellar parameters for the metal-poor sample observed with HARPS to search planets. In Sect. 2 we describe the observations with the HARPS spectrograph. Section 3 describes the procedure used to derive precise spectroscopic stellar parameters and to estimate for the stellar masses and new spectroscopic parallaxes based on the derived parameters. In Sect. 4, we compare our temperature values with the ones obtained with an IRFM (infra-red flux method) calibration to check for consistency. In Chapter 5 we redo a calibration for the temperature as a function of \textit{B-V} and [Fe/H] using the new parameters derived in this work that were added to data from a previous work. Finally in Chapter 6 we summarize the work presented here.

In this work we presented precise stellar parameters for a sample of metal-poor stars. The stellar parameters were derived in a consistent way following the same method as in previous works. It is crucial that the spectroscopic parameters for these stars are derived precisely and in a systematic way to allow correct comparison between stars. This is very useful when searching for clues for the stellar and planet formations that are typically done by comparing the stars with detected planets to the single stars that do not show any evidence of hosting any planet. These parameters and abundances will be used to study the frequency of planets as a function of the stellar parameters, but this is beyond the scope of this paper. ARES is a very important tool for this task, not only because it is an automatic tool that allows faster and more completely analysis of the spectra of many stars, but more importantly, because it clearly allows eliminating of most of the human factor that was creating larger errors in the spectral analysis (more specifically the subjective position of the continuum located by eye when measuring the EWs of the lines with interactive routines). We also present estimations for the mass of these stars using similar procedures as in previous works. Here it was necessary to overtake the problem of the high errors in the Hipparcos parallaxes presented for most of the stars in the sample. Therefore we estimated a second value for the mass by assuming a different parallax based on the derived spectroscopic parameters. The effective temperature for these metal-poor stars was tested and compared against an IRFM calibration. The comparison between the two different approaches to derive the effective temperature are consistent, meaning that our spectroscopic method is still valid for lower metallicity stars. Finally a new calibration for the effective temperature as a function of the color index \textit{B-V} and [Fe/H] is presented where the metallicity range is now wider thanks to using the parameters derived in this work.

1012.1528

Stellar metallicity strongly correlates with the presence of planets and their properties. To check for new correlations between stars and the existence of an orbiting planet, we determine precise stellar parameters for a sample of metal-poor solar-type stars. This sample was observed with the HARPS spectrograph and is part of a program to search for new extrasolar planets. The stellar parameters were determined using an LTE analysis based on equivalent widths (EW) of iron lines and by imposing excitation and ionization equilibrium. The ARES code was used to allow automatic and systematic derivation of the stellar parameters. Precise stellar parameters and metallicities were obtained for 97 low metal-content stars. We also present the derived masses, luminosities, and new parallaxes estimations based on the derived parameters, and compare our spectroscopic parameters with an infra-red flux method calibration to check the consistency of our method in metal poor stars. Both methods seems to give the same effective temperature scale. Finally we present a new calibration for the temperature as a function of B-V and [Fe/H]. This was obtained by adding these new metal poor stars in order to increase the range in metallicity for the calibration. The standard deviation of this new calibration is ~50 K. <P />Based on observations collected at the La Silla Parana Observatory, ESO (Chile) with the HARPS spectrograph at the 3.6-m telescope (ESO runs ID 72.C-0488, 082.C-0212, and 085.C-0063).Tables 1-3 are only available in electronic form at <A href="http://www.aanda.org">http://www.aanda.org</A>

12224479

2011PhRvD..83d3523A

1012

1012.1833_arXiv.txt

\label{into} Over the past decade the cosmic microwave background (CMB) has emerged as a fundamental probe of cosmology and astrophysics. In addition to the primary fluctuations of the early Universe, the CMB contains signatures of the gravitational bending of CMB photon trajectories due to matter, called gravitational lensing. Mapping this gravitational lensing is important for a number of reasons including, but not limited to, understanding cosmic structure, constraining cosmological parameters \cite{Kaplin, Smth2006} and detecting gravity waves \cite{knox2002, Kesden, SelH}. In this paper we present a local Bayesian estimate that can accurately map the gravitational lens in high resolution, low noise measurements of the CMB temperature and polarization fields. There is extensive literature on estimating the lensing of the CMB (classic references include \cite{ZaldSel1999, HuOka2002,HiraSel2003b}) and some recent observational detections \cite{Smith2007,Hira2008}. The current estimators in the literature can be loosely characterized into two types. The first type was initiated in \cite{ZaldSel1999} (see also \cite{SelZlad1999,Guzik2000}) and utilizes quadratic combinations of the CMB and its gradient to infer lens structure. The optimal quadratic combinations were then discovered by \cite{Hu2001b, HuOka2002, OkaHu2003} and are generally referred to as \lq the quadratic estimator'. This is arguably the most popular estimate of the gravitational potential and uses a first order Taylor approximation to establish mode coupling in the Fourier domain which can be estimated to recover the gravitational potential (real space analogs to these estimators can be found in \cite{buncher, carv}). The second type is an approximate global maximum likelihood estimate and was developed in \cite{HiraSel2003a, HiraSel2003b}. Our method, in contrast, locally approximates a quadratic form for the gravitational potential and estimates the coefficients locally using Bayesian methods. The locally estimated coefficients are then globally {\it stitched together} to construct an estimate of a low pass filter of the gravitational potential. The local analysis allows us to avoid using the typical first order Taylor expansion for the quadratic estimator and avoids the likelihood approximations used in global estimates. Moreover, we are able to easily handle missing pixels, problems with partial sky observations (E and B mode mixing, for example), and spatially varying or asymmetric beams. The motivation for developing this estimate stems, in part, from current speculation that likelihood methods will allow superior mapping of the lensing structure (compared to the quadratic estimator) under low noise levels, and that global likelihood methods can be prohibitively computational intensive---indeed intractable---without significant approximation. \begin{figure*} \includegraphics[height=6.9cm]{PotEst1}% \includegraphics[height=6.9cm]{PotTrue1} \caption{{\it Left:} Estimated gravitational potential on a $17^o\times 17^o$ patch of the simulated flat sky. {\it Right:} Input gravitational potential used in the simulation. See Section \ref{into} and Appendix \ref{SimDets} for the simulation details.\label{fig1}} \end{figure*} We illustrate our mapping methodology on a high resolution, low noise simulation of the CMB temperature and polarization field on a $17^\text{\,o}\times 17^\text{\,o}$ patch of the flat sky. This simulation is used throughout the paper to demonstrate findings and techniques. To get an overview of the performance of our method, Fig.~\ref{fig1} shows the estimated potential (left) from the simulated lensed CMB temperature and polarization field (observational noise levels are set at $1$ $\mu K$-arcmin for the temperature field, $\sqrt{2}$ $\mu K$-arcmin for the polarization fields, with a beam FWHM of $0.25$ arcmin). The input gravitational potential is shown in the right diagram in Fig.~\ref{fig1}. The details of the simulation procedure can be found in Appendix \ref{SimDets}. It is clear from Fig.~\ref{fig1} that the mapping accurately traces the true, unknown gravitational potential. To get an idea of the noise of this reconstruction for different realizations of the CMB $+$ noise we present Fig.~\ref{Sect} which shows the different estimates of the projected matter power spectrum using the estimated projected mass---with the local likelihood approach---for 10 different CMB $+$ noise realizations (dashed lines) while keeping the gravitational potential in Fig.~\ref{fig1} fixed. The blue curve shows the estimated projected mass power spectrum if one had access to the true gravitational potential used in our simulations. Finally we plot the theoretical ensemble average projected mass power spectrum in red to get an idea of the magnitude of the errors in the mass reconstruction. \begin{figure} \includegraphics[height=6.9cm]{Sect}% \caption{Plot of the projected mass power spectrum (red) along with the estimated power spectrum using the true, but unknown, projected mass (blue). The dashed lines correspond to different estimates of the power spectrum using the estimated projected mass---with the local likelihood approach---for different CMB realizations but the same lensing potential realization. See Section \ref{into} and Appendix \ref{SimDets} for the simulation details. \label{Sect} } \end{figure}

We have demonstrated the feasibility of using a local Bayesian estimate to accurately map the gravitational potential and displacement fields under low noise, small beam experimental conditions. The motivation for developing this estimate stems, in part, from speculation that likelihood methods will allow superior mapping of the lensing structure (compared to the quadratic estimator) under low noise levels. The main difference between the global estimates of \cite{HiraSel2003a,HiraSel2003b} and the local estimate presented here is the nature of the likelihood approximation. In \cite{HiraSel2003a,HiraSel2003b} the global likelihood is defined as a functional on the unknown gravitational potential $\phi$ and approximations are made to this functional. Our method, in contrast, uses a nearly exact likelihood---exact up to approximation (\ref{CovFun}) in Appendix \ref{Newt}---but under a local modeling approximation that assumes a quadratic $\phi$. One advantage is the added precision available to model instrumental and foreground characteristics. For example, the local analysis models the beam convolved CMB rather than the deconvolved CMB. Deconvolution induces spatial correlation in the additive instrumental noise which is potentially nonstationary if the beam spatially varies. Since this noise is not invariant under warping it complicates the global likelihood. Another advantage is that the local estimates are relatively easy to implement and parallelize. In addition, the local estimate automatically uses the highest signal to noise combinations of $Q, U$ and $T$ so there is no need to re-derive the optimal quadratic combinations for different experimental conditions. The local analysis is not free from disadvantages however. A global analysis is presumably much better suited for estimating long wavelengths in the gravitational potential and wavelengths that are shorter than the local neighborhood size. Moreover, since our estimates are defined implicitly---as the maximum of the posterior density---it is difficult to derive expected error magnitudes. However, the results presented here show that under some experimental conditions the advantages overcome the disadvantages. Moreover our local estimate uses an approximation that is inherently different from the Taylor approximation used to derive the quadratic estimator. This leaves open the possibility that the local estimate may have different bias and error characteristics which could compliment the quadratic estimator, rather than replace it. % \appendix

1012.1833

We present a new estimation method for mapping the gravitational lensing potential from observed CMB intensity and polarization fields. Our method uses Bayesian techniques to estimate the average curvature of the potential over small local regions. These local curvatures are then used to construct an estimate of a low pass filter of the gravitational potential. By utilizing Bayesian/likelihood methods one can easily overcome problems with missing and/or nonuniform pixels and problems with partial sky observations (E- and B-mode mixing, for example). Moreover, our methods are local in nature, which allow us to easily model spatially varying beams, and are highly parallelizable. We note that our estimates do not rely on the typical Taylor approximation which is used to construct estimates of the gravitational potential by Fourier coupling. We present our methodology with a flat sky simulation under nearly ideal experimental conditions with a noise level of 1μK-arcmin for the temperature field, 2μK-arcmin for the polarization fields, with an instrumental beam full width at half maximum (FWHM) of 0.25 arcmin.

12167745

2011ApJ...728..107A

1012

1012.3188_arXiv.txt

\par The strength of the metagalactic ultraviolet background (UVB) has great impact on theoretical models of structure formation \cite[e.g.][]{Haa96} and a variety of physical processes such as the inhibition of small halo collapse \cite[e.g.][]{Efs92}, the intergalactic temperature and ionization state of the intergalactic medium (IGM) \cite[e.g.][]{Hui97}, and IGM metallicity determinations \cite[e.g.][]{Rau97a}. The likely contributors to the UVB are active galactic nuclei and star formation in galaxies \citep{Sch03,Fau08a} with appear compatible with observed populations \citep{Gal95,Hop04,Hop07,Bou09} under reasonable corrections for dust attenuation, low luminosity extrapolations, redshift evolution, and escape fractions. The strength of the UVB, especially at low redshift \citep{Dav01}, is still highly uncertain despite its importance. Most recent efforts have focused on high redshifts, $z>2$, where the strongest UVB measurements exist. For instance, the detailed history of star formation \citep{Mad99,Fau08b} and the potential to measure individual active galactic nuclei (AGN) host halo masses \citep{Loe95,Fau08c} have been explored. Measurements of the photoionization rate have used three methods: observations of H$\alpha$ such as described in this paper, the line-of-sight proximity effect method \cite[e.g.][]{Car82,Bat88}, and the flux decrement method \cite[e.g.][]{Cen94,Rau97b}. The latter two require backlighting quasars and are therefore difficult or impossible at low redshift. We are motivated to constrain the current model with a different, low redshift measurement. Instead of using Lyman-$\alpha$ forest features, we pursue a measurement of the UVB powered, H$\alpha$ emission that should occur in the outskirts of local disk galaxies. As a secondary motivation, the kinematics of H$\alpha$ at distances beyond HI data are important probes to the total dark halo masses in nearby disk galaxies \citep{Chr08}. \par Galactic disks are optically thick to Lyman limit photons and maintain their observed HI distributions through self-shielding against the UVB. As recognized for decades \citep{Sun69,Fel69,Boc77}, the influence of the UVB may be investigated in the extreme outskirts of disks where the self-shielding begins to fail. These early works sought to measure this effect through disk truncation in HI. However, there appear to be cases with \citep{Cor89,vGo93} and without \citep{Wal97,Car98,Oos07} HI truncations above the critical column density predicted using current UVB estimates, implying that other processes may strip gas and mimic the result. Moreover, reaching the UVB implied truncation thresholds in 21 cm measured HI would require rather long observations with current facilities. A more robust signature of the UVB strength would be the detection of the H$\alpha$ in these outskirt regions. H$\alpha$ has been found at such radii before in actively star forming and warped galaxies by \citet{Bla97} (hereafter BFQ) with Fabry-Perot staring measurements. However, the $\mu(H\alpha)=2.3\times$10$^{-19}$ erg/s/cm$^2$/\sq\arcsec\ detection was interpreted to be due to non-UVB sources as indicated by an abnormally high [NII]$\lambda$6548 to H$\alpha$ ratio. Searches have also yielded limits in quiescent systems \citep{Vog95,Wey01,Mad01} with an upper limit for the UVB photoionization rate, $\Gamma$, of $\Gamma(z=0)<2.4-9.5\times10^{-14}$s$^{-1}$(2$\sigma$) being the deepest. The wide range due on this limit is due to gas cloud geometrical uncertainty. Despite the numerous theoretical implications and the efforts of numerous groups, a UVB powered H$\alpha$ detection still awaits discovery. \par The tactical advantages we bring to this problem are deep surface brightness limits, a large two dimensional field of view through integral field spectroscopy compared to the previous longslit and Fabry-Perot staring data, and target selection of very high inclinations to maximize signal and minimize contamination uncertainty. Our targets are edge-on, low surface brightness Sd galaxies that are rather isolated and minimally warped in order to avoid density distribution uncertainties and exposure to internally generated ionization from smaller radii. Indeeed, our most constraining target, UGC~7321, has a gas surface density below that required for significant star formation \citep{Ken89} at all radii, as well as being unusually isolated with no known companions and minimal ($<3\arcdeg$) warping \citep{Uso03}. \par In this paper we begin with a description of the simple ionization state and density model of disk galaxies that will be used to link a measured H$\alpha$ surface brightness with a particular UVB photoiozation rate in \S \ref{sec_mod_HI}. In \S \ref{sec_mod_fit}, we give disk parameter constraints based on fits to existing 21 cm data. In \S \ref{sec_alt_lim}, we argue that UGC~7321 in particular is likely to extend its HI profile beyond the current 21 cm limits without truncation. In addition, the HI observations of UGC~7321 are amongst the most sensitive such measurements published to-date. The 21 cm data allow a very precise model to be made for the gas distribution in the galaxy outskirts at the locations where we search for H$\alpha$ emission. Next, in \S \ref{sec_data}, we present deep integral field spectroscopy observations at radii corresponding to the outermost detections of 21 cm emission and beyond. We describe the choices made to stack spectra on various spatial scales. The stacked spectra are searched for H$\alpha$ detections and upper limits are derived. Particular focus is given to systematic errors. Finally, in \S \ref{sec_dis}, we discuss the context, the likely cause of the unexpectedly low limit, and further observations that can confirm our conclusions. The Appendix \ref{sec_full_mu} provides the analytic details necessary to construct the full and general H$\alpha$ surface brightness distribution model. We will quote most of the surface brightness limits in units of erg s$^{-1}$ cm$^{-2}$ arcsec$^{-2}$, but for easy comparison to alternative units we note the conversion at the wavelength of H$\alpha$ of 1 millirayleigh (mR)$=5.66\times 10^{-21}$ erg s$^{-1}$ cm$^{-2}$ arcsec$^{-2}=2.8\times 10^{-3}$ cm$^{-6}$ pc in emission measure assuming the case B coefficient we adopt.

\label{sec_dis} \par The flux decrement method is currently the most widely used method to estimate the UVB strength at high redshift. Under the fluctuating Gunn-Peterson approximation \citep{Cro98}, the Lyman-$\alpha$ forest optical depth distribution should have a normalization that depends only on well constrainted cosmological parameters and the UVB strength. The IGM temperature and density distributions may have some systematic uncertainties that propagate into knowledge of the UVB, but they are not likely the leading uncertainties. The more likely dominant uncertainties in flux decrement modeling are the source emissivities. At $z\lessapprox1$, the Lyman limit mean free path becomes larger than the horizon, so the UVB strength at z=0 is influenced by source evolution across this redshift range. AGN and stellar population luminosity functions, both observed and modeled, generally agree to better than an order of magnitude over these redshifts. The least constrained input to flux decrement modeling is the escape fraction for ionizing photons in galaxies, particularly at low redshift and low luminosity. We believe our measurement is best interpreted as an indicator of a low escape fraction. \par Our most constraining (5$\sigma$) spectral limits are $\Gamma<1.7\times 10^{-14}$ s$^{-1}$ in UGC~7321 and $\Gamma<13.5\times 10^{-14}$ s$^{-1}$ in UGC~1281 again assuming $\beta=1.8$. Several benchmarks, both empirical and theoretical, exist with which to compare these limits. Figure \ref{fig_zGam} shows the UVB strength against redshift determined by many groups. The lowest redshift proximity effect limit comes from \citet{Kul93} with analysis of 13 quasars from \citet{Bah93} between $0.16 \le z \le 1.00$ at $\Gamma(\bar{z}=0.5)=2.0^{+10}_{-1.3}\times 10^{-14}$ s$^{-1}$. However, the proximity effect method has been shown to have a high bias that depends on halo mass \citep{Fau08c} and should be interpreted with care. The theoretical model of \citet{Fau09} gives a drop in the UVB strength by a factor of 3.4 between z=0.5 and z=0.0 leaving this measurement consistent with our current limit. This agreement is interesting and somewhat unexpected given the bias of proximity effect measurements. The only existing low-z flux decrement limit is $\Gamma(\bar{z}=0.17)=5.0^{+20.}_{-4.0}\times 10^{-14}$ s$^{-1}$ \citep{Dav01}. The theoretical model itself, normalized by the flux decrement method, predicts $\Gamma(z=0)=3.8\times 10^{-14}$ s$^{-1}$ which is much higher than our new limit. There exists a second set of unpublished theoretical predictions from F. Haardt and P. Madau discussed in \citet{Fau09} giving $\Gamma(z=0)=1\times 10^{-13}$ s$^{-1}$. The latter model used a constant 10\% escape fraction of ionizing photons and an unspecified star formation history while the former used a completely theoretical and simulation-based star formation history \citep{Her03} and a scaling of the stellar UV emissivity based on high redshift flux decrement measurements that contains the escape fraction. A comparison to Lyman-break galaxy (LBG) luminosity functions led that group to require only $f_{esc,abs}\approx0.5\%$ \citep{Fau08a}. The direct measurement of galactic escape fractions is difficult due to the low values involved. While UV bright samples can range up to $\approx 3\%$ in absolute Lyman limit escape fraction \citep{Sha06}, a presumably lower-mass sample yielded $(2\pm2)\%$ \citep{Che07}. Theoretical work shows a strong decrease in $f_{esc}$ with star formation rate and halo mass \citep{Gne08} below $M_{tot}\approx10^{11}M_{\odot}$, and lower redshift observations of populations similar to LBGs show a potential redshift evolution \citep{Sia10} with $f_{esc,abs}<0.8\%$. There is no reason yet to suppose a lower bound to the escape fraction. If we interpret our limit as a scaling of the escape fraction from the models in \citet{Fau08a} at low redshift, we find $f_{esc,abs}<0.2\%$. \par It is unlikely that systematics from the model assumptions in our analysis can cause the disagreement. Contaminating ionization from the galaxies' forming stars would bias our measurement high, only making the disagreement more severe. We further note that the degree of contamination can be measured by anomalous [NII]$\lambda$6548 to H$\alpha$ ratios (BFQ) and should not, in principle, limit this type of measurement. There has been a large body of work on low strength star formation beyond the optical radii in local galaxy disks, usually labelled extended UV disks (XUV), fostered by far UV (FUV,1350-1750\AA) and near UV (1750-2750\AA) Galaxy Evolution Explorer (GALEX) data \cite[e.g.][]{Thi07}. Narrowband H$\alpha$ imaging and spectroscopy have revealed that $\sim10$\% of gas rich disks \citep{Wer10a,Wer10b,Her10} host outlying H$\alpha$ emitting complexes as either compact HII regions or dwarf satellite companions. The common H$\alpha$ fluxes observed so far are of the order of a few times $10^{-16}$ erg/s/cm$^2$. Any such systems would have been found in our data as strong detections limited in size to a few fibers. The expectation of large-scale, diffuse UVB H$\alpha$ emission should discriminate reliably against compact XUV H$\alpha$ emission. We have also visually inspected the target galaxies' GALEX data which have not yet been analyzed in any XUV focused work. UGC~1281 has only been covered in the rather shallow all-sky survey mode. UGC~7321 has been covered for 2.8ks in the NUV and 1.7 ks in the FUV under guest investigator cycle 4 proposal ID 095 (PI: J.~Lee) as part of the 11HUGS project \citep{Lee09}. In neither system is there evidence for an extended UV disk beyond the DSS2-red\footnote{The Digitized Sky Survey was produced at the Space Telescope Science Institute under U.S. Government grant NAG W-2166. The images of these surveys are based on photographic data obtained using the Oschin Schmidt Telescope on Palomar Mountain and the UK Schmidt Telescope. The plates were processed into the present compressed digital form with the permission of these institutions.} limiting contours. Finally, these contamination issues are speculative until a putative UVB H$\alpha$ detection is made. The only possible systematics that could have made a low bias to our limit are unaccounted for dust or gas distribution parameters, such as inclination, far beyond the range we have constrained. \par We have made our first analysis under the assumption that the gas distribution extends beyond the HI data limits with the same exponential form as at smaller radii. This assumption, motivated by the thin and regular HI distributions and lack of nearby companions, has the strongest impact on our interpretation. An alternative estimate without this assumption, taking only fibers that overlap with the observed HI signal, yields a very comparable limit of $\Gamma<2.3\times 10^{-14}$ s$^{-1}$ at 5$\sigma$ significance in UGC~7321. This agreement essentially comes about because our original model predicts only a minor H$\alpha$ contribution at the discarded positions under the modeled UVB strength. Nevertheless, there is no reason to assume the presence of an HI edge since the radio observations detect the gas up to the column densities where the sensitivity runs out. This result raises the question whether a redshift-dependent escape fraction is manifesting in galaxies. Alternatively, our new limits may be saying that the UVB strength, as estimated through flux decrement measurements, has been overestimated at all redshifts. The latter choice would upset the apparent agreement between current models and reionization constraints. Either case will require some modification to the UVB strength model and its implementation in structure formation simulations. We intend to pursue our measurements of these and other superthin galaxies to greater depth in order to arrive at a detection of $\Gamma(z=0)$.

1012.3188

We present new integral-field spectroscopy in the outskirts of two nearby, edge-on, late-type galaxies to search for the Hα emission that is expected from the exposure of their hydrogen gas to the metagalactic ultraviolet background (UVB). Despite the sensitivity of the VIRUS-P spectrograph on the McDonald 2.7 m telescope to low surface brightness emission and the large field of view, we do not detect Hα to 5σ upper limits of 6.4 × 10<SUP>-19</SUP> erg s<SUP>-1</SUP> cm<SUP>-2</SUP> arcsec<SUP>-2</SUP> in UGC 7321 and of 25 × 10<SUP>-19</SUP> erg s<SUP>-1</SUP> cm<SUP>-2</SUP> arcsec<SUP>-2</SUP> in UGC 1281 in each of the hundreds of independent spatial elements (fibers). We fit gas distribution models from overlapping 21 cm data of H I, extrapolate one scale length beyond the H I data, and estimate predicted Hα surface brightness maps. We analyze three types of limits from the data with stacks formed from increasingly large spatial regions and compare to the model predictions: (1) single fibers, (2) convolution of the fiber grid with a Gaussian, circular kernel (10'' full width at half-maximum), and (3) the co-added spectra from a few hundred fibers over the brightest model regions. None of these methods produce a significant detection (&gt;5σ) with the most stringent constraints on the H I photoionization rate of Γ(z = 0) &lt; 1.7 × 10<SUP>-14</SUP> s<SUP>-1</SUP> in UGC 7321 and Γ(z = 0) &lt; 14 × 10<SUP>-14</SUP> s<SUP>-1</SUP> in UGC 1281. The UGC 7321 limit is below previous measurement limits and also below current theoretical models. Restricting the analysis to the fibers bound by the H I data leads to a comparable limit; the limit is Γ(z = 0) &lt; 2.3 × 10<SUP>-14</SUP> s<SUP>-1</SUP> in UGC 7321. We discuss how a low Lyman limit escape fraction in z ~ 0 redshift star-forming galaxies might explain this lower than predicted UVB strength and the prospects of deeper data to make a direct detection. <P />This paper includes data taken at the McDonald Observatory of the University of Texas at Austin.

12213920

2011MNRAS.413..878S

1012

1012.0311_arXiv.txt

\label{sec:intro} In the current paradigm of structure formation, large objects, such as galaxies or clusters, are believed to form hierarchically, through a 'bottom-up' \citep{white1978} process of merging. About a decade ago, N-body simulations attained sufficient dynamic range to reveal that, in Cold Dark Matter (CDM) models, all haloes should contain a large number of embedded subhaloes that survive the collapse and virialization of the parent structure \citep{klypin1999,moore1999}. The properties of subhaloes on different scales has been the subject of many recent studies that have pushed the resolution of dissipationless simulations \citep[e.g.][]{springel2001,deLucia2004,kravtsov2004,gao2004, reed2005,diemand2007,zentner2005,springel2008}. The kinetic properties of subhaloes are now well understood - they make up a fraction of between 5 and 10\% of the mass of virialized haloes, on scales relevant to observational cosmology. Most of these previous studies used dissipationless cosmological simulations; although non-baryonic dark matter exceeds baryonic matter by a factor of $\Omega_{dm}/\Omega_b \simeq 6$ on average \citep[e.g.][]{komatsu2009}% , the gravitational field in the central region of galaxies is dominated by stars and gas. The cooling baryons increase the density in the central halo region mainly because of the extra mass associated with the inflow, but also because of the adiabatic contraction of the total mass distribution \citep[e.g.][]{gnedin2004}. Since this process is active for {\it both} the host halo and its subhaloes, it might be expected that subhaloes formed within hydrodynamical simulations (including gas and stars) will experience a different tidal force field and will themselves be more robust to tidal effects. Recently, a number of authors have examined the impact of baryonic physics (gas cooling, star formation and feedback) on both the central object and the satellite population in galaxy and cluster-sized haloes \citep[e.g.][]{bailin2005,nagai2005,maccio2006,weinberg2008,romano2009,romano2010,libeskind2010, sommer2010,duffy2010}. \cite{maccio2006} simulated a Galactic mass halo twice - once considering pure DM and once including baryons modeled with smoothed particles hydrodynamics, stopping the gas cooling at $z=1.5$. They found that the hydro run produced an overabundance of subhaloes in the inner regions of the halo as well as an increase by a factor of 2 in the absolute number of subhaloes with respect to the DM run. \\ \indent The issue of the distribution and properties of galactic satellites in hydro and DM simulations has been more recently revisited by \cite{libeskind2010} and \cite{romano2010}. \cite{libeskind2010} found results very similar to the work of \cite{maccio2006}, with subhaloes in the hydro simulation being more radially concentrated than their dark matter counterparts. They ascribe this effect to the higher central density of subhaloes in hydro simulation (due to the collapse of baryons into stars in the central region) that makes them more resilient to tidal forces. The increased mass in hydrodynamic subhaloes with respect to dark matter ones causes dynamical friction to be more effective, dragging the subhalo towards the centre of the host. The overall properties of the satellite population in the study of \cite{romano2010} are also consistent with \cite{maccio2006} and \cite{libeskind2010}. But, perhaps counter-intuitively, they find that satellites in the hydro run are depleted at a faster rate than the pure DM one within the central 30 kpc of the prime halo. According to their analysis, \cite{romano2010} suggest that although the baryons provide a substantial glue to the subhaloes, the main halo exhibits the same trend. This would assure a more efficient tidal disruption of the hydro subhalo population in the inner region of the halo ($\approx 0.1$ of the virial radius). Studies concerning the different results of pure dark matter and hydrodynamical simulations are especially compelling within the context of satellite effects on disk stability. Several theoretical and numerical studies have been devoted to quantifying the resilience of galactic disks to infalling satellites \citep[e.g.][]{toth1992,quinn1993,velazquez1999,font2001,read2008,villalobos2008,moster2010}. Recently \cite{kazantzidis2009} performed a very detailed study of the dynamical response of thin galactic disks to bombardment by cold dark matter substructure. They used pure Nbody simulations of the formation of a Milky Way-like dark matter halo to derive the properties of substructures and subsequently as initial conditions in subsequent high resolution satellites-disk merger simulations. Clearly, understanding {\it if} and {\it how} different possible orbital parameters and mass loss rates expected for satellites in hydro runs modify results previously obtained using pure Nbody simulations is of importance. In this work we revisit the issue of the effects of baryonic physics on the satellite population in galactic-size dark matter haloes. We improve the original study by \cite{maccio2006} in several aspects, namely with better parametrization of the baryonic physics, a full hydrodynamical approach down to redshift zero and a more extensive analysis of the satellite properties, including the time evolution of mass loss, radial position, and orbital parameters (peri and apo-centre distances). We start hydro and pure dark matter simulations from the same initial conditions and we focus our analysis on ''twins'' satellites, i.e. (sub)structures that are formed from the same Lagrangian region for the initial conditions, which should therefore share the same formation history in both simulation types. The remainder of the paper is organized as follows: in Section~\ref{sec:sims} we describe our simulations, and provide a brief summary of the numerical codes we use, including the technique employed to match satellites in the different runs. In Section~\ref{sec:res} we present our main results, focusing on several satellite properties like radial position, orbital parameters, mass loss. Finally in Section~\ref{sec:conc} we summarize and discuss our results.

\label{sec:conc} With this work we aim to provide a detailed description of the effects of baryonic physics on the properties of galactic satellites and their evolution with redshift. For this purpose we analyse three different cosmological simulations of galaxy formation, each run twice: once as pure dark matter and another with the addition of gas physics, including gas cooling, star formation and feedback. For each run we create a comprehensive catalogue of the subhalo population at $z=0$ which we trace back in time in order to study mass loss, dynamical friction and evolution of satellite orbital parameters. Within this population we focus on a sub-sample of corresponding DM and hydro satellite pairs (``twins'') and their individual evolutionary histories. Satellites are found to be more radially concentrated in the hydro simulation than in the pure DM one, confirming earlier results by \cite{maccio2006}. This bias persists also for the twin population, even if slightly less pronounced. When we restrict our analysis to the twin sub-sample we find that hydro satellites tend to enter the virial radius of the parent halo later than the corresponding DM subhaloes, with an average delay of 0.7 Gyrs. This difference cannot be ascribed to a difference in the evolution of $R_{vir}$ in the two simulations, and we speculate that it is instead related to the pressure support of the hot gas that acts against collapse, which, in low density regions, is not counterbalanced by cooling. Nevertheless further studies are need to confirm our hypothesis. Given the delay in the accretion time, the orbits of twin satellites are often weakly correlated. As a consequence, we find that the radial position at a given redshift is not sufficient to describe the satellite orbit. For this reason we define an average radius $R_{avg}$ for each satellite equal to the median between the apocenter and pericenter distances of the satellite orbit, where these last two quantities are obtained by integrating each satellite orbit in a fixed potential resembling that of the halo. We find that the ratio of the average positions in the hydro and DM cases measured at redshift zero correlates with the accretion time of the satellite and its mass at that time. Moreover, we find that both the absolute mass loss experienced by the satellite and the difference in mass loss in the hydro and DM simulations also correlate with the subhalo average distance at $z=0$. We arrive at a final picture in which more massive satellites at the time of accretion are {\it closer} to the center and {\it more} heavily stripped in the hydro simulation than in the pure DM. The situation is reversed for satellites with low $M_{infall}$ that are on average {\it further} way from the center and {\it less} stripped in the hydro simulation. This bimodality can be understood in the following way: in the outer region of the halo, dynamical friction is less important, and the distance from the center of a given satellite mainly depends on its accretion time. Since hydro satellites tend to be accreted later, they experience lower dynamical friction and mass stripping than their pure DM counterparts. In the inner regions, on the other hand, dynamical friction is stronger in the hydro simulation (due to the central stellar body) and drags satellites toward the center in a more effective way compared to the pure DM case. Although the baryons provide a substantial glue to the subhaloes, the main halo exhibits the same trend. This assures a more efficient tidal stripping of the hydro subhalo population, resulting in a larger mass loss. During the making of this work, two other groups, i.e. \cite{libeskind2010} and \cite{romano2010}, independently pursued the same subject. Both studies compared hydro and pure DM simulations using an approach similar to ours, including a focus on corresponding pairs of satellites. The publication of Libeskind et al. mainly considers the radial distribution of the satellites and a comparison of the retained masses. Like us, they find a difference in the orbits of twins. But in contrast to our approach, \cite{libeskind2010} use the statistics of the radial orbit position, rather then, e.g. determine the orbital parameters and the resulting average distance as we do. They find a larger mass loss for the satellite in the pure DM simulation and interpret this as a sign of the expected higher stability of the hydrodynamical satellite, which our results are only able to confirm in the external region of the halo. \cite{romano2010} on the other hand, find an increased mass loss for the hydrodynamical satellite in the central region of the halo, in accordance with our results. They are also able to detect the final disruption of a satellite, and have shown that the life expectancy of the satellites in the hydrodynamical simulation is indeed shorter than in the pure DM simulation, as our findings suggest. We extend this result further in the last part of our study by investigating the possible impact of different orbital parameters and mass loss in hydrodynamical simulations, i.e. whether this translates into an increase, or a reduction, in the danger that these satellites pose for the stability of a possible stellar disk at the center of the parent halo \citep[e.g.][]{kazantzidis2009,moster2010}. While the peri-centric distance does not substantially differ in hydro and DM runs, hydrodynamical satellites in the central region face a stronger mass loss. As the danger for galactic disks is linked to the mass of the satellite, this faster mass depletion for hydro satellites leads to less dangerous perturbers for galactic disks, possibly easing the problem of the existence of thin disks in a Cold Dark Matter Universe. We emphasize, though, that, as correctly pointed out by \cite{romano2010}, the effects of baryonic cooling in the center of dark matter (sub)haloes can be altered (if not reversed) for a more efficient feedback from stellar evolution and possibly central super-massive black holes, which will expel baryons from the center and decrease the central concentration of the prime halo \citep[e.g.][]{elzant2004,mashchenko2006,governato2010}. At this time, the direct comparison of our study of galactic satellites with observations is not possible. The Sloan Digital Sky Survey, which greatly contributed to the discovery and study of several Milky Way satellites \citep[e.g.][]{koposov2008,dejong2010}, covers only a single patch of the sky. Future surveys like Pan-Starrs, which will provide a more comprehensive map of the (northern) sky, promise a better understanding of the properties and spatial distribution of satellite galaxies orbiting around the Milky Way.

1012.0311

In this work, we examine the different properties of galactic satellites in hydrodynamical and pure dark matter simulations. We use three pairs of simulations (collisional and collisionless) starting from identical initial conditions. We concentrate our analysis on pairs of satellites in the hydro and N-body runs that form from the same Lagrangian region. We look at the radial positions, mass-loss as a function of time and orbital parameters of these ‘twin’ satellites. We confirm an overall higher radial density of satellites in the hydrodynamical runs, but find that trends in the mass-loss and radial position of these satellites in the inner and outer region of the parent halo differ from the pure dark matter case. In the outskirts of the halo (≈70 per cent of the virial radius), satellites experience a stronger mass-loss and higher dynamical friction in pure dark matter runs. The situation is reversed in the central region of the halo, where hydrodynamical satellites have smaller apocentre distances and suffer higher mass stripping. We partially ascribe this bimodal behaviour to the delayed infall time for hydro satellites, which on average cross the virial radius of the parent halo 0.7 Gyr after their dark matter twins. Finally, we briefly discuss the implications of the different set of satellite orbital parameters and mass-loss rates in hydrodynamical simulations within the context of thin-disc heating and destruction.

12101858

2010AIPC.1273...75W

1012

1012.5228_arXiv.txt

Searching for iron in PG1159 stars poses a difficult problem. The atmospheric parameter range (\Teff\,=\,75\,000 -- 200\,000\,K, \logg\,=\,5.5 -- 8.0 \cite{we06}) requires to look for lines from high ionization stages, at least \ion{Fe}{vii}. Numerous \ion{Fe}{vii} lines are seen in UV spectra of hot hydrogen-rich central stars of planetary nebulae (e.g., \cite{scho85}), but the search for these lines in PG1159 stars was unsuccessful. Model calculations show that \ion{Fe}{vii} lines with detectable strength can only be expected in the cooler PG1159 stars. They vanish with increasing \Teff\ because the ionization balance of iron shifts to higher stages. The same tendency is seen with decreasing gravity (i.e., increasing luminosity). Based on the lack of \ion{Fe}{vii} lines, Fe-deficiency of $\geq 1$~dex was concluded for four stars: \pgvier\ (\Teff\,=\,110\,000\,K, \logg\,=\,7.5 \cite{rei08}), the hybrid-PG1159 stars (i.e., H-Balmer lines are visible) NGC\,7094 and Abell\,43 (both have \Teff\,=\,100\,000\,K, \logg\,=\,5.5 \cite{zi09}), the PG1159 -- [WC] transition type star Abell\,78 (\Teff\,=\,110\,000\,K, \logg\,=\,5.5 \cite{we03}). Currently, these stars represent the strongest cases for Fe-deficiency. Also, Fe-deficiency was claimed in some [WC]-type central stars, the putative progenitors of PG1159 stars. This finding is not explained by evolutionary models. We speculated that Fe nuclei could have been transformed to heavier elements by excessive s-processing. \begin{figure}[t] \includegraphics[width=0.8\textwidth]{werner_poster_fig3.ps} \caption{\label{fig:ion} Vertical run of iron ionization fractions in a model atmosphere with \Teff=\,170\,000\,K and \logg=5.8.} \end{figure} \begin{figure}[t] \includegraphics[width=0.9\textwidth]{werner_poster_fig4.ps} \caption{\label{fig:var} Dependence of \ion{Fe}{x} $\lambda$979.3 on atmospheric parameters. \emph{Left and middle panels:} \Teff\ variation at two fixed surface gravities. \emph{Right panel:} variation of iron abundance at fixed \Teff\ and \logg. } \end{figure}

1012.5228

The lack of Fe VII lines in PG1159 stars had led to the conclusion that in some objects iron must be strongly depleted. We have now detected Fe X lines in FUSE spectra of the very hottest PG1159 stars (Teff = 150 000-200 000 KRX J2117.1+3412, K1-16, NGC 246, Longmore4). Surprisingly, we derive a solar iron abundance. It is conspicuous that they are among the most massive PG1159 stars (0.71-0.82 M<SUB>solar</SUB>), in contrast to those objects for which strongest Fe-deficiency was claimed (0.53-0.56 M<SUB>solar</SUB>). Based on new Fe VIII line identifications in SOHO/SUMER UV spectra of the Sun, we were able to detect these lines in FUSE spectra of several ``cooler'' (Teff&lt;~150 000) objects, among them is the prototype PG1159-035. An abundance analysis is in progress.

1018057

2010arXiv1012.2064T

1012

1012.2064_arXiv.txt

% This energy region $\sqrt{s} \approx $ 2 TeV is of particular relevance because of high energy physics and astrophysics issues. The {\it pp} total cross section, \sigmappt, and \sigmapairin ~are related and can be inferred from each other by means of the Glauber theory. The whole procedure is model dependent, the results \cite{gaisser,d&p,kopel,block06,bell} differing of about 20\% for $\sqrt{s}$ values in the TeV energy range. Available accelerators' measurements at the highest energies, are themselves affected by systematic uncertainties of difficult evaluation. The {\it pp} ($\bar pp$) cross section measurements at energies of $\sqrt{s} = $ 1.8 TeV \cite{cdf,d0,d01} differ of about 10\%, exceeding the statistical uncertainties of the measurements. It is therefore of primary interest to have experimental measurements of \sigmapairin ~and \sigmappt ~at the same center of mass energy, i.e. around $\sqrt{s} \approx $ 2 TeV, at which collider data are available. The interpretation of Extensive Air Shower measurements relies on simulations that use hadronic interaction models based on theoretically guided extrapolations of the accelerator data obtained at lower energies and restricted to limited kinematical regions. A \sigmapairin \space direct measurement and the comparison of observables as obtained from measurements and model based simulations, in the same conditions, is therefore highly recommended in order to confirm the validity of the whole analysis procedure. Measurements of the {\it p}-air inelastic cross section performed in EAS have been reported. Since air shower detectors cannot observe the depth of the first interaction of the primary particle, indirect methods have to be used. Two main techniques are used: the constant \ne-\nmu ~cuts \cite{akeno2,eastop1,argo} by means of particle arrays, and the study of the shower longitudinal profiles using fluorescence detectors \cite{fly,hires} at higher energies. Following the particle array technique, the primary energy is first selected by means of the muon number (\nmu). Proton induced showers at the same development stage are then selected by means of the shower size (\ne). The cross section of primary particles is obtained by studying the absorption in the atmosphere ($ \lambda_{obs}$) of such showers, through their angular distribution at the observation level. The rate of showers decreases exponentially with zenith angle $\theta$ (i.e. atmospheric depth of first interaction) as: \begin{equation} f(\theta)= G(\theta) f(0) \exp [- x_0 (\sec \theta - 1 )/ \lambda _{\rm obs}] \label{flu} \end{equation} where $x_{\rm 0}$ is the vertical atmospheric depth of the detector, and $ G(\theta)$ the angular acceptance. With air fluorescence detectors, the absorption length $ \lambda_{obs}$ is obtained fitting the atmospheric depth of the maximum shower development stage ($X_{max}$) distribution tail. The observed absorption length in both cases is affected by the fluctuations in the longitudinal development of the cascades and in the detector response. Such fluctuations can be studied through simulations, providing the conversion factor {\it k} between the observed absorption length and the interaction length of primary protons ({\it k}=$\lambda_{obs}/ \lambda_{int} $). This factor is then used to convert the observed experimental absorption length \Lambdaobsexp ~into the interaction one \Lambdaintexp. In this paper we will report on the measurement of the {\it p}-air inelastic cross section at primary energy $E_0 \approx 2 \cdot 10^{15}$ eV (i.e. $E_0~=~(1.5~\div~2.5)\cdot 10^{15}$~eV$, \sqrt{s} \approx $ 2 TeV) with the EAS-TOP experiment. Primary energies are below the steepening (\emph{knee}) of the primary spectrum (i.e. E$_0 < 3\cdot 10^{15}$ eV) above which the proton flux is strongly reduced. The constant \ne-\nmu ~analysis has been optimized \cite{eastop1} selecting showers at the maximum development stage where fluctuations are lower and heavier primaries rejection is improved by the request of the highest \ne values at a given primary energy. The constant \ne-\nmu ~method with the selection of cascades' maximum developement stage is equivalent to the study of $X_{max}$ distribution tail. As shown in Fig.~\ref{fig:Xmax}, the accessible part of the $X_{max}$ distribution tail depends from the vertical atmospheric depth ($x_{\rm 0}$) of the detector. This is a limitation on the possibility to exploit this method at different atmospheric depths and on the maximum zenith angle $\theta$ (i.e. atmospheric depth) that can be considered in the analysis without running out of statistics. Systematic uncertainties of the measurement and the effect of possible contributions of heavier primaries are discussed and evaluated. \begin{figure}[h] \vspace{-5mm} \includegraphics[width=80mm] {longi_qgs2.eps} \caption{\label{fig:Xmax} Depth of shower maximum ($X_{max}$) distribution for proton showers simulated with QGSJET~II in the selected energy range ($E_0~=~(1.5~\div~2.5)\cdot 10^{15}$ eV) . The shaded area represents the interval of atmospheric depths (i.e.~$1.0\leq \sec \theta \leq 1.2$) % considered in the analysis.} \vspace{-4mm} \end{figure}

Combining the results obtained with the two considered interaction models and including the systematic uncertainties % the {\it p}-air inelastic cross section is: \vspace{2mm} $\sigma^{\rm inel}_{p- \rm air} = 338 \pm 21_{stat} \pm 19_{syst} -29_{syst(He)} {\rm mb}$ \vspace{2mm} As shown in Fig.~\ref{fig:inpair}, this value is about 15\% smaller than the values in use within QGSJET II and SIBYLL and in better agreement with Refs. \cite{block06,hor,hdpm}. Predicted $\sigma^{\rm inel}_{p- \rm air}$ values, obtained from different \sigmappt~Tevatron measurements at $\sqrt{s} = $ 1.8 TeV by using different calculations based on the Glauber theory, are reported in Fig.~\ref{fig:ppair}. The present measurement is consistent with smaller values of the $\bar pp$ total cross section (\sigmappt=72.8$\pm$3.1 mb \cite{d0}, and \sigmappt =71$\pm$2 mb \cite{d01}), and the {\it pp} to {\it p}-air calculations predicting for a given value of \sigmappt, a smaller value of $\sigma_{\rm in}^{p- \rm air}$ % \cite{block06,bell}. \begin{figure}[h] \vspace{-1mm} \includegraphics[width=80mm]{pair_pp.eps} \vspace{-4mm} \caption{\label{fig:ppair} {\it p}-air inelastic vs. {\it $ \bar pp $} total cross section data. The present result ($\pm$ 1 s.d., solid lines) is shown together with the results of different calculations% derived from {\it $ \bar pp $} measurements reported at $ \sqrt{s} = $ 1.8 TeV. } \vspace{-4mm} \label{fig:ppair} \end{figure} Independently from the cross section analysis, the measured values of the absorption length $\lambda_{obs}$ are about 15\% higher than the simulated ones for both the considered interaction models. This can be probably ascribed to the fact that the measured cascades penetrate deeper into the atmosphere than predicted by the interaction models.

1012.2064

The proton-air inelastic cross section value \sigmapairin=338$\pm$21({\it stat})$\pm$19({\it syst})-28({\it syst}) mb at $\sqrt{s} \approx $ 2 TeV has been measured by the EAS-TOP Extensive Air Shower experiment. The absorption length of cosmic ray proton primaries cascades reaching the maximum development at the observation level is obtained from the flux attenuation for different zenith angles (i.e. atmospheric depths). The analysis, including the effects of the heavier primaries contribution and systematic uncertainties, is described. The experimental result is compared with different high energy interaction models and the relationships with the {\it pp} ($\bar pp$) total cross section measurements are discussed.

16077757

2011NatCo...2..175C

1012

1012.0850_arXiv.txt

1012.0850

The origin of ultra-high-energy cosmic rays (UHECRs) remains an enigma. They offer a window to new physics, including tests of physical laws at energies unattainable by terrestrial accelerators. They must be accelerated locally, otherwise, background radiations would severely suppress the flux of protons and nuclei, at energies above the Greisen-Zatsepin-Kuzmin (GZK) limit. Nearby, gamma ray bursts (GRBs), hypernovae, active galactic nuclei and their flares have all been suggested and debated as possible sources. A local sub-population of type Ibc supernovae (SNe) with mildly relativistic outflows have been detected as sub-energetic GRBs, X-ray flashes and recently as radio afterglows without detected GRB counterparts. Here, we measure the size-magnetic field evolution, baryon loading and energetics, using the observed radio spectra of SN 2009bb. We place such engine-driven SNe above the Hillas line and establish that they can readily explain the post-GZK UHECRs.

12201867

2011ASPC..448..269C

1012

1012.0061_arXiv.txt

In his seminal 1972 paper, Andrew Skumanich showed that stellar rotation decreases over time such that $v_{rot} \propto t^{-0.5}$ --- as does CaII emission, a measure of chromospheric activity and proxy for magnetic field strength. This relationship between age, rotation, and activity has been a cornerstone of stellar evolution work over the past 40 years, and has generated almost as many questions as applications. For example, angular momentum loss due to stellar winds is generally thought to be responsible for the \citet{skumanich72} law, but the exact dependence of $v_{rot}$ on age is not entirely clear, and relies on the assumed stellar magnetic field geometry and degree of core-envelope coupling \citep{kawaler1988, krishnamurthi1997}. Furthermore, later-type, fully convective stars appear to have longer active lifetimes than their early-type brethren \citep[e.g.,][]{andy08}, indicating that the lowest mass stars are capable of generating significant magnetic fields even in the absence of a standard solar-type dynamo \citep{Browning2008}. The lack of a comprehensive theoretical understanding of the age-rotation-activity relation has not prevented the development and use of gyrochronology, however, which attempts to determine the precise ages of field stars based on a presumed age-rotation relation \citep[e.g.,][]{barnes2007, mamajek2008, collier2009, barnes2010}, nor of empirical age-activity relations, which do not always find activity decaying with time quite as simply as predicted by the Skumanich law \citep[e.g.,][]{feigelson04,pace2004,giampapa2006}. Fully mapping out the dependence of stellar rotation and activity on age requires the study of stars ranging in both mass and age. Statistical constraints on the age-rotation-activity relation can be derived via analysis of Galactic field stars \citep[e.g.,][]{feigelson04,Covey2008,Irwin2010}, but the homogeneous, coeval populations in open clusters provide an ideal environment for studying time-dependent stellar properties. There are relatively few nearby open clusters, however, and fewer still have the high quality optical data needed to characterize their rotations --- in part because of the sheer difficulty involved in systematically monitoring a large number of stars over several months or more. As a result, our current view of the age-rotation-activity relation depends on observations of handfuls of stars in the field and in a small number of well-studied clusters \citep[with the Hyades being a particularly key cluster, e.g.,][]{Radick1987,jones1996,stauffer1997,terndrup2000}. The advent of time-domain surveys, with their emphasis on wide-field, automated, high-cadence observing, makes it possible to monitor stellar rotation in clusters on an entirely new scale \citep[e.g.,][]{Irwin2007,Meibom2009,Hartman2010}. Many ongoing time-domain surveys are primarily designed to identify transiting exoplanets, however, and thus aim to cover the widest area possible to a relatively modest depth. Deep targeted surveys of open clusters are therefore still required to measure rotation periods, particularly for the lowest mass stars in older, more distant open clusters. The Palomar Transient Factory provides deep, multi-epoch photometry over a wide field-of-view, and our Columbia/Cornell/Caltech PTF survey is leveraging this capability to map $v_{rot}(t)$ in open clusters of different ages. Our first CCCP target, Praesepe (the Beehive Cluster, M44, $08^h 40^m 24^s +19^{\circ} 41\arcmin$), is a nearby ($\sim$170 pc), intermediate-age ($\sim$600 Myr), and rich \citep[its membership has recently been expanded to $\sim$1200 stars;][]{adam2007} open cluster that shares many characteristics with the Hyades. Here we report the stellar rotation periods for Praesepe members derived from our first season of PTF observations. Our campaign produced $\sim$200 distinct observations of four overlapping fields covering a $3.75\times 3.30$ deg area designed to include a large number of Praesepe members identified by \citet{adam2007}. In Section~\ref{observations} we describe our PTF observations. We discuss our period-finding algorithm in Section~\ref{periods} and our results in Section \ref{results}. Finally, in Section~\ref{future} we outline the current status and future of the CCCP survey.

1012.0061

Over the past 40 years, observational surveys have established the existence of a tight relationship between a star's age, rotation period, and magnetic activity. The age-rotation-activity relation is essential for understanding the interplay between, and evolution of, a star's angular momentum content and magnetic dynamo. It also provides a valuable age estimator for isolated field stars. While the age-rotation-activity relation has been studied extensively in clusters younger than 500 Myr, its subsequent evolution is less constrained. Empirically measured rotation periods are scarce at intermediate ages (i.e., Hyades or older), complicating attempts to test reports of a break in the age-activity relation near 1 Gyr (e.g., Pace &amp; Pasquini 2004; Giardino et al. 2008). Using the Palomar Transient Factory (PTF), we have begun a survey of stellar rotation to map out the late-stage evolution of the age-rotation-activity relation: the Columbia/Cornell/Caltech PTF (CCCP) survey of open clusters. The first CCCP target is the nearby ∼600 Myr Hyades-analog Praesepe; we have constructed PTF light curves containing &gt;150 measurements spanning more than three months for ∼650 cluster members. We measure rotation periods for 40 K &amp; M cluster members, filling the gap between the periods previously reported for solar-type Hyads (Radick et al. 1987; Prosser et al. 1995) and for a handful of low-mass Praesepe members (Scholz &amp; Eislöffel 2007). Our measurements indicate that Praesepe's period-color relation undergoes a transition at a characteristic spectral type of ∼M1 -- from a well-defined singular relation at higher mass, to a more scattered distribution of both fast and slow-rotators at lower masses. The location of this transition is broadly consistent with expectations based on observations of younger clusters and the assumption that stellar-spin down is the dominant mechanism influencing angular momentum evolution at ∼600 Myr. Combining these data with archival X-ray observations and Hα measurements provides a portrait of the ∼600 Myr age-rotation-activity relation (see contribution by Lemonias et al. in these proceedings). In addition to presenting the results of our photometric monitoring of Praesepe, we summarize the status and future of the CCCP survey.

12201882

2011ASPC..448..347B

1012

1012.1856_arXiv.txt

Low--mass dwarfs ($0.08~M_{\odot} < M < 0.8~M_{\odot}$) are the dominant stellar component of the Milky Way, composing $\sim 70\%$ of all stars \citep{1997AJ....113.2246R,2010AJ....139.2679B} and nearly half the stellar mass of the Galaxy. However, despite their abundance, M dwarfs were studied in relatively small samples, due to their dim intrinsic brightnesses ($L \lesssim 0.05~L_{\odot}$). However, the situation has been radically altered in the last decade, as deep surveys covering large areas of the sky were carried out. These projects, such as the Sloan Digital Sky Survey \citep[SDSS,][]{2000AJ....120.1579Y} and the Two--Micron All Sky Survey \citep[2MASS,][]{2006AJ....131.1163S}, can trace their roots back to photographic surveys, epitomized by the National Geographic Society - Palomar Observatory Sky Survey \citep[POSS-I,][]{1963bad..book..481M} and its successor, POSS-II \citep{1991PASP..103..661R}. However, the photometric precision achieved by modern surveys distinguishes them from their photographic predecessors. For example, SDSS has imaged 1/4 of the sky to $r \sim 22$ and 2MASS imaged the entire sky to $J \sim 16.5$, with typical uncertainties of a few percent. The resulting databases contain accurate multi-band photometry of tens of millions of low--mass stars, enabling exciting new science. While there is a rich heritage of historical investigations, this article will focus on results derived from SDSS data. A multitude of studies have focused on the intrinsic properties of low--mass stars using SDSS observations. The field luminosity function (LF) and corresponding mass function (MF) was measured using over 15 million stars \citep{2010AJ....139.2679B}. The mass--radius relation of M dwarfs has been studied with eclipsing binary systems \citep{2008ApJ...684..635B, 2008MNRAS.386..416B}. Average photometric colors and spectroscopic features have also been quantified \citep{2002AJ....123.3409H, 2007AJ....133..531B, 2007AJ....134.2430D,2008AJ....135..785W, west10}. Multiple studies have attempted to estimate the absolute magnitude of low--mass stars in SDSS \citep{2002AJ....123.3409H, 2008ApJ...673..864J, 2008ApJ...689.1244S, bochanskithesis, 2009MNRAS.396.1589B, bochanski10}. However, due to the lack of precise trigonometric parallaxes for many of the M dwarfs in SDSS, absolute magnitudes and distances are derived by secondary means. Chromospheric activity, driven by magnetic dynamos within the stars, has been observationally traced with H$\alpha$ emission measured in SDSS spectroscopy \citep{2004AJ....128..426W, 2008AJ....135..785W, west10,2010ApJ...722.1352K}. Flare rates have been measured for thousands of stars \citep{2009AJ....138..633K, 2010AJ....140.1402H} and will be incorporated into predicting the flaring population observed by next--generation surveys \citep{hilton10}. While interesting objects in their own right, M dwarfs are also powerful tools for studying the Milky Way. Their ubiquity, combined with lifetimes much greater than a Hubble time \citep{1997ApJ...482..420L}, make them ideal tracers of Galactic structure and kinematics. Using photometry of over 15 million M dwarfs, \cite{2010AJ....139.2679B} measured the scale heights and lengths of the thin and thick disks. A complementary study by \cite{2008ApJ...673..864J} used higher mass stars to measure the stellar density profiles. The local gravitational potential has been probed using the kinematics of $\sim$ 7000 M dwarfs along one line of sight with proper motions, SDSS spectroscopy and photometry \citep{2007AJ....134.2418B}. This study has been expanded to the entire SDSS footprint, using a database of $\sim$ 25,000 M dwarfs with velocities and distances \citep{pineda}. Furthermore, \cite{2009AJ....137.4149F} employed proper motions and distances of $\sim$ 2 million M dwarfs to estimate velocity distributions. Below, I highlight several studies that have used SDSS observations to study both the intrinsic properties of low--mass stars \emph{and} use them to study the Milky Way. In \S \ref{sec:observations}, the technical details of SDSS photometric and spectroscopic observations are briefly described. In \S \ref{sec:lf}, a new investigation measuring the M dwarf field LF and MF and Galactic structure parameters is detailed. Kinematics within the Milky Way are explored in \S \ref{sec:kinematics}. The importance of placing large, deep surveys of M dwarfs in a Galactic context is discussed in \S \ref{sec:case_study}. Concluding remarks and avenues for future investigations are discussed in \S \ref{sec:conclusions}.

\label{sec:conclusions} I have briefly detailed some of the investigations that examine both intrinsic M dwarf properties and Galactic structure and kinematics. The requisite data for such studies, namely precise, deep, multi-band photometry over large areas of the sky, will become more common in the future, as additional large surveys come online. Within the next decade, surveys such as PanSTARRS, LSST and GAIA \citep{2002SPIE.4836..154K, 2008arXiv0805.2366I, 2001A&A...369..339P}, will produce new data sets that are ideal for similar investigations. While databases containing millions of low--mass stars will be the norm for future studies, it is important to note that systematic errors will be the dominant source of error in many areas. In SDSS, nearly all the observed M dwarfs do not have reliable trigonometric parallaxes, forcing the use of photometric parallax estimates. Reliably calibrating and testing these photometric parallaxes will only become more important as surveys push to larger distances. Other systematics, such as how metallicity affects the absolute magnitude of M dwarfs, will also be crucial to future investigations. Finally, the importance of SDSS spectroscopy can not be overstated. The large spectroscopic samples of SDSS M dwarfs have enabled many novel investigations. Significant spectroscopic followup of the next generation of surveys should be a high priority.

1012.1856

Modern sky surveys, such as the Sloan Digital Sky Survey and the Two-Micron All Sky Survey, have revolutionized the study of low-mass stars. With millions of photometric and spectroscopic observations, intrinsic stellar properties can be studied with unprecedented statistical significance. Low-mass stars dominate the local Milky Way and are ideal tracers of the Galactic potential and the thin and thick disks. Recent efforts, driven by SDSS observations, have sought to place the local low-mass stellar population in a broader Galactic context. I highlight a recent measurement of the luminosity and mass functions of M dwarfs, using a new technique optimized for large surveys. Starting with SDSS photometry, the field luminosity function and local Galactic structure are measured simultaneously. The sample size used to estimate the LF is nearly three orders of magnitude larger than any previous study, offering a definitive measurement of this quantity. The observed LF is transformed into a mass function and compared to previous studies. Ongoing investigations employing M dwarfs as tracers of Galactic kinematics are also discussed. SDSS spectroscopy has produced databases containing tens of thousands of low-mass stars, forming a powerful probe of the kinematic structure of the Milky Way. SDSS spectroscopic studies are complemented by large proper motion surveys, which have uncovered thousands of common proper motion binaries containing low-mass stars. Additionally, the SDSS spectroscopic data explore the intrinsic properties of M dwarfs, including metallicity and magnetic activity. The highlighted projects demonstrate the advantages and problems with using large data sets and will pave the way for studies with next-generation surveys, such as PanSTARRS and LSST.

12104221

2010ApJ...725L.101A

1012

1012.1584_arXiv.txt

1012.1584

We present results of 2 hr non-interrupted observations of solar granulation obtained under excellent seeing conditions with the largest aperture ground-based solar telescope—the New Solar Telescope (NST)—of Big Bear Solar Observatory. Observations were performed with adaptive optics correction using a broadband TiO filter in the 705.7 nm spectral line with a time cadence of 10 s and a pixel size of 0farcs0375. Photospheric bright points (BPs) were detected and tracked. We find that the BPs detected in NST images are cospatial with those visible in Hinode/SOT G-band images. In cases where Hinode/SOT detects one large BP, NST detects several separated BPs. Extended filigree features are clearly fragmented into separate BPs in NST images. The distribution function of BP sizes extends to the diffraction limit of NST (77 km) without saturation and corresponds to a log-normal distribution. The lifetime distribution function follows a log-normal approximation for all BPs with lifetime exceeding 100 s. A majority of BPs are transient events reflecting the strong dynamics of the quiet Sun: 98.6% of BPs live less than 120 s. The longest registered lifetime was 44 minutes. The size and maximum intensity of BPs were found to be proportional to their lifetimes.

12292125

2012PhLB..706..352D

1012

1012.4642_arXiv.txt

Elko spinor fields are unexpected spin one-half matter fields endowed with mass dimension $1$ \cite{elko,elko2}. Since its recent theoretical discovery, it has attracted much attention, in part by the wide range of possibility opened by such peculiar matter fields in cosmology and physics \cite{COSMO} and in part from the mathematical point of view \cite{MATE}. The word Elko is the acronym for {\it Eigenspinoren des Ladungskonjugationsoperators} or Dual-helicity eigenspinors of the charge conjugation operator (see Eq. (\ref{e2})). The two aforementioned characteristics of Elko (namely, spin one-half and mass dimension $1$) makes quite reduced the possible coupling to the Standard Model fields. In fact, keeping in mind that interaction terms with mass dimension greater than four should be severely suppressed by some fundamental mass scale and focusing in simple power counting renormalizable arguments, it turns out that Elko spinor fields may have quartic self-interaction and an Elko-Higgs (doublet) interaction\footnote{We shall emphasize that Elko does not carry standard $U(1)$ gauge invariance \cite{elko}.}. In this vein, such spinor field may act as a dark matter candidate. Another interesting feature about Elko is its non-locality. Elko spinor fields do not belongs to a standard Wigner's class \cite{WIG}. It was demonstrated, however, that Elko breaks Lorentz symmetry (in a subtle way) by containing a preferred direction \cite{ALS}. It is worth to note that the existence of a preferred direction -- the so-called `axis of evil' -- (as well as a self interaction) is believed to be a property of dark matter \cite{MAG}. We also remark, for completeness, that the quantum field associated to the Elko spinor is now better understood in the scope of Very Special Relativity (VSR) framework \cite{VSR}. In fact, it is possible to describe, or construct, Elko spinor fields as the spinor representation of $SIM(2)$ subgroup of VSR \cite{SAL}. In this vein, since $SIM(2)$ is the largest subgroup of VSR encompassing all the necessary physical symmetries except some (violated) discrete symmetry, the tension between Elko and Lorentz symmetries disappears. On the other hand, it is well known that accelerators will test, in a incontestable way, theories in the scope of physics beyond the Standard Model as well as shed some light to the mass generation problem\cite{Nath:2010zj,ATLAS,Martin:1997ns,AguilarSaavedra:2005pw}. Candidates of dark matter predicted in particle physics theories, like supersymmetry, are on the focus of such studies and the answers will provide additional information for a deeper level of our understanding on astrophysics and cosmology. In such way, the CERN Large Hadron Collider (LHC) results are fundamental for any study connecting high energy physics and astrophysics/cosmology. The LHC will provide center-of-mass energy enough to probe directly the weak scale and the origin of mass. Therefore, since we still have the open question of the dark matter nature, it is possible the study of the origin of mass as well as the candidate to the dark matter in the search of Elko. In considering some specific process for Elko production, radiative corrections must be taken into account. In this case, as we will see, the Elko non-locality is manifest leading to an exclusive output in the final signature. At phenomenological grounds, such a behavior suggests a different analysis for the search of Elko at accelerators. So, we consider in some detail a tree level process (where the non-locality is absent) concerning to the Elko production at the LHC, whose signature is $ \mu^{+}+\mu^{-}+2\varsigma$. Such process includes the quartic self-interaction and a coupling with the Higgs scalar field. This paper is organized as follows: In the next Section we introduce some formal aspects of the Elko spinor fields calling attention to the main characteristics that will be relevant in the subsequent analysis. In the Section III we explore the Elko non-locality, when considering radiative corrections. In the Section IV we analyze the tree level case of a viable cross-section for Elko production at the LHC. Then, we move forward investigating some peculiar aspects of our signal. In the last Section we conclude.

1012.4642

We study the prospects of observing the presence of a relatively light Elko particle as a possible dark matter candidate, by pointing out a typical signature for the process encompassing the Elko non-locality, exploring some consequences of the unusual Elko propagator behavior when analyzed outside the Elko axis of propagation. We also consider the production of a light Elko associated to missing energy and isolated leptons at the LHC, with center of mass energy of 7 and 14 TeV and total luminosity from 1 fb<SUP></SUP> to 10 fb<SUP></SUP>. Basically, the Elko non-locality engenders a peculiar signal in the missing energy turning it sensible to the angle of detection.

12168273

2011ApJ...729...39O

1012

1012.1251_arXiv.txt

Currently, $>$5,000 objects are regarded as planetary nebulae (PNe) in the local group galaxies. Of them, BoBn1 (Otsuka et al. 2008; Otsuka et al. 2010), Wray16-423, StWr2-21 (Kniazev et al. 2008; Zijlstra et al. 2006), and Hen2-436 (Walsh et al. 1997; Dudziak et al. 2000; this paper) belong to the Sagittarius (Sgr) dwarf galaxy. Sgr dwarf galaxy PNe are interesting objects as they provide direct insight into old, low-mass stars and also information on the history of the Galactic halo. In some galactic evolution scenarios, the galactic halo is partly built up from the tidal destruction and assimilation of dwarf galaxies. Sgr dwarf galaxy PNe are ideal laboratories and appropriate references to study the evolution of metal-deficient stars found in the Galactic halo and low-mass stars ($\lesssim$3.5 $M_{\odot}$) in the LMC. The metallicity of Sgr dwarf galaxy PNe is relatively low, e.g., $<$0.01 $Z_{\odot}$ in BoBn1 or $\sim$0.5 $Z_{\odot}$ in the others, which is equal to the typical metallicity of the LMC, and all the Sgr dwarf galaxy PNe are C-rich ([C/O]$\gtrsim0$; e.g., Zijlstra et al. 2006) based on gas-phased C and O abundances. The Sgr dwarf galaxy PNe and LMC PNe would complement each other in a study of metal-deficient PNe. In addition, since the distance to the Sgr dwarf galaxy is well determined (24.8 kpc; Kunder \& Chaboyar 2009) and the interstellar reddening to the Sgr dwarf galaxy PNe is relatively low ($E(B-V)$ $\lesssim$0.1), we can accurately estimate intrinsic flux densities. Moreover, by using the spatially highly resolved images taken by Hubble Space Telescope ($HST$) and 8-m class ground based telescopes, we can estimate the size and shape of the nebulae. The nebular size is an important parameter in building spectral energy distributions (SED). We have collected data on several Sgr dwarf galaxy PNe to investigate elemental abundances, dust mass, and evolutionary status of the central star in a metal-deficient environment. In this paper, we analyze the most recently secured spectral data of the Sgr dwarf galaxy PN Hen2-436 (PN G004.8--22.7). Hen2-436 is an interesting object in terms of chemical abundances and dust production in metal-poor environments. Zijlstra et al. (2006) argued that dust exists within the nebula. Sterling et al. (2009) detected [Kr\,{\sc iii}]$\lambda$2.19 $\mu$m and [Se~{\sc iv}]$\lambda$2.29 $\mu$m lines in the Gemini/GNIR spectra, although the amounts of these elements are not estimated yet. In the extra-Galactic PNe, these slow neutron capture elements ({\it s}-process elements) have so far only been detected in BoBn1 (Otsuka et al. 2010) and this nebula. The detection of such rare elements in PNe is highly interesting. For Hen2-436, we performed a comprehensive chemical abundance analysis based on optical ESO/VLT FORS2 spectra, near-IR spectra from Magellan/MMIRS, and mid-IR $Spitzer$/IRS spectra. From FORS2 spectra, we found candidate detections of fluorine (F), phosphorus (P), and krypton (Kr) forbidden lines, and we estimate abundances of these elements. The estimations of F and P are done for the first time. Both elements are synthesized by neutron capture in the He-rich intershell of AGB stars. Through spectral energy distribution (SED) modeling with the photo-ionization code {\sc Cloudy} (Ferland 2004), we infer the evolutionary status of the central star and try to estimate the dust mass. We further investigate the evolutionary status of the progenitors of Hen2-436 and the other Sgr dwarf galaxy PNe with similar metallicity. The observed chemical abundances are compared with theoretical nucleosynthesis model predictions.

We estimated elemental abundances in the Sgr dwarf galaxy PN Hen2-436 based on the archived ESO/VLT FORS2 and $Spitzer$/IRS spectra. We detected candidates of [F\,{\sc ii}]$\lambda$4790, [Kr\,{\sc iii}]$\lambda$6826, and [P\,{\sc ii}]$\lambda$7875 for the first time, which indicates that these elements are largely enhanced. We found a co-relation between C and F, P, and Kr abundances among PNe and C-rich stars. The detections of F, P and Kr in Hen2-436 support that F, P, and Kr together with C are certainly synthesized in the same layer and brought to the stellar surface by the third dredge-up. We detected some N~{\sc ii} and O~{\sc ii} ORLs and derived the ionic abundances from these lines. The discrepancy between O ORL and CEL abundances is $>$1 dex. We constructed a SED model considering dust and estimated the initial mass of the progenitor to be $\sim$1.5-2.0 $M_{\odot}$ with $Z$=0.008 and the age to be $\sim$3000 yr after the AGB phase. Hen2-436 shares its evolutionary status with the Sgr dwarf galaxy PNe StWr2-21 and Wray16-423. The observed elemental abundances of the three Sgr dwarf galaxy PNe could be explained by a theoretical nucleosynthesis model with initial mass 2.25 $M_{\odot}$, $Z$=0.008, and LMC chemical abundances. The SED model predicted that $>$2.9(--4) $M_{\odot}$ of carbon dust co-exists in the ionized nebula. Hen2-436 seems to have experienced evolution similar to LMC PNe. Based on the assumption that most of the observed dust is formed during the last two thermal pulses and the dust-to-gas mass ratio is 5.58(--3), the dust mass-loss rate and the total mass-loss rate are $<$3.1(--8) $M_{\odot}$ yr$^{-1}$ and $<$5.5(--6) $M_{\odot}$ yr$^{-1}$, respectively. Our estimated dust mass-loss rate is comparable to a similar metallicity and luminosity Sgr dwarf galaxy AGB star.

1012.1251

We have estimated elemental abundances of the planetary nebula (PN) Hen2-436 in the Sagittarius (Sgr) spheroidal dwarf galaxy using ESO/VLT FORS2, Magellan/MMIRS, and Spitzer/IRS spectra. We have detected candidates of fluorine [F II] λ4790, krypton [Kr III] λ6826, and phosphorus [P II] λ7875 lines and successfully estimated the abundances of these elements ([F/H] = +1.23, [Kr/H] = +0.26, [P/H] = +0.26) for the first time. These elements are known to be synthesized by the neutron capture process in the He-rich intershell during the thermally pulsing asymptotic giant branch (AGB) phase. We present a relation between C, F, P, and Kr abundances among PNe and C-rich stars. The detections of these elements in Hen2-436 support the idea that F, P, Kr together with C are synthesized in the same layer and brought to the surface by the third dredge-up. We have detected N II and O II optical recombination lines (ORLs) and derived the N<SUP>2+</SUP> and O<SUP>2+</SUP> abundances. The discrepancy between the abundance derived from the oxygen ORL and that derived from the collisionally excited line is &gt;1 dex. To investigate the status of the central star of the PN, nebula condition, and dust properties, we construct a theoretical spectral energy distribution (SED) model to match the observed SED with CLOUDY. By comparing the derived luminosity and temperature of the central star with theoretical evolutionary tracks, we conclude that the initial mass of the progenitor is likely to be ~1.5-2.0 M <SUB>sun</SUB> and the age is ~3000 yr after the AGB phase. The observed elemental abundances of Hen2-436 can be explained by a theoretical nucleosynthesis model with a star of initial mass 2.25 M <SUB>sun</SUB>, Z = 0.008, and LMC compositions. We have estimated the dust mass to be 2.9×10<SUP>-4</SUP> M <SUB>sun</SUB> (amorphous carbon only) or 4.0×10<SUP>-4</SUP> M <SUB>sun</SUB> (amorphous carbon and polycyclic aromatic hydrocarbon). Based on the assumption that most of the observed dust is formed during the last two thermal pulses and the dust-to-gas mass ratio is 5.58 × 10<SUP>-3</SUP>, the dust mass-loss rate and the total mass-loss rate are &lt;3.1×10<SUP>-8</SUP> M <SUB>sun</SUB> yr<SUP>-1</SUP>and &lt;5.5×10<SUP>-6</SUP> M <SUB>sun</SUB> yr<SUP>-1</SUP>, respectively. Our estimated dust mass-loss rate is comparable to a Sgr dwarf galaxy AGB star with similar metallicity and luminosity.

2946054

2010arXiv1012.0707H

1012

1012.0707_arXiv.txt

When the first extrasolar planets (``exoplanets'') were discovered, it became necessary to find names for individual objects. Given the large numbers of exoplanets, the Solar System convention of using mythological or other real names was utterly impractical and astronomically uninformative. There already exist several naming conventions which have grown out of the historical needs of the visual and spectroscopic binary communities. For instance, the components of visual binaries tend to be labeled with capital letters (e.g. $\xi$\,UMa\,A \& B), whereas spectroscopic binaries tend to be labeled with small-case letters or numbers (e.g. $\xi$\,UMa\,A consists of two stars, ``Aa'' and ``Ab'', ``Aa'' being the primary). However, this system is neither officially defined and sanctioned by the I.A.U. nor is its use in the literature uniform. Indeed, the notation used in earlier literature was often the opposite convention, i.e. capital letters for primary stars and the matching lower-case letters for the secondaries. For example, the names of the components of $\xi$\,UMa used by Aitken (\cite{aitken}; p. 249), `A'' \& ''a'' for the two components of the ``A'' system and ``B'' \& ``b'' for the two components of the ``B'' system, are carefully renamed by Griffin (\cite{griffin}; p. 276) to ``Aa'', ``Ab'', ``Ba'', and ``Bb''. Most authors relieve themselves from the obviously onerous task of giving the components of spectroscopic or astrometric binaries names by using the terms ``primary'' and ``secondary'' or designating them as "1" or "2", particularly as indices of dynamical parameters. This problem with the naming of stars in multiple systems is well-known and the source of many discussions in various commissions within the I.A.U. (see Hartkopf \& Mason \cite{HandM}). The provisional working standard adopted during the XXIV I.A.U. convention is that of the Washington Mulitplicity Catalog (WMC), which uses the following system: \begin{itemize} \item the brightest component is called ``A'', whether it is initially resolved into sub-components or not; \item subsequent distinct components not contained within ``A'' are labeled ``B'', ``C'', etc.; \item sub-components are designated by the concatenation of one or more suffixes with the primary label, starting with lower-case letters for the 2nd hierarchical level and then with numbers for the 3rd. \end{itemize} This system makes no distinction between stellar, sub-stellar, and planetary objects but does express a clear hierarchical structure. One problem with this system is that the discovery hierarchy is not necessarily identical with the dynamical hierarchy: does HD\,97950\,C orbit around HD\,97950\.A or perhaps around HD\,97950\,A+HD\,97950\,B? Systems of the latter type are often expressed as HD\,97950\,AB, i.e. by concatenating the component suffixes: the WMC contains references to things like ``A-BC'', i.e. a triple system consisting of the brightest component, ``A'', orbiting around a fainter binary, ``B''+``C''. However, this nomenclature is also not adequate enough to express the dynamical state of just one of the components. Another problem with the WMC nomenclature is that the names are purely accidental and/or historical: we are used to referring to ``Sirius B'', not ``Sirius Ab''; and ``Sirius AB'' can be a reference to both stars together or a capitalized misprint of just one. Finally, the decision, what to call the ``A'' component is arbitrary: historically, visual binaries yielded the widest binaries and so defined the upper-case usage and spectroscopic binaries came later, inducing the lower-case additions, but it is equally possible to first discover a binary using some other method, which then implies ``A'' and ``B'' (instead of ``Aa'' and ``Ab'') components, and to discover a companion system later astrometrically, which then would be the ``C'' component (rather than ``B''), even though the latter might have a totally different dynamical relationship to the first two objects. Given that the situation for multiple stars is confusing enough, when it came to naming exoplanets the simplest solution was to name the planets around single stars using a variation of the WMC convention: if the host planetary system is the ``A'' component (i.e. may or may not be a member of a hierarchical stellar system), then the first exoplanet was considered to be the secondary sub-component and should have been given the suffix ``Ab''. For example, 51\,Peg\,Aa is the host star in the planetary system 51\,Peg, and the first exoplanet is then 51\,Peg\,Ab. Since most exoplanets are in single star systems, the implicit ``A'' designation was simply dropped, leaving the exoplanet name with the lower-case letter only: 51\,Peg\,b. This meant that researchers from the exoplanetary community have adopted what we will refer to as the ``lower-case b'' nomenclature, i.e. without the reference to the primary component and probably have little or no knowledge of the historical nomenclature behind it. The usage of the lower-case b notation is not universal, however: e.g. the planets around the pulsar PSR\,1257+12 were long labeled numerically starting with the index 1 but have also been labeled with lower-case letters starting with ``a'' (Currie \& Hansen \cite{currie}) and upper-case letters starting with ``A'' (Wolszczan \cite{wolszczan}). Thus, the situation is far from uniform even for exoplanets. The usual notation becomes dangerous when considering exoplanets around the stars in binary systems, e.g. $\tau$\,Boo\,b is the name given to the first planet discovered around the primary star of the $\tau$\,Boo binary system, but could $\tau$\,Boo\,c be the 2nd star around the primary or the 1st star around the secondary? Fortunately, now that there are a few planets of this kind -- 16\,Cyg\,Bb, 30\,Ari\,Bb, $\tau$\,Boo\,Ab, HD\,178911\,Bb, HD\,41004\,Ab \& Bb -- the planets around the secondary stars have to date been correctly named. The implicit system for exoplanet names utterly failed with the discovery of circumbinary planets in systems like HW\,Vir (2 planets; Lee et al. 2009), DP\,Leo (1 planet; Qian et al. 2010), and NN\,Ser (2 planets; Beuermann et al. 2010). Lee et al. tried to circumvent the naming problem in HW\,Vir by calling the two planets ``HW Vir 3'' and ``HW Vir 4'', i.e. the latter is the 4th object -- stellar or planetary -- discovered in the system HW\,Vir, which is inconsistent with a similar convention already used for pulsar planets in the literature, where the first planet was labelled, e.g. PSR\,1257+12\,\#1 (these pulsar exoplanets are now registered in exoplanet.eu using the lower-case b notation). In the case of NN\,Ser, Beuermann et al. were confronted with multiple suggestions from various offical sources and finally chose to use the designation NN\,Ser\,c and NN\,Ser\,d, i.e. implicitly NN\,Ser\,Ac and NN\,Ser\,Ad with the central very close binary system composed of NN\,Ser\,Aa and NN\,Ser\,Ab. This solution conflicts with the standard usage of ``A'' and ``B'' for the primary and secondary stars in (pre-)cataclysmic variables and places the two stars and the two planets on the same hierarchical level. The official alternative would have been either to declare NN\,Ser\,Aa+Ab as one dynamical component with the exoplanets NN\,Ser\,B and NN\,Ser\,C orbiting around it, which would have described the dynamical separation of the stars and planets more explicitly but would have placed the planets on a higher hierarchical level than the stars (at least semantically) or to adopt the standard usage of NN\,Ser\,A \& B for the close binary stars, leaving the planets as NN\,Ser\,C \& D. No matter how hard one tries, the designations for the two circumbinary planetary systems are confusing and seemingly incompatible with the common usage for the other exoplanets. Naming conventions are not physically important -- no one really cares if an object is called Sirius\,B, $\alpha$ CMa\,Ab, GJ\,244 \#2, RXF\,J064508.6-164240, or ``Rover'', but names convey both historical {\em and} physical information about the object and the naming convention used should at least not confuse. This is particularly true for the benefit of observers, who are definitely interested in knowing which object on the sky is meant by what name. Unlike the multiple star community, which is suffering from over a century of jumbled naming conventions, the exoplanet community is still sufficiently young that it is possible to adopt a uniform nomenclature which maximizes the usefulness of the names and minimizes the amount of confusion while consciously staying as close as possible to the provisional I.A.U. multiple star naming standard. The purpose of this letter is to propose a simple, maximally compatible and yet physically informative solution for this problem, in the hopes that the I.A.U. (or at least Commision 53) would eventually adopt it for universal usage.

We have proposed a slight revision of the provisional I.A.U. nomenclature standard for multiple systems with the intent of reaching an effective and simple naming convention for extrasolar planets. Our proposal is nearly 100\% compatible with both the standard and common exoplanet usage and yet permits one to distinguish clearly between the dynamical status of planets around single stars, stars in multiple systems, and circumbinary (or higher order) planets. While primarily designed for the exoplanet community, the parenthesis syntax could naturally be used to good effect for stellar multiple systems as well. Thus, we encourage a broadly based discussion on the feasibility of endorsing this (or a similar) convention in the hopes of giving our objects useful and uniform names. \begin{figure} \begin{center} \includegraphics[width=72.0mm]{nomenclature.png} \end{center} \caption{Examples of different exoplanet name suffixes in single and binary systems using the proposed system. Upper left: exoplanet around a single star (e.g. 51\,Peg) plus a moon. Upper right: double star, each with a planet (e.g. HD\,41004), plus a circumbinary planet. Lower left: two circumbinary planets (e.g. NN\,Ser). Lower right: planet around the secondary star in a binary (e.g. HD\,178911).} \end{figure}

1012.0707

The present naming convention for extrasolar planets used by the vast majority of researchers in the field is based upon an interpretation of the provisional I.A.U. standard for multiple star systems. With the existence of hundreds of exoplanets around single stars named by this convention and a handful of exoplanets around binary stars -- circumbinary planets -- it has become necessary to find a uniform and useful naming convention for the latter which is maximally compatible with the single host-star convention and which captures as much of the dynamical information about the planet as possible. We propose a simple and generic naming convention for all exoplanets which follows the provisional I.A.U. standard but more clearly indicates their dynamical status. The proposed convention is compatible with present usage and easily extendible to exoplanets around stars in systems of arbitrary multiplicity. We invite comments and discussion on the proposed convention, in the hope of a timely adoption by the I.A.U. Commissions 5, 8+24, 26, 42, 45 and 53.

12167728

2011ApJ...728...42R

1012

1012.0641_arXiv.txt

Gamma-ray bursts (GRBs) are very fascinating cosmic objects in the universe. Since its discovery in 1973 (Klebesadel et al. 1973), GRBs opened up a new domain of astrophysical research due to the rich observational characteristics of the afterglows in a vast range of electromagnetic spectrum from $\gamma$-rays to radio wavelengths (see Gehrels et al. 2009 for a review). There is a general consensus in the literature that the diverse observational characteristics are due to the interaction of relativistic matter with the surrounding medium. The nature and energy source for this relativistic matter, particularly in the context of the long GRBs, could be in the nature of bulk motion of ions (the Fire Ball Model - see e.g., Piran 2004), cannon balls emitted from a compact object newly formed in a supernova explosion (Dar 2006), or particles accelerated by magnetized winds (the electro-magnetic model - see Lyutikov 2006 and references therein). The long GRBs are associated with supernovae and it is believed that relativistic matter of very high bulk Lorentz factor is generated in a conical jet during the collapse of a massive star at cosmological distances (see for example, Meszaros 2002). The prompt gamma-ray emission has several characteristic correlations like the peak energy $E_p$ against the isotropic luminosity $E_{iso}$ (Amati et al. 2002), spectral lag against the peak luminosity (Norris et al. 2000) etc. (see Gehrels et al. 2009, for a summary of such correlations) and these relations are used even in predicting the red shift of long duration gamma-ray bursts, although with a large uncertainty (close to a factor of 2, see, for example Xiao \& Schaefer 2009). A detailed understanding of the prompt emission is necessary to put these correlations in a firm footing so that GRBs can be used as cosmological candles and also to have a clear understanding of the central engine and the basic jet/cannon-ball emission mechanism. The morphology or temporal profile of the GRBs during the prompt emission varies asymmetrically with no apparent structure among the bursts. Some GRBs show multiple pulses and the individual pulses in a burst is a separate and unique emission with varying amplitude and intensity. In the frame work of Fireball Model (see, Zhang 2007 and Meszaros 2002 for reviews) these pulses are created with different shock strengths at different locations of the jet. Observations suggest that in most of the bursts, an individual pulse profile is in the shape of fast-rise-exponential-decay (FRED), with the width decreasing with energy (Fenimore et al. 1995; Norris et al. 2005). Spectral lag is another spectro-temporal property which is crucial to understand the dynamics and energetics of GRBs and can constrain the GRB models (Ioka \& Nakamura 2001; Shen et al. 2005; Lu et al. 2006). GRB~090618 is a very interesting object for several reasons. It is bright and relatively nearby (redshift $\sim$ 0.5) making it a good candidate to expect to have a visible supernova (if it is like SN1998bw - see Dado \& Dar 2010), though no supernova has yet been associated with this GRB. Further, its intense hard X-ray and gamma-ray emission during the prompt phase enables one to make a time resolved spectral analysis (see for example, Ghirlanda et al. 2010). In this paper, we make a detailed analysis of the prompt emission in a wide band X-ray and gamma-ray region using data from Swift/ BAT and the $RT$-2 Experiment onboard the $Coronas-Photon$ satellite (preliminary results are given in Rao et al. 2009). Since this is the first result from this experiment, we describe in detail the methodology used in deriving the response matrix and spectral fitting. We augment our results by using the publicly available $Swift$ BAT data and make a combined spectral fit. We examine the spectral and temporal characteristics of the individual pulses during the prompt emission of this GRB and investigate the implications to the source emission mechanisms. In \S 2, a summary of observations on GRB~090618 is given and in \S 3 a brief description of the $RT$-2 Experiment is given. Observations and analysis results ($RT$-2 and BAT data) are given in \S 4 and finally in \S 5, a discussion of the results are presented along with relevant conclusions.

In this paper, we have presented a method to measure the spectral parameters of the prompt emission using the recently launched $RT$-2 detectors and the Swift-BAT. Though $RT$-2 has been made to primarily study solar activities, we find that above $\sim$ 50 keV, $RT$-2 essentially acts as an all sky hard X-ray monitor. Thus, it can be used to measure the spectral and timing characteristics of the prompt emission of GRBs. GRB~090618 shows multiple peaks. It shows a systematic softening of the spectrum for the successive pulses which is associated with the variations in the timing parameters. For the successive peaks in the GRB, the peak energy shifts to lower values, the width of the pulse varies sharply with energy and the delay (which is lower for the latter pulses) as a function of energy shows a flatter dependence on energy. The parameter $\xi$, characterizing the dependence of the pulse width with energy, also shows a strong pulse (and hence time) dependence. Fenimore et al. (1995) used the width of individual pulses in several GRBs detected by BATSE and showed that the dependence on energy is a power-law, with an index of $\sim$0.4. Borgonovo et al. (2007) measured the width of several GRBs of known redshifts and found that $\xi$ shows a continuous distribution. A large number of bursts in their work show $\xi$ to be peaking around 0.1 -- 0.2. It appears that GRB~090618 belongs to this class of GRB with a narrow distribution of $\xi$. Interestingly, though the value of $\xi$ changes from pulse to pulse in this GRB, it is in the range of 0 -- 0.1. Various explanations are offered to understand the lags seen in GRBs. Shen et al. (2005) have estimated the contributions from the relativistic curvature effect. These effects can explain contributions to the lag of the order of 10$^{-2}$ -- 10$^{-1}$ s. The lags observed in GRB~090618 is much larger than this and hence quite unlikely to be due to the curvature effect. Ioka \& Nakamura (2001) have computed the kinematic dependence of lag caused by the viewing angle which could produce the observed dependences. The systematic pulse-to-pulse variation of the properties detected in GRB~090618, however, would be difficult to understand in this framework. Spectral evolution, too, can reproduce some part of the lag (Kocevski \& Liang 2003; Hafizi \& Mochkovitch 2007). The results from the present observations support the conclusion of Hakkila et al. (2008) that the spectral lags are pulse rather than burst properties. The individual pulses in a GRB could be either due to emission from multiple shock locations when a jet material encounters the supernova ejecta or due to the weakening of the relativistic matter and/or the emission of a fresh relativistic matter from the central engine. Detection of properties of a GRB pulse quite different from the previous pulses has implications for the nature of the central engine. If the central engine is shown to emit multiple ejection episodes it will strongly favor a black hole as the candidate as against a highly magnetized neutron star (see for example, Metzger 2010 for a discussion on the GRB central engines). In the case of GRB~090618, there is a marginal evidence for the last pulse (pulse 4) to be different and could be a candidate for a separate emission from the central engine. A detailed study of several such multi-peaked bursts are required to draw a firm conclusion. For example, discovery of a hard GRB pulse, after a X-ray pulse, would certainly favor a black hole as a candidate for the central engine. We have calculated the isotropic energy (E$_{iso}$) for this GRB by considering a standard cosmology model for a flat universe with q$_0$ = 1/2 and H$_0$ = 70 km s$^{-1}$ Mpc$^{-1}$. Using the measured redshift of 0.54 and the measured integrated fluence in the energy range of 20 keV to 1 MeV of 2.8 x 10$^{-4}$ ergs cm$^{-2}$, we calculate E$_{iso}$ to be 2.21 $\times$ 10$^{53}$ ergs (beaming effects are neglected). The measured time-averaged peak energy (E$_p$) for the entire burst is around 164 keV, which gives the intrinsic peak energy (E$_{p,i}$) of 252 keV. Based on the measured values of E$_{p,i}$ and E$_{iso}$, it is found that the GRB~090618 closely follows the `Amati' relation with a minor deviation, which is within the 2$\sigma$ scatter. Hence, it could be concluded that GRB~090618 is a standard candle for the category of long duration GRBs alongwith various intrinsic properties that are discussed in this paper. The recent detection of polarization in GRB~090102 (Steele et al. 2009) indicates the presence of ordered magnetic field in the source of GRBs during the prompt emission. The present measurement of the spectral and temporal parameters of GRB~090618 shows that the individual pulses show distinct behaviors.

1012.0641

We present the results of an analysis of the prompt gamma-ray emission from GRB 090618 using the RT-2 Experiment on board the Coronas-Photon satellite. GRB 090618 shows multiple peaks, and a detailed study of the temporal structure as a function of energy is carried out. As the gamma-ray burst (GRB) was incident at an angle of 77° to the detector axis, we have generated appropriate response functions of the detectors to derive the spectrum of this GRB. We have augmented these results using the publicly available data from the Swift Burst Alert Telescope detector and show that a combined spectral analysis can measure the spectral parameters quite accurately. We also attempt a spectral and timing analysis of individual peaks and find evidence for a systematic change in the pulse emission characteristics for the successive pulses. In particular, we find that the peak energy of the spectrum, E<SUB>p</SUB> , is found to monotonically decrease with time, for the successive pulses of this GRB.

12163990

2011AIPC.1331..246V

1012

1012.2644_arXiv.txt

GALEX~J193156.8+011745 (GALEX~J1931+0117, thereafter) is a hydrogen-rich white dwarf \citep{ven2010} showing optical heavy-element lines and an infrared excess. The original low-resolution spectrum obtained with the New Technology Telescope (NTT) at La Silla Observatory showed a strong Mg\,{\sc ii}$\lambda$4481 doublet and weaker silicon lines. Follow-up echelle spectroscopy obtained with the Very Large Telescope (VLT)-Kueyen enabled a detailed abundance study. The near-solar ($\pm0.5$ dex) abundances of oxygen, magnesium, silicon, calcium and iron bear the signature of an external supply of material accreting onto the surface of the white dwarf. Based on available data, Vennes et al. concluded that the supply may originate from a close, sub-stellar companion or from a cool debris disc. The presence of heavy elements in hydrogen-rich white dwarfs has variously been interpreted as intrinsic to the star, or as extrinsic, i.e., supplied by the interstellar medium \citep{dup1993}, by a nearby companion as in post-common envelope systems \citep{deb2006,kaw2008}, or by a debris disc \citep{zuc2003,kil2006,far2008}. However, accretion from the interstellar medium is unlikely because of supply shortages \citep{far2010a}. In the extrinsic scenarios, the elements are accreted and diffused in the atmosphere and envelope of the star \citep[see][]{fon1979,koe2009}. An intrinsic, or internal, reservoir of heavy elements is also possible, but in either scenario a self-consistent solution of the diffusion equation must explore the effect of radiative acceleration on trace elements \citep{cha1995a,cha1995b}. As a class, the polluted DA white dwarfs, or DAZs, are often defined by the detection of the Ca\,{\sc ii} H\&K doublet in optical spectra \citep{zuc2003,koe2005}. Exceptionally, Mg\,{\sc ii}$\lambda$4481 is, so far, only detected in a handful of warm white dwarfs \citep[but see][]{kaw2010} such as EG~102 \citep{hol1997}, GALEX~J1931+0117, and two warm white dwarfs from the Sloan Digital Sky Survey (SDSS) that show evidence of dusty and gaseous discs \citep{gan2006,gan2007}. The presence of a large concentration of magnesium in the last two objects helped establish a strong link between heavy element pollution and dense circumstellar environments. Moreover, an infrared excess that cannot otherwise be explained by a cool companion, may be attributed to a dust ring as in the case of the DAZ white dwarf G29-38 \citep{gra1990}. We present new high-dispersion spectroscopic observations that help elucidate the nature of the peculiar abundance pattern in GALEX~J1931+0117. We revise our abundance measurements and explore the effect of vertical abundance inhomogeneities.

We presented a model atmosphere analysis of the high-metallicity white dwarf GALEX~J1931+0117. The abundance pattern obtained using homogeneous model atmospheres shows that magnesium, silicon, and iron are near or above solar abundances, while carbon, oxygen and calcium are below solar abundances. However, a line profile analysis performed using vertical abundance distributions obtained by solving the steady-state diffusion equation shows that the accretion flow is rich in oxygen, silicon, and iron while it is depleted in other elements. Although the oxygen abundance is below solar in the line forming region, it must be supplied in larger quantity because of its short diffusion time-scale relative to other elements. The effect of pressure shifts are apparent in several strong silicon lines. This effect predicted by \citet{ham1989} and \citet{krs1992} also impacts radial velocity measurements. A lack of radial velocity variations between two epochs that are 123 days apart rules out a close binary scenario for the origin of the accreted material. We are left with the possibility that the infrared excess belongs to a warm disc of debris material that accretes onto the white dwarf surface. In conclusion, our analysis of the abundance pattern in GALEX~J1931+0117 and the absence of a close companion, as well as a comparison with similar cases (e.g., DAZ PG1015+161) support the likely presence of a dusty disc around GALEX~J1931+0117. \begin{theacknowledgments} This research is supported by GA AV grant numbers IAA301630901 and IAA300030908, respectively, and by GA \v{C}R grant number P209/10/0967. \end{theacknowledgments}

1012.2644

We present an analysis of high-dispersion and high signal-to-noise ratio spectra of the DAZ white dwarf GALEX J1931+0117. The spectra obtained with the VLT-Kueyen/UV-Visual Echelle Spectrograph show several well-resolved Si II spectral lines enabling a study of pressure effects on line profiles. We observed large Stark shifts in silicon lines in agreement with laboratory measurements. A model atmosphere analysis shows that the magnesium, silicon and iron abundances exceed solar abundances, while the oxygen and calcium abundances are below solar. Also, we compared the observed line profiles to synthetic spectra computed with variable accretion rates and vertical abundance distributions assuming diffusion steady-state. The inferred accretion rates vary from Ṁ = 2×10<SUP>6</SUP> for calcium to 2×10<SUP>9</SUP> g s<SUP>-1</SUP> for oxygen and indicate that the accretion flow is dominated by oxygen, silicon and iron while being deficient in carbon, magnesium and calcium. The lack of radial velocity variations between two measurement epochs suggests that GALEX J1931+0117 is probably not in a close binary and that the source of the accreted material resides in a debris disc.

12163245

2011A&A...527A.148L

1012

1012.0477_arXiv.txt

\label{sec:intro} The abundance patterns of light elements (up to Al) in globular clusters are in the process of being carefully investigated. In particular, many groups are studying the origin of the larger spread in C, N, O, Na, Mg, and Al abundances compared to field stars of similar metallicity. The present status of the observed light element (O, Na, Mg, and Al) abundances in globular clusters and their possible consequences for the formation and enrichment history of these stellar populations have been presented in a series of publications by \citet[e.g.][]{Carretta09b,Carretta10b} (see also reviews by \citealt{Gratton04} and \citealt{Charbonnel05b}). The main findings, as inferred from high-resolution spectroscopy of individual globular cluster stars, are apparent enhancements of N, Na, and Al abundances and deficiencies in Li, O, and Mg. These patterns can be naturally explained in terms of the enrichment of the nucleosynthesis rest-products of H-burning at high temperatures \citep{Denisenkov89,Langer93}. Since the resulting anti-correlations between the O--Na abundances in particular have not only been seen in evolved red giant branch (RGB) stars, but also in turn-off (TO) and subgiant branch (SGB) stars \citep{Gratton01}, an intrinsic stellar evolutionary cause is very unlikely, i.e. the stars cannot have established these abundance patterns themselves. It is instead believed that the gas that formed second generation stars in globular clusters underwent early pollution by slow ejecta from intermediate or massive stars. In this context, it is significant that photometric evidence of multiplicity has been found in some clusters, e.g. the parallel main sequences identified in $\omega$ Cen and NGC\,2808. These observations seemingly necessitate that a difference in the He content be present \citep[e.g.][]{Norris04}, in qualitative agreement with the self-enrichment process responsible for the other light-element variations. Helium enrichment is also commonly invoked to explain the extended horizontal branch observed in many clusters \citep{DAntona08}. We are, however, far from building a fully consistent picture of the chemical evolution of globular clusters that can explain all the various observations simultaneously. A key unknown is the nature of the polluting objects. One possibility is that so-called hot bottom burning occurs at the base of the convective envelope in intermediate-mass stars during the asymptotic giant branch (AGB) phase, leaving nucleosynthesis products in the envelopes that are subsequently expelled \citep{Ventura08,Ventura10}. In addition, super-AGB stars have been suggested to be responsible for the most extreme anomalies \citep{Pumo08}. The main alternative scenario is a slow mechanical wind from rapidly rotating massive stars \citep[e.g.][]{Decressin07b}, whose envelopes have been enriched in H-burning products by means of deep internal mixing (see Sect.\,\ref{sec:conseq}). Yet another option is mass loss from massive binary systems, as suggested by \citet{deMink09}. By performing accurate abundance analysis of many elements (and isotopic ratios) in large stellar samples, we may be able to pin-point the nature of the progenitors \citep[e.g.][and references therein]{Charbonnel05b}. A common property of the competing scenarios is that the pollution would mainly alter the light element abundances of the second generation stars in globular clusters, thus leaving $\alpha$ and iron-peak elements unaffected. This is necessary to explain the homogeneous composition of these elements seen in most clusters. Elements created in the s-process may be affected by AGB pollution, suggesting that correlations exist between s-process and light element anomalies. In NGC\,6752, an unexplained correlation was indeed identified between Al abundances and Y, Zr, and Ba \citep{Yong05}, but the systematic heavy element variations are small (0.1\,dex) and comparable to the statistical scatter. One must also bear in mind that the yields of AGB stars are uncertain \citep[e.g.][]{Charbonnel07b,Decressin09,Ventura10}. In addition to mapping the presence of abundance trends and correlations, it is also essential to investigate, preferably with sound number statistics, the fraction of stars with normal chemical compositions, similar to those in the field, and the fraction of second and possibly third generation stars in globular clusters. Linking this information to other cluster observables, one may construct a schematic model for the episodes of star formation and evolution. \citet{Carretta10b} describe a possible general formation scenario in which a precursor population, forming from the gas assembled at a very early epoch inside a CDM halo, efficiently raises the metal content of the gas cloud via core-collapse supernova explosions. These trigger a second, large episode of star formation, the so-called primordial population (first generation). The slow winds of massive or intermediate-mass stars of the primordial population feed a cooling flow, and the intermediate (second) generation of stars are formed in the central parts of the cluster, out of material enriched in H-burning products. The remaining gas is dispersed by core-collapse supernovae of the second generation and star formation ceases. The present-day cluster is dominated by the second generation, with a smaller fraction, approximately 30\%, being left of the primordial population. Critical factors determining the outcome of this scenario are the initial mass function (IMF) of the polluting stars, the initial total mass of the cluster, and the amount of mixing between processed gas in the slow stellar ejecta and pristine cluster gas. We discuss these issues in Sect.\,\ref{sec:conseq}. HST photometry of NGC\,6397 produces a remarkably clean HR-diagram, with a very tight main sequence \citep{Richer08} and a very compact blue horizontal branch, i.e. there are no obvious photometric signs of multiple populations. The cluster is well-studied in terms of numbers of stars for which high-resolution spectra have been obtained, but only a handful of elements have previously been analysed even on the RGB \citep[most recently by][]{Castilho00,Gratton01,Thevenin01,Korn07,Carretta09b}, as summarised in Table \ref{tab:mabund}. Early studies of the strengths of the G and CN band in RGB stars in NGC\,6397 (and other clusters) suggested that there are anomalies in their C and N abundances \citep{Bell79,Briley90}. Eventually, \citet{Gratton01} also detected a significant spread in Na abundance for a sample of ten TO stars and RGB stars, findings that clearly pointed to an intrinsic, rather than evolutionary, origin. \citet{Carretta05} corroborated these findings for a larger sample of stars and also found a significant O--Na anti-correlation, as well as a large spread in C and N abundance \citep[see also][]{Pasquini04}. In the latest analysis by \citet{Carretta09a,Carretta09b}, the O--Na anti-correlation is present, although the number statistics are still rather small, oxygen measurements in particular being few in number. A Mg--Al anti-correlation has not been identified in NGC\,6397, but Mg also seems to exhibit a certain scatter \citep{Korn07}. \citet{Lind09b}, hereafter Paper I, presented Na abundances for $>100$ TO, SGB, and RGB stars, and found that the most heavily Na-enriched stars are also significantly depleted in Li. It is thus clear that NGC\,6397, like other globular clusters, should no longer be regarded as a single stellar population despite the tightness of its colour-magnitude diagram. However, even if pollution indeed seems to have taken place in the cluster, it is unclear to which extent, and which elements are affected by it. In this study, which targets red giants, we cover as many elements as possible for a large sample of stars and give a more decisive answer to this question.

\label{sec:conclusions} By studying red giant branch stars in the globular cluster NGC\,6397, we have demonstrated the possibility of distinguishing between the present stellar generations spectroscopically, by making use of a double-peaked histogram of Na abundances. A two-population fit returns a Na abundance similar to the halo field for 25\% of the stars, which we characterise as belonging to a first generation, whereas the remaining 75\% (i.e. second generation) have highly elevated Na abundances. This bimodal abundance signature should be verified for an extended sample, and similar histograms may also be possible to produce for N and Al abundances, for which large spreads are present in globular clusters. The abundance spreads are smaller for O and Mg, making the task more challenging. Highly precise abundance analysis, with small statistical error bars will be necessary to identify these patterns, if indeed present. Based on 17 different elements heavier than Al, we conclude that there is no evidence that $\alpha$, iron-peak or neutron-capture abundances are significantly different between the stellar generations. We have also estimated the difference in He abundance between two stars of each generation by enforcing the mass fraction of iron to be same within the error bars of the analysis. This small (if any) difference in He is expected from models of cluster self-enrichment and also supported by the tightness of the main sequence of NGC\,6397. The same exercise should be performed for other globular clusters, especially those displaying multiple main sequences, hopefully bringing us closer to identifying the process responsible for early cluster pollution. We have finally shown that the abundance patterns observed in NGC\,6397 can be well reproduced within the ``wind of fast rotating massive stars" scenario. On the basis of the Li-Na relation we have been able to infer the dilution factors between the ejecta of massive polluters and interstellar gas of pristine composition. The observed ratio of second to first generation stars was used to estimate the number fraction of first generation long-lived low-mass stars that must have been lost by NGC\,6397, and thus a lower limit to the initial mass of this globular cluster.

1012.0477

Context. The chemical compositions of globular clusters provide important information on the star formation that occurred at very early times in the Galaxy. In particular the abundance patterns of elements with atomic number z ≤ 13 may shed light on the properties of stars that early on enriched parts of the star-forming gas with the rest-products of hydrogen-burning at high temperatures. <BR /> Aims: We analyse and discuss the chemical compositions of a large number of elements in 21 red giant branch stars in the metal-poor globular cluster NGC 6397. We compare the derived abundance patterns with theoretical predictions in the framework of the "wind of fast rotating massive star"-scenario. <BR /> Methods: High-resolution spectra were obtained with the FLAMES/UVES spectrograph on the VLT. We determined non-LTE abundances of Na, and LTE abundances for the remaining 21 elements, including O (from the [OI] line at 630 nm), Mg, Al, α, iron-peak, and neutron-capture elements, many of which had not been previously analysed for this cluster. We also considered the influence of possible He enrichment in the analysis of stellar spectra. <BR /> Results: We find that the Na abundances of evolved, as well as unevolved, stars in NGC 6397 show a distinct bimodality, which is indicative of two stellar populations: one primordial stellar generation of composition similar to field stars, and a second generation that is polluted with material processed during hydrogen-burning, i.e., enriched in Na and Al and depleted in O and Mg. The red giant branch exhibits a similar bimodal distribution in the Strömgren colour index c<SUB>y</SUB> = c<SUB>1</SUB> - (b - y), implying that there are also large differences in the N abundance. The two populations have the same composition for all analysed elements heavier than Al, within the measurement uncertainty of the analysis, with the possible exception of [Y/Fe]. Using two stars with almost identical stellar parameters, one from each generation, we estimate the difference in He content, ΔY = 0.01 ± 0.06, given the assumption that the mass fraction of iron is the same for the stars. <BR /> Conclusions: NGC 6397 hosts two stellar populations that have different chemical compositions of N, O, Na, Mg, and probably Al. The cluster is dominated (75%) by the second generation. We show that massive stars of the first generation can be held responsible for the abundance patterns observed in the second generation long-lived stars of NGC 6397. We estimate that the initial mass of this globular cluster is at least ten times higher than its present-day value. <P />Based on data collected at European Southern Observatory (ESO), Paranal, Chile, under program IDs 077.A-0018(A) and 281.D-5028(A), as well as data collected with the Danish 1.54 m at European Southern Observatory (ESO), La Silla.Tables A.1 and A.2 are only available in electronic form at <A href="http://www.aanda.org">http://www.aanda.org</A>Two tables with line equivalent widths, chemical abundances, and stellar parameters are only available in electronic form at the CDS via anonymous ftp to cdsarc.u-strasbg.fr (130.79.128.5) or via <A href="http://cdsarc.u-strasbg.fr/viz-bin/qcat?J/A+A/527/A148">http://cdsarc.u-strasbg.fr/viz-bin/qcat?J/A+A/527/A148</A>

12273003

2012ASSP...28..151R

1012

1012.5587_arXiv.txt

\label{sec:1} The integrated galactic initial mass function (IGIMF) theory \cite{wk05} is based on the following 3 assumptions: \begin{itemize} \item most stars in galaxies form in star clusters (SCs). Within each SC, the IMF can be approximated by $\xi(m)\propto m^{-\alpha}$, with $\alpha=1.3$ for $m<0.5$ M$_\odot$ and $\alpha=2.35$ for $0.5$ M$_\odot<m<m_{max}$. The value of $m_{max}$ depends on the mass of the SC (the larger the mass, the higher the probability of forming massive stars). \item Also the SCs are distributed according to a power law, $\xi_{ecl}\propto {M_{ecl}}^{-\beta}$, where $M_{ecl}$ is the mass of the SC. Observations suggest $\beta$ to be about 2 \cite{ll03}. \item The maximum possible mass of a SC increases with the star formation rate (SFR) of the galaxy \cite{lr00}. The IGIMF in galaxies depends thus on the SFR. Galaxies with high SFRs contain larger clusters and, consequently, a larger fraction of massive stars. The IMF is therefore flatter than in galaxies with low SFRs. \end{itemize}

1012.5587

It is well established that the [α/Fe] ratios in elliptical galaxies increase with galaxy mass. This relation holds also for early-type dwarf galaxies, although it seems to steepen at low masses. The [α/Fe] vs. mass relation can be explained assuming that smaller galaxies form over longer timescales (downsizing), allowing a larger amount of Fe (mostly produced by long-living Type Ia Supernovae) to be released and incorporated into newly forming stars. Another way to obtain the same result is by using a flatter initial mass function (IMF) in large galaxies, increasing in this way the number of Type II Supernovae and therefore the production rate of α-elements. The integrated galactic initial mass function (IGIMF) theory predicts that the higher the star formation rate, the flatter the IMF. We have checked, by means of semi-analytical calculations, that the IGIMF theory, combined with the downsizing effect (i.e. the shorter duration of the star formation in larger galaxies), well reproduces the observed [α/Fe] vs. mass relation. In particular, we show a steepening of this relation in dwarf galaxies, in accordance with the available observations.

5087248

2011ASPC..444..217I

1012

1012.2472_arXiv.txt

\label{intro} Relativistic flows in association with intense gravitational and magnetic fields are commonly linked up to extremely energetic phenomena in the Universe, viz. pulsar winds, anomalous X-ray pulsars, soft gamma-ray repeaters, gamma-ray bursts, relativistic jets in active galactic nuclei, etc. The necessity to model the aforementioned astrophysical scenarios in the framework of relativistic MHD (RMHD), together with the fast increase in computing power, is pushing towards the development of more efficient numerical algorithms. In the last years, considerable progress has been achieved in numerical special RMHD (SRMHD), by extending the existing high-resolution shock-capturing (HRSC) methods of special relativistic hydrodynamics \citep[e.g.,][]{martilr:03}. In the so called Godunov-type methods, an important subsample of HRSC methods, numerical fluxes are evaluated through the exact or approximate solution of the (local) Riemann problem. Despite the fact that such an exact solution in SRMHD is known \citep{romero,GR06}, approximate algorithms are usually preferred because of their larger numerical efficiency. Several authors \citep[see, e.g.,][and references therein]{Antonetal10} have developed independent {\it Roe-type} algorithms based on linearized Riemann solvers relying on the characteristic structure of the RMHD equations. The purpose of the present paper is twofold. On one hand, the objective is to present a {\it regular} set of right and left eigenvectors of the flux vector Jacobian matrices of the RMHD equations, and span a complete basis in {\it any} physical state, including degenerate states. On the other hand, wish to evaluate numerically the performance of a RMHD Riemann solver based on the aforementioned spectral decomposition. Both the theoretical analysis and the numerical applications presented in this paper are based on the work developed by \cite{Antonetal10}, where we have characterized thoroughly all the degeneracies of RMHD in terms of the components of the magnetic field normal and tangential to the wavefront in the fluid rest frame. Our numerical method deviates in several aspects from previous works based on linearized Riemann solver approaches \citep{komissarov99,Balsara01,koldoba}. First, numerical fluxes are computed from the spectral decomposition in conserved variables. Second, we present explicit expressions also for the left eigenvectors. Third, and most important, we have extended classical MHD strategy \citep{BW88} to relativistic flows, giving sets of right and left eigenvectors which are well defined through degenerate states. Based on the full wave decomposition (FWD) provided by the renormalized set of eigenvectors in conserved variables, we have also developed a linearized (Roe-type) Riemann solver. Extensive testing against one- and two-dimensional standard numerical problems allows us to conclude that our solver is very robust. When compared with a family of simpler solvers that do not require the knowledge of the full characteristic structure of the equations in the computation of the numerical fluxes, our solver turns out to be less diffusive than HLL and HLLC, and comparable in accuracy to the HLLD solver. The amount of operations needed by the FWD solver makes it less efficient computationally than those of the HLL family in one-dimensional problems. However its relative efficiency increases in multidimensional simulations.

1012.2472

In a recent paper (Antón et al. 2010) we have derived sets of right and left eigenvectors of the Jacobians of the relativistic MHD equations, which are regular and span a complete basis in any physical state including degenerate ones. We present a summary of the main steps followed in the above derivation and the numerical experiments carried out with the linearized (Roe-type) Riemann solver we have developed, and some note on the (non-)convex character of the relativistic MHD equations.

12168042

2011ApJ...728..142L

1012

1012.4548_arXiv.txt

Investigation of the flow behavior of the accreting matter in the vicinity of a black hole is important since the spectrum and the intensity of the emitted radiation depend on the flow structure. The event horizon presents the unique inner boundary condition in which the in-falling matter crosses the horizon with the speed of light ($c$). Therefore black hole accretion has to be transonic, as a result of which existence of one sonic point or critical point is assured for black hole accretion. General relativity also ensures that matter must posses sub-Keplerian angular momentum closer to the horizon. Although within the marginally stable circular orbit ($r_{ms}$), the angular momentum at $r{\lsim}r_{ms}$ is definitely sub-Keplerian (and the value $l {\sim} l_{ms}=l|_{r=r_{ms}}$), but at larger radius the angular momentum should be generally large. Therefore, a general accretion disk should have viscosity to remove the angular momentum outwards. The first serious model of viscous accretion disk was presented by \citet{ss73}, in which the angular momentum distribution was Keplerian, the accretion disk was geometrically thin and optically thick. In Shakura-Sunyaev disk the pressure and advection terms were not properly considered and no attempt was made to satisfy the inner boundary condition around a black hole apart from the adhoc termination of the disk at $r \leq r_{ms}$. Along with this theoretical short coming, Shakura-Sunyaev disk also failed to explain power-law high energy part of a black hole candidate spectrum. Therefore, search for another component of an accretion disk which may explain the origin of the high energy radiations from black hole candidates, were undertaken by various groups. One such model which got a wide attention was ADAF [\eg \citet{i77}, \citet{ny94} hereafter NY94]. This model was first constructed around a Newtonian gravitational potential, where the viscously dissipated energy is advected along with the mass, momentum and the entropy of the flow. The original ADAF solution was self-similar and wholly subsonic, and was found to be thermally and dynamically stable. Howover, the low viscosity ADAF showed convective instability \citep{ia99}, that has no dynamical effect if the angular momentum is transported outward but it is dynamically important in case the opposite is true. The global solution of ADAF showed that the flow actually becomes transonic at around few Schwarzschild radii ($r_g$), and the self-similarity may be maintained far away from the sonic point \citep{cal97}. Simultaneous to these developments, there were some interesting research going on sub-Keplerian flows around black holes. It has been shown that sub-Keplerian flow does posses multiple sonic point in a significant range of the energy-angular momentum parameter space \citep{lt80}. One of the consequences of existence of multiple sonic points, is that the flow accreting through the outer sonic point can be slowed down by the centrifugal barrier. This slowed down matter acts as barrier to the faster fluid following it. If the strength of the barrier is strong enough then accretion shocks may form \citep{c89}. General global solutions in the advective domain incorporating viscosity and thermal effects were obtained by many independent researchers \citep{c90,c96,lgy99,lmc98,gl04}. Furthermore, it has also been shown that the global ADAF solution is a subset of the general advective solutions \citep{lgy99}. Whether a flow will follow an ADAF solution or some kind of hybrid solution with or without shock will depend on the outer boundary condition and the physical processes dominant in the disk. Although steady-state solutions are possible in a certain range of parameter space \citep{c89,cd04,mlc94,mrc96a}, but advective solutions with discontinuities such as shocks are generally prone to various kind of instabilities. Since, various flow variables across the shock surface jumps abruptly, this results in a markedly different cooling, heating and other dissipation rates across the shock. This may render the shock unstable. For example, in presence of bremsstrahlung cooling, resonance between cooling timescales and in fall timescales in the post shock part of the disk gives rise to oscillating shocks \citep{msc96b}. \citet{lmc98} showed that beyond a critical viscosity post-shock disk may oscillate. The interaction of the outflow and the inflow may also cause the bending instability in the disk \citep{makbc01}. \citet{mtk99} showed that in presence of non-axisymmetric azimuthal perturbations the shock initially becomes unstable but stabilizes within a finite radial extent into an asymmetric closed pattern. Moreover, the post-shock region may be associated with the elusive Compton cloud that produces the hard photons \citep{ct95,cm06,mc08} and may also be the base of the jet \citep{dc99,dcc01,cd07,dc08,bdl08,dbl09}. Therefore instabilities of the post-shock region may manifest itself as the variabilities observed in the emitted hard photons seen in microquasars and active-galactic nuclei \citep{msc96b}. To add a new twist, \citet{ft04} conjectured the presence of multiple shocks and \citet{gcsr10} actually reported the presence of two oscillating shocks giving rise to two quasi-periodic oscillations. In this paper, we concentrate on the study of instabilities of rotating fluid around black holes, generated by the angular-momentum transport by viscosity. Since the temperature, density etc are higher and the velocity is lower in the post-shock region compared to the pre-shock region, the angular momentum transport rate should be different in the two regions of the disk. In other words, in this paper we simulate transonic, viscous, rotating fluid around black holes. We employ a new code to study the effect of angular momentum transport in the accretion disk. Unlike other purely Eulerian codes, this new code is especially developed to strictly conserve angular momentum in absence of viscosity. In \S 2, governing equations and assumptions are presented. In \S 3, the code which was built to calculate the evolution of angular momentum as accurately as possible is described, along with tests for a rotating transonic flow and a viscous flow. In \S 4, the structure and the instability shown in simulations are presented, along with descriptions on the nature of the instability. A summary and discussion is presented in \S 5.

This paper is intended to study the time-dependent simulations of large amplitude oscillations of advective, viscous, sub-Keplerian disks, to complement earlier works of studying low amplitude oscillations undertaken by Molteni and his collaborators \citep{lmc98}. As an improvement we have employed a new code which uses the Lagrangian TVD/remap approach. This code strictly conserved the angular momentum without viscosity, and reduced the numerical dissipation considerably (\eg \S 3). Tests showed that the shock capturing capability of this code is better than both standard Eulerian code and Lagrangian SPH code (\eg Fig. 1), and followed the angular momentum transfer of the viscous, subsonic analytical solution extremely well (\eg Fig. 2). Oscillation of accretion shock was borne out by the different rates of angular momentum transfer across the shock and the heat dissipated due to the presence of viscosity. It has been shown that in presence of low viscosity parameter ($\alpha=0.003$), the shock front of a disk, with the same initial and boundary conditions as those of the inviscid case, did tend to expand and settled at a larger distance from the disk (Fig. 4b). For an even higher viscosity ($\alpha \gsim 0.005$), the rate of angular momentum transfer was higher, which caused a faster rate of shock front expansion. As the shock front exceeded a possible equilibrium position it started to oscillate (Fig. 4c). However, it is to be remembered that the value of the critical viscosity parameter ($\alpha \sim 0.005$ in the present case) is not sacrosanct, but actually depends on the initial condition. For example, it has been shown that the critical viscosity parameter will be higher for flows with lower angular momentum, while for a fluid with higher initial energy the critical viscosity parameter will be lower \citep[see,][]{cd04}. Hence, if proper initial condition is used then a stable shock is expected to form for higher viscosity parameters (\ie $\alpha \sim 0.1$ -- $0.2$) too, investigation of which, however, is not the point of interest for the present paper. A detailed study of the disk dynamics was conducted for reasonably high viscosity (\ie $\alpha=0.01,~\&~0.1$). For $\alpha=0.01$ the shock oscillation amplitude was found to be quite high $\gsim 100r_g$. This resulted in a large sound speed gradient in the post-shock subsonic flow. In case of large amplitude shock oscillation, the rate of outward angular momentum transport in a region closer to the inner sonic point was shown to be much higher compared to the rate of angular momentum transport near the shock. As a result, our simulation showed that the angular momentum to be piled up in an intermediate region between the shock and the inner sonic point. The expanding shock also increased the inflow velocity in the immediate post-shock region only to be decelerated by the extra centrifugal pressure due to the piled-up angular momentum further inside the disk (\eg Fig. 6a -- 6d). The inflow velocity in the post shock disk may be increased to the extent that it may again become supersonic, then the resistance from the excess centrifugal pressure from the piled-up angular momentum distribution may cause the formation of inner shock. In case of moderately high $\alpha$, the distance between the peak of the angular momentum distribution and the outer shock is large enough to allow for the $v_r$ to become supersonic again and enhanced the possibility to form the inner shock. It is to be noted that the amplitude of shock oscillation will possibly be lesser for multi-dimensional simulation. Viscosity is more active in the post-shock disk, and hence the extra centrifual force due to the piled angular momentum and the heat dissipated by viscosity both actively take part in shock oscillation. However, in case of realistic accretion flow, a part of viscous heat dissipated in the post-shock disk will also be spend to puff it up, which would imply less outward push on the shock surface. Hence for a flow with same injection and viscosity parameters, the oscillation amplitude for a multi-dimensional disk is expected to be lesser compared to a purely conical flow. Consequently, the critical viscosity above which the disk becomes oscillatory will also be higher. The time evolution of shocks for higher viscosity was shown to be distinctly different from that of the lower one. The inner shock was weaker and more sporadic for a disk with $\alpha=0.1$. The main reason was because of the higher rate of angular momentum transport. Even when the shock was around $100r_g$, highly efficient angular momentum transport created a smooth increase of angular momentum, which only peaked closer to $r_{sh}$. As shock expanded $v_r$ increased, but the opportunity to become supersonic was minimized since the peak of the $l(r)$ was closer to the shock. Hence the inner shock, if formed at all, was weaker. However, since shock amplitude for $\alpha=0.1$ was much larger than the case with $\alpha=0.01$, with time the formation of inner shock became more regular, and the behavior was more similar to that of $\alpha=0.01$. The oscillatory motion of the shock induced oscillation in all the disk parameters like emission, rate of matter consumed by the black hole, the rate of angular momentum consumed by the black hole, and the average angular momentum of disk. All these parameters did oscillate with the same period as that of the shock. The disc oscillation started with $\alpha \gsim 0.005$. Considering $M_{BH}=10~M_{\odot}$, for $\alpha=0.005$ the oscillation frequency of the outer shock was $5~Hz$, and the inner shock $10~Hz$, for $\alpha=0.006$ the frequencies were $1~Hz$ and $3~Hz$ respectively, and for $\alpha=0.01$ the two frequencies were $0.125~Hz$ and $0.25~Hz$. Hence one may conclude that apart from the dependence of the oscillation frequency on injection parameters, the QPO frequency definitely decreases with increasing viscosity and vice-versa. Observationally, GRO J1655-40 exhibits a rise in QPO frequency in its rising state and a fall in QPO frequency in its declining phase in 2005 \citep{cdnp08}. \citet{cdp09} plotted the QPO frequency for the object XTE J1550-564 in 1998 burst phase. They showed that in the rising phase of the outburst, the low frequency QPO increases from $0.08 Hz$ to $13.1~Hz$ and then starts to decrease in the declining phase before disappearing. Such rise and fall of QPO frequencies may be explained by the change in shock oscillation frequency due to the change of the net viscosity of the disk. In presence of viscosity a positive angular momentum gradient \ie $dl/dr \geq 0$ helps in outward transport of angular momentum. However, a negative gradient may trigger inward transport of angular momentum. The $dl/dr<0$ condition was attained in the disk in at least two locations, at the outer shock front and just behind the peak of the specific angular momentum distribution. Those regions were subject to the rotational instability. $dl/dr<0$ caused the average angular momentum $<l>$ of the disk to increase, and hence the period and the amplitude of the shock oscillation to increase too. This is less perceptible for lower $\alpha$ and the shock oscillation achieved quasi-saturation, but for $\alpha=0.1$ the shock went outside the computation domain. We repeated the simulation with $\alpha=0.3$ (not presented in the paper) and in this case too the shock went outside the domain, although formation of inner sonic point and oscillation of the two shocks were observed too. In case of multi-dimensional simulations, a part of the post shock matter would have ejected along the vertical direction in the form of winds, which would have carried away a part of the angular momentum, such that the increase of $<l>$ may have been arrested for higher $\alpha$. This would have meant that the shock oscillation may saturate for $\alpha \geq 0.1$. Hence, we conjecture that the non-saturation of shock oscillation for $\alpha\gsim 0.1$, could be an artifact of one-dimensional simulation. We will test it in a future work using multi-dimensional simulations.

1012.4548

We investigated the instability of advective accretion flow as a consequence of angular momentum transfer in one-dimensional, quasi-spherical transonic accretion flow around a non-rotating black hole. The code is designed to include the effects of viscosity; the hydrodynamics component preserves angular momentum strictly with Lagrangian and remap method in the absence of viscosity, while the viscosity component updates viscous angular momentum transfer through the implicit method. We performed two tests to demonstrate the suitability of the code for accretion study. First, we simulated the inviscid, low angular momentum, transonic accretion flow with shocks around a black hole, and then the subsonic, self-similar ADAF solution around a Newtonian object. Both simulations fitted the corresponding analytical curves extremely well. We then simulated a rotating, viscous, transonic fluid with shocks. We showed that for low viscosity parameter, stable shocks at larger distance are possible. For higher viscosity parameter, more efficient angular momentum transfer in the post-shock disk makes the shock structure oscillatory. Moreover, as the shock drifts to larger distances, a secondary inner shock develops. We showed that the inner shock is the direct consequence of the expansion of the outer shock, as well as the creation of regions with ∂l/∂r &lt; 0 due to more efficient angular momentum transfer near the inner sonic point. We showed that all disk parameters, including emissivity, oscillate with the same period as that of the shock oscillation. Our simulation may have implications for low frequency quasi-periodic oscillations, e.g., GRO J1655 - 40 and XTE J1550 - 564.

5395548

2011PhRvD..83f3506N

1012

1012.3460_arXiv.txt

\label{section:intro} The simplest cosmological model consistent with the spatial flatness expectation of inflation is the Einstein-de Sitter (EdS) universe with $\Omega_m=1$ and $\Omega_K=0$. This formed the basis of the ``standard cold dark matter'' (SCDM) cosmology which provided a good description of the large-scale structure of the universe \cite{Blumenthal:1984bp,Davis:1985rj}. It was noted however that it is inconsistent with the angular power spectrum of the clustering of galaxies in the APM survey \cite{Efstathiou:1990xe}. Subsequently large-angle anisotropies in the cosmic microwave background (CMB) were detected by COBE \cite{Smoot:1992td}, thus providing an absolute normalization of the amplitude of primordial density perturbations. The SCDM model assumes that the power spectrum of the primordial density perturbations has the scale-invariant form: $P(k) \equiv |\delta_k|^2 \propto k^n$, with $n=1$ and the predicted amplitude of matter fluctuations on small (cluster and galaxy) scales is then too high relative to observations \cite{White:1992ri}. However if the epoch of matter-radiation equality is delayed by lowering the CDM density to $\Omega_m \sim 0.3$, then the peak in the power spectrum of density fluctuations is shifted to larger scales thus decreasing the power on small scales and enabling a match to the data. To maintain spatial flatness a non-zero value of the cosmological constant was invoked, with $\Lambda \sim 2H_0^2$ corresponding to $\Omega_\Lambda \equiv \Lambda/3H_0^2 \sim 0.7$. Subsequently it was also observed that Type Ia supernovae (SNe~Ia) at redshift $z \simeq 0.5$ appear $\sim25\%$ fainter than expected in an EdS universe \cite{Riess:1998cb,Perlmutter:1998np}. Together with measurements of galaxy clustering in the 2dF survey \cite{Efstathiou:2001cw} and of cosmic microwave background (CMB) anisotropies by WMAP \cite{Spergel:2003cb}, this changed the standard cosmological model to an accelerating universe with a dominant cosmological constant term, which has been widely interpreted as a manifestation of the physical vacuum or ``dark energy''. This ``concordance'' $\Lambda$CDM cosmology (with $\Omega_\Lambda \simeq 0.7$, $\Omega_m \simeq 0.3$, $h \simeq 0.7$) has proved to be consistent with other cosmological data, in particular, baryonic acoustic oscillations detected in the SDSS \cite{Eisenstein:2005su} and measurements of mass fluctuations from clusters and weak lensing \cite{Contaldi:2003hi}. Further observations of both SNe~Ia \cite{Riess:2004nr,Astier:2005qq,WoodVasey:2007jb} and the WMAP 3-year results \cite{Spergel:2006hy} have continued to firm up the model. Embarrassingly, however, this model lacks a {\em physical} basis. There are two serious problems with the notion that the universe is dominated by some form of vacuum energy. The first is the notorious fine-tuning problem of vacuum fluctuations in quantum field theory --- the energy scale of the inferred cosmological energy density is $\sim10^{-12}$ GeV, which is many orders of magnitude below the energy scale of $\sim10^2$ GeV of the Standard Model of particle physics, not to mention the Planck scale of $M_\mathrm{P}\equiv (8\pi G_\mathrm{N})^{-1/2}\simeq2.4\times10^{18}$ GeV \cite{Weinberg:1988cp}. The second is the equally acute coincidence problem: since $\Omega_\Lambda/\Omega_m$ evolves as the cube of the cosmic scale factor $a(t)$, there is no reason to expect it to be of ${\cal O}(1)$ {\em today}, yet this is supposedly the case. Interestingly the WMAP results alone do not require dark energy if the assumption of a scale-invariant primordial power spectrum is relaxed. This is well justified given our present ignorance of the physics underlying inflation which is believed to have created these fluctuations. It has been demonstrated \cite{Hunt:2007dn,Hunt:2008wp} that the temperature angular power spectrum of an EdS universe with $h \simeq 0.44$ matches the WMAP data well if the primordial power is enhanced by $\sim 30\%$ in the region of the second and third acoustic peaks (corresponding to spatial scales of $k \sim 0.01-0.1h~\mathrm{Mpc}^{-1}$). This alternative model with {\em no} dark energy has a slightly better $\chi^2$ for the fit to WMAP-3 data than the concordance ``power-law $\Lambda$CDM model'' and, inspite of having more parameters, has an {\em equal} value of the Akaike information criterion used in model selection. Other EdS models with a broken power-law spectrum \cite{Blanchard:2003du} have also been shown to fit the WMAP data. Moreover, an EdS universe can fit measurements of the galaxy power spectrum if it includes a $\sim 10\%$ component of hot dark matter in the form of massive neutrinos of mass $\sim 0.5$~eV \cite{Blanchard:2003du,Hunt:2007dn,Hunt:2008wp}. Clearly the main evidence for dark energy comes from the SNe~Ia Hubble diagram. It should be kept in mind that the acceleration of the expansion rate is not directly measured but inferred from measurements of the apparent magnitudes and redshifts of SNe~Ia. Indeed \emph{all} the evidence for dark energy is geometrical, i.e., based on interpreting the data in an assumed homogeneous model universe. So far there is no convincing observation of dynamical manifestations, e.g., the ``late integrated Sachs-Wolfe effect''. In fact what is actually inferred from observations is {\em not} an energy density, just a value of ${\cal O}(H_0^2)$ for the otherwise unconstrained $\Lambda$ term in the Friedmann equation. It has been suggested that this may simply be an artifact of interpreting imprecisely measured cosmological data in the oversimplified framework of an universe assumed to be described by the exactly isotropic and homogeneous Friedmann-Robertson-Walker (FRW) metric, in which $H_0 \sim 10^{-42} \mathrm{GeV} \sim (10^{28} \mathrm{cm})^{-1}$ is the {\em only} scale \cite{Sarkar:2007cx}. Note that the non-zero value of $\Omega_\Lambda \equiv \Lambda/3H_0^2$ is inferred from the ``cosmic sum rule'' $\Omega_m + \Omega_K + \Omega_\Lambda=1$, which is just a restatement of the Friedmann equation. If however this equation does not describe the real universe exactly, and in order to do so \emph{other} non-zero terms ought to have been added to the sum rule, then we may mistakenly infer a value for $\Omega_\Lambda$ of ${\cal O}(1)$ if these other terms are in fact important. For example, in an inhomogeneous universe averaged quantities satisfy modified Friedmann equations which contain extra terms since the operations of spatial averaging and time evolution do not commute \cite{Buchert:1999er}. These ``backreaction'' terms depend upon the variance of the local expansion rate and hence increase as inhomogenities develop. However although backreaction behaves just like a cosmological constant, whether its expected magnitude can indeed account for the apparent cosmological acceleration is debated and remains an open question at present \cite{Wetterich:2001kr,Ishibashi:2005sj,Vanderveld:2007cq,Wiltshire:2007fg,Khosravi:2007bq,Leith:2007ay,Behrend:2007mf,Rasanen:2008it,Li:2008yj,Paranjape:2008jc}. Another possibility is that inhomogeneities affect light propagation on large scales and cause the luminosity distance-redshift relation to resemble that expected for an accelerating universe. This has been investigated for a ``Swiss-cheese'' universe in which voids modelled by patches of Lema\'{i}tre-Tolman-Bondi (LTB) space-time are distributed throughout a homogenous background. However, the results depend on the specific model: some authors find the change in light propagation to be negligible because of cancellation effects \cite{Biswas:2007gi,Brouzakis:2007zi,Brouzakis:2008uw,Valkenburg:2009iw}, whereas others claim it can partly mimic dark energy \cite{Marra:2007pm,Marra:2007gc,Kainulainen:2009sx}. It may be that observers preferentially choose sky regions with underdense foregrounds when studing distant objects such as SNe~Ia, so the expansion rate along the line of sight is then greater than average; such a selection effect may also allow an inhomogeneous universe to fit the observations without dark energy \cite{Mattsson:2007tj}. In this paper we are mainly interested in a ``local void'' (sometimes referred to as ``Hubble bubble'') as an explanation for dark energy; to prevent an excessive CMB dipole moment due to our peculiar velocity we must be located near the centre of the void. An underdense void expands faster than its surroundings, thus younger supernovae inside the void would be observed to be receding more rapidly than older supernovae outside the void. Under the assumption of homogeneity this would lead to the mistaken conclusion that the expansion rate of the universe is accelerating, although both the void and the global universe are actually decelerating. The local void scenario has been investigated by several authors using a variety of methods \cite{Moffat:1994qy,Tomita:1999rw,Celerier:1999hp,Tomita:1999qn,Tomita:2000rf,Tomita:2000jj,Tomita:2001gh,Iguchi:2001sq,Tomita:2002df,Moffat:2005yx,Moffat:2005ii,Alnes:2005rw,Mansouri:2005rf,Vanderveld:2006rb,Garfinkle:2006sb,Chung:2006xh,Biswas:2006ub,Alnes:2006uk,Caldwell:2007yu,Alexander:2007xx,Clarkson:2007pz,Uzan:2008qp,GarciaBellido:2008nz,GarciaBellido:2008gd,Clifton:2008hv,Bolejko:2008cm}. By modelling the void as a open FRW region joined by a singular mass shell to a FRW background, it was found that a void with radius 200 Mpc can fit the supernova Hubble diagram without dark energy \cite{Tomita:2001gh}. It was also shown that a LTB region which reduces to a EdS cosmology with $h=0.51$ at a radius of 1.4~Gpc can match both the supernova data and the location of the first acoustic peak in the CMB \cite{Alnes:2005rw}. Ref.\cite{Alexander:2007xx} attempted to find the smallest possible void consistent with the current supernova results --- their LTB-based `minimal void' model has a radius of 350 Mpc. Unfortunately, since this model is equivalent to an EdS universe with $h=0.44$ {\em outside} the void where the SDSS luminous red galaxies lie, as it stands it is unable to fit the measurements of the baryonic acoustic oscillation (BAO) peak at $z\sim0.35$ \cite{Blanchard:2005ev}. LTB models of much larger voids were considered in Ref.\cite{GarciaBellido:2008nz} (with radii of 2.3 Gpc and 2.5 Gpc and Hubble contrasts of 0.18 and 0.30 respectively) and it was demonstrated they can fit the BAO data, as well as the SNe\,Ia data and the location of the first CMB peak. Ref.\cite{Clifton:2008hv} found the best fit to the SNe Ia data for a void of radius $1.3 \pm 0.2$ Gpc and Ref.\cite{Bolejko:2008cm} confirmed that such a void provides an excellent fit to the ``Union'' dataset of SNe\,Ia. In this paper we demonstrate that, contrary to the results obtained in Refs.~\cite{Clifton:2009kx,Zibin:2008vk,Moss:2010jx,Biswas:2010xm}, a Gpc-sized void can simultaneously fit the SNe Ia data as well as the full CMB power spectrum, while also satisfying constraints from local Hubble measurements, primordial nucleosynthesis and the BAO data, if the primordial power spectrum is \emph{not} assumed to be nearly scale-invariant. The layout of the paper is as follows. In Sec.~\ref{section:LTB} we summarize the general relativistic framework for LTB models and describe the characterization of the void. In Sec.~\ref{section:primordial} we discuss the form of the primordial power, and present a physical model with a primordial power spectrum that is not scale-free. In order to compare observables in the void model to existing cosmological data, some formalism needs to be developed. This is done in Sec.~\ref{section:fitting}, and the statistical approach is discussed in Sec.~\ref{section:method}. Finally Sec.~\ref{section:results} presents the main results of the paper.

\label{section:discussion} An interesting proposed test of void models is the Compton $y$-distortion of the CMB spectrum that is produced by the scattering of photons by reionized gas in regions of the void that see a highly anisotropic LSS \cite{Caldwell:2007yu}. In the single-scattering and linear approximations and under the assumption that the dipole anisotropy dominates the distortion, this can be written as \cite{Moss:2010jx} \begin{equation} \label{y} y = \frac{7}{10}\int_0^{z_\mathrm{re}} \mathrm{d}z\frac{\mathrm{d}\tau}{\mathrm{d}z}\beta\left(z\right)^2\,, \end{equation} where $\tau$ is the optical depth, $\beta\left(z\right)$ is the dipole temperature anisotropy in the CMB observed at redshift $z$, and the integral is taken up to the redshift of reionization $z_\mathrm{re}$. The FIRAS instrument on COBE provides an upper bound $y<1.5\times10^{-5}$ (at 2$\sigma$) \cite{Fixsen:1996nj}. While this can constrain some void profiles, voids lacking an overdense outer shell are found not to be in tension with this upper bound \cite{Moss:2010jx}. The Gaussian void profile that we consider does not have a prominent overdense outer shell. Since we do not expect an interesting constraint, we chose not to evaluate the $y$-distortion, which would be uncertain in any case since the exact reionization history (and thus $z_\mathrm{re}$) of the void model is unknown. Another direct test of a local void is via the kinetic Sunyaev-Zel'dovich effect \cite{Sunyaev:1972eq}. This has been applied to cluster kSZ observations (\emph{e.g.}, Ref.\cite{GarciaBellido:2008gd,Yoo:2010ad}) and to full-sky data \cite{Zhang:2010fa} from the Atacama Cosmology Telescope (ACT) \cite{Das:2010ga}, and is found to place constraints on large voids. Allowing an inhomogeneous bang-time or correctly accounting for the effects of radiation (see \cite{Regis:2010iq,Clarkson:2010ej}) may potentially weaken these constraints while preserving the fits described here. Such an analysis is beyond the scope of the current paper but will be investigated in the future. In summary, in this paper we have presented a local void model which fits SNe Ia and CMB data, local $H_0$ values, nucleosynthesis constraints and BAO \emph{without} requiring dark energy and thus provides a counterexample to the claim that dark energy is necessary to fit these observations.

1012.3460

In the standard cosmological model, the dimming of distant Type Ia supernovae is explained by invoking the existence of repulsive “dark energy” which is causing the Hubble expansion to accelerate. However, this may be an artifact of interpreting the data in an (oversimplified) homogeneous model universe. In the simplest inhomogeneous model which fits the SNe Ia Hubble diagram without dark energy, we are located close to the center of a void modeled by a Lemaítre-Tolman-Bondi metric. It has been claimed that such models cannot fit the cosmic microwave background (CMB) and other cosmological data. This is, however, based on the assumption of a scale-free spectrum for the primordial density perturbation. An alternative physically motivated form for the spectrum enables a good fit to both SNe Ia (Constitution/Union2) and CMB (WMAP 7-yr) data, and to the locally measured Hubble parameter. Constraints from baryon acoustic oscillations and primordial nucleosynthesis are also satisfied.

2268900

2011A&A...525A..16P

1012

1012.0149_arXiv.txt

The group of chemically peculiar (CP) stars on the upper main sequence display peculiar lines and line strengths, in addition to other peculiar features such as a strong global stellar magnetic field \citep{Bab47}. This subclass of B to F-type stars is characterized by variable line strengths and radial velocity changes as well as photometric variability of in general the same periodicity. One can usually distinguish between He-weak/strong, HgMn, Si, SrCrEu, and Am stars. The subgroup of SrCrEu objects, for example, typically have overabundances of up to several dex for e.g., Sr, Cr, Eu, and other rare earth elements compared to the Sun. Photometric variability of the CP star $\alpha^{2}$ CVn was first reported by \citet{Guth14}. The light curves can be fitted well by a sine wave and its first harmonic with varying amplitudes for different photometric filter systems. For some CP stars, a double-wave structure of the photometric light curves is detected \citep{Mai80}. However, similar magnetic field modulus variations are rare exceptions \citep{Mathy97}. The variability of CP stars is explained in terms of the oblique rotator model \citep{Stibbs50}, according to which, the period of the observed light, spectrum, and magnetic field variations is the rotational period. Accurate knowledge of the period of variability and its evolution in time for CP stars is a fundamental step in understanding their complex behaviour, especially as far as it concerns the phase relation between the magnetic, spectral, and light variations \citep{Miku10}. In 2004, an on-line catalogue of photometric observations of magnetic CP stars (mCPod\footnote{http://astro.physics.muni.cz/mcpod/}) was initiated \citep{Miku07}. This was intended to gather all available photometric data for CP stars, into a single, freely accessible database. The archive presently contains about 150\,000 individual measurements of 151 CP stars and is being constantly updated with the latest photometric data. In this paper, we present Str{\"o}mgren $uvby$ time series for 27 CP stars within the boundaries of open clusters. In addition, we queried the Hipparcos photometric database \citep{Leeuw97} for entries of our targets and found four matches. Our observations cover a time interval of about six years with typically fifteen measurements for each object. A detailed time series analysis was performed, implementing five different methods to minimize spurious detections. Because the variability is assumed to be related to the rotation, a detailed study of the light curves is essential, not only to determine astrophysical parameters using more realistic stellar model atmospheres, but also to interpret the substantial information provided when mapping the stellar surface with Doppler imaging techniques \citep{Lehm07}. The choice of open cluster CP stars allows us to establish fundamental parameters such as the age, distance, reddening, and metallicity of numerous cluster members more accurately than usually possible for Galactic field stars. \begin{table} \begin{center} \caption[]{Log of observations with the extinction coefficients for the Str{\"o}mgren $uvby$ filters. For the majority of the observations, the standard extinction coefficients for LaSilla fit the data very well.} \begin{tabular}{llccccc} \hline \hline Teles. & Obs. & Night & $k_u$ & $k_v$ & $k_b$ & $k_y$ \\ & & [JD] & [mag] & [mag] & [mag] & [mag] \\ \hline Bochum & HMM & 2446581 & 0.570 & 0.350 & 0.230 & 0.170 \\ & & 2446582 & 0.570 & 0.350 & 0.230 & 0.170 \\ & & 2446583 & 0.570 & 0.350 & 0.230 & 0.170 \\ & & 2446584 & 0.570 & 0.350 & 0.230 & 0.170 \\ & & 2446585 & 0.570 & 0.350 & 0.230 & 0.170 \\ & & 2446587 & 0.570 & 0.350 & 0.230 & 0.170 \\ & & 2447963 & 0.570 & 0.350 & 0.230 & 0.170 \\ & & 2447964 & 0.570 & 0.350 & 0.230 & 0.170 \\ & & 2447965 & 0.570 & 0.350 & 0.230 & 0.170 \\ & & 2447966 & 0.570 & 0.350 & 0.230 & 0.170 \\ & & 2447967 & 0.570 & 0.350 & 0.230 & 0.170 \\ & & 2447968 & 0.570 & 0.350 & 0.230 & 0.170 \\ & & 2447969 & 0.570 & 0.350 & 0.230 & 0.170 \\ & & 2447970 & 0.570 & 0.350 & 0.230 & 0.170 \\ & & 2447971 & 0.570 & 0.350 & 0.230 & 0.170 \\ & & 2447972 & 0.570 & 0.350 & 0.230 & 0.170 \\ & & 2447973 & 0.570 & 0.350 & 0.230 & 0.170 \\ & & 2447974 & 0.570 & 0.350 & 0.230 & 0.170 \\ & & 2447976 & 0.570 & 0.350 & 0.230 & 0.170 \\ SAT & CT & 2448307 & 0.570 & 0.350 & 0.230 & 0.170 \\ & & 2448308 & 0.570 & 0.350 & 0.230 & 0.170 \\ & & 2448309 & 0.570 & 0.350 & 0.230 & 0.170 \\ & & 2448310 & 0.570 & 0.350 & 0.230 & 0.170 \\ & & 2448311 & 0.570 & 0.350 & 0.230 & 0.170 \\ & & 2448313 & 0.570 & 0.350 & 0.230 & 0.170 \\ & & 2448314 & 0.570 & 0.350 & 0.230 & 0.170 \\ & & 2448315 & 0.570 & 0.350 & 0.230 & 0.170 \\ & & 2448308 & 0.570 & 0.350 & 0.230 & 0.170 \\ & & 2448308 & 0.570 & 0.350 & 0.230 & 0.170 \\ SAT & HH & 2448392 & 0.570 & 0.350 & 0.230 & 0.170 \\ & & 2448393 & 0.570 & 0.350 & 0.230 & 0.170 \\ & & 2448394 & 0.570 & 0.350 & 0.230 & 0.170 \\ & & 2448398 & 0.570 & 0.350 & 0.230 & 0.170 \\ & & 2448400 & 0.570 & 0.350 & 0.230 & 0.170 \\ & & 2448405 & 0.570 & 0.350 & 0.230 & 0.170 \\ & & 2448407 & 0.592 & 0.369 & 0.244 & 0.184 \\ & & 2448408 & 0.592 & 0.369 & 0.244 & 0.184 \\ & & 2448409 & 0.592 & 0.369 & 0.244 & 0.184 \\ & & 2448410 & 0.592 & 0.369 & 0.244 & 0.184 \\ & & 2448411 & 0.735 & 0.455 & 0.320 & 0.240 \\ & & 2448413 & 0.575 & 0.355 & 0.230 & 0.170 \\ & & 2448414 & 0.575 & 0.355 & 0.230 & 0.170 \\ & & 2448415 & 0.575 & 0.355 & 0.230 & 0.170 \\ & & 2448416 & 0.575 & 0.355 & 0.230 & 0.170 \\ SAT & HMM & 2448638 & 0.720 & 0.480 & 0.350 & 0.270 \\ & & 2448639 & 0.660 & 0.420 & 0.305 & 0.240 \\ & & 2448640 & 0.630 & 0.380 & 0.280 & 0.230 \\ & & 2448641 & 0.590 & 0.360 & 0.250 & 0.200 \\ & & 2448642 & 0.600 & 0.355 & 0.250 & 0.190 \\ & & 2448643 & 0.600 & 0.360 & 0.255 & 0.205 \\ & & 2448644 & 0.590 & 0.350 & 0.240 & 0.190 \\ \hline \end{tabular} \label{obs_log} \end{center} \end{table}

We have investigated photometrically a total of 27 CP stars within the boundaries of open clusters. We established variability for fourteen CP stars with previously unknown rotation periods, and confirmed the rotation periods of eight more stars with an increase in precision for two of them. The applied open cluster parameters (Table \ref{list_clusters}) are averaged ones based on a set of literature values. Starting with the comprehensive compilation by \citet{PN06}, we searched the literature for additional (new or overseen) parameters. They were all checked for plausibility by using appropriate isochrones and available photometric data taken from the WEBDA database. Significant different results were removed and finally a strict average and standard deviations were calculated. The log\,$T_{\rm eff}$ values were calculated using the calibration of \citet{Net08} for CP stars. We made use of the GCPD catalogue\footnote{http://obswww.unige.ch/gcpd/gcpd.html} and the WEBDA database to extract all available photometric data of the programme stars. To obtain absolute magnitudes (assuming that all programme stars are indeed cluster members), we took averaged values and the mean cluster distances (see Table \ref{list_clusters}). Since the stars in NGC~2516 in particular, exhibit a strong differential reddening, individual reddening values were determined whenever possible for all objects as suggested by \citet{Net08}. Within the latter reference a bolometric correction for magnetic CP stars was also introduced, which was used to calculate the luminosity. For the remaining Am and HgMn objects, the bolometric correction by \citet{Bal94} for normal stars was applied. In Table \ref{logs}, the individual values are listed. Figure \ref{hrd} shows the location of the investigated CP stars in a log\,$L/L_{\sun}$ versus log\,$T_{\rm eff}$ diagram. The dashed line denotes the terminal-age main-sequence. The evolutionary tracks for individual masses and ages are interpolated between the solar metallicity ones listed by \citet{Schall92}, \citet{Schaer93a}, \citet{Schaer93b}, and \citet{Charb93}. Two stars, HD 89856 and HD 96729 are located below the ZAMS and deviate significantly from the apparent cluster age (Table \ref{list_clusters}). These objects are definite non-members according to our analysis and that of \citet{Land07}. We also marked the location of the other questionable cluster members as discussed in Sect. \ref{membs}. The target stars cover the typical mass range for mid B to late A type main-sequence objects from about 1.7M$_{\sun}$ to 4.5M$_{\sun}$ as other members of this group \citep{Poehn05}. The photometric periods of 0.7--4.5~days are consistent with the typical rotation velocities of CP stars \citep{North92}. There is a hint of the period decreasing with increasing age, in particular for stars with a mass below 3M$_{\sun}$, but this is not statistically significant due to poor number statistics. The same is true for a possible correlation of the period with the stellar mass and effective temperature. However, with the ongoing observations, more light will be shed on these important topics. Our observations fill an important gap in previous photometric long-time studies of CP stars. The apparent open cluster members are excellent targets for follow-up observations, based on for example polarimetry, high-resolution spectroscopy, and surface mapping techniques. Follow-up observations within our framework are already under way with the Rapid Eye Mount (REM) telescope at La Silla. This will help us to understand the apparent stellar cycles, such as that of the Sun, for this group of magnetic peculiar objects.

1012.0149

Context. Photometric variability of chemically peculiar (CP) stars of the upper main sequence is closely connected to their local stellar magnetic field and their rotational period. Long term investigations, as presented here, help us to identify possible stellar cycles (as in the Sun). Furthermore, these data provide a basis for detailed surface mapping techniques. <BR /> Aims: Photoelectric Strömgren uvby time series for 27 CP stars within the boundaries of open clusters are presented. In addition, Hipparcos photometric data (from 1989 to 1993) are used for our analysis. Our observations cover a time period of about six years (1986 to 1992) with typically fifteen measurements for each objects. These observations help us to determine the rotational periods of these objects. <BR /> Methods: A standard reduction procedure was applied to the data. When possible, we merged our data sets with already published ones to obtain a more significant result. A detailed time series analysis was performed, involving five different methods to minimize spurious detections. <BR /> Results: We established, for the first time, variability for fourteen CP stars. For additional two stars, a merging of already published data sets, resulted in more precise periods, whereas for six objects, the published periods could be confirmed. Last, but not least, no significant variations were found for five stars. Apart from six stars, all targets seem to be members of their host open clusters. <BR /> Conclusions: The present observations fill an important gap in previous photometric long-time studies of CP stars. The presented open cluster members are excellent targets for follow-up observations, employing for example polarimetric, high-resolution spectroscopic, and surface mapping techniques. <P />All photometric measurements are available in electronic form at the CDS via anonymous ftp to cdsarc.u-strasbg.fr (130.79.128.5) or via <A href="http://cdsarc.u-strasbg.fr/viz-bin/qcat?J/A+A/525/A16">http://cdsarc.u-strasbg.fr/viz-bin/qcat?J/A+A/525/A16</A>

12168085

2011ApJ...728..147C

1012

1012.5136_arXiv.txt

The magnetic field plays an important role in physical processes occurring at all relevant coronal heights. From the photosphere to the inner corona, the strength of the magnetic field can be measured by the conventional Zeeman splitting technique. However, in the outer corona beyond, say, 1.2-1.5 $R_\odot$, the field gets too weak to be measured directly. Indirect methods available at present are mostly based on numerical extrapolations or various types of radio techniques. With the extrapolation method, the coronal magnetic field distributions are resolved numerically by extrapolating the measured photospheric magnetic field, making use of potential field (Schatten et al. 1969; Schrijver {\&} Derosa 2003), linear or nonlinear force free field assumptions (e.g., Yan {\&} Sakurai 2000; Wiegelmann, 2008; He {\&} Wang 2008), or solving the full-set of magnetohydrodynamic (MHD) equations (e.g., Linker et al. 1999). There also exist various types of techniques employing radio emissions to derive the coronal magnetic field strength. The first one utilizes the well-known Faraday-rotation effect acting on a linearly polarized radio signal passing through coronal structures. Signals from both extragalactic radio sources (Sakurai {\&} Spangler 1994; Mancuso {\&} Spangler 1999, 2000; Spangler 2005; Ingleby et al. 2007) and spacecraft radio emitters (P\"{a}tzold et al. 1987) have been analyzed over the past decades. To use this method one needs to determine independently the coronal electron density distribution and the geometry of the magnetic field. The method applies to the heliocentric distance range of $3 - 10$ $R_\odot$. The second one is based on the band-splitting phenomenon observed during Type-II radio bursts related to shocks driven by coronal mass ejections (CMEs) (Smerd et al. 1974, 1975; Vrsnak, et al. 2002; Cho et al. 2007). The phenomenon is interpreted as plasma emissions from downstream and upstream of the shock front at different frequencies. To implement this method one needs to apply the MHD shock theory, and presume the value of the plasma $\beta$, the shock geometry, as well as the coronal electron density distribution. The method works for the heliocentric distance range of $1.5 - 3$ $R_\odot$. The third method makes use of observations of circularly polarized thermal radio emissions (Sastry 2009), which was recently explored to estimate the magnetic field strength in coronal streamers at heliocentric distances of 1.5 and 1.7 $R_\odot$ by Ramesh et al. (2010). In this paper, we present a novel seismological method to evaluate the strength of the magnetic field in the outer corona. Coronal seismology is a way to diagnose the physical parameters of the corona with observational and MHD theoretical analyses of waves and oscillations. Here we present a seismological study to derive the Alfv\'en speed and magnetic field strength with the use of the so-called streamer waves, reported recently by Chen et al. (2010) (referred to as paper I hereinafter). The waves were observed by the Large Angle and Spectrometric Coronagraph (LASCO), and generated as the aftermath of the CME-streamer interaction event dated on 6 July 2004. The wave properties like the wavelength, period, and phase speed, as well as the possibility of deriving the magnetic field strength with this wave, are already presented in paper I. Here we briefly discuss the theoretical basis and procedures of implementing the concerned seismological study. The waves are regarded as the fast kink body mode carried by the plasma sheet structure of a streamer. The phase speed of the mode have two contributions. One is the phase speed of the mode in the plasma rest frame, the other is the speed of the solar wind streaming along the plasma sheet. The former can be well approximated by the exterior Alfv\'en speed according to a linear wave dispersion analysis, which will be carried out in the Appendix of this paper. As a result, the radial profile of the Alfv\'en speed and that of the magnetic field strength can be deduced given the speed and density of the solar wind plasmas. In paper I, by taking these solar wind parameters from a simplified two-dimensional MHD model accounting for streamers, current-plasma sheets, and slow-fast winds developed by Chen {\&} Hu (2001), the magnetic field strengths at 5 and 10 $R_\odot$ are evaluated. It should be pointed out that the seismological study in paper I is rather preliminary and incomplete. The purpose of this paper is to further improve and complete the study. It will be improved from the following three aspects. First, in paper I the Alfv\'en speed in the region surrounding the plasma sheet is set to be equal to the phase speed of the kink mode in the plasma rest frame. In this paper, we will conduct a parameter study on dispersion relation of the fast kink body mode with a simplified slab model of magnetized plasmas given by Edwin {\&} Roberts (1982), making use of different sets of coronal parameters prescribed according to available observational diagnoses and theoretical modelings. From the study, we deduce the appropriate connection of the kink mode phase speed to the exterior Alfv\'en speed. Second, we employ available observational results to put constraints on the flow speed and number density of the solar wind along the plasma sheet. The former will be constrained by measurements of plasma blobs, which are structures released intermittently through streamer cusps and flowing outwards passively in the wind along the plasma sheet. The feasibility of using blobs to yield the wind speed has been discussed in several papers (Sheeley et al. 1997; Wang et al. 1998, 2000; Song et al. 2009; and Chen et al. 2009), and will not be repeated here. The number density will be derived by inversion of the pB (polarized brightness) measurements given by LASCO at the time prior to the CME-streamer interaction event. Third, in paper I only the magnetic field strengths at two distances were estimated, while in this paper two sets of radial profiles of both the Alfv\'en speed and magnetic field strength in the heliocentric range of 3 to 10 $R_\odot$, corresponding to two subsequent wave phases, will be presented. This will provide the information on not only spatial but also temporal variations of the two critical coronal parameters. In the following section, we show major relevant results on streamer waves of paper I, and present the radial profiles of the Alfv\'en speed. The strength of the coronal magnetic field will be deduced in the third section. Conclusions and discussion are given in the fourth section of this paper, and the associated wave dispersion analysis is presented in the Appendix.

In this paper we provide a novel method to diagnose the Alfv\'en speed and magnetic field strength in the corona based on the observations of streamer waves, which propagate with phase speeds consisting of two components. One is the speed of the wave mode in the plasma rest frame, the other is the speed of the mean solar wind. The method applies to heliocentric ranges from 3 to about 10 $R_\odot$ in the region surrounding the plasma sheet. To implement the diagnosis, we first establish the connection between the Alfv\'en speed and the observed phase speed, then we put constraints on the solar wind velocities with blob measurements, and determine the density distributions through the inversion of the LASCO pB data. The obtained profiles of the magnetic field strength are in line with other estimates in the corona. Previous studies indicate that the magnetic flux tube experiences a dramatic expansion in the neighborhood of the streamer cusp, and a possible contraction beyond (e.g., Wang 1994; Bravo {\&} Stewart 1997; Chen {\&} Hu 2001, 2002; Hu et al. 2003; Li et al. 2006). Till now, there exist no direct observational proofs of the above peculiar feature of the magnetic field near the streamer cusp. From the seismological study of this paper, we see that the magnetic field along the plasma sheet expands more or less radially starting from as near as 3 $R_\odot$. Therefore, the mentioned intriguing expanding process of the flux tube occurs, if it does, below this distance. This provides observational constraints on relevant models. According to our seismological studies on the basis of speed measurements of the two wave phases, we find that the Alfv\'en speed and magnetic field strength at a fixed distance decrease with time. This has been suggested to be a result of the recovering process of the CME-disturbed corona. In the process, the magnetic field initially stretched out by the CME ejecta may get relaxed through processes like magnetic reconnections, and the evacuated corona may get refilled gradually through plasma heating and resultant expansions. As mentioned in the above section, the occurrence of this dynamic recovering process contributes directly to the uncertainty of our results, since we have adopted identical and steady solar wind parameters for the diagnoses associated with the two phases. Due to the lack of direct measurements on these parameters in the near-Sun region, it is currently not possible to evaluate the impact of using time-dependent solar wind parameters on our results. Apart from the undeterminancy associated with solar wind densities and velocities, there exist two other major factors contributing to our diagnostic uncertainties. One is the error coming from the phase speed measurements, which was already discussed and estimated to be about $\pm10\%$ (or $\pm50$ km s$^{-1}$) in paper I. This error will be passed directly to the evaluation of the Alfv\'en speed and the magnetic field strength. Another factor stems from the approximate relationship between the kink mode and the external Alfv\'en speed. In the paper the relationship was determined with the dispersion relation given by Edward {\&} Roberts (1982) for a simplified magnetized plasma slab configuration under Cartesian coordinates. This geometry is different from the realistic spherical, inhomogeneous (both in the radial and latitudinal directions), and time-dependent streamer-plasma-sheet configuration. Therefore, future studies should investigate the properties of the kink mode under more realistic geometry. In addition, theoretical and numerical efforts should continue to explore excitation conditions and propagation properties of streamer waves in the process of CME-streamer interaction, and determine the connection between the wave properties and the background plasma properties. These works will be of great benefit to future seismological studies with streamer waves.

1012.5136

We present a novel method to evaluate the Alfvén speed and the magnetic field strength along the streamer plasma sheet in the outer corona. The method is based on recent observations of streamer waves, which are regarded as the fast kink body mode carried by the plasma sheet structure and generated upon the impact of a fast coronal mass ejection (CME) on a nearby streamer. The mode propagates outward with a phase speed consisting of two components. One is the phase speed of the mode in the plasma rest frame and the other is the speed of the solar wind streaming along the plasma sheet. The former can be well represented by the Alfvén speed outside the plasma sheet, according to a linear wave dispersion analysis with a simplified slab model of magnetized plasmas. The radial profiles of the Alfvén speed can be deduced with constraints put on the speed of the solar wind, which is done by making use of the measurements of streamer blobs flowing passively in the wind. The radial profiles of the strength of the coronal magnetic field can be depicted once the electron density distribution is specified, this is done by inverting the observed polarized brightness data. Comparing the diagnostic results corresponding to the first wave trough and the following crest, we find that both the Alfvén speed and magnetic field strength at a fixed distance decline with time. This is suggestive of the recovering process of the CME-disturbed corona.

12209169

2011JHEP...03..019C

1012

1012.5300_arXiv.txt

\label{sect:intro} The existence of the dark matter has been firmly established and it constitutes about 23\% of the total energy density in the universe. The nature of the dark matter is one of the most outstanding questions in cosmology and particle physics. Many different types of experiments are deployed to detect the dark matter and to measure its properties, including direct detections from the recoils of the nuclei hit by the dark matter particle, indirect detections of the cosmic rays from dark matter annihilations or decays, and collider searches by direct production of dark matter particles. Recently, there has been an interesting observation of anomalous $e^+/e^-$ excess in the energy range of 1 -- 100 GeV measured by the PAMELA collaboration \cite{Adriani:2008zr}, which may be interpreted as indirect dark matter signals, coming from dark matter annihilations or decays inside the galactic halo. In addition, the $e^+ +e^-$ spectrum measured by the Fermi-LAT experiment between 20 GeV and 1 TeV is harder than that inferred from previous experiments~\cite{Abdo:2009zk}, which may also be attributed to the contribution from the dark matter. In this paper, we consider a new dark matter candidate which could naturally produce the excess of the electron/positron flux observed in these experiments. The cosmic positrons are one of the prominent signals for indirect dark matter detections. For the most popular dark matter candidate, a weakly interacting massive particle (WIMP), electrons and positrons can be produced from annihilations of the WIMPs in the galactic halo. However, to account for the PAMELA excess, a large boost factor at the order of 100 or larger is required to increase the annihilation rate~\cite{Bergstrom:2008gr,Barger:2008su,Cholis:2008hb,Cirelli:2008pk}. In addition, large flux of gamma rays will be produced in the dark matter annihilations, which is severely constrained by the observed gamma ray spectrum~\cite{Bertone:2008xr,Nardi:2008ix,Essig:2009jx,Meade:2009iu}. As a result, explaining the PAMELA excess by annihilating dark matter has a hard time to satisfy the constraints from the annihilation cross section and cosmic gamma ray data. Another possibility to generate the observed electron/positron spectrum is that if the dark matter particle is not absolutely stable, but decays with a very long lifetime. A small fraction of the dark matter particles has decayed, producing electrons and positrons in the decay products~\cite{Chen:2008dh}. The decaying dark matter has an easier time to satisfy the gamma ray constraints, but to explain the PAMELA excess the lifetime needs to be of the order of $10^{26}$--$10^{27}$ seconds~\cite{Nardi:2008ix,Essig:2009jx,Meade:2009iu}, which seems to be an additional arbitrary parameter coming from nowhere. For the decaying dark matter, it is usually assumed that the symmetry that protects the stability of the dark matter particle is not exact, but violated by some highly suppressed interactions. It has been argued that the required lifetime can be obtained from a TeV scale particle decaying through dimension-6 operators suppressed by the grand unification scale $m_{\rm GUT} \sim 2\times 10^{16}$ GeV \cite{Nardi:2008ix, Arvanitaki:2008hq}. In this paper we consider another possibility that an exact symmetry is carried by two sequestered sectors, which interact indirectly only through the visible sector (standard model). The lightest particle charged under this symmetry is absolutely stable. However, the dark matter is made of the next to the lightest particle charged under the same symmetry, which is only approximately stable due to sequestering. The dark matter particle decays to the truly stable particle with a long lifetime because of the highly suppressed interactions between the two sequestered sectors. The standard model (SM) particles produced in the decays can be observed, and could be responsible for the anomalies in the cosmic ray experiments. We show that such decaying dark matter can arise naturally in the goldstini scenario proposed recently~\cite{Cheung:2010mc}. In this scenario, supersymmetry (SUSY) is spontaneously broken in more than one sequestered sectors. There is a goldstino associated with the spontaneously broken SUSY in each sector. The SUSY in different sectors are connected by supergravity and only one combination of the goldstini is eaten and becomes the longitudinal mode of the gravitino. The other combinations of the goldstini acquire a mass of twice the gravitino mass at the lowest order due to the supergravity effect. Assuming $R$-parity is exactly conserved, and if the gravitino and an uneaten goldstino are the lightest and the next to the lightest supersymmetric particles (LSP and NLSP), respectively, the cosmic electrons and positrons can be produced from decays of the goldstino dark matter to the gravitino. If the two SUSY breaking sectors only interact indirectly through the visible supersymmetric standard model (SSM) sector, the interactions responsible for the goldstino decays are highly suppressed and the required lifetime for the observed electron/positron excess can be naturally obtained. A distinct feature of this scenario is that the dark matter decays dominantly through three-body processes, producing a pair of SM particles and another invisible massive particle. Most of the studies of decaying dark matter before assumed that the dark matter particle decays through a two-body process to a pair of SM particles or a pair of portals to four SM particles without additional missing particles. Some exceptions are in Ref.~\cite{Ibarra:2009dr} where the three-body decays including a neutrino, as well as from internal bremsstrahlung, are considered. The constraint on the anti-proton flux, which shows no excess in the PAMELA experiment \cite{Adriani:2008zq}, requires that the decays of dark matter particles dominantly produce leptons. In the case of two-body decays, the muon and tau final states are preferred \cite{Meade:2009iu} and the direct decay to the electron and positron pair would give a sharp edge on the energy spectrum at half the mass of the dark matter particle, which is not seen by Fermi-LAT. On the other hand, the electrons and positrons coming from the three-body decays will have a softer and smooth spectrum which may still be consistent with other observations. As will be shown, the goldstino couplings to the SM particles, unlike the universal coupling of the gravitino, are governed by the fractions of the soft-SUSY breaking masses coming from different SUSY breaking sectors for the corresponding superpartners. It is easy for the goldstino to have preferential decays to leptons if different superpartners receive different soft masses from different SUSY-breaking sectors. This paper is organized as follows. In Sec.~\ref{sect:effgold} we derived the goldstini interactions with SM fermions using the method of constrained superfields developed recently by Komargodski and Seiberg~\cite{Komargodski:2009rz}. From the interactions we can calculate the decay rate of a goldstino to the gravitino and a pair of SM fermions. The interactions with other SM fields are collected in the Appendix. In Sec.~\ref{sect:decayingDM} we discuss the model of the decaying goldstino dark matter and the parameters which can give rise to the PAMELA signal and satisfy other astrophysical and cosmological constraints. In Sec.~\ref{sect:pheno} we perform fits to the electron/positron energy spectra observed by PAMELA and Fermi-LAT experiments with the decaying goldstino scenario, and identify decay modes and parameters which can be consistent with the observation data. We then briefly discuss the collider phenomenology of this scenario. Conclusions are drawn in Sec.~\ref{sect:conclusions}. Throughout this paper, we use ``goldstini'' when we refer to the goldstino fields coming from different SUSY-breaking sectors, and ``goldstinos'' to represent the plural form of the same-species goldstino.

\label{sect:conclusions} In this work we proposed a new scenario for supersymmetric decaying dark matter in theories with goldstini, where the uneaten goldstino dominantly decays into gravitino, which shows up as missing energy, and two SM particles. In this scenario it is not necessary to introduce $R$-parity violations, since the goldstino decays through dimension-8 operators and naturally has a long lifetime suitable to explain the positron excess observed by the PAMELA collaboration. We derive low-energy effective interactions of the goldstini and show that the couplings can be non-universal, while the gravitino coupling remains universal as expected. The non-universality of the goldstini coupling is crucial for the dark matter to be leptophilic, so as to avoid the lack of excess in the anti-proton spectrum measured by PAMELA. To obtain the correct goldstino relic density for the dark matter, however, seems to require some fine tuning of the reheating temperature in the early universe. A distinct feature of this scenario is the three-body decay of the dark matter, which results in softer energy spectra for the electrons and positrons, as opposed to a sharp edge in the case of the more conventional two-body decay. Consequently, it is possible to fit both the positron excess in the PAMELA data and the hardening feature in the $e^++e^-$ flux measured by the Fermi-LAT. We find decays into $e^++e^-$ with 100\% branching fraction, which is disfavored if the dark matter decays into two or four SM particles, could still provide reasonably good fits to PAMELA and Fermi-LAT. In addition, universal coupling of the dark matter with all three lepton flavors, which may be favored from other considerations, could also fit the data well. In this work we have assumed the hadronic decay modes of the dark matter, as well as prompt decays into photons, are suppressed in order to satisfy constraints from anti-proton and gamma ray measurements. However, it is worth pointing out that most studies on these constraints are based on the assumption that the dark matter decays into two-body final states, while the decay proceeds through three-body channel with a missing particle in our scenario. It would be interesting to re-evaluate these constraints in a more model-independent fashion for the case of three-body decays with missing particles.

1012.5300

We consider a new scenario for supersymmetric decaying dark matter without R-parity violation in theories with goldstini, which arise if supersymmetry is broken independently by multiple sequestered sectors. The uneaten goldstino naturally has a long lifetime and decays into three-body final states including the gravitino, which escapes detection, and two visible particles. The goldstini low-energy effective interactions are derived, which can be non-universal and allow the dark matter to be leptophilic, in contrast to the case of a single sector supersymmetry breaking. In addition, the three-body decay with a missing particle gives a softer spectrum. Consequently, it is possible to fit both the e <SUP>+</SUP>/ e <SUP>-</SUP> excess observed by the PAMELA and the e <SUP>+</SUP> + e <SUP>-</SUP> measurements by the Fermi-LAT using universal couplings to all three lepton flavors or 100% branching fraction into electrons/positrons, both of which are disfavored in the conventional scenario of dark matter decays into two or four visible particles without missing energy.

12225158

2011PhRvD..83j3510B

1012

1012.4656_arXiv.txt

Simple inflationary models can be broadly but conveniently classified into three main categories, see e.g.~\cite{FieldRev,StringRev} for reviews. One category comprises hybrid inflation models~\cite{hybrid}, where the inflaton is responsible for the slow-roll dynamics while another scalar field is responsible for the end of inflation. Interesting aspects of these models include the possible production of topological defects at the end of inflation and possible realisations in supergravity as F-term or D-term inflation~\cite{FieldRev} and in string theory in the framework of brane inflation~\cite{StringRev}. Another category comprising chaotic inflation~\cite{chaotic} includes the large field inflation models. They have the simplest field theoretic origin as a single massive scalar field is enough to generate inflation. Another relevant feature of these models is the large, super-Planckian value of the inflaton and therefore the possibility to generate a significant amount of primordial gravity waves during inflation. Unfortunately, such large values of the inflaton may render these models more difficult to realize in supergravity or string theory. In this respect, the category of small field or hilltop inflation models~\cite{hilltop}, which includes for instance the original new inflation model~\cite{new}, may be more natural. In these models, inflation occurs near a maximum or an inflection point of the potential, where the curvature is negative. This makes the slow-roll conditions somewhat easier to achieve and leads to cosmological perturbations with a negative spectral index as favoured by observations. Another interesting aspect of these models is that inflation can occur at very low energy scales, while still generating an acceptable spectrum of primordial density perturbations~\cite{lowscale}. These models offer also string theoretic realisations, for instance in the form of racetrack models~\cite{StringRev} where the imaginary part of a K\"ahler modulus acts as the inflaton. Another consequence of small field inflation follows from the Lyth bound on the inflaton excursion~\cite{lythbound}, which implies that gravitational waves generated during inflation are highly suppressed in this setting. At the end of inflation, the inflaton condensate must decay and reheat the universe. It has become clear during the last twenty years or so~\cite{traschen} that, in most models, reheating starts with an explosive and non-perturbative production of large, non-thermal fluctuations of the inflaton and other bosonic fields coupled to it, in the process of preheating~\cite{KLS}. The subsequent dynamics are characterized by a highly non-linear and turbulent-like evolution, before the system eventually settles into thermal equilibrium. Preheating may have many interesting consequences in cosmology, like the production of stochastic backgrounds of gravitational waves~\cite{GWcha, GWhyb, GWpre, GWflat, GWvec}, primordial black holes~\cite{malik, bassett, kudoh1, kudoh2}, non-gaussian curvature perturbations~\cite{rajantie, bond} or primordial magnetic fields~\cite{magnetic}. Preheating after chaotic and hybrid inflation has been abundantly studied in the literature, both analytically and with numerical lattice simulations~\cite{KTlatt, latticeeasy, defrost}. In models of chaotic inflation, preheating occurs typically via broad parametric resonance~\cite{KLS}, where fields coupled to the inflaton are produced non-adiabatically as the inflaton oscillates around the minimum of its potential, or via tachyonic resonance~\cite{tachres} in the presence of trilinear or non-renormalizable bosonic interactions for the inflaton. In models of hybrid inflation, field fluctuations are amplified by a tachyonic (spinodal) instability~\cite{tachyonic} as the fields roll towards the true minimum of the potential at the end of inflation. By contrast, preheating after small field inflation remains less explored. On the analytical side, the problem is complicated by the fact that closed-form solutions for the evolution of the homogeneous background during preheating are in general not available, while standard approximations like WKB for the evolution of the perturbations are not applicable in this context. On the numerical side, lattice simulations of preheating in these models must cover scales that range from the Hubble rate at the end of inflation up to the inflaton mass at the minimum of the potential, which typically differ by several orders of magnitude. Preheating after new inflation was studied in \cite{PreNew}\footnote{See \cite{traschen} for earlier work.}, where it was described as a combination of both tachyonic amplification and non-adiabatic resonance. In this sense, it is somehow intermediate between preheating after chaotic inflation and preheating after hybrid inflation. Some aspects of preheating in this model where also studied numerically in \cite{kudoh2} and in the context of K\"ahler moduli / Roulette inflation in \cite{PreRoul}. The combination of tachyonic and non-adiabatic effects occurs also in preheating via trilinear or non-renormalizable interactions after chaotic inflation~\cite{tachres}. However, we will see that preheating after small field inflation occurs in a qualitatively different way. The main purpose of this paper is to develop the analytical understanding of preheating in the class of small field inflation models. This is a first necessary step before a further study of the dynamics with lattice simulations and of the cosmological consequences of preheating in these models. To illustrate some of the qualitative features of preheating after small field inflation, consider the inflaton potential shown in Fig.~\ref{smallpot}. Slow-roll inflation ends when the second slow-roll parameter $|\eta|$ becomes of order one, where the curvature of the potential is negative and of the order of the Hubble rate squared. The inflaton then rolls down the potential, passes through an inflection point where the curvature vanishes and oscillates around the minimum with an amplitude controlled by Hubble friction. Before the inflaton reaches the inflection point, its fluctuations have a negative effective mass squared and modes with momentum $k^2 < -V''(\phi)$ are amplified by a tachyonic instability. This process may affect a very wide range of scales, from the Hubble rate at the end of inflation up to the maximum of $|V''(\phi)|$, but it is much more efficient for the low-momentum modes. The tachyonic amplification stops when the inflaton crosses the inflexion point of the potential, oscillates around the minimum and goes back towards the tachyonic region. During this interval of time, the mass of the inflaton varies very non-adiabatically with time. This can lead to a further growth of the perturbations, but we will see that this effect is typically negligible. However, the evolution of the modes during this time interval when the inflaton oscillates around the minimum of the potential has dramatic consequences for the fate of the fluctuations when the inflaton goes back into the region where the curvature is negative. Indeed, during this second tachyonic episode, the amplitude of the fluctuations starts to decrease exponentially with time. For the modes with sufficiently low momenta, this exponential decrease occurs during the amount of time the inflaton condensate climbs back along its potential. If the inflaton climbed back up to the point from where it started, this effect would exactly compensate the growth of these modes during the first tachyonic episode. By contrast, the modes with higher momentum, which were amplified at a slower rate during the first tachyonic episode, continue to be amplified during the second one. The whole process may then repeat itself several times as the inflaton oscillates around the minimum of the potential. The net effect then follows from a competition between low-momentum modes which grow a lot but also decrease a lot and modes with higher momenta which grow less but also decrease less. As we will see, the analytical description of this process is very similar to the problem of tunneling through a volcano-shaped potential (two potential barriers separated by a deep potential well) in quantum mechanics. Alternatively, the initial tachyonic growth may be so efficient that preheating ends during the first tachyonic episode, depending on the parameters. \begin{figure}[htb] \begin{center} \includegraphics[width=13cm]{smallpot.eps} \end{center} \vspace*{-5mm} \caption{Schematic form of the inflaton potential in a model of small field inflation.} \label{smallpot} \vspace*{-3mm} \end{figure} As we will see, another specificity of small field models is that the spectrum of the inflaton fluctuations amplified at the beginning of preheating is peaked around the Hubble rate, because the modes with $k/a \lesssim H$ are amplified from the beginning of preheating when $V'' \sim -H^2$ while the modes with higher momenta are amplified only later and at a slower rate. The situation is different for preheating after chaotic inflation, where the fluctuations amplified by parametric resonance are typically peaked at scales $k / (aH) > q^{1/4} \gg 1$ with the initial resonance parameter $q \gg 1$ for preheating to be efficient~\cite{KLS}. Similarly, for preheating after hybrid inflation, the typical scale of the fluctuations amplified by the tachyonic instability is set by the curvature of the potential in the direction of the symmetry-breaking field and must be sub-Hubble to satisfy the so-called waterfall condition~\cite{hybrid}. In fact, in models of small field inflation, the modes amplified by the initial tachyonic instability leave the Hubble radius, because preheating starts when the the second slow-roll parameter $|\eta|$ becomes of order unity while the first slow-roll parameter $\epsilon$ is still much smaller than one, so that the universe is still inflating at the beginning of preheating. When preheating ends in less than one oscillation of the inflaton condensate, large density perturbations of the Hubble size may lead to an abundant production of primordial black holes, see also~\cite{kudoh2}, which may put strong constraints on models of small field inflation. This is in contrast with preheating after chaotic and hybrid inflation, where the field inhomogeneities are peaked at scales that are much shorter than the Hubble radius~\cite{kudoh1,defrost}. Furthermore, large field fluctuations at the Hubble scale and the fact that small field inflation can occur at very low energy scales may lead to the production of gravitational waves from preheating with frequencies today that are sufficiently small to fall into the range accessible by high-sensitivity ground-based and even space-based interferometric experiments~\cite{GWpre}. A detailed study of these cosmological consequences of preheating after small field inflation will appear elsewhere~\cite{paper2}. The rest of the paper is organized as follows. In section \ref{SecBack}, we specify the models of small field inflation that we will consider and we describe the inflationary dynamics and the subsequent evolution of the inflaton condensate. Section \ref{Sec1} is dedicated to the first stage of preheating, when the inflaton condensate first rolls towards the minimum of the potential. In Section \ref{Sec2}, we study the second stage of preheating, when the inflaton condensate oscillates around the minimum. Some of the results used in that Section are derived in the Appendix. In Section~\ref{SecPert}, we first consider other possible non-perturbative decay channels for the condensate. We then discuss the perturbative decay of the inflaton after preheating and the resulting reheat temperature. Finally, Section \ref{SecConclu} contains a summary of our results and directions for future work.

\label{SecConclu} We have studied the linear stage of preheating in the class of small field inflation models, where the curvature of the inflaton potential is negative during inflation. Although this is one of the most common classes of inflationary models, preheating after small field inflation remained much less studied than preheating after chaotic and hybrid inflation. On the analytical side, the problem is complicated by the fact that closed form solutions for the background evolution of the inflaton condensate are in general not available, while standard approximation like WKB for the evolution of perturbations are not accurate. Nevertheless, we saw that a detailed analytical description of preheating in this class of models is possible. The analytical methods that we developed in this paper may also be applied to the study of preheating in more complicated models. We showed that preheating after small field inflation is usually dominated by the tachyonic amplification of inflaton fluctuations in the intervals of time when the inflaton condensate rolls in the region where the curvature of its potential is negative. A peculiar feature of this process is that the inflaton fluctuations experience a succession of exponential growths and decreases, so we called it "tachyonic oscillation". The exponential decreases of the fluctuations in these intervals of time arise because the coefficient of the decreasing mode is much larger than the coefficient of the growing mode, by an amount that depends on the scale of the fluctuation. Despite this temporary exponential decrease of the fluctuations, the full process is very efficient and we showed that preheating completes typically after less than five oscillations of the inflaton condensate. The range of scales of the fluctuations amplified in the course of preheating is very wide, extending from the Hubble scale to the curvature of the potential at the minimum, with a peak given by Eq.~(\ref{peakp2}) when preheating completes in more than one oscillation of the condensate. When the condition (\ref{cond1v}) is satisfied, the first tachyonic instability is so efficient that preheating completes in less than one oscillation of the condensate. In that case, the spectrum of the inflaton fluctuations at the end of preheating is peaked around the Hubble scale, or even slightly outside the Hubble radius because of the residual amount of inflation during preheating. Density fluctuations at the Hubble scale may then lead to an abundant production of primordial black holes, see also~\cite{kudoh2}, which can put constraints on small field inflation models, see e.g.~\cite{PBH} for a recent update on observational constraints on primodrial black holes. The large field fluctuations amplified by preheating lead also to the production of gravitational waves (GW), with a peak frequency today~\cite{GWpre} \be \label{fstarGW} f_* \approx \left(\frac{k_* / a}{H_p}\right)\,\left(\frac{\rho_p^{1/4}}{10^{11}\,\mathrm{GeV}}\right)\,10^3\,\mathrm{Hz} \ee where $H_p$ and $\rho_p$ are the Hubble rate and the energy density during preheating, and $k_*/a$ is the charactersitic momentum amplified by preheating~\footnote{Eq.~(\ref{fstarGW}) is valid when the universe becomes quickly radiation-dominated after preheating. Note also that, in models where gauge fields play an important role in the dynamics of preheating, extra peaks may appear at well-distinct frequencies in the final GW spectrum~\cite{GWvec}.}. GW from preheating after chaotic and hybrid inflation tend to have a frequency today above $10^3$ Hz, which is too high to be observable by high-sensitivity interferometric experiments, either because $\rho_p^{1/4} \gg 10^{11}$ GeV or because $k_* / a \gg H_p$~\cite{GWpre}. GW from the non-perturbative decay of condensates different from the inflaton, in particular super-symmetric flat directions, are more promising in this respect although the GW amplitude may be suppressed in that cases~\cite{GWflat}. On the other hand, GW from preheating after small field inflation may fall naturally in the frequency range accessible by ground-based and even space-based interferometers when $k_* / a \sim H_p$ and the energy scale of inflation is small enough, which is common in this class of models. We also saw that preheating after small field inflation may be followed by a long matter-dominated stage before the universe thermalizes, which would further redshift and dilute these GW. A detailed study of the production of primordial black holes and GW from preheating after small field inflation is currently under way~\cite{paper2}. Finally, although an analytical understanding of the linear stage of preheating is often very useful to study its cosmological consequences, lattice simulations are eventually necessary to further study the non-linear dynamics. This is a difficult task in general for preheating after small field inflation, because of the wide range of scales that appear in the problem. In such cases, lattice simulations can be performed only for a limited range of parameters and the analytical results that we derived in this paper may be used to extrapolate the results to other regions of the parameter space.

1012.4656

Whereas preheating after chaotic and hybrid inflation models has been abundantly studied in the literature, preheating in small field inflation models, where the curvature of the inflaton potential is negative during inflation, remains less explored. In these models, a tachyonic instability at the end of inflation leads to a succession of exponentially large increases and decreases of the inflaton fluctuations as the inflaton condensate oscillates around the minimum of its potential. The net effect is a competition between low-momentum modes which grow and decrease significantly, and modes with higher momenta which grow less but also decrease less. We develop an analytical description of this process, which is analogous to the quantum mechanical problem of tunneling through a volcano-shaped potential. Depending on the parameters, preheating may be so efficient that it completes in less than one oscillation of the inflaton condensate. Preheating after small field inflation may also be followed by a long matter-dominated stage before the Universe thermalizes, depending on the energy scale of inflation and the details of the inflaton interactions. Finally, another feature of these models is that the spectrum of the inflaton fluctuations at the end of preheating may be peaked around the Hubble scale. In fact, because preheating starts when the second slow-roll parameter |η| becomes of order unity while the first slow-roll parameter ɛ is still much smaller than 1, the Universe is still inflating during preheating and the modes amplified by the initial tachyonic instability leave the Hubble radius. This may lead to an abundant production of primordial black holes and gravitational waves with frequencies today which are naturally small enough to fall into the range accessible by high-sensitivity interferometric experiments.

12164010

2011AIPC.1331...41B

1012

1012.1559_arXiv.txt

The first dusty disc around a main sequence star was observed in 1984 around Vega \citep{vega1984}. Since then our knowledge of such systems has improved significantly, and it is now known that 32\% of A stars exhibit excess emission in the infrared, over and above the stellar photosphere \citep{su06}. This is thermal emission from dust particles orbiting the star in a debris disc. Debris discs are collisionally dominated in that the smallest bodies in the system are continuously replenished by collisions between larger objects and are subsequently removed by radiation pressure. The long term evolution of such systems can be modelled by the steady state collisional models of \cite{wyatt07} and is expected to be a slow decline in brightness as the disk mass is depleted due to collisional erosion. A decrease in brightness with age is indeed observed \citep[e.g.][]{su06} and can be well fitted with the models of \cite{wyatt07}, allowing such models to characterize the population of main sequence A stars debris discs reasonably accurately. Dust is also seen around some post-main sequence stars. In some cases this dust can be a result of the evolution of the star, for example material emitted in the stellar wind form spherical shells of dust that are observed around AGB stars (e.g. \citealt{AGBshell2010}) or even stable discs observed around post-AGB stars, possibly linked to binarity (e.g.~\citealt{winkel09postAGBbinary}). Infrared excess observed around giant stars, e.g. \cite{jura99}, and the helix nebula \citep{helix}, on the other hand, has been interpreted as a disc similar to debris discs on the main sequence (although alternative interpretations exist, see e.g. \citealt{kimzuckerman01}). Hot dust is also observed in small radii ($<$0.01AU) discs around white dwarfs, e.g. \cite{farihi09} or \cite{2007Kuchner}, again inferred to originate from a debris disc. However, in contrast to main sequence debris discs, these discs cannot be in steady state since material at such small radii has a short lifetime. Rather models suggest that these discs are formed when an asteroid approaches close to the star where it is tidally disrupted \citep{jurasmallasteroid}. There are not yet enough observations of discs around post-main sequence stars to fully understand the population and it is not clear how the few discs that have been discovered around post-main sequence stars relate to the progenitor population of debris discs on the main sequence. In this contribution we discuss the evolution of debris discs beyond the main sequence, with particular reference to models presented in \cite{bonsor10}. These models intend to be a theoretical framework in which all of the effects of stellar evolution on a disc are investigated. We use these models to evolve the population of debris discs observed around main sequence A stars \citep{rieke05} \citep{su06} and discuss the detectability of such discs in terms of the observations of dust around evolved stars. \begin{figure} \includegraphics[width=0.5\textwidth] {_rsst05_sufig9_24um.ps} \includegraphics[width=0.5\textwidth] {_rsst05_sufig9_70um.ps} \caption{The fit to the observations of \citet{su06} at 24 (left) and 70$\mu$m (right), comparable to Fig.2 of \citet{wyatt07}. The plots show the fraction of stars with flux ratios in different age bins ($<$30Myr, 30-190Myr, 190-400Myr), at 24$\mu$m $\frac{F_{disc}}{F_*}=$ 1-1.25 (diamond: small excess), 1.25-2 (square: medium excess), $>$2 (triangle: large excess) and similarly at 70$\mu$m $\frac{F_{disc}}{F_*}$= 1-5 (diamond:small excess), 5-20 (square: medium excess), $>$20 (triangle: large excess). Observed values are shown with $\sqrt{N}$ error bars, whilst model values are joined with dotted, dash and solid lines, for small, medium and large excess.} \label{fig:fit} \end{figure}

\label{sec:conclusions} We have presented a piece of that work \citep{bonsor10} that provides a theoretical framework in which all of the effects of stellar evolution on debris discs are considered. Our models find that it is significantly harder to detect debris discs around evolved stars. The fraction of discs with detectable excess decreases significantly on the giant branch, yet further on the horizontal branch and discs around white dwarfs are very hard to detect. Approximately 10\% of giant stars should have detectable discs, whilst the highest chances of detecting a debris disc around a white dwarf are for very young, nearby white dwarfs. This fits with the absence of such detections apart from the helix nebula \cite{helix} and a handful of other systems also presented in this volume (see entries by Chu and Bilikova). According to our models the helix nebula, although not anticipated to have a detectable disc is at the optimum age and distance for such a detection, should for some reason our models underpredict the disc flux. Evolved main sequence debris discs cannot directly explain the near-IR observations of hot dusty discs around white dwarfs e.g. \cite{farihi09}. The minimum disc radius in our population is 10AU, whilst these sytems have radii less than R$_{\odot}$. However the fact that we predict a population of cold, undetectable planetesimal belts around white dwarfs has great implications for such discs. These belts could potentially provide the reservoir of material that replenishes such systems. \begin{theacknowledgments} I would like to acknowledge the support of an STFC funded PhD (Bonsor) and various fruitful discussions with colleagues. \end{theacknowledgments}

1012.1559

The population of debris discs on the main sequence is well constrained, however very little is known about debris discs around evolved stars. In this work we provide a theoretical framework that considers the effects of stellar evolution on debris discs; firstly considering the evolution of an individual disc from the main sequence through to the white dwarf phase, then extending this to the known population of debris discs around main sequence A stars. It is found that discs around evolved stars are harder to detect than on the main sequence. In the context of our models discs should be detectable with Herschel or Alma on the giant branch, subject to the uncertain effect of sublimation on the discs. The best chances are for hot young white dwarfs, fitting nicely with the observations e.g the helix nebula (Su et al.. 2007) and 9 systems presented by Chu &amp; Bilikova (this volume). Although our baseline models do not predict such a high rate of detectable discs.

12163279

2011A&A...526A..64R

1012

1012.0769_arXiv.txt

4U 1538$-$52, discovered by the \emph{Uhuru} satellite (\cite{giacconi74}), is an X-ray pulsar with a B-type supergiant companion, QV Nor. It has an orbital period of $\sim$3.728 days (\cite{DWP77}; \cite{clark00}), with eclipses lasting $\sim$0.6 days (\cite{becker}). This X-ray persistent system produces this radiation when the neutron star captures matter from the wind of the B supergiant star. Assuming a distance to the source of 5.5 kpc (\cite{becker}) and isotropic emission, the estimated X-ray luminosity is $\sim (2-7)\times 10^{36}$ erg s$^{-1}$ in the 3$-$100 keV range (\cite{rodesPhD}). Therefore, the size of the ionization zone may be a relatively small region in the stellar wind (\cite{HM77}; \cite{vLoon01}). Fluorescence iron emission lines from X-ray pulsars are produced by illumination of neutral or partially ionized material by X-ray photons with energies above the line excitation energy. Possible sites of fluorescence emission may be: (i) accretion disk (mostly seen in low mass X-ray binary sources); (ii) stellar wind (in high mass X-ray binary pulsars); (iii) material in the form of a circumstellar shell; (iv) accretion column; (v) material in the line of sight; or some combination of these locations. In this sense, fluorescence lines in the X-ray spectrum are an interesting tool for studying the surrounding wind regions and elemental abundance in X-ray sources. In 4U 1538$-$52 the emission line at $\sim$6.4 keV can usually be described within the uncertainties either by a single narrow Gaussian line or by a multiplet of narrow Gaussian lines (\cite{WSH83}). Observations carried out with \emph{Tenma} detected it at 6.3$\pm$0.2 keV and EW 50$\pm$30 eV (\cite{makishima87} 1987), while the Rossi X-ray Timing Explorer (\emph{RXTE}) saw it at 6.25$\pm$0.06 keV and EW 61 eV (\cite{coburn2}). Other X-ray observatories such as \emph{BeppoSAX} (\cite{robba}) and \emph{RXTE} have used a Gaussian line at $\sim$6.4 keV for describing the fluorescence of iron in a low-ionization region (\cite{mukherjee06} 2006; \cite{rodesPhD}). The variability of this line was studied by \cite{jjrrXrU05}. In this paper, we present a spectral analysis based on the observation of 4U 1538$-$52 performed with data from the \emph{XMM-Newton} satellite. The observation covers the orbital phase interval 0.75--1.00 and we detect the presence of the K$_\alpha$ iron line at $\sim$6.4 keV and some blended emission lines below 3 keV. In Sect.~\ref{data} we describe the observation and data analysis. We present in Sect.~\ref{timing} timing analysis; in Sect.~\ref{analyse} spectral analysis, and in Sect.~\ref{conclusion} we summarize our results.

\label{conclusion} We presented the spectral analysis of the HMXB 4U 1538$-$52 using an \emph{XMM-Newton} observation. The X-ray continuum is well fitted by three absorbed power-laws with a photon index $\sim$1.13, describing the hard, scattered and soft excess, respectively. The inferred unabsorbed flux is $\sim 2.1 \times 10^{-10}$ erg s$^{-1}$ cm$^{-2}$ in the 0.3--11.5 keV energy band, corresponding to a luminosity of $\sim 7.5 \times 10^{35}$ erg s$^{-1}$, assuming an isotropic emission and a distance to the source of 5.5 kpc (\cite{becker}). The flux found by \emph{RXTE} in the 3--11.5 keV energy band and an orbital phase of 0.85 was 1.3$\times$10$^{-10}$ erg s$^{-1}$ cm$^{-2}$. Using the spectrum obtained by \emph{XMM-Newton} over the same range of phases, the flux we obtained in this work was two times lower, 6.3$\times$10$^{-11}$ erg s$^{-1}$ cm$^{-2}$. The soft excess is present in the spectrum and can be modelled with different absorbed components. We simply modelled the soft emission with an absorbed power law component, although Fig.~\ref{PNMOS} still showed spectral residuals at lowest energies. \textbf{Our results show that a blackbody component could also be the physical origin of this soft excess, taking the associated errors into account}.The soft excess in other HMXBs has been explained by a blend of Gaussian emission lines only (\cite{boroson03}). We also tried to fit the soft emission using Gaussian profiles, but we did not obtain a significant improvement of the fit because of the low level of counts below 0.6 keV. We detected an iron K$\alpha$ line at $\sim$6.41 keV, with an EW of $\sim$50 eV. The \emph{BeppoSAX} observation of this system obtained an EW of 57 eV in the same orbital phase range 0.75--1.00 (\cite{robba}) and showed an increase in the post-egress phase to 85 eV. The phase-averaged spectrum obtained by \emph{RXTE} reported an EW of 62 eV (\cite{coburn2}). In addition this iron line is detected in all orbital phases (\cite{robba}; \cite{rodesPhD}), therefore the K$\alpha$ iron line is not only produced by fluorescence from less ionized iron near the neutron star's surface but also a fraction of the observed line flux must originate from more extended regions. \textbf{We have also detected a number of emission lines which we interpret as recombination lines from highly ionized species. Since these lines are detected in eclipse, they must be produced in an extended halo. Likewise, we have found an absorption feature at 2.1 keV. Whether it is produced by physical properties of the source or it is due to calibration effects, is still an open issue. } We compared the phase-average spectrum to the eclipse spectrum. In the phase-average spectrum, we found no evidence of any other iron line apart from that at 6.4 keV, and the absorption edge at $\sim$7.1 keV was well described by the X-ray absorption model. The 6.4--7.2 keV energy band showed a complex structure, but we did not find a proper model to describe it or detect other iron features significantly. We also detected discrete recombination lines in both \emph{EPIC/PN} and \emph{EPIC/MOS} spectra. The emission lines reported in Sect.~\ref{fluorescence} can be identified with He- and H-like ions. Although we could identify some of them, the energy resolution of the \emph{EPIC} instruments is such that the lines could be blended. However, in the eclipse spectrum, we detected the fluorescence iron line at 6.4 keV and one more iron emission line at 6.6 keV. The discrete recombination lines detected are listed in Table~\ref{emlieclparam}. The presence of these lines in an eclipse spectrum implies that the formation region extends beyond the size of the B supergiant. Moreover, the ionization state was estimated to range from $10^{2.1}$ to $10^{3.2}$ erg cm s$^{-1}$, due to the simultaneous detection of elements with both low and high ionization levels. \textbf{This broad range of $\xi$ also suggests either that the emitting material is present over a wide range of distances from the compact object or has a large range of densities. } \textbf{The pulse phase-resolved spectroscopy showed significant variability of the photon index and the unabsorbed flux, but no clear correlation or anticorrelation between them. Significant variations with the pulse phase were also observed in the different column absorption values, but again did not show a clear relationship to the other parameters. } \textbf{Future observations with high spectral resolution instruments will be needed to unambiguously resolve possible blended lines found in this study allowing the full use of their diagnostic capability.}

1012.0769

Context. The properties of the X-ray emission lines are a fundamental tool for studying the nature of the matter surrounding the neutron star and the phenomena that produce these lines. <BR /> Aims: The aim of this work is to analyse the X-ray spectrum of 4U 1538-52 obtained by the XMM-Newton observatory and to look for the presence of diagnostic lines in the energy range 0.3-11.5 keV. <BR /> Methods: We used a 54 ks PN &amp; MOS/XMM-Newton observation of the high-mass X-ray binary 4U 1538-52 covering the orbital phase between 0.75 to 1.00 (the eclipse ingress). We modelled the 0.3-11.5 keV continuum emission with three absorbed power laws and looked for the emission lines. <BR /> Results: We found previously unreported recombination lines in this system at ~2.4 keV, ~1.9 keV, and ~1.3 keV, which is consistent with the presence of highly ionized states of S XV Heα, Si XIII Heα, and either Mg Kα or Mg XI Heα. On the other hand, in spectra that are both out of eclipse and in eclipse, we detect a fluorescence iron emission line at 6.4 keV, which is resolved into two components: a narrow (σ ≤ 10 eV) fluorescence Fe Kα line plus one hot line from highly photoionized Fe XXV. <BR /> Conclusions: The detection of new recombination lines during eclipse ingress in 4U 1538-52 indicates that there is an extended ionized region surrounding the neutron star.

12167926

2011ApJ...729L..28P

1012

1012.0243_arXiv.txt

Supernova remnants (SNRs) have long been considered to be the primary source of Galactic cosmic-rays (CRs) below the {\it knee} of the cosmic-ray spectrum, $\sim$ 10$^{15}$ eV. TeV $\gamma$-ray observations of SNRs such as RX J1713.7-3946 and RX J0852.0-4622 provide evidence for the acceleration of ions \citep{aharonian07a,aharonian07b}. However the TeV emission can also be attributed to inverse-Compton scattering by the same electron population that produces the X-ray synchrotron emission. Viewed in X-rays, the young ($\sim$ 330 yr; \citealt{fesen06}) Galactic SNR Cassiopeia A (Cas A) consists of a shell whose emission is dominated by emission lines from O, Si, S, and Fe \citep[e.g.,][]{vink96,hughes00,will02,will03,hwang03,laming03}. Exterior to this shell are faint X-ray filaments which mark the location of the forward shock. The emission found here is nonthermal X-ray synchrotron emission from shock accelerated electrons \citep{allen97,gotthelf01,vink03}. These forward shock filaments are observed to expand with a velocity of $\simeq$ 5000 km s$^{-1}$ \citep{delaney03,patnaude09}, assuming a SNR distance of 3.4 kpc \citep{reed95}. Nonthermal emission filaments are also observed in the interior of the SNR and are believed to be either forward shock filaments seen in projection \citep{delaney04,patnaude09} or associated with efficient acceleration of electrons at the SNR reverse shock \citep{uchiyama08,helder08}. Fluctuations in both exterior and interior nonthermal filaments have also been reported \citep{patnaude07,uchiyama08,patnaude09}, and the variability is cited as evidence for rapid synchrotron cooling of TeV electrons in mG--scale fields. A two to four year timescale for variations is evidence for efficient diffusive shock acceleration in SNR shocks, or alternatively the variations are seen as evidence for magnetic field fluctuations due to plasma waves behind the shock \citep{bykov08}. Emission from Cas A has been seen at energies up to $\sim$ 40 keV with the {\it Suzaku} HXD PIN detector \citep{maeda09}, up to 100 keV with {\it CGRO} OSSE and {\it Integral} IBIS \citep{the96,renaud06}, and GeV emission has been detected using {\it Fermi}--LAT \citep{abdo10}. The {\it Fermi} observations do not rule out either a leptonic origin to the GeV emission from a combination of nonthermal bremsstrahlung and inverse--Compton emission or a hadronic origin from neutral pion decay. Finally, Cas A has been detected at even higher TeV energies with HEGRA, MAGIC and Veritas \citep{aharonian01,albert07,humensky08}. Interestingly, the centroids for the GeV--TeV emission are located in the western region of Cas A, where the nonthermal X-ray emission is brightest \citep{helder08,maeda09}. Here we present {\it Chandra} ACIS-S3 observations of Cas A taken in 2009 and 2010 which, when compared to ACIS-S3 observations taken between 2000 and 2007, show the remnant's nonthermal X-ray emission in the 4.2--6.0 keV band to have decreased at a rate of $\simeq$ 1.5--2.0\% yr$^{-1}$. In \S~2 we discuss our observations, data reduction, and spectral analysis and in \S~3 we discuss our results and offer some conclusions about the current and future evolution of the nonthermal emission in Cas A.

Our analysis of the 4.2--6 keV flux of Cas A shows a decline of 1.5\% yr$^{-1}$ in the nonthermal X-ray emission across the entire SNR over 11 years, with a slightly larger decline rate of 1.9\% yr$^{-1}$ from regions along the remnant's western limb. We find that qualitatively, the observed spectral steepening and decline in flux can be explained by a simple model for changes in the electron cutoff energy which are brought about by a natural deceleration of the shock. We estimate an average deceleration of Cas A's forward shock velocity $\approx$ 10 -- 40 km s$^{-1}$ yr$^{-1}$. The predicted decline in the nonthermal X-ray emission is about 4\% yr$^{-1}$, which is nearly twice that observed. The difference between the predicted and observed decline might be explained by the fact that the 4.2--6.0 keV continuum emission is not entirely due to synchrotron emission from shock accelerated electrons, but some of it is from thermal continuum which does not evolve on the same timescale as the nonthermal emission. We have compared our results to models where the decline is a natural consequence of either a decrease in the number of particles entering the shock or a decrease in the efficiency of the shock to amplify the magnetic field, and find that these models predict a decline of $\sim$ 0.1--0.5\% yr$^{-1}$, which is significantly less than the observed decline of 1.5--1.9\% yr$^{-1}$.

1012.0243

We present new Chandra ACIS-S3 observations of Cassiopeia A which, when combined with earlier ACIS-S3 observations, show evidence for a steady ~1.5%-2% yr<SUP>-1</SUP> decline in the 4.2-6.0 keV X-ray emission between the years 2000 and 2010. The computed flux from exposure corrected images over the entire remnant showed a 17% decline over the entire remnant and a slightly larger (21%) decline from regions along the remnant's western limb. Spectral fits of the 4.2-6.0 keV emission across the entire remnant, forward shock filaments, and interior filaments indicate that the remnant's nonthermal spectral power-law index has steepened by about 10%, with interior filaments having steeper power-law indices. Since TeV electrons, which give rise to the observed X-ray synchrotron emission, are associated with the exponential cutoff portion of the electron distribution function, we have related our results to a change in the cutoff energy and conclude that the observed decline and steepening of the nonthermal X-ray emission is consistent with a deceleration of the remnant's sime5000 km s<SUP>-1</SUP> forward shock of ≈30-70 km s<SUP>-1</SUP> yr<SUP>-1</SUP>.

2034709

2011A&A...527A..59C

1012

1012.0419_arXiv.txt

Several features of colour-magnitude diagrams (CMDs) and luminosity functions (LFs) of Galactic Globular Clusters (GCs) can be employed to test the accuracy of low-mass, metal-poor stellar models \citep[see, e.g.][]{rfs88}. The bump appearing in the GC Red Giant Branch (RGB) LF is one of these important benchmarks. It is produced by the encounter of the H-burning shell with the H-abundance discontinuity left over by the outer convection at its maximum depth \citep{tho, ibe} reached during the first dredge-up. The sharp increase of the H-abundance causes a sudden decrease of the mean molecular weight ($\mu$), that affects the efficiency of the H-burning shell, proportional to a high power of $\mu$ \citep[see][]{kipwei,salcas}. This occurrence causes a temporary drop of the surface luminosity, before it starts to increase again. As a consequence, a low-mass RGB star crosses the same luminosity interval three times, and a bump (over-density) appears in the RGB differential LF (star counts per magnitude bin) of GCs \citep[for a detailed discussion we refer to][]{salaris02}. Given that the RGB-bump brightness depends on the maximum depth attained by the convective envelope, and the chemical profile above the advancing H-burning shell, the comparison between predicted and observed luminosity of the RGB-bump, provides valuable information about the internal structure of low-mass stars at the beginning of their RGB evolution. Since its first detection in the LF of NGC104 \citep[47Tuc --][]{kdd} the RGB bump has been the subject of several theoretical and observational investigations \citep{fusipecci90, cassisi97, alves99, zoccali99, bono01, riello03, bjork06, dicecco10}. Thanks to these works, we have now accurate measurements of its brightness in many GCs as well as in Local Group dwarf galaxies \citep[see][and references therein]{monelli10}. The parameter routinely adopted to compare observations with theory is the quantity $\Delta V_{\rm HB}^{\rm Bump}= V_{Bump}-V_{HB}$, that is, the V-magnitude (or bandpasses similar to Johnson V) difference between the RGB-bump and the horizontal branch (HB) at the RR Lyrae instability strip level \citep{fusipecci90, cassisi97}. This has the advantage of being formally independent of distance and reddening, and not affected by any uncertainty in the zero point of the photometry. The most recent comparisons between \vhbb models and observations \citep[see, e.g., Fig.~10 in][]{dicecco10} seem to confirm a discrepancy (at the level of $\sim$0.20~mag or possibly more) for GCs with total metallicity [M/H] below $\sim -$1.5, in the sense that the predicted RGB-bump luminosity is too high. The quantitative estimate of the discrepancy depends on the adopted metallicity scale. At the upper end of the GC metallicity range, the existence of a discrepancy depends on the adopted metallicity scale. One drawback of using \vhbb as diagnostic of the RGB-bump luminosity, is that uncertainties in the placement of the observed HB level for GCs with blue HB morphologies, and in theoretical predictions of the HB luminosity \citep[i.e., due to uncertainties in the calculations of the He-core mass at the He-flash, see e.g.][]{cas10}, hamper the interpretation of discrepancies between theory and observations. An alternative avenue explored in this paper is offered by measuring the magnitude difference between the Main Sequence (MS) Turn-Off (TO) and the RGB-bump brightness \vtob$= V_{TO}-V_{bump}$, that bypasses the HB. Observationally, an accurate estimate of the TO brightness requires both very high quality photometric datasets, and a detailed analysis of the uncertainty associated with the presence of binary stars. To the best of our knowledge, so far only \citet{caputo02} and \citet{meissner06} have studied the \vtob parameter. \citet{caputo02} used \vtob in combination with $\Delta V_{\rm HB}^{\rm TO}= V_{TO}-V_{HB}$ to investigate the metallicity scale of a large sample of galactic GCs, but did not attempt to assess the level of agreement between predicted and observed \vtob values. More recently, \citet{meissner06} used the \vtob together with other CMD age indicators, to check their mutual self-consistency. As a result, they found that the GC ages estimated from \vtob were younger by about 2~Gyr, in comparison with estimates based on the $\Delta V_{\rm HB}^{\rm TO}$ parameter. This occurrence was interpreted as an evidence that stellar models predict a too bright RGB-bump, by $\sim$0.2-0.3~mag. We wish to reanalyze this issue employing new, accurate photometry of a large sample of GCs, that enabled us to determine both TO and RGB-bump magnitudes for 12 GCs, covering a large metallicity range. Our methodological approach is the following. We have first determined the apparent magnitudes of both TO and RGB-bump in our GC sample, and employed the cluster relative distances from a theoretical MS-fitting technique. As a second step, we have estimated individual cluster ages from the TO absolute magnitudes, obtained assuming the empirical MS-fitting distance to NGC6752 by \citet{grat03} as zero point of our relative distance scale. Another set of ages for each cluster is then determined from their observed \vtob, and compared with the TO ages. The outcome of this comparison constrains the level of agreement between predicted and observed RGB-bump luminosity, independently of the HB. The plan of this paper is as follows: section~\ref{frame} presents briefly the observational dataset and the theoretical models adopted in our analysis; estimates and comparisons of TO and \vtob ages are described in section~\ref{comp}, followed by a final discussion.

The main result of our analysis is summarized by Fig.~\ref{difteo}, discussed in the previous section. The values of \vacstob predicted by theoretical models for cluster ages estimated from the TO absolute magnitudes, are larger than observed. Given that the observed TO magnitude is by definition matched by the theoretical isochrones to determine the TO age, this discrepancy implies that the absolute magnitude of the RGB-bump in the models is too bright. An extension of this type of analysis to a larger, homogeneous sample of GC photometries is obviously desirable; hovewer our results based on a limited sample of clusters provide already clear evidence of a real \lq{over-luminosity}\rq\ of the predicted absolute magnitude of the RGB-bump, irrespective of problems with HB modelling and placement of the reference HB level in clusters with only blue HB stars. The simplest explanation for this discrepancy envisages a systematic underestimate of the cluster metallicities by $\sim$0.2~dex. An higher [M/H] would eliminate the discrepancy, because it causes a lower TO age and a lower theoretical RGB-bump brightness for each cluster. There is of course no indication that the metallicity scale we adopted is affected by this type of systematics, but this is a point to be considered. In the following we expand our discussion to see how improvements in the micro- (e.g. opacities, nuclear reaction rates) and macro-physics (e.g., element transport meachanisms) employed in stellar evolution calculations, and the recently extablished presence of multiple stellar populations with varying chemical patterns in individual GCs, can account for this discrepancy. \subsection{Improved micro- and macro-physics} A straightforward explanation for the discrepancy highlighted in Fig.~\ref{difteo} could be an underestimate of the radiative opacity at temperatures around a few $10^6$~K -- typical temperatures at the lower boundary of the convective envelope. Higher opacities would shift the convection boundary -- hence the H-abundance discontinuity -- to deeper layers, causing a fainter RGB-bump. However, this solution does not seem plausible, for the following reasons: i) radiative opacities in this temperature range should not be affected by an uncertainty larger than $\sim5$\% \citep[see, e.g.][]{guzik08} and this small change is not able to reconcile theory with observations; ii) the discrepancy theory-observations increases with decreasing [M/H], and it does not seem very likely that radiative opacities become less accurate when the metal content decreases. The isochrones employed to determine both the cluster ages from the TO brightness, and the theoretical values of \vacstob, do not account for the effect of atomic diffusion (including radiative levitation). Although current spectroscopic observations of globular cluster stars show that atomic diffusion is at least partially inhibited by additional turbulence/mixing \citep[see, i.e.][]{korn07} -- induced for example by rotation \citep[see, i.e.,][]{egge10} -- we summarize here the effect on TO ages and \vacstob values in case of full efficiency. According to the results by \cite{vand02} and \cite{mich10} -- that expand upon previous studies by \cite{cds97, cassisi98} where the effect of radiative levitation was not considered -- atomic diffusion makes the RGB-bump magnitude brighter by 0.03-0.06~mag at fixed age, and also decreases the cluster TO ages by at most $\sim$1.5~Gyr for the lowest metallicity clusters. The combined effect on \vacstob would decrease the discrepancy for the most metal poor cluster in our sample by $\sim$0.05~mag at most. The effect becomes less significant with increasing metallicity. On the other hand, the recent redetermination of the $^{14}N(p,\gamma)^{15}O$ reaction rate -- not included in our adopted models -- would increase the cluster ages by $\sim$ 1~Gyr, and at the same time make the RGB-bump brighter by $\sim$0.06~mag at fixed age \citep{weiss05, pietrinferni10}. The net result would be an increase of the discrepancy by $\sim 0.10$~mag or more, that would move the mean value of $\Delta$(\vacstob) up to $\sim$0.30~mag. Overall, the combined effect of the new $^{14}N(p,\gamma)^{15}O$ reaction rate and inclusion of atomic diffusion (plus radiative levitation) would exacerbate the discrepancy between theory and observations, that would become on average of the order of 0.25~mag. Another possibility to mitigate the discrepancy is to include overshooting beyond the formal boundary of the convective envelope \citep[see, e.g.,][]{alongi91}. Calculations by \cite{cassisi02} show that the inclusion of convective overshooting decreases the RGB-bump brightness by ${\rm \sim 0.8 mag/H_P}$ (where $H_P$ denotes the local pressure scale height); the discrepancy between theory and observations would disappear with the inclusion of convective overshooting of the order of $\sim 0.25$ below the Schwarzschild boundary of the convective envelope. Besides overshooting from the convective boundary, \cite{cassisi02} have investigated also the effect on the RGB-bump shape and brightness, of a smoother chemical discontinuity left over by the first dredge-up. A smoother chemical discontinuity could be produced, for example, by turbulent mixing counteracting the efficiency of atomic diffusion. \cite{cassisi02} results show that the bump luminosity decreases by $\sim$0.25 mag/${\rm H_p}$, where the smoothing length is expressed in units of the local pressure scale height. Given that smoothing the discontinuity alters also the shape of the RGB luminosity function in the bump region, this hypothesis is potentially testable. As estimated by \cite{cassisi02}, a sample of more than 120 RGB stars within $\pm$0.2~mag of the peak of the RGB-bump, and random photometric errors smaller than 0.03~mag can potentially disclose this effect in the RGB luminosity function. \subsection{The role of GC multipopulations} A very important issue to be considered, is the effect on the cluster RGB-bump luminosity and TO ages of subpopulations with varying degrees of the CNONa anticorrelation and the -- likely -- associated increased He abundance, as observed in individual GCs \citep[see, e.g.,][for a review]{gcs04}. If the sum of the CNO abundance stays constant among all stars in a given cluster -- as observed, within the measurement errors -- the RGB-bump magnitude is affected only by the possible increase of helium. As shown by, e.g., \cite{cassisi97} and \cite{salaris06}, increasing the initial He abundance increases the bump brightness at fixed age and [Fe/H]. In a 'real' cluster the size of this effect depends on the exact amount of He-enhancement and the fraction of stars involved, but the main point is that this can only exacerbate the discrepancy displayed in Fig.~\ref{difteo}. As for the ages from the TO luminosity, one has to notice that within the individual clusters analyzed in this paper, there are no clear signs of large spreads of the initial He abundance, in terms of a split of the MS in the CMD. A reasonable upper limit to the He spread of 0.05 in mass fraction, would decrease the TO age by not more than $\sim$0.5~Gyr \citep[see, i.e.,][]{salaris06}. As a conclusion, the effect of subpopulations with enhanced He within individual clusters in our sample would not solve the discrepancy highlighted by Fig.~\ref{difteo}. Only NGC~1851 shows a clear split of the subgiant branch in our adopted CMD, whose origin is still debated \citep[see, e.g.,][]{cassisi08, carretta10b}. The TO measurement has been obtained considering only the most populated SGB, that should harbour stars with a 'standard' He and metal distribution \citep{cassisi08} so that also in this case our TO age estimates should be reliable.

1012.0419

We present new measurements of the magnitude of the main sequence turn off and the red giant branch bump in the luminosity function of a sample of Galactic globular clusters with updated estimates of [Fe/H] and [α/Fe], employing photometric data collected with the Advanced Camera for Survey onboard the HST. We compare measured and predicted values of the magnitude difference between these two features, a rarely employed diagnostic of the internal structure of low-mass stars at the beginning of their red giant evolution. Our analysis discloses a clear discrepancy between theory and observations, because the theoretical red giant branch bump magnitudes are too bright by on average ~0.2 mag. This corroborates results from the more widely studied magnitude difference between horizontal branch and red giant bump, avoiding the well known problems associated with determining the horizontal branch level from colour-magnitude diagrams and with uncertainties in the luminosity of horizontal branch stellar models. We briefly discuss several potential solutions of this discrepancy.

2302276

2011MNRAS.413.2314B

1012

1012.0305_arXiv.txt

\label{introduction} Although rare today, ultraluminous infrared galaxies (ULIRGs) -- galaxies with infrared (IR) luminosities exceeding $10^{12}$\,L$_{\odot}$ -- were extremely common in the early Universe, signposting systems undergoing intense, dust-obscured star formation. Moreover, they contribute a significant fraction of the submillimetre (submm) background \citep{fixsen98}. This important high-redshift population was first discovered in the form of bright submm sources behind massive, lensing clusters \citep*{smail97}, and in blank fields \citep[e.g.][]{hughes98,barger98,eales99}, using the Submm Common User Bolometer Array \citep[SCUBA;][]{holland99} on the 15-m James Clerk Maxwell Telescope (JCMT); a number of surveys with a variety of instruments have now brought the number of known submm galaxies (SMGs) to several hundred \citep[e.g.][]{coppin06, bertoldi07, greve08, scott08}. Cross-identifying the submm sources with emission at other wavelengths is made difficult by the poor spatial resolution of even the largest submm telescopes. For example, the combination of JCMT and SCUBA resulted in a resolution of 14\,arcsecond (arcsec; {\sc fwhm}) at 850\,$\umu$m. The best way to overcome this would be with mm/submm interferometric observations -- capable of locating the submm emission directly, with arcsec accuracy \citep[e.g.][]{downes99, gear00, iono06, wang07, younger07, ivison08, cowie09}. Such observations, however, require a large investment of observing time with the few existing facilities that are capable, although the advent of the Atacama Large Millimeter/Submillimeter Array (ALMA) will make this strategy much easier in the future. In the meantime, attaining higher resolution is possible using radio interferometric and IR observations, where the empirical correlations between the far-IR and radio wavebands \citep{condon92} or the bolometric IR/mid-IR \citep{elbaz02} make it much easier to identify the submm emitter, particularly given the low source densities in the radio \citep{ivison98, ivison00, ivison02, smail00, dannerbauer04}. This work has typically relied on data from the Very Large Array (VLA) at 1.4\,GHz and {\it Spitzer} using the 24-$\umu$m channel of the MIPS instrument \citep{werner04, rieke04}. In addition, high-redshift SMGs can be identified through their IR colours as measured by {\it Spitzer's} IRAC camera \citep[e.g.][]{pope06}. Here we present radio, mid-IR (24-$\umu$m) and IRAC counterparts to the 126 SMGs that have been detected in the Large APEX Bolometer Camera (LABOCA) Extended Chandra Deep Field South [ECDFS] Submm Survey (LESS), a deep blank-field 870-$\umu$m survey, down to a 3.7-$\sigma$ limit of 4.4\,mJy \citep{weiss09}. The ECDFS is an exceptional area for multi-wavelength, wide-field studies of galaxy evolution due to deep X-ray \citep{giacconi01,lehmer05,luo08}, optical \citep{giavalisco04,beckwith06}, IR (Dickinson et al., in preparation) and radio \citep{miller08,ivison10} data. The CDFS portion of the field has also been surveyed \citep{scott10} with the AzTEC 1.1-mm bolometric camera \citep{wilson08} on the Atacama Submillimeter Telescope Experiment. The paper is organised as follows: in Section~\ref{catalogues} we describe the submm, radio, 24-$\umu$m and IRAC data that have been used to identify counterparts to the submm sources, with particular emphasis on the techniques used to extract source fluxes and positions from the radio map. Section~\ref{strategy} contains details of our counterpart identification strategy and in Sections~\ref{results} and \ref{irac} we present lists of the likely counterparts and their properties. Section~\ref{discussion} discusses these results in detail, ascertaining the effectiveness of our strategy. We also derive the redshift distribution of the radio-detected robust counterparts using the radio-submm spectral index relation of \citet{carilli99,carilli00} before drawing our conclusions in Section~\ref{conclusions}. In an appendix we present detailed notes on some of the sources as well as multi-wavelength maps with the counterparts marked. We assume a flat $\Lambda$CDM cosmology of $\Omega_{\Lambda} = 0.73$, $\Omega_{m} = 0.27$ and $H_0 = 70.5$~km\,s$^{-1}$\,Mpc$^{-1}$ \citep{hinshaw09}.

\label{conclusions} Using a probabilistic approach, we have attempted to identify reliable counterparts to the 126 SMGs recently discovered at a wavelength of 870\,$\umu$m in the LESS survey of the ECDFS using the LABOCA camera on the APEX telescope \citep{weiss09}. Taking values of the corrected Poissonian probability (the so-called $p$-statistic, $p$) that are less than or equal to 0.05 to indicate a secure identification, i.e.\ a highly unlikely chance coincidence, we have found reliable radio and/or 24-$\umu$m counterparts to 62 SMGs. A further 17 SMGs were identified using IRAC sources that fell within a colour-flux cut that was constructed from the results of the radio and MIPS analysis. In contrast to most previous work of a similar nature, we have based our identifications on rigorously constructed catalogues of 1.4-GHz and MIPS/IRAC sources. In total we find that 79 out of the 126 SMGs have secure counterparts, an identification fraction of 63~per~cent. This is not as high as some other studies, partly due to the relatively shallow radio map and somewhat larger submm beam. In several cases it is obvious that multiple submm emitters are blended and consequently difficult to identify. Finally, in creating our radio catalogue we have performed simulations in order to correct the flux densities for `flux boosting'. This has particular relevance to the calculation of source redshifts based on the radio-submm spectral index, a technique which often uses de-boosted submm fluxes, but ignores the corresponding effect in the radio. With the systematic shift towards lower redshifts removed, the median redshift of the radio-detected SMGs in our sample is $\overline{z}=2.2^{+0.8}_{-0.7}$ (1-$\sigma$ errors). This is identical to that found by both \citet{chapman05} and \citet{wardlow10}, the latter using the sample of SMGs identified in this paper, but using a different technique (optical to mid-IR multi-band photometry) for measuring the source redshifts. The median redshift of the full sample is likely to be rather higher as the unidentified SMGs by definition have weak radio emission. The current generation of submm cameras produce maps with such poor resolution that a probabilistic approach to identifying submm galaxies is inevitable. Ideally, identification work such as that presented in this paper would be done with telescopes offering sub-arcsecond resolution, such as the IRAM Plateau de Bure Interferometer and the Submillimeter Array (SMA). However, due to their limited sensitivity (small numbers of antennas and relatively poor atmospheric transmission), many hours are required for a reliable detection of a typical submm galaxy. In the future, the Atacama Large Millimeter/Submillimeter Array (ALMA) will revolutionise the study of high-redshift star formation with its order of magnitude increase in sensitivity and imaging fidelity which will make pinpointing the origin of the submm emission in surveys such as LESS a relatively trivial exercise, requiring only minutes per source to achieve a high dynamic range image.

1012.0305

We present radio and infrared (3.6-24 μm) counterparts to submillimetre galaxies (SMGs) detected in the Extended Chandra Deep Field-South with the Large APEX Bolometer Camera (LABOCA) 870-μm bolometer camera on the 12-m Atacama Pathfinder Experiment. Using the Very Large Array at 1.4 GHz and Spitzer, we have identified secure counterparts to 79 of the 126 SMGs [signal-to-noise ratio (S/N) &gt; 3.7, S<SUB>870</SUB> &gt; 4.4 mJy] in the field, 62 via their radio and/or 24-μm emission, the remainder using a colour-flux cut on Infrared Array Camera 3.6- and 5.8-μm sources chosen to maximize the number of secure, coincident radio and 24-μm counterparts. In constructing our radio catalogue, we have corrected for the effects of 'flux boosting', then used the corrected flux densities to estimate the redshifts of the SMGs based on the radio/submm spectral indices. The effect of the boosting correction is to increase the median redshift by 0.2 resulting in a value of ? (1σ errors) for the secure radio counterparts, in agreement with other studies, both spectroscopic and photometric.

12167709

2011ApJ...728..104S

1012

1012.3394_arXiv.txt

\hspace{1.0em} The Sun is located roughly at the center of a large interstellar feature in the Local spiral arm (the Orion Spur). This conspicuous structure, known as Local Cavity (LC) or Local Bubble, consists of a very irregular, low density volume ($n_{HI} < 0.005$ cm$^{-3}$) of the interstellar medium (ISM), whose borders extend from $\approx 65$ to $250$pc, depending on the direction to which we observe it \citep{coxreynolds1987,welsh1994,lallement2003,vergely2010}. This cavity is surrounded by several other interstellar bubbles that are sometimes associated to strong star-forming activity, and are frequently believed to have been generated by supernovae (SN) explosions and very intense stellar winds from massive OB stars \citep{quigley1965,berkhuijsen1971,weaver1979}. Several efforts have been made to build a tridimensional view of the LC, mainly by using Na I observations in the direction of nearby stars, which is a suitable tracer of the neutral gas \citep{welsh1994,sfeir1999,lallement2003,vergely2010}. Information on the shape and size of this structure can also be inferred from ultraviolet interstellar absorption lines \citep{frisch1981,frisch1983,paresce1984,centurion1991,welsh1994,redfield2000,sallmen2008,welsh2005}, interstellar reddening \citep[][Reis et al. 2010, in preparation]{franco1990,corradi1997,corradi2004,frisch2007,reis_corradi2008,vergely2010} and interstellar polarization \citep{tinbergen1982,reiz1998,heiles1998,heiles2000,leroy1999}. The overall shape of the LC suggests that it is being compressed due to the expansion of the neighboring bubbles, with a narrower dimension along the Galactic Plane (GP), and probably opened in the direction of the Galactic halo. Furthermore, this ``chimney" structure is slightly tilted relative to the GP ($\approx 20^{\circ}$), and perpendicular to the Gould's Belt, a large complex of young massive OB stars surrounding the Local ISM \citep{welsh1994,sfeir1999,lallement2003,vergely2010}. In the direction of the Galactic Center, we find a particularly interesting structure known as the Loop I superbubble, which is centered at the Scorpio-Centaurus OB association (Sco-Cen, $l\approx329^{\circ},b\approx+17.5^{\circ}$), and probably created due to its intense stellar activity \citep{blaauw1964,degeus1992,sartori2003,preibisch2008}. This neighboring interstellar bubble, located at $d\approx 130$pc and defined by a large sky-projected diameter ($\sim115^{\circ}$), has been better revealed by radio continuum observations, which show several arc-shaped shells of interstellar material \citep{berkhuijsen1971,iwan1980,haslam1982,heiles1998}. The proximity between the Local and Loop I bubbles led some authors to believe that some kind of interaction could be taking place between them \citep{coxreynolds1987,centurion1991}. In fact, by analyzing shadowing effects at the wide-angle soft X-ray survey from ROSAT ($0.25$keV), \citet{egger_aschenbach1995} proposed that the collision between both bubbles led to the formation of a wall of neutral gas surrounded by a dense interstellar ring feature at the interaction zone. This conclusion was inspired by collisional models of spherical shock fronts, which revealed that a dense interacting wall would arise, encompassed by an even denser annular feature, if at least one of the interacting shells have reached the radiative stage before the collision occurred \citep{yoshioka1990}. This idea was supported by an anticorrelation between the shadows from the soft X-ray maps and the neutral hydrogen (HI) emission from the local ISM \citep{egger_aschenbach1995,breit2000}. Up to date, several attempts have been made to determine the distance to the supposed interacting region, leading to widely different results. \citet{centurion1991} suggested a distance of $40\pm 25$ pc, from the analysis of ultraviolet spectra to eight stars at the region defined by $315^{\circ}<l<330^{\circ}$ and $15^{\circ}<b<25^{\circ}$. \citet{egger_aschenbach1995} used HI column densities data from \citet{fruscione1994} to determine a distance of $70$ pc, defined by a jump in $N_{H}$ from $\leq 10^{20}$ cm$^{-2}$ to $\geq 7\times 10^{20}$ cm$^{-2}$ at this distance. Using $E(b-y)$ color excess data and high resolution spectroscopy in the direction of the Southern Coalsack, Chamaeleon, and Musca dark clouds \citet{corradi1997,corradi2004} suggested that the interaction zone is twisted and folded, located at $120-150$ pc along this line-of-sight. As previously pointed out by \citet{dame2001}, several dark clouds ($\rho$ Oph, Lupus, R CrA, G317-4, Southern Coalsack, Chamaeleon, and Musca) are located at the same mean distance of $150$ pc, in the direction of Loop I. However, \citet{welsh2005} identified the presence of an interstellar cloud in the direction of $(l,b)\approx(330^{\circ},+18^{\circ})$, at approximately $90$ pc from the Sun, suggesting that this could be part of the interface between the Local and Loop I bubbles. Recently, \citet{reis_corradi2008} used a larger sample of $E(b-y)$ color excess data distributed along the entire interface region to map the interstellar dust distribution. The analysis led to the conclusion that the expected transition from nearly $E(b-y)=0\fm015$ to $E(b-y)\approx 0\fm070-0\fm100$ \citep[which corresponds to the ring's density, as proposed by][]{egger_aschenbach1995}, occurs at the western (left) side at $110\pm 20$pc while the eastern (right side) transition cannot be clearly defined before $280\pm 50$ pc. Moreover, the structure of the interstellar magnetic field along the borders of the LC have been previously studied by several polarimetric surveys \citep{mathewson_ford1970,tinbergen1982,reiz1998,leroy1999,heiles2000}. It is known that the local ISM is filled by an irregular, large-scale magnetic flux of average intensity $\langle B\rangle\approx 2.2\mu$G \citep{heiles1998,beck2001,heiles2005}. Although no final conclusion has been reached in relation to the dominant physical mechanism responsible for the alignment of the interstellar dust particles, it is generally accepted that in the majority of cases, grain alignment occurs with the grain's major axis perpendicular to the magnetic field direction \citep{hall1949,hiltner1949,davis_greenstein_1951,jones_spitzer1967,codina1976,purcell1979,lazarian1995a,lazarian1995b,draine_wein_1996,draine_wein_1997,fosalba2002,lazarian2007}. This configuration of dust grains gives rise to an anisotropic extinction which results in a partially polarized transmitted light beam from a distant star, with position angle in the same direction as the field $\mathbf{\bf{B}}$. Therefore, there is a strong correlation between the direction of the plane-of-sky projected component of $\mathbf{\bf{B}}$ and the polarization vectors, which can be used as a powerful tool to map the Galactic magnetic field, as well as to probe the nature of the interstellar dust particles. All-sky polarization surveys exhibit a large-scale vectors distribution pattern which is generally correlated to the direction of the local interstellar structures. Specifically, it is frequently found that polarization vectors may be aligned roughly perpendicular or parallel to the interstellar filaments, depending on several physical and geometrical factors, including projection effects \citep{heiles1998,fosalba2002,whittet2003,heiles2005}. Particularly along the Galactic plane it is noted an overall distribution of polarization angles which is mainly horizontal (i.e., parallel to the plane). Such trend reflects the morphology of the local magnetic field, directed mainly along the local spiral arm. In fact, this predominant orientation along the Galactic plane vanishes when the local magnetic field is viewed face-on along its ``poles" at $l\approx(80^{\circ},260^{\circ})$, which is roughly coincident with the direction of the local spiral arm. Large-scale mappings of the magnetic field structure in other spiral galaxies show that the orientation parallel to the galactic plane is a general trend which may be attributed to differential rotation of the galactic disk and magnetic flux freezing with the interstellar matter \citep{zweibel1997,beck2002}. In this work, we present an optical polarimetric survey in the direction of the interface between the LC and Loop I. Section \ref{obsdata} provides a description of the observational data, as well as the reduction process. The results and analysis of the correlation between polarimetric and color excess data, as well as of the spatial distribution of the polarimetric vectors and polarization degree as a function of distance are shown on sections \ref{correlation_p_colourexcess}, \ref{polvec} and \ref{poldist}. Discussion of the results is carried out on section \ref{discussion} and the conclusions are shown on section \ref{conclusions}.

\label{conclusions} We have carried out a polarimetric survey of $878$ Hipparcos stars in the general direction of the Loop I superbubble's interface with the LC. Our sample was complemented with the data from \citet{heiles2000} catalogue. The main results of the analysis are summarized below: \begin{itemize} \item Along the ring structure proposed by \citet{egger_aschenbach1995}, the left side rise from $P\sim 0.2\%$ to $P\sim0.9-1.3\%$, which correspond to the ring's column density, occurs at $\approx100$ pc, while at the right side regions the same transition occurs only beyond $\approx250$ pc. This trend corroborates the color excess analysis by \citet{reis_corradi2008}; \item A gradual decrease in the polarization values is observed in the left-right direction along the contour of the interaction ring; \item The analysis of the polarization vectors direction along the interface revealed that along the right side the predominant direction of the vectors do not present correlation with the direction of the ring contour, in contrast to the situation observed at the left side, where the vectors run parallel to the ring structure. \end{itemize} Altogether, these evidence confirm the distorted nature of the interstellar interface between the LC and Loop I. The low polarization values toward some areas along the interaction ring suggest that, if it really exists, it is probably highly fragmented and twisted. However, the configuration of the sky-projected component of magnetic field in relation to the direction of the ring contour reveals markedly different behavior in the left and right sides. This fact casts some doubt on the existence of the ring-like structure as a unique large-scale feature. Our methods were mainly the analysis based on the distances to the interstellar structures and the shape of magnetic field lines along the interface. These studies were used in a qualitative comparison with the existent models for the local interstellar medium. The existence (or not) of the ring is closely related to the true nature and associated origin models of the Local Cavity, and shall deserve a particularly special attention.

1012.3394

The Sun is located inside an extremely low density and quite irregular volume of the interstellar medium, known as the Local Cavity (LC). It has been widely believed that some kind of interaction could be occurring between the LC and Loop I, a nearby superbubble seen in the direction of the Galactic center. As a result of such interaction, a wall of neutral and dense material, surrounded by a ring-shaped feature, would be formed at the interaction zone. Evidence of this structure was previously observed by analyzing the soft X-ray emission in the direction of Loop I. Our goal is to investigate the distance of the proposed annular region and map the geometry of the Galactic magnetic field in these directions. On that account, we have conducted an optical polarization survey of 878 stars from the Hipparcos catalog. Our results suggest that the structure is highly twisted and fragmented, showing very discrepant distances along the annular region: ≈100 pc on the left side and 250 pc on the right side, independently confirming the indication from a previous photometric analysis. In addition, the polarization vectors' orientation pattern along the ring also shows a widely different behavior toward both sides of the studied feature, running parallel to the ring contour on the left side and showing no relation to its direction on the right side. Altogether, these evidences suggest a highly irregular nature, casting some doubt on the existence of a unique large-scale ring-like structure.

12205356

2011IAUS..275..280F

1012

1012.0218_arXiv.txt

1012.0218

The observed X-ray luminosity of SS 433 is ~10<SUP>36</SUP> erg/s, it is known that all the radiation is formed in the famous SS 433 jets. The bolometric luminosity of SS 433 is ~10<SUP>40</SUP> erg/s, and originally the luminosity must be realized in X-rays. The original radiation is probably thermalized in the supercritical accretion disk wind, however the missing more than four orders of magnitude is surprising. We have analysed the XMM-Newton spectra of SS 433 using a model of adiabatically and radiatively cooling X-ray jets. The multi-temperature thermal jet model reproduces very well the strongest observed emission lines, but it can not reproduce the continuum radiation and some spectral features. We have found a notable contribution of ionized reflection to the spectrum in the energy range from ~3 to 12 keV. The reflected spectrum is an evidence of the supercritical disk funnel, where the illuminating radiation comes from deeper funnel regions, to be further reflected in the outer visible funnel walls (r &gt;= 2 . 10<SUP>11</SUP> cm). The illuminating spectrum is similar to that observed in ULXs, its luminosity has to be no less than ~10<SUP>39</SUP> erg/s. A soft excess has been detected, that does not depend on the thermal jet model details. It may be represented as a BB with a temperature of T<SUB>bb</SUB> ~ 0.1 keV and luminosity of L<SUB>bb</SUB>~3 . 10<SUP>37</SUP> erg/s. The soft spectral component has about the same parameters as those found in ULXs.

12167725

2011ApJ...728...79A

1012

1012.1925_arXiv.txt

Standard stellar evolution theory predicts that only one mixing event will change the surface composition of a low-mass star as it ascends the red giant branch. That event is the so-called first dredge-up (FDU, see \citealt{1967ApJ...147..624I}) associated with the inwards migration of the base of the convective envelope into regions where hydrogen burning via the CNO-bicycle has occurred. Relatively modest changes in surface C and N abundance are predicted, and once the convective envelope recedes outwards these changes are brought to a halt. However, observations of low-mass red giants ($M < 2.5 M_\odot$, see for example \citealt{1998A&A...336..915C}) for essentially all compositions show trends among light element abundances that cannot be accounted for by the FDU. Some form of non-convective mixing seems to occur whereby greater amounts of the products of partial hydrogen burning are cycled into the convective envelope over a much longer timescale, and during more advanced phases of RGB evolution, than can be explained by the FDU. The observational results summarised below imply that current canonical models do not include essential physics of this so-called ``extra mixing''. Regardless of what the physical mechanism is, extra mixing is required to conform to the following observational criteria: \begin{enumerate} \item It commences after the hydrogen burning shell has erased a composition discontinuity in the radiative zone that marked the innermost limit of the convective envelope during the FDU event, and may continue to at least the tip of the RGB \citep{1991ApJ...371..578G, 1998A&A...332..204C, 2000A&A...354..169G, 2003PASP..115.1211S, 2003ApJ...585L..45S, 2004MmSAI..75..347W, 2008AJ....136.2522M}. The onset of the extra mixing is thus thought to coincide with a local maximum (the so-called ``bump'') observed in the RGB luminosity function of globular clusters. \item It must occur over a range of masses and metallicities \citep[and references therein]{2009A&A...502..267S}, being active in giants of all metallicities from solar to at least [Fe/H] $\sim -2.5$ \citep{2000A&A...354..169G} and masses less than $\sim 2.5 M_\odot$ \citep{1977ApJ...217..508L}, although not necessarily with equal efficiency throughout these mass and metallicity ranges. \item It must deplete \el{7}{Li} \citep{1998A&A...332..204C, 2009A&A...502..267S,2009A&A...503..545L}. \item It must decrease the \el{12}{C}/\el{13}{C} ratio \citep{1994A&A...282..811C, 1996ASPC...98..213C}, since values lower than predicted by the FDU are found among Population I field giants \citep{1976ApJ...210..694T, 1981ApJ...248..228L, 1998A&A...332..204C}, open cluster giants \citep{1989ApJ...347..835G, 1991ApJ...371..578G,2009A&A...502..267S, 2010MNRAS.407.1866M}, globular cluster giants \citep{2003ApJ...585L..45S, 2007A&A...461L..13R} and halo field giants \citep{1986ApJ...311..826S,2000A&A...354..169G}. \item It must decrease the total carbon abundance since systematic decreases with advancing luminosity on the upper half of the red giant branch are seen both among globular clusters and halo field giants \citep{1981ApJS...47....1S,1989asgc.conf...71S,1982ApJS...49..207C, 1983ApJ...266..144T,1986PASP...98..473L,2000A&A...354..169G,2001PASP..113..326B,2003PASP..115.1211S,2008AJ....136.2522M,2010AJ....140.1119S}. In Population II giants the behaviour of the carbon abundance can serve as an even more potent probe of the extent of extra mixing than the \el{12}{C}/\el{13}{C} isotope ratio, because the latter can attain near-equilibrium values for only moderate amounts of mixing that would otherwise cause only small ($\sim$0.1 dex) changes in [C/H] \citep{1986ApJ...311..826S}. \item As a consequence of the previous point it must increase the nitrogen abundance. The results of CN cycling are observed on the upper half of the red giant branch. Halo field stars on the upper RGB were found by \citet{2000A&A...354..169G} to show an excess of nitrogen compared to those on the lower RGB. \end{enumerate} It is expected that the mechanism(s) will also destroy \el{3}{He} inside the star \citep{1986ApJ...302...35D, 1995PhRvL..75.3977H, 1996ApJ...465..887D, 1999ApJ...510..217S, 2007A&A...467L..15C, 2007A&A...476L..29C}. As we cannot observe \el{3}{He} in stellar atmospheres directly this is not an observational constraint, but it is a significant requirement from the study of chemical yields and galactic evolution. The importance of \el{3}{He} is discussed in Section 3. In the last four decades many candidate extra-mixing mechanisms have been suggested. These include: rotational mixing \citep{1979ApJ...229..624S, 2005ApJ...631..540C, 2006A&A...453..261P}, magnetic fields \citep{2009PASA...26..161P, 2008ApJ...684L..29N, 2007ApJ...671..802B}, and internal gravity waves \citep{2000MNRAS.316..395D}. Individually, none of these have been proven to be satisfactory. Due to its promising ability to account for the above requirements, and the necessity of its occurrence in low mass giants just after the FDU, in this study we focus on \el{3}{He}-driven ``thermohaline mixing'' (\citealt[EDL06 hereafter]{2006Sci...314.1580E}, \citealt[CZ07a hereafter]{2007A&A...467L..15C}, \citealt[EDL08 hereafter]{2008ApJ...677..581E}).\footnote{One must concede that multiple processes and indeed interactions between them can affect the transport. Models have been made that include multiple processes \citep{2010A&A...521A...9C, 2010A&A...522A..10C} by simply adding the diffusion coefficients for each process. This does not allow for the interaction between the processes as discussed by \citet{2008ApJ...684..626D}.} The name comes from a phenomenon seen in oceans, and is taken from the two major determinants of the density of sea water - its heat content (``thermo'') and its salinity (the salt or ``haline'' content). It is common to find warm salty water overlying cool fresh water. Although the higher salinity of the warm water makes it denser, the higher heat content acts to stabilise the stratification. The subsequent evolution of the system is determined by the competition between two diffusion processes and their associated timescales - the time for the (stabilizing) heat to diffuse away compared to the time for the (destabilizing) salt to do the same. Hence the process is often called ``doubly-diffusive'', and it has been studied in the oceanographic context for many years (see recent reviews, theoretical modelling, observations and laboratory experiments in Progress of Oceanography Volume 56, 2003; e.g. \citealt{ruddick}). Within oceans it is now well known that the rapid diffusion of heat from the warm salty layer produces an over-dense layer that begins to sink into the cooler fresh water below. The temperature stays roughly the same as the surrounds, and ``salt fingers'' form which extend downward delivering the saltier water to deeper regions. Reciprocal fresh-water fingers move upward and replace the salty water with fresher water. Laboratory experiments have also played a role in helping to characterise the instability. Work by Stommel and Faller published in \citet{stern} as well as more recently \citet[and references there in]{Krishnamurti} have helped elucidate the instability. A similar process can occur in stars. Here it is not salt but the mean molecular weight that is the ``destabilising agent'', in the words of \citet{2010arXiv1006.5481D}. Usually, the molecular weight increases as we move toward the centre of the star, as a result of nuclear burning and fusion reactions. If it were to decrease, then the plasma would be buoyantly unstable, just as is the case in convection. However, just as in the oceanic case, we must include the rapid thermal diffusion which can act to stabilise the motion. In this paper we consider the situation where some local event causes a decrease in the molecular weight in an otherwise stable region within a star. We investigate the effects of the resultant ``thermohaline mixing'' or doubly diffusive process that is initiated by a molecular weight ($\mu$) inversion \citep*{1972ApJ...172..165U,1980A&A....91..175K}\footnote{This has been referred to as ``$\delta \mu$ mixing'' by EDL06 and EDL08, to distinguish it from a separate occurrence of ``thermohaline mixing'' in stars. In that case we may have mass transfer in a binary system, where material of a higher mean molecular weight is accreted on an envelope of lower molecular weight. This situation is unstable and some thermohaline circulation will take place to redistribute the composition of the star to result in a stable stratification (exchanging energy with the thermal content of the material in doing so). This case is not relevant to the discussions in this paper, and more information may be found in \citet{2004MNRAS.355.1182C} and \citet{2007A&A...464L..57S}.}. Recently thermohaline mixing has featured prominently in the literature. EDL06, CZ07a, EDL08, \citet{2010A&A...521A...9C} and \citet{2010A&A...522A..10C} have discussed in detail the important consequences of its inclusion during the RGB. The dichotomy in RGB carbon abundance between metal poor stars and carbon-enhanced metal-poor stars has been explained by \citet{2009MNRAS.396.2313S} using thermohaline mixing. \citet{2010A&A...521A...9C}, \citet{2010MNRAS.403..505S} and \citet{2010A&A...522A..10C} have shown that its operation beyond the giant branch may affect the subsequent asymptotic giant branch (AGB) evolution. This mechanism may be a crucial part of stellar physics that has been missing from the models. We thus believe it pertinent to investigate the effects of the mechanism in some detail. In this paper we will take the approach of trying to model the mixing in spherically symmetric stellar models and investigate its effect on observable surface abundances. In this respect our approach follows that of CZ07a and \citet{2010A&A...522A..10C}. We will examine the change of the abundances of carbon and nitrogen on the red giant branch of globular clusters, focusing on the case of M3. Currently a range of approaches is taken to include thermohaline mixing in evolution codes, especially if the \el{12}{C}/\el{13}{C} ratio is the constraint used to determine the extent of extra mixing. The \el{12}{C}/\el{13}{C} ratio was one of the first indicators that extra mixing must operate on the RGB. It is classically used to probe the results of FDU \citep{1975MNRAS.170P...7D,1976ApJ...210..694T, 1994A&A...282..811C}. It naturally follows that the \el{12}{C}/\el{13}{C} ratio could be used to trace the extent of extra mixing and constrain any mechanism. The change in \el{12}{C}/\el{13}{C} ratio following FDU will depend on the efficiency of mixing and allow us to explore the mixing velocity via a diffusion approximation. EDL08 estimated the mixing speed with their formula for the diffusion coefficient and found that a range of three orders of magnitude in their free parameter can lead to the observed levels of \el{12}{C}/\el{13}{C} and \el{3}{He} depletion. \citet{1972ApJ...172..165U} and \cite*{1980A&A....91..175K} both use essentially the same formula (UKRT formula hereafter) for the diffusion coefficient but their choice of the free parameter varies by two orders of magnitude. In an attempt to constrain the parameter space we address the following questions: \begin{enumerate} \item The \el{12}{C}/\el{13}{C} ratio is generally used as a tracer to probe the extent of mixing. This ratio saturates near the CN equilibrium value rather quickly in low metallicity stars, and hence is of limited utility. Is there a better way to constrain the mixing? \item Which formalism should be used? Here we will limit our investigation to the EDL08 and UKRT prescriptions for the diffusion coefficient. \item Once the preferred formalism is identified, what diffusion co-efficient (or mixing velocity) is needed to match observations? What values do we use for any free parameters? \end{enumerate} As has been the practice for many years, we turn to globular clusters to test our understanding of stellar theory. \citet{2002PASP..114.1097S}, \citet{2003PASP..115.1211S} and \citet{2008AJ....136.2522M} have compiled observations of carbon and nitrogen along the giant branch of M3. This has provided us with a valuable alternative to the \el{12}{C}/\el{13}{C} probe. By matching our models to the carbon depletion (as a function of absolute magnitude) observed in this cluster we can attempt to constrain both the form of the thermohaline diffusion coefficient and the values of any parameters contained therein. As carbon and nitrogen are intrinsically linked in the CN burning cycle we include observations of nitrogen as an additional tracer. Furthermore, we identify when extra mixing begins in the models and compare this to the observed luminosity function bump (LF bump) in the cluster.

M3 is a well studied system that demonstrates the abundance patterns we commonly associate with globular clusters. Along with many other clusters it displays significant [C/Fe] depletion along the RGB, the implication being that some form of internal, non-canonical mixing must be occurring. Our models with thermohaline mixing show that the carbon and nitrogen observations can be explained if we adopt the hybrid theory outlined in \citet{2002PASP..114.1097S}, where stars in the cluster are undergoing extra mixing as they ascend the RGB and the presence of primordial abundance inhomogeneities due to ON cycling are needed to explain the initial carbon and nitrogen abundances. We have used observations of M3 to investigate our theoretical understanding of thermohaline mixing. Our findings are summarised below: \begin{enumerate} \item The variation of {$\rm{[C/Fe]}$} with magnitude provides a much more stringent test of any proposed extra-mixing mechanism than simply matching the final \el{12}{C}/\el{13}{C} ratio. When data of sufficient quality is available this can constrain the details of any proposed extra-mixing formulation. In the present case the UKRT formulation of thermohaline mixing is a far better fit than the phenomenological prescription given by EDL08, although both fit the constraint provided by the carbon isotope ratio. \item The UKRT prescription of thermohaline mixing with \textit{C$_{t}$}=1000 seems to best fit the data for M3. This is consistent with the results of \citet{2010A&A...522A..10C} for higher metallicities and CZ07a for a range of metallicities. \item We infer that there is a spread of $\sim$0.3 to 0.4 dex in [C/Fe] in the stars in M3 from their birth. Without this initial difference in [C/Fe] between the two populations, thermohaline mixing cannot reproduce the change in [C/Fe] seen on the giant branch. That there are two populations is absolutely required, because once \textit{C$_{t}$} is sufficiently large, an increase of this coefficient doesn't produce a bigger $\Delta$[C/Fe]. Primordial C and N inhomogeneities have been directly observed as abundance differences among main sequence stars in globular clusters such as M13 and NGC 6752 \citep{2004AJ....127.1579B, 2005A&A...433..597C}, which have similar metallicity to M3. \item Thermohaline\footnote{Note, when we refer to ``thermohaline mixing" we are referring to the linear theory as proposed by \citet{1972ApJ...172..165U} and \citet{1980A&A....91..175K}.} mixing can produce the observed values of the carbon isotopes seen in M3. \item To reproduce the entire spread of [C/Fe] values seen in the giants of M3 it is essential that thermohaline mixing operate in both the CN-strong and CN-weak populations identified in (3). In this case we can explain the full spread in [C/Fe] seen near the tip of the giant branch in M3. A similar exercise was carried out by \citet{2003ApJ...593..509D} for the case of M92, modelled with a simple parameterized extra-mixing formulation. The data for M92 are from many different sources and make it difficult to estimate precisely where the extra mixing begins. For this reason we do not discuss M92 further in this paper. \item Thermohaline mixing can produce a significant change in [N/Fe] as a function of M$_V$ on the RGB for initially CN-weak stars but not for initially CN-strong stars, which have so much N to begin with that any extra mixing does not significantly affect the surface composition. \item The level of depletion of carbon is dependent on the depth to which the material is mixed and how fast it is mixed. \item Mass loss has little effect on the surface abundances. \item Both the predicted and observed composition changes take place at a luminosity that is higher than the LF bump in \citet{1990A&A...238...95F}. The observed abundances begin to decrease at a luminosity lower than the LF bump preferred by \citet{2003PASP..115.1211S}. Uncertainties in the distance modulus make it difficult to draw further conclusions. \end{enumerate} We have seen that the linear theory of \citet{1972ApJ...172..165U} and \citet{1980A&A....91..175K} provides a fit to the carbon and nitrogen abundances in the giants in M3 (assuming there are two different populations initially). CZ07a have shown that the same theory (and indeed the same parameter \textit{C$_{t}$}=1000) seems to fit field stars of a range of metallicities (see also \citealt{2010A&A...522A..10C}). It is important to note that thermohaline mixing is more than a theory with an adjustable parameter. For example, its beginning is determined clearly by the fusion of \el{3}{He} which produces a molecular weight inversion. Also, the hydrodynamics provides the physical formulation for the diffusion co-efficient used, and hence its variation throughout the star, to within a constant which depends on the geometry of the fingers expected in the mixing process. Nevertheless, until we have a complete theory that also determines this geometric factor, or at the least some numerical simulations, there is a gap between understanding the fundamental physics and the sort of work presented here, to match the observations. To this end we note the work of \citet{2010arXiv1006.5481D} which addresses this issue. There has been recently considerable progress in modelling the oceanic case \citep{2010arXiv1008.1808S,2010arXiv1008.1807T}. We are ourselves working on 3D hydrodynamic models in the stellar context, which will be the subject of another paper. Although our models can explain M3 very well the question remains whether this work can be extended to all clusters. The UKRT prescription for thermohaline mixing appears to model the internal mixing of a young metal rich cluster. Old metal-poor clusters such as M92 will have undergone a different mixing history. Furthermore \citet{1978IAUS...80..333S} suggests changing the metallicity and helium content will drastically alter the location and size of the LF bump. Preliminary work by \citet{2010arXiv1006.5828A} will be the subject of subsequent studies. This work further supports the ability of thermohaline mixing to explain extra mixing on the RGB.

1012.1925

We review the observational evidence for extra mixing in stars on the red giant branch (RGB) and discuss why thermohaline mixing is a strong candidate mechanism. We recall the simple phenomenological description of thermohaline mixing and aspects of mixing in stars in general. We use observations of M3 to constrain the form of the thermohaline diffusion coefficient and any associated free parameters. This is done by matching [C/Fe] and [N/Fe] along the RGB of M3. After taking into account a presumed initial primordial bimodality of [C/Fe] in the CN-weak and CN-strong stars, our thermohaline mixing models can explain the full spread of [C/Fe]. Thermohaline mixing can produce a significant change in [N/Fe] as a function of absolute magnitude on the RGB for initially CN-weak stars, but not for initially CN-strong stars, which have so much nitrogen to begin with that any extra mixing does not significantly affect the surface nitrogen composition.

12213567

2011MNRAS.413..386B

1012

1012.1322_arXiv.txt

Within the last decade, three foundational observational probes have been well recognized as opening up observational windows to reveal some of the most valuable secrets of the Universe. These are: the temperature fluctuations of the cosmic microwave background radiation (CMB), recently measured with high precision by the Wilkinson Microwave Anisotropy Probe (\emph{WMAP}) satellite \citep{spergel07, komatsu10}; the Hubble diagram inferred from the Type Ia supernovae (SNe) observations \citep[e.g.][]{perlmutter, riess} and the measurement of the large scale structure (LSS) of the Universe as traced by the spatial distribution of galaxies \citep{percival02,teg_sdss,sanchez06,percival_sdss,percival_last,reid}. With the completion of large redshift surveys, such as the Two-degree Field Galaxy Redshift Survey ($2$dFGRS)\footnote{http://msowww.anu.edu.au/2dFGRS/} and the Sloan Digital Sky Survey (SDSS)\footnote{http://www.sdss.org/} it has been possible to push the level of accuracy of LSS studies. The recent detection of the baryon acoustic oscillations (BAO, \citealp[e.g.][]{ei, cole2df,gaz, ariel_last,percival_last}) in redshift surveys have been key to this progress. Over the last few years these observations have established the concordance cosmological model, based on a flat space-time in a current phase of accelerated expansion due to the presence of a dominating dark energy component, whose equation of state is compatible with Einstein's cosmological constant $\Lambda$. This is the so-called $\Lambda$CDM cosmological model. Galaxy cluster samples also provide valuable information to obtain constraints on cosmological parameters. Galaxy clusters are the largest bounded structures in the Universe. They are associated with the highest peaks in the matter density field and are recognized as biased tracers of the underlying matter distribution \citep[e.g.][]{bardeen}. Their deep potential wells make them the largest astrophysical laboratories in the Universe, where the combination of gravitation and baryonic physics has been intensively studied through the analysis of cluster properties such as scaling relations \citep[e.g.][]{lm0,pratt,mantzI}, density profiles \citep[e.g.][]{makino}, pressure profiles \citep[e.g.][]{pressure}, baryon fractions \citep[e.g.][]{baryon_frac}, etc. The abundances of galaxy clusters determined by their luminosity function \citep[e.g.][]{reflex_lf} can also be used to constrain parameters like the matter content in the universe $\Omega_{\rm m}$, and the amplitude of the density fluctuations characterized by $\sigma_{8}$, the rms linear perturbation theory variance in spheres of radius $8\,$ Mpc $h^{-1}$ \citep[e.g.][]{reflex1_s8}. The spatial distribution of galaxy clusters, characterized by its power spectrum (or correlation function), provides useful information about the cosmological model of the Universe. The shape of this measurement is particularly sensitive to the parameter combination $\Omega_{\rm m}h$, which complements the constraints on these parameters obtained from the analysis of fluctuations in the CMB. Furthermore, the amplitude of the galaxy cluster power spectrum also contains important information that can be related to theoretical models in a more direct way than for measurements based on galaxy samples \citep[e.g.][]{moscardini}. \begin{figure*} \includegraphics[width=14cm]{reflex_mask_pres_2.epsi} \caption[]{Distribution of REFLEX II clusters in equatorial coordinates. Open circles represent the position of the REFLEX II clusters. Filled circles represent the REFLEX II clusters without redshift. The points represent $5$ per cent of the random catalogue constructed with the REFLEX II selection function. The empty regions in the Southern hemisphere corresponds to the cut in galactic coordinates $|b_{II}|<20^{o}$ of the Milky Way (band) and the Magellanic clouds. } \label{dsky} \end{figure*} In the past few years the ROSAT-ESO Flux Limited X-ray (REFLEX) catalogue \citep{hb_catalogo} has been used to measure fundamental cosmological quantities. The REFLEX catalogue is based on \emph{ROSAT} All Sky Survey (RASS) observations \citep{truemper}, complemented with follow-up observations as described by \citet{guzzo_optical}, yielding spectroscopic redshifts for 447 clusters with flux limit of $3 \times 10^{-12}\, {\rm erg}\,{\rm s}^{-1}{\rm cm}^{-2}$ (in the \emph{ROSAT} energy band $0.1-2.4$\,keV). The REFLEX catalogue was, to date, the largest statistically complete X-ray detected cluster sample. The clustering properties of this survey were analyzed by means of the power spectrum \citep{reflex1}, the cluster correlation function \citep{collins_cf}, cluster-galaxy cross-correlation functions \citep{ariel_reflexI} and Minkowski functionals \citep{kerscher}. Sub-samples of the REFLEX catalogue complemented by detailed follow-up observations have been used to constrain cluster scaling relations \citep[e.g.][]{lm0,stanek,pratt,mantzI}. In this paper we present the analysis of the power spectrum of the new REFLEX II catalogue. The REFLEX II is an extension of the REFLEX catalogue to a lower limiting flux ($1.8 \times 10^{-12}\, {\rm erg}\,{\rm s}^{-1}\,{\rm cm}^{-2}$) allowing the inclusion of $464$ new clusters over the original sample and yielding a total of 911 clusters with spectroscopic redshifts for $\sim 95$ per cent of the sample. In addition to the enlarged sample size of REFLEX II, several improvements were made to the data reduction: (i) we use the RASS survey product RASS III which gives a few percent more sky exposure in formerly underexposed areas due to improved attitude solutions, consequently recovering a few more clusters at higher flux, (ii) for the count rate to flux conversion an estimated temperature has to be applied, which is now obtained with up-to-date scaling relations based on the REXCESS Survey \citep{bohringer_07} with the L-T relation described in \citet{pratt}; (iii) the total flux and X-ray luminosity is now estimated inside the radii of $r_{500}$ and $r_{200}$ (based on relations described in \citealt{pratt} and \citealt{arnaud_05}). These calculations now involve less extrapolation than the estimates for the previously used fiducial radius. Besides the advantages provided by a larger cluster sample, the power spectrum analysis presented here represents an improvement over that of \citet{reflex1} in a number of ways. In particular, our analysis is complemented with a set of $N$-body simulations, the L-BASICC II \citep{angulo,ariel1}, from which we constructed a suit of $100$ REFLEX II mock catalogues. These catalogues were calibrated to reproduce the measured REFLEX II X-ray luminosity function. Selection criteria of the REFLEX II sample were applied in their construction, yielding a large suit of mocks that can be used to analyze the statistical methods applied to the data. The details of the construction of these mock catalogues will be described in a forthcoming paper (S\'anchez et. al. in preparation). Our ensemble of mock catalogues allowed us to show that it is possible to construct an accurate model of the shape and amplitude of the REFLEX II power spectrum. This model includes the effects of the non-linear evolution of density fluctuations, redshift-space distortions and halo bias, which introduce deviations in the clustering signal with respect to the simple predictions of linear perturbation theory. This will allow us to use the full information contained in the REFLEX II power spectrum to obtain constraints on cosmological parameters. This paper is organized as follows. In Section \ref{sec_sel} we describe the REFLEX II sample and the survey selection function, followed by a brief description of the construction of the mock catalogues in Section \ref{sec_mocks}. In Section \ref{sec_power} we describe the power spectrum estimator and show the measurements of the REFLEX II window function and the covariance matrix. In Section~\ref{lumi_bias} we model the amplitude of the power spectrum measured from the mock catalogues. In Section ~\ref{sec_malm} we explore the sensitivity of the REFLEX II sample to distortions induced by flux-selection effects. In Section \ref{sec_shape} we model the shape of the power spectrum. The model of the shape and the amplitude is applied to the REFLEX II sample in Section ~\ref{sec_q}. We end with our conclusions in Section \ref{sec_conc}. Our fiducial cosmological model consist of a flat $\Lambda$CDM Universe with a matter energy density parameter of $\Omega_{\rm m}=0.25$, a dark energy equation of state $w=-1$, a dimensionless Hubble parameter $h=0.7$ \footnote{The Hubble constant $H_{0}$ in units of $100 \,{\rm km}\,{\rm s}^{-1}{\rm Mpc}^{-1}$.} and a spectral index of primordial scalar fluctuations $n_{s}=1$. Throughout this paper we always refer to the X-ray luminosity in the \emph{ROSAT} hard energy band $0.1-2.4\, {\rm keV}$ and, whenever it is not explicitly written, its units are given in $10^{44}{\rm erg}\,{\rm s}^{-1}h^{-2}$.

\label{sec_conc} In this paper we presented the measurement and analysis of the power spectrum from the new REFLEX II catalogue, which is an extension of the original REFLEX sample \citep{hb_catalogo} to a lower limiting flux ($1.8 \times 10^{-12}\, {\rm erg}\,{\rm s}^{-1}\,{\rm cm}^{-2}$). The new sample contains $911$ X-ray detected galaxy clusters of which $860$ have measured redshifts in the range $0\leq z\lesssim 0.6$ and X-ray luminosities in the range $4.9\times 10^{40}\leq L_{X}/({\rm erg}\,\,{\rm s}^{-1}h^{-2})\leq 1.96\times 10^{45}$. The total flux and X-ray luminosities are estimated using the up-to-date scaling relations based on the REXCESS Survey \citep{bohringer_07,pratt}. The new sample allowed us to perform a detailed study of the full shape and amplitude of the power spectrum of X-ray detected galaxy clusters. We complemented this analysis by using a set of 100 independent mock catalogues constructed to match the selection function of the REFLEX II survey. The clustering properties of these mock catalogues are in good agreement with those measured in the REFLEX II sample. Thus, this ensemble provides a reliable tool to test the statistical methods applied to the data. In particular, we used the mock catalogues to test a model for the luminosity dependence of bias, to construct covariance matrices of the the REFLEX II power spectrum and to analyze the possible systematic effects that might affect this measurement. Due to the flux-limited selection of the REFLEX II survey, the clustering pattern of galaxy clusters might be affected by scale-dependent distortion, as has been observed in galaxy surveys \citep[e.g.][]{teg_sdss,percival_sdss}. Using the mock catalogues, we have shown that these distortions might affect the clustering in configuration space (i.e., when measured with the cluster-correlation function) on scales $r \geq 150\,{\rm Mpc}\, h^{-1}$, which would naively correspond to scales $k\leq 0.04 h {\rm Mpc}^{-1}$ in Fourier space. In order to test the impact of this flux-selection effect on the final measurements of power spectrum, we implemented the luminosity dependent estimator of \cite{pvp}, which is designed to correct for this distortion. We observed that the shape of the power spectrum measured by means of the FKP estimator does not show significant distortions compared to the results from the PVP estimator. This implies that the flux-selection of the REFLEX II sample does not introduce a significant systematic effect in the measurement of the power spectrum of this catalogue. \begin{figure} \includegraphics[width=8cm, angle=0]{power_reflex_qmodel2.epsi} \caption[]{Best-fitting $Q$-model for the REFLEX II mock power spectrum (points with error bars). See Fig.~\ref{reflex2_power_qmodel} for description.}\label{reflex2__qmodel} \end{figure} The shape of the mean power spectrum from our ensemble of mock catalogues is in good agreement with the measured power spectrum from the REFLEX II sample, and is statistically distinguishable from the linear perturbation theory predictions on intermediate scales. This implies a clear signature of non-linear evolution in the X-ray cluster spatial distribution. Nevertheless, given the level of accuracy of the measurements of power spectrum in the REFLEX II sample, it is sufficient to model these distortions using the $Q$-model of \citet{cole2df}. We find that this prescription provides a good description of the measurements from the mock catalogues on intermediate scales ($0.02\leq k\,/(h\,{\rm Mpc}^{-1}) \leq 0.25$). This model can also be used to describe the shape of the measured REFLEX II power spectrum, providing a valuable tool to extract the cosmological information contained in the shape of this statistic. The next generation of X-ray galaxy clusters surveys, such as {\it eROSITA}\footnote{http://www.mpe.mpg.de/heg/www/Projects/EROSITA/main.html} and \emph{WFXT}\footnote{http://wfxt.pha.jhu.edu/}, will provide measurements of the two-point statistics of the cluster population with higher accuracy than present-day samples, for which a more detailed modelling of non-linearities will be required \citep[e.g.][]{rpt, montesano}. Our measurements of the REFLEX II power spectrum are compatible with the prediction of the $\Lambda$CDM cosmological model and shows good agreement with the previous results from the REFLEX sample \citep{reflex1}, save the expected differences due to the lower limiting flux of the REFLEX II sample. We showed that our measurements cannot provide a statistically significant detection of BAO, which is mainly due to the moderate volume probed by the survey (compared to the volume probed by current galaxy redshift surveys). We found that the power spectra measured from the REFLEX II sample and the mock catalogues are compatible with a scale-independent effective bias in the range of wavenumbers $0.01\leq k\, /(h\,{\rm Mpc}^{-1})\leq 0.1$, and that a simple theoretical prediction, based on the halo-mass bias, the halo mass function and the mass-luminosity relation, is able to describe these measurements. This, together with the modeling of the shape of the power spectrum given by the $Q$-model, provides a link to the cosmological models and allows our measurements to reach their full constraining power.

1012.1322

We present the power spectrum of galaxy clusters measured from the new ROSAT-ESO Flux-Limited X-Ray (REFLEX II) galaxy cluster catalogue. This new sample extends the flux limit of the original REFLEX catalogue to 1.8 × 10<SUP>-12</SUP> erg s<SUP>-1</SUP> cm<SUP>-2</SUP>, yielding a total of 911 clusters with ≥94 per cent completeness in redshift follow-up. The analysis of the data is improved by creating a set of 100 REFLEX II-catalogue-like mock galaxy cluster catalogues built from a suite of large-volume Λ cold dark matter (ΛCDM) N-body simulations (L-BASICC II). The measured power spectrum is in agreement with the predictions from a ΛCDM cosmological model. The measurements show the expected increase in the amplitude of the power spectrum with increasing X-ray luminosity. On large scales, we show that the shape of the measured power spectrum is compatible with a scale-independent bias and provide a model for the amplitude that allows us to connect our measurements with a cosmological model. By implementing a luminosity-dependent power-spectrum estimator, we observe that the power spectrum measured from the REFLEX II sample is weakly affected by flux-selection effects. The shape of the measured power spectrum is compatible with a featureless power spectrum on scales k &gt; 0.01 h Mpc<SUP>-1</SUP> and hence no statistically significant signal of baryonic acoustic oscillations can be detected. We show that the measured REFLEX II power spectrum displays signatures of non-linear evolution.

3825980

2012MNRAS.423.1463A

1012

1012.4441_arXiv.txt

The Sunyaev--Zel'dovich (SZ) effect is the inverse-Compton scattering of cosmic microwave background (CMB) photons from the hot plasma within clusters of galaxies (\citealt{SZE}, see e.g. \citealt{BIRK_SZ_REVIEW} and \citealt{CARL_SZ_REVIEW} for reviews). The surface brightness of an SZ signal does not depend on the redshift ${z}$ of the cluster and the integrated signal is only weakly dependent on $z$ via the angular diameter distance. Hence an SZ-effect flux-density-limited survey can provide a complete catalogue of galaxy clusters above a limiting mass (see e.g. \citealt{BART_SILK}, \citealt{AMI_EXPECTED_RESULTS}, \citealt{ACT_KOSO} and \citealt{SPT_INTRO}). Detecting and imaging the SZ effect has gradually become routine since it was first securely detected by \cite{Birkinshaw_1981} and first imaged by \cite{Jones_1993}. Until recently, SZ observations have been directed almost entirely towards clusters selected optically or in X-ray, for example with AMI (\citealt{7CLUSTERS}), AMiBA (\citealt{AMiBA_WU}), APEX (\citealt{APEX_HALV}), CBI (\citealt{CBI_UDOM}), CBI-2 (\citealt{CBI_PEAR}), OCRA (\citealt{OCRA_LANC}), OVRO/BIMA (\citealt{CARL_JOY}), RT (\citealt{RT_GRAINGE}), SuZIE (\citealt{SuZIE_HOL}), SZA (\citealt{Muchovej_2011}) and the VSA (\citealt{VSA_LANC}). Now, however, SZ blind surveying is underway, with ACT and SPT having produced initial results (\citealt{ACT_CLUSTERS}, \citealt{ACT_CLUSTERS2}, \citealt{SPT_2009}, \citealt{SPT_CLUSTERS2} and \citealt{SPT_CLUSTERS3}). The Arcminute Microkelvin Imager (AMI) is conducting a blind cluster survey at 16\,GHz in twelve regions, each typically one deg$^2$, which contain no previously recorded clusters. The AMI cluster survey focuses on depth, aiming to detect weak SZ-effect signals from clusters of galaxies with a mass above $M_{T,200}$ $=$ 2 $\times$ $10^{14}$$M_{\odot}$, where $M_{T,200}$ corresponds to the total cluster mass within a spherical volume such that the mean interior density is 200 times the mean density of the Universe at the current epoch. The outline of this paper is as follows. In Section \ref{sec:OBS_INFO}, we give a brief description of the instrument, observations, data reduction and map making techniques. Identifying cluster candidates is described in Section \ref{sec:METHOD} -- we stress that some readers will wish to jump to the start of Section \ref{sec:METHOD} which is an important overview of the three analysis methods and of their assumptions. We discuss how we apply a Bayesian analysis to the AMI data in Section \ref{sec:BAYES} and present the results in Section \ref{sec:RESULTS}. We assume a concordance $\rm{\Lambda}$CDM cosmology, with $\rm{\Omega_{m}}$ = 0.3, $\rm{\Omega_\Lambda}$ = 0.7 and H$_{0}$ = 70 km\,s$^{-1}$Mpc$^{-1}$. The dimensionless Hubble parameter $h_{70}$ is defined as $h_{70}$ = H$_{0}$/(70 km\,s$^{-1}$Mpc$^{-1}$). All coordinates are given at equinox J2000.

\label{sec:CONC} \begin{itemize} \item We have presented a large, complex Sunyaev--Zel'dovich structure in an AMI blind field. The structure may be two separate components or be a single merging system. \item A Bayesian analysis using a physical model for the cluster (including assumed priors on the number density of clusters) was used to constrain cluster parameters such as $\beta, r_{c}, M_{T,200}$ and $z$. Using the Bayesian evidences we have calculated formal probabilities of detection taking into account point sources, receiver noise and the statistical properties of the primary CMB anisotropy. For the deeper component we find a formal probability of detection ratio of 7.9 $\times$ $10^{4}$:1 when assuming the \cite{Evrard_02} cluster number count and 2.1 $\times$ $10^{5}$:1 when assuming \cite{Jenkins} as the true prior. We derive a cluster mass of ${M_{T,200}}=5.5^{+1.2}_{-1.3}$$\times$ $10^{14}{h_{70}^{-1}M_{\odot}}$. \item A Bayesian analysis using a phenomenological model of the gas distribution was also used to quantify the significance of the detection and again taking into account point sources, receiver noise and the statistical properties of the primary CMB anisotropy. For the deeper component we find $\Delta T_{0} = -295^{+36}_{-15} \mu \rm{K}$. \item In our pointed follow-up observation the cluster system is detected with a high significance, with each map indicating that there is a 0.6mJy/beam peak decrement ($8\sigma_{\rm{SA,pointed}}$) towards the deeper component and an integrated decrement flux density ($\rm{S_{SZ,integrated}}$) of $\rm{\approx1.2mJy\rm/beam}$. The other component has a 0.5mJy peak decrement and an integrated decrement of 0.7mJy. \item Using the approximation $\rm{M_{T}^{5/3} \propto S_{\rm{SZ,integrated}}}$ we anticipate that the AMI blind cluster survey will detect clusters with $M_{T,200} >2$ $\times$ $10^{14}{h_{70}^{-1}M_{\odot}}$ at $4\sigma_{\rm{SA,survey}}$. \end{itemize}

1012.4441

We present an interesting Sunyaev-Zel’dovich (SZ) detection in the first of the Arcminute Microkelvin Imager (AMI) ‘blind’, degree-square fields to have been observed down to our target sensitivity of ?. In follow-up deep pointed observations the SZ effect is detected with a maximum peak decrement greater than eight times the thermal noise. No corresponding emission is visible in the ROSAT all-sky X-ray survey and no cluster is evident in the Palomar all-sky optical survey. Compared with existing SZ images of distant clusters, the extent is large (≈10 arcmin) and complex; our analysis favours a model containing two clusters rather than a single cluster. Our Bayesian analysis is currently limited to modelling each cluster with an ellipsoidal or spherical β model, which does not do justice to this decrement. Fitting an ellipsoid to the deeper candidate we find the following. (a) Assuming that the Evrard et al. approximation to Press &amp; Schechter correctly gives the number density of clusters as a function of mass and redshift, then, in the search area, the formal Bayesian probability ratio of the AMI detection of this cluster is 7.9 × 10<SUP>4</SUP>:1; alternatively assuming Jenkins et al. as the true prior, the formal Bayesian probability ratio of detection is 2.1 × 10<SUP>5</SUP>:1. (b) The cluster mass is ?. (c) Abandoning a physical model with number density prior and instead simply modelling the SZ decrement using a phenomenological β model of temperature decrement as a function of angular distance, we find a central SZ temperature decrement of ?K - this allows for cosmic microwave background primary anisotropies, receiver noise and radio sources. We are unsure if the cluster system we observe is a merging system or two separate clusters. We request that any reference to this paper cites ‘AMI Consortium: Shimwell et al. 2012’.

12273929

2012CQGra..29a5010G

1012

1012.4488_arXiv.txt

The Laser Interferometer Space Antenna (LISA) \cite{LISA} is expected to see thousands of low-frequency gravitational wave sources when it is launched around $2025$. LISA's detection of extreme mass ratio inspirals (EMRIs) -- a white dwarf, neutron star or small black hole (collectively referred to as small compact object, SCO) inspiraling toward a supermassive black hole of mass $10^5 - 10^7$ solar masses -- is expected to provide an unprecedented level of insight into the structure of spinning black hole spacetimes as well as into the dynamics and populations of EMRIs in galactic nuclei. EMRI signals will be extracted with matched filtering techniques from LISA's data stream. Once the presence of a signal is established one can use the resulting roughly-estimated parameters to refine the search using waveforms calculated with accuracies better than one cycle in the roughly $10^5$ cycles that are expected to accumulate during the last year of inspiral \cite{Gair:CQG21}. These high-accuracy templates are expected to be sufficiently accurate that the masses can be determined to about one part in $10^4$ \cite{BarackCutler:PRD69} and will require modeling the motion of the SCO with fractional accuracy of $10^{-5}$ or better. The SCO's motion and corresponding gravitational wave emission may be described using a perturbative treatment since the mass ratio for EMRIs is very small, between about $10^{-4}$ and $10^{-7}$ for EMRIs in LISA's detectable bandwidth. Gravitational waves from EMRIs carry energy, linear momentum and angular momentum that results in a force on the compact object, called the self force, which causes the SCO to inspiral toward the supermassive black hole. The self force on the SCO is conceptually different than, for example, the radiation reaction on a point charge in flat spacetime. Radiation reaction is local in time and is typically proportional to the time-derivative of the charge's acceleration. However, the self force also accounts for the history-dependent force arising from the interactions of the SCO with waves emitted in the past that have backscattered off the background spacetime curvature. As such, self force is nonlocal in time and depends on the past motion of the SCO. The self force can be computed perturbatively in powers of the mass ratio or, more generally, in powers of the size of the SCO to the curvature scale of the background spacetime it moves on. To develop templates with sufficient accuracy to achieve LISA's science goals necessitates corrections through at least second order in the mass ratio. We motivate this last statement with a more detailed discussion below. In this paper, we take first steps towards building high-accuracy EMRI waveforms by deriving the expressions for the self force on a SCO due to the emission of scalar perturbations. Specifically, we calculate the self force through third order in the mass ratio within a class of nonlinear scalar models that is constructed to be analogous to the kinematical structure of the perturbative General Relativistic description of EMRIs. \subsection{High-accuracy EMRI waveforms} Current theoretical techniques for generating high-accuracy waveforms are significantly underdeveloped when compared with those methods producing less accurate waveform templates, which are more useful for detecting EMRIs (for a review see \cite{Drasco:CQG23, Vallisneri:CQG26}). The former waveforms include ``Capra'' waveforms (see e.g., \cite{Drasco:CQG23}) and ``two-timescale'' waveforms \cite{Mino:PRD67, HindererFlanagan:PRD78}. Capra waveforms constitute the highest standard of accurate source modeling for EMRIs since the underlying calculations are based on a minimum number of assumptions and are thought to represent the binary's evolution and gravitational wave emission most accurately. These waveforms are sourced by the solutions to the self force equations of motion describing the perturbed motion of the SCO in the background supermassive black hole spacetime. The self force equation through first-order in the (very small) mass ratio is a complicated integro-differential equation, called the MiSaTaQuWa equation \cite{MinoSasakiTanaka:PRD55, QuinnWald:PRD56}, for the SCO's worldline coordinates. Accordingly, not a single Capra waveform has been computed despite recent advances and progress with numerical computations from several research groups; see Ref.\,\cite{Barack:CQG26} for a recent review. It was argued in \cite{Burko:PRD67, Rosenthal:PRD73} that if only the first-order (dissipative part of the) self force drives the quasi-circular inspiral of a SCO of mass $m$ moving in a Schwarzschild spacetime with mass $M$, then the accumulated phase of the gravitational waveform over an inspiral time $\sim M/\epsilon$, where $\epsilon = m/M \ll 1$, is, schematically, \begin{align} \Phi \sim \frac{1}{\epsilon} + O(\epsilon^0) ~. \label{phase1} \end{align} The last term represents the error in the phase from not including second and higher-order self force effects, which represents an $O(1)$ correction. Thus, to produce waveforms accurate to less than a cycle requires that second-order self force corrections be included in determining the SCO's motion. Two-timescale waveforms are based on a systematic adiabatic expansion in which the typical orbital period $T_{\rm orb}$ is small compared to the inspiral timescale $T_{\rm insp}$. To leading order in $T_{\rm orb}/T_{\rm insp}$, Mino \cite{Mino:PRD67} showed that the waveform phase depends only on the time average of the dissipative part of the first-order self force. Hinderer and Flanagan \cite{HindererFlanagan:PRD78} extended Mino's work by placing it within a systematic two-timescale expansion to calculate ``post-adiabatic'' (PA) corrections to Mino's result. They found at 1PA that the time averaged dissipative part of the second-order self force is just as important as fluctuations in the conservative part of the first-order self force. Therefore, second-order self force corrections are important to maintain the consistency of the inspiral's adiabatic evolution. \subsection{Transient resonances} Recently, Flanagan and Hinderer \cite{Flanagan:2010cd} discovered that the SCO may undergo transient resonances during the course of its inspiral. Resonances occur only when the SCO evolves on a non-equatorial and eccentric orbit in a spinning supermassive black hole spacetime. Such orbital configurations are expected to be generic for EMRI sources detectable by LISA. During a transient resonance, the frequency of the true gravitational wave signal undergoes sudden jumps and an adiabatic treatment of the inspiral breaks down. Thus, the corresponding signal-to-noise ratio may be significantly diminished if using a template bank of two-timescale waveforms. One way to address this problem is to provide increasingly accurate waveforms that are capable of tracking the phase evolution of the SCO even through the transient resonances. There are at least two ways to do this. One may patch a kludge waveform \cite{GairGlampedakis:PRD73, Babaketal:PRD75} (describing the resonant phases) into a very accurate two-timescale waveform (describing the system at all other times, during the adiabatic inspiral phases) or one may develop Capra waveforms that incorporate higher-order self force corrections to the SCO's motion. The former approach is a reasonable possibility since the durations of the transient resonances are short compared to the radiation reaction time scale but will be only as accurate as the kludge model used. The latter method is the most direct and accurate way but may be difficult to realize given that even first order Capra inspiral waveforms have not yet been computed. Nevertheless, in the presence of resonances the accumulated phase of the waveform is, schematically, \begin{align} \Phi \sim \frac{1}{\varepsilon} + \frac{ 1 }{ \sqrt{\varepsilon} } + O ( \varepsilon^0 ) \label{phase2} \end{align} and the $O(\varepsilon^{-1/2})$ contribution, which originates from passing through a resonance, requires knowing a part of the self force at second order in $\varepsilon$ \cite{Flanagan:2010cd}. Therefore, for EMRIs that pass through at least one resonance, the second and possibly higher-order contributions to the self force are especially important. Estimates of the effects on the inspiral waveform phase from passing through a resonance indicate that the phase may change by $\sim 20$ rad for an EMRI with a mass ratio of $10^{-6}$ \cite{Flanagan:2010cd}. Not only is this a significant change for one resonant crossing but the effect accumulates for each resonance encountered during the inspiral. As a result, even detecting EMRIs with LISA could be affected by the dephasing from transient resonances. Thus, the motion of the SCO is needed with (possibly very) high accuracy in order to sufficiently describe its evolution before, during and after each transient resonance. Based on these previous works, there are several indications suggesting a need to model EMRI sources with high accuracy, which will require incorporating the effects on the SCO from second (and possibly higher) order self force corrections. \subsection{Intermediate mass ratio inspirals} Higher order self force corrections may also be needed to model binaries with less extreme mass ratios and could be useful even for those with comparable mass ratios. It is natural to think that including higher-order self force corrections in the SCO equations of motion will allow for the mass ratio to be relaxed to higher values. Doing so may offer the only way of describing binaries with intermediate mass ratios (IMRs), which have mass ratios in the range of $\sim 10^{-1} - 10^{-4}$, since neither the post-Newtonian approximation nor numerical relativity are particularly good tools for studying the inspirals of IMRs (however, see the recent work of \cite{Lousto:2010ut}). In fact, an alternative approach based on self force methods for binaries with IMRs may be useful for calibrating semi-analytical models (e.g., Effective One Body \cite{BuonannoDamour:PRD59}) and phenomenological hybrid waveforms \cite{Ajithetal:CQG24}. \subsection{Self force in scalar models} In this paper, we introduce a class of nonlinear scalar theories that is developed with a structure very similar to the perturbative General Relativistic description of EMRIs and will serve as a scalar analog of these sources. We calculate the finite (or regular) part of the scalar self force on the SCO through {\it third} order in the ratio of the size of the SCO -- $R_m$ -- to the background curvature length scale -- ${\cal R}$ -- and denoted by $\varepsilon = R_m / {\cal R}$, which is just the mass ratio in the strong-field regime of a supermassive black hole. We take the background to be specified for all time and, for simplicity, do not include the effects from the scalar field's stress-energy on the spacetime so that the scalar field is analogous to the propagation of metric perturbations on a fixed background spacetime. In regularizing the formally divergent self force expressions we use the standard tools of renormalization borrowed from the fields of high energy physics and condensed matter. The methods used in this paper will be of direct use for calculating the second order gravitational self force for EMRIs in a future paper. Historically, scalar models offer a simpler framework for studying the underlying issues of self force regularization and for developing practical self force computational schemes. Indeed, the most useful regularization scheme (both in terms of physical intuition and practical computations) was first developed and understood in the context of the self force on a scalar charge from a linear scalar field in \cite{DetweilerWhiting:PRD67}. The first numerical computation of the self force (which was performed for a circular geodesic in Schwarzschild spacetime) was accomplished in a linear scalar theory in \cite{Burko:PRL84} and predated, by about seven years, the corresponding computation in the gravitational case \cite{BarackSago:PRD75}. Because of the relative simplicities that scalar models afford, it seems likely that the physics of higher-order self force effects can be investigated more easily and quickly than in the gravitational EMRI context. Specifically, one can address the qualitative and quantitative effect that higher-order self force corrections have on the waveforms themselves (i.e., the change in the phase compared to first-order accurate waveforms, parameter estimation, etc.); how transient resonances and the number of resonances encountered during the inspiral affect the waveform and the SCO's motion; and how much the mass ratio can be relaxed to higher values while still having a reasonably accurate description of the system (the accuracy can be addressed using the next higher order to bound or estimate the errors). It is likely that these questions can begin to be studied using the results from this paper and with small modifications to some self force codes currently in use (particularly those using $3+1$ methods as in \cite{Vegaetal:PRD80} since these do not rely on a mode decomposition of the field). \subsection{Previous work in higher-order self force corrections} The first work (that we are aware of) regarding higher-order self force computations was carried out by Burko in Ref.\,\cite{Burko:PRD67}. He computed the second-order self force on a scalar charged particle in quasi-circular orbit in a Schwarzschild background. Burko derived a formal expression for the equations of motion on the particle for this scenario. The force on the particle included contributions from the second-order self force and also from the product of two first-order pieces. Since the second-order expressions were unknown, Burko retained only the latter contributions and estimated the change in the accumulated phase due to these self force corrections when compared to the first-order accurate waveform phase. He found that there was a relevant correction of $O(1)$ cycles to the first-order accurate phase. The change to the phase from amplitude corrections was found to be about a tenth of a cycle and hence irrelevant. Later, Rosenthal developed a rigorous program to regularize scalar \cite{Rosenthal:CQG22} and gravitational \cite{Rosenthal:PRD72, Rosenthal:PRD73} perturbations through second order in perturbation theory. Using these regularized perturbations, he derived formal expressions for the second-order gravitational self force in Ref.\,\cite{Rosenthal:PRD74} in a gauge different from the standard Lorenz gauge, which he called the Fermi gauge \cite{Rosenthal:PRD72}. The Fermi gauge is a coordinate system in which the regular part of the field perturbation and its covariant derivative are made to vanish on the worldline. As such, the first-order self-force corrections vanish in this gauge and the background spacetime is composed of the original background plus the first-order metric perturbations generated by the motion of the SCO. Unfortunately, Rosenthal's approach does not seem practical for numerical self force computations for EMRIs since constructing the Fermi gauge requires the first-order self force on the SCO to be made to vanish, which is accomplished by integrating the MiSaTaQuWa equation -- a feat that has yet to be accomplished. Recently\footnote{Private communication with S.\,Detweiler.}, Detweiler \cite{Detweiler:CQGcapra} has shown with matched asymptotic expansions that the motion of the SCO can be described by a geodesic in a perturbed background spacetime having a smooth metric at the location of the particle through second order in $\varepsilon$. Detweiler's results imply that one may continue to interpret, at higher orders in $\varepsilon$, the SCO's motion as either being perturbed by self force corrections on a fixed background or as being geodesic in a perturbed spacetime. \subsection{Organization} This paper is organized as follows. In Section \ref{sec:eft} we give a brief overview of the effective field theory (EFT) approach as it applies to EMRIs. The EFT framework provides an efficient way to systematically calculate the self force at higher orders in $\varepsilon$. In Section \ref{sec:actionprinciple} we discuss how to consistently implement outgoing boundary conditions for the field within a variational principle. In Section \ref{sec:linearsf} we demonstrate how to isolate and evaluate the singular part of the well-known linear scalar self force on a charge in a non-vacuum spacetime so that we may compare the result from our formalism and methods to the standard result in Ref.\,\cite{Quinn:PRD62}. In Section \ref{sec:nonlinear} we develop a class of nonlinear scalar models that is designed to have a structure analogous to the perturbation theory used to describe EMRIs in General Relativity. In Section \ref{sec:nonlinearsf} we then compute the formally divergent self force expressions in the nonlinear scalar model through third order in $\varepsilon$. In Section \ref{sec:renormalization} we renormalize these expressions by introducing counter terms into the action to cancel those divergences. In Section \ref{sec:eom} we write down the finite, third-order self force equations of motion. In Section \ref{sec:conclusion} we conclude with a discussion. The Appendices are devoted to deriving the quasilocal expansions used in regularizing the self force expressions in Section \ref{sec:nonlinearsf}, to proving that power-divergent integrals vanish in dimensional regularization and to listing the Feynman rules for the nonlinear scalar model introduced in Section \ref{sec:nonlinear}. The regular part of the self force in the nonlinear scalar model valid through $O(\varepsilon^3)$ is given in (\ref{renormalizedsf1}) and (\ref{renormalizedsf3}) using the Detweiler-Whiting decomposition \cite{DetweilerWhiting:PRD67} for the retarded Green's function and in (\ref{renormalizedsf4}) and (\ref{effmass4}) using the Hadamard decomposition \cite{Hadamard}. Using the Detweiler-Whiting decomposition we find evidence suggesting that the self force through $O(\varepsilon^3)$ can be written solely in terms of the regular part of the field and its derivatives when evaluated on the worldline. In a later paper in this series we will explicitly show that this is indeed the case by calculating the radiative scalar perturbations and computing their effect on the SCO's motion \cite{Galley:Nonlinear2}. We use units where $c=1$ and define the gravitational constant $G$ in terms of a mass parameter ($m_{pl}$) to be $32\pi G \equiv m_{pl}^{-2}$. The metric signature is $(-,+,+,+)$. We frequently use the notation where the worldline coordinates at proper times $\tau$ and $\tau'$ are denoted by the shorthand $z^\mu$ and $z^{\mu'}$, respectively, so that $z^\mu = z^\mu (\tau)$ and $z^{\mu'} = z^\mu (\tau')$. The same goes for tensors evaluated on the worldline at some proper time. We also use a mixed notation where a quantity such as $V(x; z]$ indicates that $V$ is a function of $x^\mu$ but is a functional of $z^\mu (\tau)$.

\label{sec:conclusion} In this paper we introduced a nonlinear scalar model for extreme mass ratio inspirals that is a natural analogue of the corresponding perturbative General Relativistic description. This model should be useful for studying the role of higher-order self force corrections for building high-accuracy waveforms for precise parameter estimation, for quantifying the effect of transient resonances on the phase evolution of waveforms, and for providing a sufficiently simple context to develop numerical methods for computations of higher-order self force corrections that may then be applied to the gravitational problem. This last program should be particularly useful for calibrating semi-analytic models \cite{BuonannoDamour:PRD59}, for building phenomenological-based hybrid waveforms \cite{Ajithetal:CQG24}, and for constructing gauge-invariant EMRI observables at higher orders in $\varepsilon$. This model has the interesting feature that, despite being constructed to be as similar as possible to perturbed General Relativity, one can always perform a field redefinition using (\ref{psi1}) so that the initially nonlinear scalar field theory is transformed to a linear one (with suitable changes in the field-particle interactions). Such a transformation provides a very clean and simple derivation of the self-force corrections at higher orders since all contributions coming from the self-interaction of the field are subsequently removed. It is natural then to wonder if such a transformation can be made when describing gravitational EMRIs to remove some, if not all, of the nonlinear self-interactions of the metric perturbations. Indeed, it may be that a combination of a field redefinition and a gauge transformation (one that is $O(\varepsilon^2)$ so as to keep the linearized metric perturbations in the Lorenz gauge at leading order) successfully removes all nonlinear self-interaction terms at a given order. If this is the case then calculating higher order self force corrections will be much easier to derive analytically and solve numerically. We will discuss these issues further in a future paper. In order to calculate the self force on the SCO we demonstrated how the usual action principle fails to retain the outgoing boundary conditions when integrating out the field -- a conservative dynamics for the worldline ensues because the full Lagrangian is time-reversal invariant. However, using the difference in the actions for two worldline and field histories (one evolving forward in time and the other backward) allows one to ``break'' the time-reversal symmetry naturally contained in the usual action principle and incorporate the time-asymmetric outgoing boundary conditions into the effective action. In this way, one can describe the open classical system dynamics of the SCO from an action principle and derive the self force (and waveforms -- see the next paper in this series) in a self-consistent manner. Using this new action principle for open classical systems together with the effective field theory formalism we derived the self force on the SCO through {\it third} order in $\varepsilon$. Our main results are given in (\ref{renormalizedsf1}) and (\ref{renormalizedsf3}). We separated the singular parts of the self force from those that are regular on the worldline using the Detweiler-Whiting decomposition for the retarded Green's function. For completeness and comparison, we also presented the more complicated third-order self force expressions in terms of Hadamard's tail propagator in (\ref{renormalizedsf4}) and (\ref{effmass4}). Despite the appearance of higher order poles and history-dependent divergent forces -- see (\ref{regselfforce1}) -- we found that only three {\it local} and {\it time-independent} counter terms are needed to absorb all of these divergences, which amounts to a renormalization of the SCO mass $m$ and the coupling constants $c_1$ and $c_2$ appearing in (\ref{nonlinear2}). In addition, we showed that this nonlinear scalar model only generates power-divergent integrals that automatically vanish in a regulariziation scheme that manifestly preserves the symmetries of the theory (e.g., dimensional regularization). No logarithmic divergences, which represent physical effects of screening by the scalar perturbations, can be generated in this model and thus none of the coupling constants appearing in (\ref{nonlinear2}) or, equivalently, (\ref{nonlinear1}) undergo a classical renormalization group running. Our results in (\ref{renormalizedsf3}) suggest an ambiguity in defining the effective mass of the SCO since, in this model, the effective mass is proportional to the SCO's inertial mass ($m_{\rm eff} = \Gamma_{\rm DW} m$) and one can associate the proportionality factor $\Gamma_{\rm DW} (z^\mu)$ with a contribution to the effective mass or with an additional contribution to the self force. Both interpretations yield the same acceleration on the SCO and the same waveform and physical observables. In the next paper in this series we shall derive the expressions for the scalar perturbations emitted by the SCO through third order in $\varepsilon$. We shall also show that one can more easily derive the self force by calculating the regular part of the radiative field and then evaluating it in the worldline equations of motion (see the discussion below (\ref{renormalizedsf3})). This suggests that current 3+1 self force codes \cite{Vegaetal:PRD80} may need only minimal modifications to compute higher-order self force corrections. In regularizing the third-order waveforms (which also contain divergences since the field is sourced by an effectively point-like object) we will show that exactly the same counter terms used in regularizing the third-order self force expressions -- see (\ref{counterterms10})-(\ref{counterterms12}) -- also regularize the scalar perturbations and the self force expressions derived from them. Therefore, the use of point particle treatments for the SCO in conjunction with our renormalization program will be explicitly self-consistent.

1012.4488

The motion of a small compact object in a background spacetime is investigated in the context of a model nonlinear scalar field theory. This model is constructed to have a perturbative structure analogous to the general relativistic description of extreme mass ratio inspirals (EMRIs). We apply the effective field theory approach to this model and calculate the finite part of the self-force on the small compact object through third order in the ratio of the size of the compact object to the curvature scale of the background (e.g. black hole) spacetime. We use well-known renormalization methods and demonstrate the consistency of the formalism in rendering the self-force finite at higher orders within a point particle prescription for the small compact object. This nonlinear scalar model should be useful for studying various aspects of higher-order self-force effects in EMRIs but within a comparatively simpler context than the full gravitational case. These aspects include developing practical schemes for higher-order self-force numerical computations, quantifying the effects of transient resonances on EMRI waveforms and accurately modeling the small compact object’s motion for precise determinations of the parameters of detected EMRI sources.

12235282

2012ApJ...749..119B

1012

1012.1787_arXiv.txt

The flux variability of very high energy (VHE) gamma-rays on minute timescales detected from the BL Lac object \pks{} \citep{ah07pks} and Mkr 501 \citep{mkr501_magic} challenges the standard scenarios suggested for the explanation of the nonthermal properties of TeV blazars \citep{bfr08,gub09}. The extremely short duration of the flares impose severe constraints on the size of the gamma-ray producing region, of \be l'\le c \tau'\simeq 3\times10^{13}\tau'_{3}\,\rm cm, \label{eq_size_prop} \ee where $l'$ and $\tau'_{3}=\tau'/10^3 $~s are the proper production size and the variability timescale in the frame of the jet respectively and $c$ is the light speed. The proper variability time-scale, $\tau'$, is connected to the variability in the observer frame, $\tau$, by the relation \be \tau= \frac{\tau'}{ \delta} \, , \label{eq_var_tran} \ee where $\delta$ is the Doppler factor of the moving source (the blob): \begin{equation} \delta=\frac{1}{\Gamma\left(1-\beta\cos(\alpha)\right)} \ . \label{db} \end{equation} Here the bulk Lorentz factor, $\Gamma$, accounts for the relativistic transformation of time, and $\left(1-\beta\cos(\alpha)\right)$ is responsible for the kinematic shrinking of the duration of the radiation and $\beta=v/c$. Relativistic jets ejected from the central engines are common phenomena for different types of active galactic nuclei (AGN). In particular, apparent superluminal speeds ${\beta_{\rm app}=\beta\sin{\alpha}/(1-\beta\cos{\alpha})}$ (in units of the speed of light~$c$) as high as $\sim 40$ have been detected for radio components on (projected) scales of $\sim 1-10\;$pc \citep[see e.g.][]{J01,J05,mjl10} in blazars - AGN with jets directed to the observer. This implies very large Lorenz factors of bulk motion given that $\Gamma \geq \beta_{\rm app}$. Because of the large bulk Lorenz factors of TeV blazars, the condition Equation~(\ref{eq_var_tran}) allows significant relaxation of the requirement supplied by Equation~(\ref{eq_size_prop}). In particular, \citet{lb08} argued that if a perturbation is produced by the central engine, its size should exceed the gravitational radius of the black hole in the observer frame. Consequently, the proper size of the production region is expected to be larger than $\Gamma r_ {\rm g}$, where $r_{\rm g}=GM_{\rm BH}/c^2 \approx 1.5 \times 10^{13} M_{\rm BH,8}$ cm is the gravitational radius of a black hole (BH) of mass $M_{\rm BH,8}=M_{\rm BH}/10^8 M_\odot$. In this case, the variability time-scale $\tau_2=\tau/10^2$s imposes a strict upper limit on the gravitational radius: \be r_{\rm g}<3\times10^{12} \tau_2 \frac{\delta}{ \Gamma}\, \rm cm \ . \label{rg} \ee Thus, the detection of gamma-ray flux variability $\tau \sim 200$~s constrains the BH mass, to be $M_{\rm BH,8} <1$. In reality, since the main energy release occurs in the inner parts of the accretion disk of radius $r \sim 10 r_{\rm g}$, the mass of the black hole should be close to $10^7 M_\odot$. Generally, in the case of an extreme Kerr black hole, the energy release takes place at the gravitational radius. However, even in this case one needs an entire rotation period for an effective energy release, i.e. the characteristic time cannot be much shorter than $2 \pi r_{\rm g}/c$. This implies that even in this case the upper bound of the BH mass of $10^7 M_\odot$ cannot be significantly relaxed. This conclusion is true, in particular, for the model of internal shocks. Note, however, that it is based on the assumption that the perturbation (the reason of the flare) originates in the central engine. Therefore it cannot be unconditionally extended to other possible scenarios as is claimed by \citet{dfkb09}. Indeed, Equation~(\ref{rg}) is not valid if perturbations are produced by an external source, e.g. by plasma condensations (often called "blobs") which do not have a direct link to the central black hole. Such blobs can be produced, in particular, by interactions of stars with the base of the jet, as proposed by \citet{bab10} to explain the TeV flares of M~87 on scales of days \citep[see also][]{abr10}, where the interaction of gas clouds from broad line regions are discussed). The jet power of M~87 is relatively modest, $L_{\rm j} \simeq 10^{44}\,\rm erg\,s^{-1}$. The results of \citet{bab10} show that, while this power is sufficient to blow-up the envelope (atmosphere) of the star which initially has been pulled out by the tidal force of the BH, such a jet appears to be not sufficiently powerful for acceleration of the gas cloud to relativistic velocities. Actually this works in a positive direction for M~87, given the large aspect angle of the jet. Otherwise, the gamma-ray flux could not be observed because of the Doppler de-boosting effect. On the other hand, the suggested mechanism of formation of hadronic blobs in the jets cannot apply to powerful gamma-ray blazars unless the blobs are accelerated to Lorentz factors $\Gamma \geq 10$. Remarkably, this can be realized in a quite natural way in powerful jets with $L_{\rm j} \geq 10^{46}\,\rm erg\,s^{-1}$. Interestingly, such powerful jets can ablate the star atmosphere without help of tidal forces (the interesting implications of this effect are discussed below). Moreover, the powerful jets can drag and disrupt the star's envelope into an ensemble of blobs moving with large Lorentz factors which, from the point of view of explanation of very short time variability is a quite comfortable situation. Another important aspect of the short time variability is related to the efficiency of acceleration and radiation mechanisms. Currently the most conventional approach for modeling of VHE emission production in active galaxies is based on the inverse Compton (IC) scattering of relativistic electrons, the soft target photon field being either the synchrotron radiation of same electrons (the so-called Synchrotron-self Compton (SSC) model), or provided by external sources (EIC model). The apparent advantage of IC models is the combination of two factors: (1) the acceleration of electrons to relatively modest energies ($\leq 1$~TeV) can be effectively realized within different acceleration scenarios; and (2) these electrons radiate readily in interactions of ambient radiation and magnetic fields. Nevertheless, while the IC models allow rather satisfactory explanations of the energy spectra and variability patterns of many blazars in general, the parameters used to fit some specific objects appear incompatible with the parameters defined from observations. Moreover, the observed short variability time-scale demands conditions which appear to be quite uncomfortable, in terms of the strength of magnetic field and related consequences concerning the strong deviation from equipartition between the energy density of relativistic electrons and magnetic fields. The requirement of weak, (generally less than 1 G) magnetic fields is one of the key postulates of the IC paradigm of gamma-ray production in blazars. Moreover, in the case of some objects with unusually hard source spectra (after correction for intergalactic absorption), such as 1ES~0229+200, the magnetic field is required to be as small as 1mG \citep{tgg09}. The magnetic field in the blazar jets can be reduced to such small values only at very large distances from the central engine, namely $\gtrsim 10^{18} \ \rm cm$. Although this idea has some observational support related to the transparency of blobs in the radio band \citep{mjd08}, it is likely that regions of highly variable gamma-ray emission are located at smaller distances from the central engine \citep{tgb10}. In particular, the EIC models require the location of any gamma-ray emitter to be located closer to the BH into the so-called Broad Line Regions (BLR), i.e typically at distances $R \sim 10^{17} - 10^{18}$~cm. This implies that any IC model can be realized only if one finds a way dramatically reduce the magnetic field in the jet. Although this cannot be excluded (e.g. because of reconnection of the B-field \citep{kbl07,gub09,mu10} or due to the effective bulk acceleration of plasma \citep{kvk10,tnm10}), strong magnetic fields exceeding 1~G remain a more favored option, as long as we deal with strong jets on sub-parsec scales. In this regard, the models which invoke high energy protons for the production of gamma-rays passes certain advantages despite a quite popular view that they are not effective emitters \citep[see e.g.][]{s10psyn}. Actually this is true only for proton-proton and proton-photon interactions. What concerns the synchrotron radiation of protons, with a key assumption on the acceleration of particles with a rate close to $t^{-1} \sim ceB/E$, where $E$ is the proton energy, coupled with a strong magnetic field between 10 to 100 Gauss, and a large Doppler factor, $\delta \geq 10$, it is that can provide relevant acceleration and radiation timescales, as well as explain the extension of gamma-ray spectra to TeV energies \cite[see e.g.][]{ah00}.

The ultra-short TeV gamma-ray flares of blazars detected in the case of \pks{} \citep{ah07pks} and Mkr 501 \citep{mkr501_magic} on 100~s timescales represent a serious challenge for current models of blazars. This challenge concerns the origin and the sites of formation of these flares, the acceleration and radiation mechanisms, the hydrodynamics of relativistic outflows, amongst others. Since the upper limit on the size of production region, of $3 \times 10^{12} \tau_2 \ \rm cm$, is smaller by an order of magnitude than the gravitational radius of a black hole of mass $10^{8} M_\odot$ (which is required to power distant blazars), the only way to avoid the situation of invoking quite uncomfortable upper limits on the mass of the central black hole (as small as $10^7 M_\odot$), is to invoke the Doppler boosting. However, this can be realized only in the case of an external origin of the processes which cause these ultra-short flares. If the flares are initiated by disturbances originating from the central black hole (e.g., due to internal shocks), the linear size (in the observer's frame) of the flare production region cannot be smaller than the gravitational radius of the black hole, independent of the Doppler factor of the jet. The model suggested in this work readily solves the problem of connecting the flares to the interactions of the red giant stars with the powerful jets. Due to these interactions the red giant loses a significant fraction its atmosphere. The cloud, accelerated by the magnetically driven jet up to a Lorentz factor of $\Gamma \sim 30$, likely separates into many small fragments. These ``blobs'' represent the ideal sites for the production of flares, provided that a significant fraction of jet energy absorbed by the cloud is converted (e.g. due to relativistic shock acceleration or magnetic reconnection) to relativistic particles. The effective acceleration of particles is a necessary, but not a sufficient condition for the interpretation of the gamma-ray radiation of blazars. Any successful model of TeV blazars require adequate cooling times through gamma-radiation; they should be comparable, or often even shorter compared to the characteristic timescales of other radiative and non-radiative processes. Generally, leptonic models of gamma-ray loud blazars, through the realization of SSC or external IC scenarios, do provide adequate radiation timescales, but at the expense of the assumption of a rather weak magnetic field, typically less than 1 G, which in powerful blazars ($L_{\rm j} \geq 10^{46} \mbox{ erg s}^{-1}$), is well below the magnetic field in the jet as long as it's concerned with sub-parsec distances (see Figure~\ref{magjet}). This is a quite challenging requirement of the discussion of feasibility of, which is generally ignored in the literature. In the JRGI scenario suggested here the problem can be formally solved assuming that the magnetic field inside the blob is much smaller than in the jet. However, in the case of the SSC models, this assumption still does not allow a relaxation of the second requirement of an extremely large jet Lorentz factor, $\Gamma \geq 1000$. Although such Lorentz factors for the bulk motion cannot be excluded\footnote{ We should note that it is rather difficult to reach such a high value of the bulk Lorentz factor, e.g. due to the so called ``photon breeding mechanism'' \citep{sp06}, which does not allow AGN jets with bulk Lorentz factors exceeding 50.}, in particular at large, $\geq 1$pc distances from the BH (see Figure~\ref{magjet}), in the proposed JRGI model it hardly can work. At such distance the jet ram pressure is not sufficient able to ablate the atmosphere of the star. The requirements on the magnetic field and the jet Lorentz factor are more relaxed in the external IC model. Nevertheless, one should note that within the JRGI scenario, the external Compton model has some specific features. In order to avoid severe gamma-gamma absorption, the Compton scattering should proceed in the Thomson regime. This can be fulfilled if the radiation region is located at large distances, i.e. regions still with quite a large Lorenz factors for the jet, $\Gamma_{\rm j} \sim 100$. One of the main postulates of the JRGI scenario is the effective star-jet interaction. This requires the location of the blobs that emit gamma-rays to be at small distances from the BH, typically $z \sim 10^{17}$~cm. In the case of powerful jets, $L_{\rm j} \geq 10^{47} \mbox{ erg s}^{-1}$, this implies a very large magnetic field, $B \sim 100$~G and a moderate Lorenz factor, $\Gamma_{\rm j} \sim 20$. Both parameters match nicely with the interpretation of the TeV gamma-ray flares as a result of proton-synchrotron radiation by highly magnetized blobs, formed and accelerated in jet-star interactions. This model demands the acceleration of protons to energies of $10^{19}$eV, and implies the acceleration of protons with a rate close to the maximum (theoretically possible) rate of, $t_{\rm acc} \sim r_{\rm L}/c$. This is quite a robust requirement, which however, can be provided, in principle, by certain acceleration mechanisms. Another challenge of the proposed scenario is related to the power of the jet. Namely, the proton synchrotron model of TeV gamma-rays can be effective, provided that: (i) the mass of BH does not significantly exceed $M \sim 10^8 M_{\odot}$ and; (ii) the jet power is not significantly below $10^{47} \mbox{ erg s}^{-1}$. An unambiguous implication of these two requirements (working in two different directions) is that the jet should have a super-Eddington luminosity. Although this could seem like quite a dramatic assumption, we note the requirement of super-Eddington luminosities seems to be an unavoidable, model-independent conclusion for GRBs and also likely for powerful gamma-ray blazars \citep{g11texas}. \subsection{Stellar density in the vicinity of a SMBH.} An important question in the suggested scenario is the expected rate of the flaring events, which is related to the number density of RGs at the relevant jet scales. The jet region suitable for the production of the powerful flares (similar to the burst detected from \pks), can be defined as $z<1\rm pc$, and the corresponding side cross section of the jet is $S_{\rm j}\approx z^2 \theta\sim 10^{33}\theta_{-1} z^2_{17}\rm cm^2$. Thus, the number of flaring events per year can be estimated as $\Upsilon \approx S_{\rm j} V_{\rm orb} n$. Equation~(\ref{sv}) provides an estimate for the density of RGs required to produce $\Upsilon$ flaring events per year: \begin{equation} n\sim10^6\Upsilon M_{\rm BH,8}^{-1/2}\theta^{-1}_{-1}z^{-3/2}_{17}\rm pc^{-3}\,. \label{eq:density_RG} \end{equation} Unfortunately, there are no direct measurements of the stellar density in the vicinity of BHs. Thus, depending on the assumed distribution law, the number of RGs in the vicinity of the BH may or may not be sufficient. However, we note that studies of possible stellar density profiles in the vicinity of the BH in AGNs (see e.g. \citep{bkck82,mcd91}) show that densities similar to the one required ($\sim 10^6$~pc$^{-3}$) are rather feasible. Moreover, under the influence of X-ray radiation close to BHs, normal stars can drastically increase the rate of their stellar wind. Thus, wind-fed clouds can be formed during the jet -- star interaction. This interaction can mimic the interaction of RG atmosphere with the jet (Dorodnitsyn, private communications). Since, the stellar density of normal stars is higher up to 2 orders of magnitude than the density of RGs, this effect can significantly relax the requirement imposed by Equation~(\ref{eq:density_RG}) on the stellar density in the vicinity of BHs. \appendix

1012.1787

We propose a new model for the description of ultra-short flares from TeV blazars by compact magnetized condensations (blobs), produced when red giant stars cross the jet close to the central black hole. Our study includes a simple dynamic model for the evolution of the envelope lost by the star in the jet and its high-energy nonthermal emission through different leptonic and hadronic radiation mechanisms. We show that the fragmented envelope of the star can be accelerated to Lorentz factors up to 100 and effectively radiate the available energy in gamma rays predominantly through proton synchrotron radiation or external inverse Compton scattering of electrons. The model can readily explain the minute-scale TeV flares on top of longer (typical timescales of days) gamma-ray variability as observed from the blazar PKS 2155-304. In the framework of the proposed scenario, the key parameters of the source are robustly constrained. In the case of proton synchrotron origin of the emission, a mass of the central black hole of M <SUB>BH</SUB> ≈ 10<SUP>8</SUP> M <SUB>⊙</SUB>, a total jet power of L <SUB>j</SUB> ≈ 2 × 10<SUP>47</SUP> erg s<SUP>-1</SUP>, and a Doppler factor of the gamma-ray emitting blobs of δ &gt;= 40 are required. For the external inverse Compton model, parameters of M <SUB>BH</SUB> ≈ 10<SUP>8</SUP> M <SUB>⊙</SUB>, L <SUB>j</SUB> ≈ 10<SUP>46</SUP> erg s<SUP>-1</SUP>, and δ &gt;= 150 are required.

12205464

2011Icar..211..876C

1012

1012.0324_arXiv.txt

The origin of fine-grained dust rims surrounding chondrules in carbonaceous chondrites is still debated. Analyses of CM chondrite thin sections by optical and scanning electron microscopy (Metzler et al. 1992), as well as detailed theoretical models (Morfill et al. 1998, henceforth MDT; Cuzzi 2004; Ormel et al. 2008, henceforth OCT) suggest that these dust mantles could very well have formed in the gaseous environment of the primitive solar nebula, as a result of accretion processes. Nevertheless, advocates of a non-nebular origin of these rims favor their formation on the CM parent asteroids, through the impact and compaction of matrix material around chondrules (Trigo-Rodr\'{\i}guez et al. 2006), followed by aqueous alteration. This interpretation arises mostly because the porosity of observed dust mantles that envelop chondrules is lower than what, presumably, could be produced by agglomeration of dust in the solar nebula. The non-nebular scenario still requires further quantitative analysis. For example, numerical techniques similar to those employed in the study of collisions between Kuiper belt objects (Leinhardt and Stewart 2009) could be used to measure shock stresses in asteroidal collisions. Theoretical modeling of the formation of fine-grained dust rims assumes that two populations, chondrules and dust, both free floating in the gas of the primitive solar nebula, come into contact. MDT put forward the notion that sweep-up of dust by chondrules in a confined volume of the nebula, prior to the formation of meteorite parent bodies, reproduces the near-linear relation between chondrule radius and rim thickness (e.g. Metzler et al. 1992, Paque and Cuzzi 1997). MDT do not make specific assumptions about the nature of the gas flow. Cuzzi (2004) goes further and incorporates the role of turbulence in the sweep-up process, showing that the initial relative abundance of chondrules is an important parameter that determines both the rate at which dust is depleted locally, and the rate of growth of rimmed particles. The analytical model of Cuzzi (2004) predicts that chondrule-sized particles will acquire their observed rim volumes in $\sim$10$^{2}$--10$^{3}$ years, subject to the values of the dimensionless turbulent viscosity parameter $\alpha$ (Shakura and Sunyaev 1973, Balbus and Hawley 1998), which is a measure of the stresses in an accretion disk. The precise value of $\alpha$ in protoplanetary disks (including the solar nebula) is uncertain, and authors often consider values ranging from $\sim 10^{-5}$ to 0.1. For disks around T Tauri stars, $\alpha \sim 10^{-4}$--$10^{-2}$ has been inferred (Hartmann et al. 1998). Assuming high values of $\alpha$, the Cuzzi (2004) model yields chondrule rimming times as low as a few years. The effect of parametrized turbulence has been included by OCT in a numerical scheme to study inter-chondrule sticking via dust rims. Their Monte Carlo algorithm contains a detailed treatment of the collisions among porous dust grain aggregates, and their results evidently show a difference in the evolution of rimmed-chondrule size for two values of $\alpha$, $10^{-6}$ and $10^{-4}$. The lower $\alpha$ value allows for rimmed-particle radii growth by a factor of 7, compared to a factor of $\sim$ 2 for $\alpha=10^{-4}$, before the effects of dust compaction set in. This occurs in $\sim 10^{3}$ years in the former case, and in $\sim 10^{2}$ years in the latter case. The final relation between chondrule and rim masses that OCT obtain is, once again, nearly linear. A key ingredient in the above calculations is the collision velocity between solid particles of different sizes. In a protoplanetary disk, the following sources of relative velocities may operate: 1) thermally-induced Brownian motion, important for micron-sized dust grains; 2) drift towards the central star and difference in azimuthal (orbital) velocities (the latter also known as ``transverse'' velocity), both due to a radial pressure gradient in the disk gas (e.g. Weidenschilling 1977); 3) settling towards the disk midplane, as a result of the vertical component of the star's gravity; and 4) turbulence. It was only recently that accurate analytical expressions for turbulent relative velocities were obtained (Ormel and Cuzzi 2007). These formulae show that turbulence can dominate mutual velocities for chondrule- and smaller-size particles in the solar nebula at 3 AU, over systematic radial drift, provided $\alpha \gtrsim 10^{-5}$ (OCT, their Figure 1). Our understanding of the evolution of protoplanetary disks is significantly improved when their vertical structure is taken into account. For example, numerical simulations show that the gas velocity dispersion (Fromang and Papaloizou 2006), viscous stress (Turner et al. 2007), and accretion rate (Turner et al. 2010) depend on height $z$ above the disk midplane. Furthermore, this vertical dependence bears directly on the dynamical behavior of solid bodies. For instance, the increase of gas velocities at high $z$ (where gas densities are lower) produce higher particle relative velocities. While it is recognized that protoplanetary disks must be in a turbulent state, there is no general consensus regarding the mechanism responsible for the production of turbulence. However, workers in the field acknowledge that a magnetic field, coupled to the differential rotation of the disk gas, could be the crucial driver of disk turbulence. The outcome of the ensuing magnetorotational instability (or MRI; Balbus and Hawley 1998) is self-sustained turbulence that supplies the necessary viscous stresses for the disk angular momentum to be transported outwards. Hydrodynamic and magnetic stresses (the latter greater than the former by a factor of a few) both contribute to the turbulent viscosity in a disk. There is a caveat associated with the saturation of the MRI in the solar nebula: gas ionization may be too low, and electrical resistivity too high, in dense regions of the nebula to allow coupling of the gas to the magnetic field. These regions remain nearly laminar and are sandwiched between conducting layers where turbulence does develop, owing to ionization by cosmic rays and stellar X-rays that are absorbed before they reach the midplane. Furthermore, recombination of free electrons on micron-sized dust grains is efficient and reduces the ionization fraction below that necessary to sustain turbulence (Turner et al. 2010). Dust abundance is therefore important in establishing levels of turbulent activity. Notwithstanding the significance of gas resistivity, in this study it is set to zero to focus on the growth of rimmed, chondrule-sized solid particles as they accrete dust in a local neighborhood of the solar nebula, in which all of the gas flow is turbulent due to the MRI and the turbulent intensity varies with time and position. The ideal magnetohydrodynamic (MHD) numerical set-up used allows us not only to determine the final size distribution of dust-mantled objects, but in the process it performs a self-consistent calculation of the dust abundance as a function of time and spatial location, with turbulence and vertical settling as the sources of relative velocities between dust and chondrules. Section 2 presents a semi-analytical argument regarding the importance of turbulence in dust-chondrule dynamics. Section 3 describes the numerical technique used to model the solar nebula and the solids, including the interaction between chondrules and dust. This Section also lists the simulations that were carried out. Section 4 presents the results of the MHD calculations, and a discusion follows in Section 5. A summary and conclusions are given in Section 6.

In this work, accretion of dust by chondrule-sized particles is modeled using a MHD representation of a turbulent solar nebula. MHD processes are a viable mechanism to produce and sustain turbulent viscosity in protoplanetary disks, and the dynamical evolution of the ionized gas in the solar nebula must depend crucially on its ability to couple to different magnetic field configurations. Relative velocities between chondrules and micron-sized dust grains can be excited by several sources, and here we focused on the effect imparted by MRI turbulence and vertical settling due to stellar gravity, in a local neighborhood of the nebula. An important parameter in determining the rate of depletion of dust due to sweep-up by chondrules is the latter's volume density. The MHD simulations show that chondrules accrete dust mostly within $\sim$ 1 scale-height ($H$) of the nebula midplane, since the growth of dust rims decreases aerodynamic coupling of particles to the gas, and increases their response to vertical gravity. Time scales to deplete dust to 1\% of its initial abundance vary from 10 to $\sim$ 800 years, roughly the same times in which chondrules reach their mass-averaged, asymptotic radius. If the sticking efficiency of dust to chondrules is below $10^{-2}$, the growth time scale could be $\sim 10^{3}$ years. Low sticking eficiencies (i.e., very little dust rim growth) can be common due to very high, MRI--driven, turbulent relative velocities between chondrules and dust. Vertical stratification of the nebular gas density allows to compare the dust rimming process at different nebula heights, and this has been done by initially placing one chondrule population at the nebula midplane and another population at a height of 2.3$H$. The growth rate of dust rims in these two cases is very similar, at least for the first 300 years of evolution, with the midplane population having a mass--averaged radius $\sim$ 1 \% larger than the upper layer population during the course of the numerical run. Size distributions of rimmed chondrules obtained from the MHD simulations could be used for comparison with data from actual chondrites. The aerodynamic properties of chondrules suspended in nebular gas are reasonably well understood, and such comparison would provide further insight into the MHD conditions in the primitive solar nebula. In particular, the level of turbulence, and hence the ionization state of the gas, could be constrained. Sweep-up of dust by chondrules is likely to influence the evolution of a laminar dead zone in the solar nebula. Future MHD calculations should incorporate explicit ionization processes (such as those due to stellar X-rays and radionuclide decay) to obtain a more complete, layered flow structure in which chondrule rimming at different locations can be characterized. The vertical variation of turbulent intensity and dust abundance leads to a modest dependence of rimmed-chondrule size on disk height. Whether this variation could be imprinted and identified on the meteoritic record will require further analysis to account for processes such as fragmentation and compaction of dust structures. Compaction of initially fractal dust aggregates can result from collisions, and compact dust rims can play a role in the sticking between rimmed chondrules. Monte Carlo (MC) methods (OCT, Zsom and Dullemond 2008) provide a robust procedure to investigate coagulation of chondrule compounds, since they treat interactions \textit{between} dust--rimmed chondrules, providing an extension to the present work. These codes take as one of their inputs particle--particle relative velocities $\Delta v$, which could be taken directly from a MHD simulation (e.g., by measuring $\Delta v$ between Lagrangian particles located inside an Eulerian grid cell, as in Carballido et al. 2010). Although a MC code does not resolve the spatial structure of the protoplanetary disk, it contains all the physics of particle--particle collisions. On the other hand, the MHD run provides the spatial structure of the nebula plus position--dependent relative velocities, which can be plugged back into the MC code to provide self--consistency to rimmed chondrule--chondrule collisions. Furthermore, if non--ideal MHD effects are allowed for (e.g., a dead zone), the growth of chondrule compounds in both turbulent and nearly laminar regions could be characterized and related to the time-- and space--dependent ionization state of the gas, which may be useful in chemical studies of chondrule compounds. The use of larger shearing boxes now allows to study locations of protoplanetary disks that reach up to 30 scale heights in radial extent (Stone and Gardiner 2010). Thus, in addition to turbulence and vertical settling, the contribution of radial drift to relative velocities between chondrules and dust could be incorporated self-consistently in a MHD-Monte Carlo numerical scheme. The combination of these two techniques is likely to set the nebular origin of dust mantles on a stronger footing.

1012.0324

Numerical magnetohydrodynamic (MHD) simulations of a turbulent solar nebula are used to study the growth of dust mantles swept up by chondrules. A small neighborhood of the solar nebula is represented by an orbiting patch of gas at a radius of 3 AU, and includes vertical stratification of the gas density. The differential rotation of the nebular gas is replaced by a shear flow. Turbulence is driven by destabilization of the flow as a result of the magnetorotational instability (MRI), whereby magnetic field lines anchored to the gas are continuously stretched by the shearing motion. A passive contaminant mimics small dust grains that are aerodynamically well coupled to the gas, and chondrules are modeled by Lagrangian particles that interact with the gas through drag. Whenever a chondrule enters a region permeated by dust, its radius grows at a rate that depends on the local dust density and the relative velocity between itself and the dust. The local dust abundance decreases accordingly. Compaction and fragmentation of dust aggregates are not included. Different chondrule volume densities ρ<SUB>c</SUB> lead to varying depletion and rimmed-chondrule size growth times. Most of the dust sweep-up occurs within ∼1 gas scale-height of the nebula midplane. Chondrules can reach their asymptotic radius in 10-800 years, although short growth times due to very high ρ<SUB>c</SUB> may not be altogether realistic. If the sticking efficiency Q of dust to chondrules depends on their relative speed δv, such that Q &lt; 10 <SUP>-2</SUP> whenever δv &gt; v<SUB>stick</SUB> ≈ 34 cm/s (with v<SUB>stick</SUB> a critical sticking velocity), then longer growth times result due to the prevalence of high MRI-turbulent relative velocities. The vertical variation of nebula turbulent intensity results in a moderate dependence of mean rimmed-chondrule size with nebula height, and in a ∼20% dispersion in radius values at every height bin. The technique used here could be combined with Monte Carlo (MC) methods that include the physics of dust compaction, in a self-consistent MHD-MC model of dust rim growth around chondrules in the solar nebula.

12231077

2011PThPS.190..120L

1012

1012.5737_arXiv.txt

Primordial scalar perturbations in the Universe are approximately Gaussian and have approximately flat power spectrum~\cite{Komatsu:2010fb}. The first property suggests that these perturbations originate from amplified vacuum fluctuations of weakly coupled quantum field(s). The flatness of the power spectrum may be due to some symmetry. The best known is the symmetry of the de~Sitter space-time under spatial dilatations supplemented by time translations. This is the approximate symmetry of the inflating Universe~\cite{inflation}, which ensures approximate flatness of the scalar spectrum generated by the inflationary mechanism~\cite{infl-perturbations}. Inflation is not the only option in this regard, however. Indeed, the flat scalar spectrum is generated also in the scalar theory with negative exponential scalar potential in flat space-time~\cite{minus-exp} (see also ref.~\cite{minus-old}). Its equation of motion is invariant under space-time dilatations supplemented by the shifts of the field. This symmetry remains approximately valid in slowly evolving, e.g., ekpyrotic~\cite{ekpyrosis} or ``starting''~\cite{starting} Universe, hence the flatness of the resulting perturbation spectrum in these models. It is worth noting that there are other mechanisms capable of producing flat or almost flat scalar spectrum~\cite{Wands:1998yp, Mukohyama:2009gg}. In some cases, there is no obvious symmetry that guarantees the flatness, i.e., the scalar spectrum is flat accidentally. In search for alternative symmetries behind the flatness of the spectrum, one naturally comes to conformal invariance~\cite{vrscalinv, Creminelli:2010ba}. In the scenario of ref.~\cite{vrscalinv}, it is supplemented by a global symmetry. The simplest model of this sort has global symmetry $U(1)$ and involves complex scalar field $\phi$, which is conformally coupled to gravity and for long enough time evolves in negative quartic potential \be V(\phi) = - h^2 |\phi|^4 \; . \label{jul22-1} \ee The theory is weakly coupled at $ h < 1 $. One assumes that the background space-time is homogeneous, isotropic and spatially flat, $ds^2 = a^2(\eta)(d\eta^2 - d{\bf x}^2)$. Then, due to conformal invariance, the dynamics of the field $ \chi (\eta, {\bf x})= a \phi$ is independent of the evolution of the scale factor and proceeds in the same way as in Minkowski space-time. One begins with the homogeneous background field $\chi_c (\eta)$ that rolls down the negative quartic potential. Its late-time behavior is completely determined by conformal invariance, \be \chi_c (\eta) = \frac{1}{h (\eta_* - \eta)} \; , \label{jul22-2} \ee where $\eta_*$ is an arbitrary real parameter, and we consider real solution, without loss of generailty. As we review in section~\ref{Review}, at early times the linear perturbations about this background oscillate in conformal time as modes of free massless scalar field, while at late times the perturbations of the phase \[ \theta =\sqrt{2}\mbox{Arg}~\phi \] freeze out. Somewhat unconventional normalization of the phase is introduced for future convenience. At the linear level, their power spectrum is flat, \be \sqrt{{\cal P}_{\delta \theta}} = \frac{h}{2\pi} \; . \label{jul21-1} \ee As discussed in ref.~\cite{vrscalinv}, this property is a consequence of conformal and global symmetries. \begin{wrapfigure}{r}{6.6cm} % \includegraphics[width=6.6cm]{potential1a} \caption{The scalar potential. Bullets show the evolution of the scalar field. Arrows at the end point at the bottom of the potential indicate that there are perturbations of the phase. \label{fig1}} \end{wrapfigure} The scenario proceeds with the assumption that the scalar potential $V(\phi)$ has, in fact, a minimum at some large value of $|\phi|$, and that the modulus of the field $\phi$ eventually gets relaxed to the minimum, see figure~\ref{fig1}. The simplest option concerning further evolution of the perturbations $\delta \theta$ is that they are superhorizon in the conventional sense by the time the conformal rolling stage ends. We proceed under this assumption. The phase perturbations remain frozen out,\footnote{For contracting Universe, this property of superhorizon modes holds if the dominating matter has stiff equation of state, $w>1$. This appears to be necessary for the viability of the bounce scenario anyway, see the discussion in refs.~\cite{smooth, ekpyro}.} and their power spectrum remains flat. At some much later cosmological epoch, the perturbations of the phase are converted into the adiabatic scalar perturbations by, e.g., the curvaton mechanism~\cite{Linde:1996gt} (in that case $\theta$ is a pseudo-Nambu-Goldstone curvaton, and reprocessing occurs as discussed in ref.~\cite{Dimopoulos:2003az}) or modulated decay mechanism~\cite{Dvali:2003em,Dvali:2003ar}. In either case, the power spectrum is not distorted, so the resulting adiabatic perturbations have flat primordial power spectrum. If conformal invariance is not exact at the rolling stage, the scalar power spectrum has small tilt, which depends on both the strength of the violation of conformal invariance and the evolution of the scale factor at the rolling stage~\cite{Osipov}. A peculiar property of the model is that the modulus of the rolling field also acquires perturbations. At late times, modes of the modulus (i.e., radial direction) have red power spectrum (see section~\ref{modulusperturbations} for details), \be \sqrt{{\cal P}_{|\phi|}(k)} \propto k^{-1} \; . \label{jul23-2} \ee One consequence is that there are perturbations of the energy density with red spectrum right after the conformal rolling stage, but before the modulus freezes out at the minimum of $V(\phi)$. These are not dangerous, provided that the energy density of the field $\phi$ is small compared to the total energy density at all times before the modulus freezes out, i.e., the cosmological evolution is governed by some other matter at that early epoch. In this paper we assume that this is indeed the case. The second consequence is that the infrared radial modes interact with the perturbations of the phase, and in principle may have strong effect on the latter. This is one of the issues we address in this paper. We show that to the linear order in $h$, the infrared effects can be absorbed into field redefinition, so there is no gross modification of the results of the linear analysis due to the effect of the infrared modes. The large wavelength modes of $\delta |\phi|$ are not entirely negligible, however. The modes whose present wavelengths exceed the present Hubble size $H_0^{-1}$ induce statistical anisotropy in the perturbations of the phase $\delta \theta$, and hence in the resulting adiabatic perturbations: the power spectrum of the adiabatic perturbation $\zeta$ has the following form, \be {\cal P}_\zeta ({\bf k}) = {\cal P}_0 (k) \left(1 + c_1 \cdot h \cdot \frac{H_0}{k} \cdot \hat{k}_i \hat{k}_j w_{ij} - c_2 \cdot h^2 \cdot ({\bf \hat k u})^2\right)\; . \label{jul22-6} \ee The first non-trivial term, linear in $h$, is free of the infrared effects; $w_{ij}$ is a traceless symmetric tensor of a general form with unit normalization, $w_{ij} w_{ij} =1$, ${\bf \hat k}$ is a unit vector, ${\bf \hat k}= {\bf k}/k$, and $c_1$ is a constant of order 1 whose actual value is undetermined because of the cosmic variance. In the last term, ${\bf u}$ is some unit vector independent of $w_{ij}$, and the positive parameter $c_2$ is logarithmically enhanced due to the infrared effects. This is the first place where the deep infrared modes show up. Clearly, their effect is subdominant for small $h$. The statistical anisotropy encoded in the last term in \eqref{jul22-6} is similar to that commonly discussed in inflationary context~\cite{aniso}, and, indeed, generated in some concrete inflationary models~\cite{soda}: it does not decay as momentum increases and has special tensorial form $({\bf \hat{k} u})^2$ with constant ${\bf u}$. On the other hand, the first non-trivial term in \eqref{jul22-6} has the general tensorial structure and decreases with momentum. The latter property is somewhat similar to the situation that occurs in cosmological models with the anisotropic expansion before inflation~\cite{Peloso}. Overall, the statistical anisotropy \eqref{jul22-6} may be quite substantial, since there are no strong bounds on $h$ at least for the modulated decay mechanism of conversion of the phase perturbations into adiabatic ones. The non-linearity of the scalar potential gives rise to the non-Gaussianity of the perturbations of the phase $\delta \theta$, and hence the adiabatic perturbations in our scenario, over and beyond the non-Gaussianity that may be generated at the time when the phase perturbations get reprocessed into the adiabatic perturbations. In view of the result outlined above, this non-Gaussianity is not plagued by the infrared effects at the first non-trivial order in $h$. Therefore, the gradient expansion of the effective background is useless for the study of the non-Gaussianity, and we have to perform a hard-core calculation. Because of the symmetry $\theta \to - \theta$, the bispectrum of the phase vanishes, while the general expression for the trispectrum is rather cumbersome. It simplifies in the folded limit (according to the nomenclature of ref.~\cite{Bartolo:2010di}); the explicit form of the trispectrum in this limit is given by eq.~(\ref{Eq/Pg12/2:in-in}). The paper is organized as follows. In section~\ref{Review} we review the linear analysis of the model. In section~\ref{twoorders} we study the effect of infrared modes of the modulus $\delta |\phi|$ on the perturbations of the phase $\delta \theta$ at the leading and subleading orders of the gradient expansion, and to the linear order in $h$. Statistical anisotropy is analysed in section~\ref{anisotropy}. Non-Gaussianity is considered in section~\ref{Section/Pg27/1:Ykis2010/Non-Gaussianity}. We conclude in section~\ref{conclude}.

\label{conclude} We conclude by making a few remarks. First, our mechanism of the generation of the adiabatic perturbations can work in any cosmological scenario that solves the horizon problem of the hot Big Bang theory, including inflation, bouncing/cyclic scenario, pre-Big Bang, etc. In some of these scenarios (e.g., bouncing Universe), the assumption that the phase perturbations are superhorizon in conventional sense by the end of the conformal rolling stage may be non-trivial. It would be of interest to study also the opposite case, in which the phase evolves for some time after the end of conformal rolling. Second, we concentrated in the first part of this paper on the effect of infrared radial modes, and employed the derivative expansion. The expressions like \eqref{jul25-32}, which we obtained in this way, must be used with caution, however. Bold usage of \eqref{jul25-32} would yield, e.g., non-vanishing equal-time commutator $[\theta ({\bf x}), \theta ({\bf y})]$, which would obviously be a wrong result. The point is that the formula \eqref{jul25-32} is valid in the approximation ${\bf v}=\mbox{const}$; with this understanding, the equal-time commutator vanishes, as it should. Finally, the non-linearity of the field equation gives rise to the intrinsic non-Gaussianity of the phase perturbations and, as a result, adiabatic perturbations. The non-Gaussianity emerges at the order $O(h^2)$, and may be sizeable for large enough values of the coupling $h$. The form of the non-Gaussianity is rather peculiar in our scenario. Unlike in many other cases, the three-point correlation function vanishes, while the four-point correlation function of $\delta \theta$ (and hence of adiabatic perturbations) involves the two-point correlator of the independent Gaussian field $\delta \eta_*$. In view of the results of section~\ref{twoorders}, the correlation functions of $\delta \theta$ are infrared-finite, at least to the order $O(h^2)$. This is confirmed by the direct calculation in section~\ref{Section/Pg27/1:Ykis2010/Non-Gaussianity}. The authors are indebted to A.~Barvinsky, S.~Dubovsky, A.~Frolov, D.~Gorbunov, E.~Komatsu, V.~Mukhanov, S.~Mukohyama, M.~Osipov, S.~Ramazanov and A.~Vikman for helpful discussions. We are grateful to the organizers of the Yukawa International Seminar ``Gravity and Cosmology 2010'', where part of this work has been done, for hospitality. This work has been supported in part by Russian Foundation for Basic Research grant 08-02-00473, the Federal Agency for Sceince and Innovations under state contract 02.740.11.0244 and the grant of the President of Russian Federation NS-5525.2010.2. The work of M.L. has been supported in part by Dynasty Foundation.

1012.5737

Primordial scalar perturbations may be generated when complex conformalscalar field rolls down its negative quartic potential. We begin with the discussion of peculiar infrared properties of this scenario. We then consider the statistical anisotropy inherent in the model. Finally, we discuss the non-Gaussianity of scalar perturbations. Because of symmetries, the bispectrum vanishes identically. We present a general expression for the trispectrum and give its explicit form in the folded limit.

12134247

2010ASPC..438..216B

1012

1012.2932_arXiv.txt

Electromagnetism is one of the four fundamental forces in nature. Electric and magnetic fields affect us in more ways than it is possible to list. This is true across all scales, from the very small, such as the electric impulses that control our bodies, to the very large, such as the magnetic field of the Earth that protects us from the ionized Solar wind. Of course, electric and magnetic fields are important throughout the cosmos. Magnetic fields are a fundamental component of the interstellar medium (ISM), which also includes gas (atomic, molecular, and ionized), dust and cosmic rays \citep{spitzer}. They are essential in the formation of stars, and they provide pressure balance which prevents gravitational collapse of our Galaxy \citep{BC90}. Magnetic fields undoubtedly play a role in the creation of galaxies as well as the formation of galaxy clusters. Understanding the origin and evolution of cosmic magnetism is one of the key science drivers for the next-generation radio telescopes \citep[including ASKAP and SKA; see][and section \ref{revolutions}]{bg04}. Our Galaxy provides a natural laboratory for exploring cosmic magnetic fields. It is reasonable to anticipate that determining how the field within our own Galaxy was generated, how it is sustained, and how it is evolving, will all contribute significantly towards understanding cosmic magnetism. The history of geomagnetism began around 1000 AD in China with the discovery of a magnetic compass, though it wasn't until 1600 that the Earth's magnetic field was first scientifically described by William Gilbert and then formally measured in the 1800s by Carl Gauss \citep[see][and references therein]{Stern}. By contrast, the idea that the Galaxy also has a magnetic field was proposed only recently by \citet{fermi49}, who suggested it plays an important role as a generator of observed high energy cosmic ray particles. In the 1950s, the observation of polarised radiation was attributed to synchrotron radiation \citep{kiepen1, kiepen2} and observations of polarised dust allowed for initial investigations of the local magnetic field \citep{davis51}. However, it was not until the late 1970s that studies of the Galactic magnetic field beyond our local spiral arm really began \citep[e.g.][]{Ruz77,sk79}.

1012.2932

Cosmic magnetic fields are an integral component of the interstellar medium (ISM), having influence on scales ranging from star formation to galactic dynamics. While observations of external galaxies offer a ‘birds-eye-view' of magnetic fields within galaxies, it is equally important to explore the magnetic field of our own Milky Way Galaxy, which offers a more detailed, albeit more complicated view. Over the past decade there has been a significant increase in interest in the Galactic magnetic field, fueled largely by innovations developed through the Canadian Galactic Plane Survey. In this paper, I review the current state of understanding of the Galactic magnetic field, and discuss briefly new and future observations that will provide exciting new insights about the field.

12137814

2010idm..confE..62P

1012

1012.2335_arXiv.txt

\label{introduction} The existence of dark matter (DM) is inferred from many different astrophysical and cosmological observations, which indicate that it constitutes about 80\% of the mass content of the Universe. However, aside from its gravitational interactions, very little is known about its nature. Among the many proposed particle candidates, the most popular one are weakly interacting massive particles (WIMPs) with masses in the range 10~GeV--10~TeV. Although most proposed WIMPs are stable and are produced thermally in the early Universe with an annihilation cross section (times relative velocity) of $\langle\sigma v\rangle \sim 3 \times 10^{-26}$~cm$^3$~s$^{-1}$, DM may be unstable but long-lived, with a lifetime $\tau_\chi$ much longer than the age of the Universe $t_{\rm U} \simeq 4 \times 10^{17}$~s. Moreover, DM might have been produced via non-thermal processes and have a larger annihilation cross section than the canonical value for WIMP thermal relics. Among the different ways to detect DM, indirect searches look for the products of DM annihilation or decay, which include antimatter, neutrinos and photons. During the last years, different approaches have been proposed to constrain dark matter properties by using indirect measurements~\cite{indirectnuprop, indirectgamma, Hensley:2009gh}. However, to extract the properties of the DM particle from the detection of an indirect signal requires several pieces of information. There exist many different degeneracies among the different parameters which determine the energy spectrum of the signal. In general, this prevents accurate reconstruction of the DM properties from the energy spectrum alone. In particular, the sole measurement of the energy spectrum would make it impossible to know if the indirect signal from DM is produced by annihilation or decay. The spectrum of the former is characterized by a cutoff at an energy equal to the DM mass, while the cutoff in the spectrum from the latter is at an energy equal to half of the DM mass. In this talk (see also Ref.~\cite{PalomaresRuiz:2010pn}), we address the question whether annihilation and/or decay can be identified as the origin of a DM signal in gamma rays. We note that if DM is unstable and produces an observable signal from decay, an annihilation signal will also be present. We show that there is a range of parameters for which the two signals would be comparable, and in this case, angular information could help to determine their presence and their relative contribution to the total signal. Although very challenging, this would identify DM as an unstable particle. In particular, in order to tackle this problem, we suggest a strategy to distinguish between these scenarios using future gamma-ray observations of Milky Way dwarf galaxies. We show that, in the case that a gamma-ray signal is clearly detected, the origin could be identified as DM decay, annihilation, or both by examining the dependence of the intensity and energy spectrum on the angular distribution of the emission. Furthermore, if annihilation and decay each contribute significantly to the signal, we show how these observations could be used to extract information about the DM mass, lifetime, annihilation cross section, and dominant annihilation and decay channels. In addition, as a byproduct of this analysis, one might also establish or limit the contribution to the signal from substructure in the dwarf galaxy's halo.

In this talk we have outlined a strategy to constrain DM properties in the event of the clear detection of an indirect signal from gamma-ray observations of dwarf galaxies. We addressed the question of how scenarios of DM annihilation, decay, or both could be distinguished, and what information could be obtained about the intrinsic properties of the DM particle and its small-scale distribution from this type of indirect measurement. In summary, we have shown that a DM particle with an annihilation cross-section and lifetime just beyond the limits currently established could produce a clear spectral change on an angular scale within the reach of future experiments. Ongoing observations by current and future experiments will continue to improve the prospects for detecting and mapping a DM signal in the coming years.

1012.2335

Although most proposed dark matter candidates are stable, in order for dark matter to be present today, the only requirement is that its lifetime is longer than the age of the Universe, t_U ~ 4 10^17 s. Moreover, the dark matter particle could be produced via non-thermal processes and have a larger annihilation cross section from the canonical value for thermal dark matter, &lt;sigma v&gt; ~ 3 10^{-26} cm3/s. We propose a strategy to distinguish between dark matter annihilation and/or decay in the case that a clear signal is detected in future gamma-ray observations of Milky Way dwarf galaxies with gamma-ray experiments. The discrimination between these cases would not be possible in the case of the measurement of only the energy spectrum. We show that by studying the dependence of the intensity and energy spectrum on the angular distribution of the signal, the origin of the signal could be identified, and some information about the presence of substructure might be extracted.

1506019

2011ApJS..193...17P

1012

1012.5665_arXiv.txt

Understanding the formation of complex molecules and dust grains in space is a major problem in modern astrophysics. Stars on the asymptotic giant branch (AGB) efficiently produce C, N, O and s-process elements. Mass loss leads to the formation of expanding circumstellar envelopes (CSEs) whose molecules and dust grains are major sources of replenishment of the interstellar medium \citep{herwig2005,BussoEtAl1999}. The radius of the C/O core, after the exhaustion of core He burning, is R$_{c} \approx 10^{9}$ cm, and the temperature is $\sim10^{8}$ K. He burning continues in a shell around the core and the size of the stellar photosphere expands to R$_{*}\approx 10^{13}$ cm. During the phase of intermittent burning of H and He (in the so-called ``thermal pulse''), products of nuclear burning are dredged up by convection and brought to the stellar surface. The zone within a few stellar radii is dynamically important for mass-loss; in this zone molecules produced in the stellar atmosphere are moved to a region away from the star that is cool enough ($\sim 1000$~K) so that dust grains can condense. Radiation pressure on the grains drives the CSE's expansion. The outer circumstellar region is comparatively less dense by several orders of magnitude and richer in molecular gas, steadily expanding outwards with velocities $\sim$ 10 km s$^{-1}$. The chemistry in the outermost part of the circumstellar shell is driven by the interstellar UV radiation \citep{glassgold1996}. Due to the clumpy nature of the shell, photochemistry may be important in the inner regions as well \citep{DecinEtAl2010}. IRC+10216 (CW Leo) is a well known AGB carbon star ([C]$>$[O] and presence of s-process elements) with a high mass-loss rate (several $\times 10^{-5}$ M$_{\odot}$ yr$^{-1}$) at a distance of 150 pc \citep[e.g. ][]{YoungEtAl1993, CrosasMenten1997}. Owing to its closeness to the Sun, it has been possible to study the physical and chemical processes in its large circumstellar envelope in great detail \citep[e.g. ][]{Olofsson1999}. There are nearly 60 molecules observed in the circumstellar shell of IRC+10216 as a result of previous single-dish line surveys \citep{KawaguchiEtAl1995, CernicharoEtAl2000, AveryEtAl1992, GroesbeckEtAl1994,HeEtAl2008, ZiurysEtAl2002, CernicharoEtAl2010,TenenbaumEtAl2010}. Mapping the spatial distribution of molecules in the circumstellar envelope of IRC+10216 is important for several reasons: (1) molecular (and isotopic) abundances can be accurately determined, since the excitation temperature can be inferred from the spatial location of the molecules in the envelope; (2) such data can be readily and quantitatively compared with chemical models predicting abundances as a function of radial distance from the star ; (3) parent molecules can be distinguished from product molecules given their distribution in the envelopes; (4) molecules important for the creation of dust can be identified (e.g., the distribution of SiO, SiS, SiN and silicon carbides appear to be centrally concentrated near the region of dust formation and probably are dominant constituents of grains forming in this region, their abundance decreases with distance from the star); (5) multiple transitions of the same molecule allows mapping of physical conditions (e.g., temperature) in the envelope. Interferometric maps of NaCN, SiO, SiS, CS, HC$_{5}$N, SiCC, NaCl, MgNC, CN, HNC, C$_{2}$H, C$_{3}$H, C$_{4}$H have been presented by \cite{GuelinEtAl1996} and \cite{DayalAndBieging1995}. Except for the SiO J=5--4 line at 217 GHz \citep{SchoierEtAl2006} and the CS J=14--13 line at 685 GHz \citep{YoungEtAl2004} mapped with the SMA, all the other maps were obtained using the IRAM Plateau de Bure Interferometer (PdBI) or the Berkeley-Illinois-Maryland Array (BIMA) at around 100 GHz. We selected the 345 GHz band for our survey primarily because very little data exists in this frequency range, and it contains transitions from many astrochemically important molecules, including various salts, the cyanopolyyne HC$_{3}$N, and cyclic molecules such as C$_{3}$H$_{2}$. Two line surveys in the 345 GHz band have been published by \cite{AveryEtAl1992} in the frequency range of 339.6--364.6 GHz with a sensitivity of 0.3 K rms made with the 15 m diameter James Clerk-Maxwell Telescope (JCMT) and by \cite{GroesbeckEtAl1994}, in the frequency range of 330.2-358.1 GHz, with the Caltech Submillimeter Observatory (CSO) 10.4 m telescope with an rms noise level of 65 mK (4.6 Jy). Our frequency range goes well beyond these surveys, including the range of 300--330 GHz which is almost unexplored.

The SMA line survey of IRC+10216 has yielded 442 lines, 293 of which have been assigned to known transitions. Most are from molecules known to exist in the circumstellar envelope of IRC+10216, including, SiCC, SiS, SiO, CS, C$_{4}$H, CH$_{3}$CN, HCN, HC$_{3}$N, and their isotopic species. Also detected are several lines from salts and metals, including, NaCl, KCl, AlCl, and AlF. More than 100 lines remain unidentified. Maps of these U-lines typically show very compact emission, suggesting vibrationally excited lines of known simple molecules, produced very close to the star or even with its photosphere. Assignment of the substantial number of unidentified lines observed in the inner envelope awaits laboratory measurements of rotational transitions from high lying levels in the ground and vibrationally excited states of polyatomic molecules.

1012.5665

We report a spectral-line survey of the extreme carbon star IRC+10216 carried out between 293.9 and 354.8 GHz with the Submillimeter Array. A total of 442 lines were detected, more than 200 for the first time; 149 are unassigned. Maps at an angular resolution of ~3'' were obtained for each line. A substantial new population of narrow lines with an expansion velocity of ~4 km s<SUP>-1</SUP> (i.e., ≈30% of the terminal velocity) was detected. Most of these are attributed to rotational transitions within vibrationally excited states, emitted from energy levels above the v = 0, J = 0 ground state with excitation energy of 1000-3000 K. Emission from these lines appears to be centered on the star with an angular extent of &lt;1''. We use multiple transitions detected in several molecules to derive physical conditions in this inner envelope of IRC+10216.

12201819

2011ASPC..437..147C

1012

1012.3796_arXiv.txt

The Solar Dynamics Observatory (SDO) is NASA's first mission of the Living with a Star program, which is designed to study the causes of solar variability and its impacts on life and humanity's technological development. Solar variability is intimately related to magnetic activity, so the main goals of the program are to understand the mechanisms that produce these fields and drive them to the surface, and be able to predict when and where the energy stored in them is eventually going to be released in the form of particle ejections and changes in the solar irradiance. \noindent SDO was launched from Cape Canaveral on February 11, 2010, carrying three instruments on board: the Atmospheric Imaging Assembly (AIA), the Extreme ultraViolet Experiment (EVE) and the Helioseismic and Magnetic Imager (HMI). The nominal lifetime of the mission is just over 5 years, with an extension of up to 10 years. The spacecraft follows a geosynchronous orbit (24h period at 36000 km) passing over over the ground station in White Sands Missile Range in New Mexico once a day, to where it downloads the more than 1 TB of data a day that the three instruments produce altogether.

To summarize the reasons for the discrepancies it is necessary to understand the different nature of the data provided by both instruments. To start with, HMI and SP observe different spectral lines, and the sensitivity of these to the physical and thermodynamical properties of the atmosphere are not identical. The operation mechanisms of the instruments are such that simultaneity is never properly achieved, and although there are ways around it (like building a composite of HMI images that are chosen to be very close in time to each scanning step of an SP map) this was not done in this study. The large differences in spatial and spectral resolution are probably the biggest caveat for an inter-instrument comparison. The alignment process (rotation and pixel scaling) required an interpolation of the HMI data, which also introduces some degree of smearing. Residual misalignments remain because the variable scanning step size of the SP --and other second order effects-- were not considered All in all, the comparison is very promising and it suggests that HMI, given its limitations, produces comparable results to SP. A more detailed and accurate comparison is being carried out in order to understand the different sources of the disagreements. The bottom-line is that, for any given ground or space-based observation obtained with any instrument on any part of the solar disk, there will always be reliable photospheric vector-magnetograms at a high cadence and a consistent 1\arcsec\ spatial resolution available. HMI data are useful on their own for studies of the evolution of photospheric magnetic fields and helioseismology, but can also fill up the gaps and/or complement observations from other instruments

1012.3796

The Helioseismic and Magnetic Imager (HMI) has just started producing data that will help determine what the sources and mechanisms of variability in the Sun's interior are. The instrument measures the Doppler shift and the polarization of the Fe I 6173 Å line, on the entire solar disk at a relatively-high cadence, in order to study the oscillations and the evolution of the full vector magnetic field of the solar Photosphere. After the data are properly calibrated, they are given to a Milne-Eddington inversion code (VFISV, Borrero et al. 2010) whose purpose is to infer certain aspects of the physical conditions in the Sun's Photosphere, such as the full 3-D topology of the magnetic field and the line-of-sight velocity at the solar surface. We will briefly describe the characteristics of the inversion code, its advantages and limitations -both in the context of the model atmosphere and the actual nature of the data-, and other aspects of its performance on such a remarkable data load. Also, a cross-comparison with near-simultaneous maps from the Spectro-Polarimeter (SP) onboard Hinode will be made.

2775592

2011A&A...527A..37B

1012

1012.3872_arXiv.txt

\label{sect:intro} The source $\beta$~Hydri (\bhyi, HD 2151, HR 98, HIP 2021) is a single, bright subgiant star (m$_{V}$=2.80) that is clearly visible to the naked eye about $12^{\circ}$ from the South Pole. It is the closest subgiant star, with a spectral and luminosity type between G2\,IV \citep{hoffleit95,evans57} and G0\,V \citep{gray06}, and it is one of the oldest stars in the solar Galactic neighbourhood. It is frequently regarded as representing the future of the Sun \citep{dravins93i,dravins93ii,dravins93iii}, making it a particularly interesting object of study. Improvements to the fundamental parameters of \bhyi~have been presented in a number of recent papers. Recent interferometric measurements of \bhyi~have yielded an accurate (0.8\%) angular diameter for this star \citep{north07}. Also, the \emph{Hipparcos} parallax of \bhyi~has been improved from an uncertainty of 0.4\% \citep{esa97} to 0.08\% \citep{van07}. The combination of these two values gives a direct measure of \bhyi's radius with high accuracy. Moreover, since the bolometric flux of this star is known \citep{blackwell98}, its position in the Hertzprung-Russell (HR) diagram is, in principle, well-constrained. \cite{frandsen87} and \cite{edmonds95} made unsuccessful attempts to detect stellar oscillations in \bhyi, placing upper limits on the p-mode amplitudes. \cite{bedding01} and \cite{carrier01} finally confirmed the presence of solar-like oscillations in \bhyi~and estimated the large frequency separation $\delta\nu$ to be about $55\,\mu {\rm Hz}$, but were unable to identify individual mode frequencies. Subsequently, \cite{bedding07} observed \bhyi~during more than a week with the high-precision spectrographs HARPS and UCLES. Besides confirming the oscillations detected in their previous observations in 2000, they were able to identify 28 oscillation modes that included some mixed modes of spherical degree $l=1$. Mixed modes occur in stars that have left the main-sequence stage of their evolution \citep[e.g.][]{osaki75,aizenman77}, and they provide useful information about the core. The presence of mixed modes, together with the availability of very precise non-seismic and seismic data for \bhyi, places the star in a privileged position for asteroseismic studies \citep[e.g.][]{cunha07}. Theoretical models of \bhyi~based on its seismic and non-seismic data have been published by \cite{fernandes03}, \cite{dimauro03}, and \cite{guln09}. \cite{fernandes03} examined the position of \bhyi~in the HR diagram by first considering the non-seismic data of the star. In order to estimate the mass of \bhyi, they used available seismic data, namely the large frequency separation, to remove partially the helium-content vs mass degeneracy that exists when only non-seismic observational constraints are used. They also emphasized the usefulness of individual frequencies to constrain the age of \bhyi~due to the presence of mixed modes in its observed oscillation spectrum. \cite{dimauro03} computed models of \bhyi, also based on its global parameters. They used the oscillation frequencies of \bhyi~to compare with the model frequencies. Their theoretical models reproduced the observed oscillation spectrum of \bhyi well, as well as the observed large and small frequency separations, after they applied an ad-hoc shift to the computed frequencies. In fact, when comparing the computed and the observed frequencies, one should bear in mind that there may be an offset between them. This offset is well known from helioseismology and is also present when comparing the observed and computed frequencies for other stars. It arises from improper modelling of the surface layers of stars. \cite{kjeldsen08} used solar data to derive an empirical correction for the near-surface offset which can be applied to stars other than the Sun. In our work, we apply this empirical correction to the model frequencies of \bhyi~before comparing to the observed ones. We extend the analysis of \cite{guln09}, and also present a detailed discussion on the application of the near-surface correction. \begin{figure*} \begin{center} $\begin{array}{cc} \includegraphics[width=9cm,angle=0]{tst.ps} & \includegraphics[width=9cm,angle=0]{tst1.ps}\\ \end{array}$ \end{center} \caption{Left panel: The position of \bhyi~in the HR diagram. The constraints on the fundamental parameters ($T_{\rm{eff}}$, $L/\rm{L_{\odot}}$) are indicated by the 1-$\sigma$ error box (solid) and on the radius by diagonal solid lines. We also show the corresponding 3-$\sigma$ uncertainties by dashed lines. Two evolutionary tracks for the best models found by method 2 (cf. Table\,\ref{tab:res} ) are plotted with dash-dotted and solid curves, representing the models with and without gravitational settling and diffusion, respectively. Right panel: the same as in the left panel but zoomed in. The selected models are marked by filled squares.} \label{fig:hr} \end{figure*}

\label{sect:end} We computed two grids of evolutionary models, using the ASTEC code, in order to find the model that most closely reproduces the seismic and non-seismic data of \bhyi. The parameters used for each grid are given in Table\,\ref{tab:param}. We computed the oscillation frequencies for the models that lie inside the 3-$\sigma$ error box of the star's position in the HR diagram with the ADIPLS code, and compared them with the observed frequencies. There is an offset needed to be taken into account in this comparison, due to improper modelling of the near-surface layers of the stars. We used the approach proposed by \cite{kjeldsen08} to correct the computed frequencies of \bhyi~from this offset. We used two methods in order to find the model that reproduces the observed oscillation frequencies of \bhyi. In our analysis, we argue that the method involving the $\chi^2$ test, method 2, is the most robust way to find the best model, since it takes into account all the individual frequencies, the mixed modes, and also the uncertainties on the observed frequencies. Analysing the \'echelle diagrams of the representative models found with method 2 (cf. Sect.\,\ref{sect:res_dis}), we see that the surface correction works very well for $l=0$ modes, and for $l=1$ and 2 modes with frequencies lying in the frequency range of the observed radial modes. This was expected since the correction term was computed using only those radial modes. Observed $l=0$ modes with higher frequencies are thus needed in order to improve the surface correction. Our best models give $M$=1.04 M$_\odot$ and an age of 6.1 -- 7.3 Gyr for \bhyi, depending on the inclusion of gravitational settling and diffusion of helium. In either case, the radius is found to be $R\sim1.785\,\rm{R}_{\odot}$, which is in good agreement with the one determined by interferometry, $R = 1.809 \pm 0.015$ R$_{\odot}$. However, there are other models fitting the data similarly well. We used the parameters of those models (with $\chi^2 < 10$) to determine the internal error regarding our analysis. We calculated the mean value, and the uncertainties were taken as the standard deviation. We found $M=1.08\pm0.03$ M$_\odot$, age = $6.40 \pm 0.56$ Gyr, and $R = 1.811 \pm 0.020 $ R$_{\odot}$. These results are also consistent with the results of \cite{fernandes03}, who derived, $M=1.10_{-0.07}^{+0.04}$ M$_\odot$ and $M=1.09\pm 0.22$ M$_\odot$, through the HR diagram analysis and $\Delta\nu_0$, respectively, and a stellar age between 6.4 and 7.1 Gyr.

1012.3872

Context. Comparing models and data of pulsating stars is a powerful way to understand the stellar structure better. Moreover, such comparisons are necessary to make improvements to the physics of the stellar models, since they do not yet perfectly represent either the interior or especially the surface layers of stars. Because β Hydri is an evolved solar-type pulsator with mixed modes in its frequency spectrum, it is very interesting for asteroseismic studies. <BR /> Aims: The goal of the present work is to search for a representative model of the solar-type star β Hydri, based on up-to-date non-seismic and seismic data. <BR /> Methods: We present a revised list of frequencies for 33 modes, which we produced by analysing the power spectrum of the published observations again using a new weighting scheme that minimises the daily sidelobes. We ran several grids of evolutionary models with different input parameters and different physics, using the stellar evolutionary code ASTEC. For the models that are inside the observed error box of β Hydri, we computed their frequencies with the pulsation code ADIPLS. We used two approaches to find the model that oscillates with the frequencies that are closest to the observed frequencies of β Hydri: (i) we assume that the best model is the one that reproduces the star's interior based on the radial oscillation frequencies alone, to which we have applied the correction for the near-surface effects; (ii) we assume that the best model is the one that produces the lowest value of the chi-square (χ<SUP>2</SUP>), i.e. that minimises the difference between the observed frequencies of all available modes and the model predictions, after all model frequencies are corrected for near-surface effects. <BR /> Results: We show that after applying a correction for near-surface effects to the frequencies of the best models, we can reproduce the observed modes well, including those that have mixed mode character. The model that gives the lowest value of the χ<SUP>2</SUP> is a post-main-sequence model with a mass of 1.04 M<SUB>⊙</SUB> and a metallicity slightly lower than that of the Sun. Our results underscore the importance of having individual frequencies to constrain the properties of the stellar model.

2519208

2011ApJ...732...70L

1012

1012.1877_arXiv.txt

There has been an-ongoing effort to understand the evolution of low-mass X-ray binaries (LMXBs) since their basic nature was understood back in the late 1960's and early 1970's \citep[see, e.g.,][]{faul71,rapp83,webb83,joss84,nels85,pyly88,pyly89,bhat91,iben95,pods02,pfah03,nels04,belc08}. There are two rather distinct parts of the evolution to consider: (i) the formation of a neutron star in orbit with a low-mass donor star, and (ii) the subsequent portion of the evolution when mass is transferred from the companion donor star to the neutron star. The difficulty with the first part of the evolution involves the conceptual problem of keeping the binary system bound while the massive progenitor of the neutron star (NS) explodes in a supernova event. A part of this problem was addressed by invoking a common envelope phase during which the lower-mass secondary ejects the envelope of the massive NS progenitor \citep{pacz76,meye79,webb79,bhat91}. However, even this solution is not so straightforward in that there may be insufficient gravitational energy release between the inspiraling low-mass star and the core of the NS progenitor to successfully eject the envelope \citep[see, e.g.,][]{dewi00,pfah03}. This difficulty could be overcome by invoking secondaries that are of intermediate mass, e.g., $2-4\, M_\odot$ to increase the gravitational energy release \citep[see, e.g.,][]{pfah03}. In turn, there has been a persistent conceptual misunderstanding that the transfer of mass from a star of $2-4\,M_\odot$ onto a $1.4 \,M_\odot$ NS was dynamically unstable. That was shown to be untrue (\citealp{pyly88,pyly89,taur99,pods00,king01}; PRP02; see also \citealp{davi98,king99}). Starting in the 1970's and continuing until the present, there have been numerous evolution studies of a limited number of LMXBs and IMXBs exploring the various paths which lead to very different intermediate and end-stage products. One of the more systematic was the study by \citet{pods02} and \citet{pfah03} which involved about 150 LMXB and IMXB systems. Given the two-dimensional parameter space of initial $P_{\rm orb,i}$ and $M_{\rm 2,i}$ (where $P_{\rm orb,i}$ and $M_{\rm 2,i}$ are the initial orbital period and mass of the donor star), the types of evolutionary paths for LMXBs or IMXBs is really quite large and varied. Possibilities include evolution to a minimum orbital period of $\sim$70 min with very low-mass H-rich donor stars; evolution to an ultracompact state with $P_{\rm orb}$ in the range of $6-50$ minutes and He-rich donors; and evolution to wide orbits of days-to-months with low-mass giant donor stars. In order to better understand the many possible evolution paths of LMXBs and IMXBs in a more systematic way, we have calculated an extensive grid of binary models encompassing 42,000 initial combinations of $P_{\rm orb,i}$ and $M_{\rm 2,i}$. This is two orders of magnitude larger than the study we conducted in 2002. We took advantage of a newly developed stellar evolution code whose equations of state allow for the evolution of very low mass stars with cold dense interiors, and a cluster of computers which speeds up the overall calculation beyond what was readily available a decade ago. With the recent discovery of PSR J1614-2230, with the most massive neutron star known ($1.97\pm 0.04\,M_\odot$), a relatively close 8.7-day orbit, and a fairly massive white dwarf companion ($0.5\,M_\odot$), we decided to first apply our evolution calculations to understanding the origins of this system. The specific goals are to understand whether a natal NS with canonical mass of $\sim$$1.4\,M_\odot$ can accrete sufficient material to grow to nearly $2\,M_\odot$, to see whether the donor star mass is consistent with the observed $P_{\rm orb}$, and to investigate how this system is related to possible progenitors in the guise of the LMXB -- Cyg X-2. In all, we found $\sim$500 of our evolution tracks which produce systems that are at least generically related to PSR J1614-2230. In this paper we introduce our set of binary evolution calculations which cover an extensive grid of initial orbital periods and donor masses (\S \ref{sec:overview}). In \S \ref{sec:PSR1614} we show how our binary evolution calculations apply directly to PSR J1614-2230. In \S \ref{sec:Discuss} we discuss a number of the general lessons we have gleaned from this study.

\label{sec:Discuss} We have used {\tt MESA} to compute a large grid of 42,000 binary evolution models for LMXBs and IMXBs. We showed the broad sweep of possible evolutions, from systems which attain orbital periods as short as 6 minutes to those which grow to long-period binaries with giant donor stars. We leave a detailed discussion of the bulk of these results for a future paper. Here we have focused on what we can learn about the evolutionary paths to the newly discovered binary pulsar, PSR J1614-2230. We have selected a subset of the evolution models (515 in total) which lead to systems like PSR J1614-2230 and its evolutionary cousin, Cyg X-2, to examine in more detail. In particular, we show how an orbital period of 8.7-days can be easily understood, as can the $0.5\,M_\odot$ companion star. We show that for the proposed scenario, the degenerate companion star is mostly (i.e., 90\%) C and O (with a surrounding shell of He which comprises 10\% of the white-dwarf mass), but there is also a thin outer envelope that is composed of $\sim$15\% H. We found, however, that starting with neutron stars of canonical mass $1.4\,M_\odot$, we were not able to produce the observed $1.97$~M$_{\odot}$ NS with the requisite combination of orbital period {\it and} white dwarf mass to match the PSR J1614-2230 system. We were able to evolve neutron stars as massive as $2\,M_\odot$, with the correct orbital period, but not in combination with a massive enough white dwarf. Also, the orbital period and white dwarf mass combination was easy to generate, but not with the correct $P_{\rm orb}$. As discussed above, the difficulty with evolving high-mass NSs in orbit with massive white dwarf companions is that the progenitors of these white dwarfs are initially substantially more massive than the NS, resulting in very rapid, thermal-timescale mass transfer (greatly in excess of the Eddington limit), and the neutron star is thereby prevented from accreting a significant fraction of the donor star's mass. At this time, we are therefore tentatively led to conclude that the initial mass of the NS would have to have been higher than the canonical value of $1.4\,M_\odot$ in order to declare that the binary evolution of this system is fully understood. By running some 700 supplementary models with initially higher-mass NSs, we found that to successfully produce a system like PSR J1614-2230 requires a minimum initial neutron star mass of at least $1.6 \pm 0.1 \, M_\odot$, as well as initial donor masses and $P_{\rm orb}$ of $\sim$$4.25 \pm 0.10 \,M_\odot$ and $\sim$$49 \pm 2$ hrs, respectively. For completeness in regard to producing high neutron-star masses, we point out that a number of the dynamically stable case B evolution tracks that we have generated yield mass transfer rates in excess of several times $10^{-4}\,M_\odot$ yr$^{-1}$ that last for intervals of several thousand years. If `hypercritical accretion', where the gravitational energy is carried off in neutrinos and the Eddington limit is thereby circumvented, is able to occur during these intervals \citep[see, e.g.,][]{houc91,brow00}, then perhaps there is a chance for the neutron stars to grow during this phase of the evolution. However, it is not clear to us that the hypercritical accretion scenario, which was developed for spherical accretion onto a neutron star, would be applicable to accretion via a disk \citep[but, see][]{more08}. A more general conclusion from this L/IMXB study is that there is a subclass of systems which start with intermediate mass donor stars (of $\gtrsim 2.2 \, M_\odot$), with $P_{\rm orb}$ well above the `bifurcation period' (in the range of $2-4$ days), which leads to systems like Cyg X-2 and PSR J1614-2230. Such systems terminate their mass transfer before the donor star develops a degenerate core, and they end up with $P_{\rm orb}$ well below the $P_{\rm orb}(M_c)$ relation. (See also \citealp{iben85}; PRP02; \citealp{han00,taur00}). Another important general lesson that we can take from this broad look at L/IMXB evolution is that {\it intermediate}-mass donor stars can evolve to virtually all the known types of LMXB systems that exist at the current epoch. These include CV-like evolution paths, ultracompact X-ray binaries, and systems with giant donor stars. This is an important finding because it is significantly easier for intermediate-mass stars to successfully eject the envelope of their massive companion progenitors of the NSs, and then remain bound during the ensuing supernova explosion. Finally, we have provided a perspective on where in the $P_{\rm orb}$ and white dwarf mass plane we can expect to find the He white dwarfs that follow the $P_{\rm orb}(M_c)$ relation, as well as where the systems with the most massive neutron stars should be found. Depending on the upper mass limit to a neutron star, $M_{\rm ns, max}$, Fig.\,\ref{fig:NSmass} shows that there should be a fair number of black holes with mass between $M_{\rm ns, max}$ and $\sim$$2.8\,M_\odot$ in binaries with He or He/CO white dwarfs that range between $\sim$0.2 and 0.4 $M_\odot$ and with orbital periods in the range of $1-40$ days. One way to detect such important relics of stellar evolution is to search for white dwarfs with interestingly high orbital velocities (i.e., $v \simeq 300 \,(P_{\rm orb}/{\rm days})$ km s$^{-1}$) and unseen companions. An important caveat to this is that if such low-mass black holes form in significant numbers in LMXBs that are still undergoing mass transfer, they would be directly detected by their accretion. However, there is no evidence for LMXBs with low-mass black holes (see, e.g., \citealp{farr10}; \citealp{ozel10}).

1012.1877

We have computed an extensive grid of binary evolution tracks to represent low- and intermediate-mass X-ray binaries (LMXBs and IMXBs). The grid includes 42,000 models which cover 60 initial donor masses over the range of 1-4 M <SUB>sun</SUB> and, for each of these, 700 initial orbital periods over the range of 10-250 hr. These results can be applied to understanding LMXBs and IMXBs: those that evolve analogously to cataclysmic variables, that form ultracompact binaries with P <SUB>orb</SUB> in the range of 6-50 minutes, and that lead to wide orbits with giant donors. We also investigate the relic binary recycled radio pulsars into which these systems evolve. To evolve the donor stars in this study, we utilized a newly developed stellar evolution code called "MESA" that was designed, among other things, to be able to handle very low mass and degenerate donors. This first application of the results is aimed at an understanding of the newly discovered pulsar PSR J1614-2230 which has a 1.97 M <SUB>sun</SUB> neutron star, P <SUB>orb</SUB> = 8.7 days, and a companion star of 0.5 M <SUB>sun</SUB>. We show that (1) this system is a cousin to the LMXB Cyg X-2; (2) for neutron stars of canonical birth mass 1.4 M <SUB>sun</SUB>, the initial donor stars which produce the closest relatives to PSR J1614-2230 have a mass between 3.4 and 3.8 M <SUB>sun</SUB>; (3) neutron stars as massive as 1.97 M <SUB>sun</SUB> are not easy to produce in spite of the initially high mass of the donor star, unless they were already born as relatively massive neutron stars; (4) to successfully produce a system like PSR J1614-2230 requires a minimum initial neutron-star mass of at least 1.6 ± 0.1 M <SUB>sun</SUB>, as well as initial donor masses and P <SUB>orb</SUB> of ~4.25 ± 0.10 M <SUB>sun</SUB> and ~49 ± 2 hr, respectively; and (5) the current companion star is largely composed of CO, but should have a surface H abundance of ~10%-15%.

1675543

2011ApJ...727L..22Z

1012

1012.3333_arXiv.txt

Classical T~Tauri stars (CTTS) are known to be magnetically active protostars showing clear observational signatures of accretion (circumstellar disks) and ejection (jets and outflows). The stellar magnetic fields measured by spectropolarimetric observations \citep[up to a few kG,][]{JK07} can deeply affect the dynamics of the circumstellar region: truncating the disk and channeling the accretion flow along the magnetic surfaces down to the stellar surface; providing an acceleration mechanism for different types of outflows, stellar winds along the opened magnetospheric fieldlines \citep{Matt08a} and ejections associated with the magnetic star-disk interaction \citep{Shu94,Ferreira00,Romanova09,Zanni09}. Their spin represents a controversial issue. While a wide range of rotation periods is observed among low-mass young stars \citep[0.2 up to 20 days,][]{Irwin09}, around half of them slowly rotate, much below the break-up limit. Besides, many slow rotators show clear accretion signatures \citep{Herbst07}, as in the case of CTTS, which have an average rotation period around $\sim8$ days, corresponding to $\sim10\%$ of their break-up speed \citep{Bouvier93}. Since CTTS are accreting mass and angular momentum from the surrounding accretion disk and they are still contracting, they would be expected to noticeably spin-up in a few million years: conversely, there are indications that solar-mass slow rotators are prone to keep their rotation period constant for $\sim5$ Myr \citep{Irwin09}. Therefore, some mechanism must act to efficiently remove angular momentum from these slowly rotating stars. Grounded on models originally developed to explain the period changes of pulsars \citep{GL79}, one of the most widespread scenarios foresees that a significant spin-down torque is provided along the magnetospheric fieldlines connecting the star and the disk region rotating slower than the star \citep{ko91}. On the other hand, both analytical \citep{matt05a} and numerical models \citep{ZF09} have questioned the efficiency of this mechanism, due to the limited extent of the connected region and the weakness of the magnetic connection. \begin{deluxetable*}{c l c c c c c c c c} \tablecaption{Star sample\label{table:sample}} \tablewidth{0pt} \tablehead{ \multicolumn{2}{c}{Object} & \colhead{$M_\star$} & \colhead{$R_\star$} & \colhead{$B_\star$} & \colhead{$P_\star$} & \colhead{$\delta$} & \colhead{$L_\mathrm{UV}$} & \colhead{$\dot{M}_\mathrm{obs}$} & Ref. \\ \multicolumn{2}{c}{Name} & \colhead{[$M_\odot$]} & \colhead{[$R_\odot$]} & \colhead{[G]} & \colhead{[days]} & & \colhead{[$L_\odot$]} & \colhead{[$M_\odot \; \mathrm{\mbox{yr}}^{-1}$]} & \\ } \startdata \multirow{2}{*}{BP~Tau} & (a) & \multirow{2}{*}{0.7} & \multirow{2}{*}{1.95} & \multirow{2}{*}{600} & \multirow{2}{*}{7.6} & \multirow{2}{*}{0.05} & 0.179 & $2 \times 10^{-8}$ & 1,3 \\ & (b) & & & & & & 0.023 & $2.5 \times 10^{-9}$ & 3,5 \\ \multirow{2}{*}{V2129~Oph} & (a) & \multirow{2}{*}{1.35} & 2.4 & 175 & \multirow{2}{*}{6.5} & 0.06 & 0.143 & $ 10^{-8}$ & 2 \\ & (b) & & 2.1 & 450 & & 0.05 & 0.01 & $6.3 \times 10^{-10}$ & 6 \\ CV~Cha & & 2.0 & 2.5 & 300 & 4.4 & 0.07 & 0.61 & $3 \times 10^{-8}$ & 4\\ CR~Cha & & 1.9 & 2.5 & 200 & 2.3 & 0.14 & 0.02 & $10^{-9}$ & 4\\ \multirow{2}{*}{AA~Tau} & (a) & \multirow{2}{*}{0.7} & \multirow{2}{*}{2} & \multirow{2}{*}{1500} & \multirow{2}{*}{8.2} & \multirow{2}{*}{0.05} & 0.025 & $2.8 \times 10^{-9}$ & 1,5 \\ & (b) & & & & & & 0.006 & $6.3 \times 10^{-10}$ & 5 \\ \enddata \tablerefs{[1] \citet{Gullbring98}, [2] \citet{Donati07}, [3] \citet{Donati08}, [4] \citet{Hussain09}, [5] \citet{Donati10a}, [6] \citet{Donati10b}} \end{deluxetable*} \citet{matt05b} have therefore proposed that stellar winds could efficiently remove angular momentum directly from the star along the opened fieldlines of the magnetosphere. Besides, they suggested that these outflows could derive their energy directly from the accretion power. This would be also consistent with the fact that accreting CTTS seem to have on average longer rotation periods than their non-accreting counterparts (weak-lined T~Tauri stars, WTTS), indicating a connection between spin-down and accretion \citep{Lamm05}. It is commonly assumed that the accretion power is liberated in a shock due to the impact of the accretion streams with the stellar surface. While a fraction of the accretion energy can be converted \citep[e.g. into Alfv\'en waves,][]{Sch88} and possibly injected into the wind, the emission of the shocked material can explain the observed optical excess and UV continuum \citep{Gull00}. Observations of the accretion shock luminosity can be therefore used to constrain the accretion energy which is available to power the stellar wind. In this letter we try to estimate the spin-down efficiency and the energy requirement of accretion powered stellar winds (APSW) compatible with measurements of magnetic fields and accretion luminosities of several CTTS. In Section \ref{sec:model} we describe a simple analytical APSW model and we apply it to a specific CTTS example in Section \ref{sec:BP}. We determine the stellar parameters which are compatible with a spin equilibrium situation in Section \ref{sec:spin} and we summarize our conclusions in Section \ref{sec:dicsussion}.

\label{sec:dicsussion} We applied the accretion powered stellar wind model \citep{matt05b} to a sample of CTTS to verify if stellar winds are a viable mechanism to spin-down the rotation of accreting and contracting protostars. According to this scenario, a fraction of the energy deposited by the magnetospheric accretion flow onto the surface of the star could be used to drive the stellar wind: we added the additional constraint that the same accretion energy must be used to power the emission of the accretion shock. In Section \ref{sec:BP} we showed that, for a given spin-down efficiency $f_\mathrm{J}\neq0$, a maximum accretion luminosity can be attained: when the accretion power and, consequently, the spin-up torque become too large, the stellar wind consumes too much energy and a smaller and smaller fraction is left to support the emission. In Table \ref{table:limits} we showed that the stars in our sample characterized by a high accretion luminosity ($L_\mathrm{UV}\geq0.1L_\odot$) impose severe limits on the accretion power which is available to drive the stellar wind so that the spin-down torque is not strong enough to achieve a spin equilibrium. This is consistent with Eq.~(\ref{eq:blim}): in the range of parameters covered in our sample, a dipolar component of kG intensity is required by the wind to spin-down a star characterized by such a high UV emission. In Fig.~\ref{fig:LD} we also showed that an important fraction of our sample (around $50\%$) would require such a strong dipolar component to be compatible with a zero-torque condition: at the moment of writing a dipolar component of kG intensity has been measured only in the case of AA~Tau. Equation (\ref{eq:blim}) and Fig.~\ref{fig:LD} clearly show that lower UV luminosity and/or faster spinning stars require weaker fields, more consistent with the dipolar intensities currently measured, to be in spin equilibrium with an APSW. In fact, stars in the sample with a low accretion luminosity ($L_\mathrm{UV}\ll0.1L_\odot$) are compatible with a $f_\mathrm{J}\geq1$ situation, but the corresponding APSW at maximum spin-down efficiency is energetically very demanding, as confirmed by Eq.~(\ref{eq:efflim}). Some low luminosity cases at spin equilibrium are less demanding, see e.g. the CR~Cha or AA~Tau examples. Still, the mass fluxes would correspond roughly to the entire mass flux of T~Tauri jets \citep[1-20$\%$,][]{Cabrit09}: this would imply that stellar winds are the primary ejecting component of young stars, which seems unlikely \citep{Ferreira06}. Besides, even for relatively low ejection rates ($f_\mathrm{M}\leq0.1$) the energy input is still an issue. It is already known that the wind can not be thermally driven \citep{matt07}: a temperature close to virial ($\sim10^6$ K) determines a too high emission and is incompatible with observations \citep{Dupree05}. Turbulent Alfv\'en waves represent another possible pressure source \citep{Dec81}. Furthermore, it has been suggested that the amplitude of the waves generated by the impact of the accretion streams onto the surface of the star is greater than interior convection-driven wave amplitude \citep{Cranmer09}. In this case, the accretion/ejection energy coupling is not easy to determine: recent models suggest anyway that the wind mass loss rates due to this mechanism are generally very low \citep[$10^{-5}<f_\mathrm{M}<10^{-2}$,][]{Cranmer09}. Besides, it is important to remark that in our sample, when the APSW ejection efficiency $f_\mathrm{M}$ at spin equilibrium becomes $\leq 0.1$, the star-disk system approaches a propeller regime ($R_\mathrm{t}\gtrsim R_\mathrm{co}$, see Eq.~\ref{eq:fmmin}), as in the typical AA~Tau case \citep{Donati10a}. When $R_\mathrm{t}\gtrsim R_\mathrm{co}$, the spin-down torque due to the star-disk interaction, which has not been taken into account here, is in principle enough to slow down the stellar rotation \citep{matt05a, Ustyu06}. We therefore conclude that accretion powered stellar winds are unlikely to be the sole mechanism to provide an efficient spin-down torque for accreting classical T~Tauri stars. Our study suggests that a conservative limit on the wind spin-down torque ($f_\mathrm{J}<0.1-0.2$) reduces the mass flux and power requirements to values more compatible with models of wave-driven winds from T~Tauri stars \citep[$f_\mathrm{M}\sim f_\mathrm{E}<1\%$,][]{Cranmer08}. The problem of the spin of accreting and contracting stars like T~Tauri has still many open issues. It is likely that diverse mechanisms contribute at the same time with different degrees: stellar winds, magnetospheric star-disk angular momentum exchanges \citep{ZF09}, magnetospheric ejections driven by the star-disk interaction \citep{Ferreira00,Zanni09}.

1012.3333

The rotation period of classical T Tauri stars (CTTS) represents a longstanding puzzle. While young low-mass stars show a wide range of rotation periods, many CTTS are slow rotators, spinning at a small fraction of breakup, and their rotation period does not seem to shorten, despite the fact that they are actively accreting and contracting. Matt &amp; Pudritz proposed that the spin-down torque of a stellar wind powered by a fraction of the accretion energy would be strong enough to balance the spin-up torque due to accretion. Since this model establishes a direct relation between accretion and ejection, the observable stellar parameters (mass, radius, rotation period, magnetic field) and the accretion diagnostics (accretion shock luminosity) can be used to constrain the wind characteristics. In particular, since the accretion energy powers both the stellar wind and the shock emission, we show in this Letter how the accretion shock luminosity L <SUB>UV</SUB> can provide upper limits to the spin-down efficiency of the stellar wind. It is found that luminous sources with L <SUB>UV</SUB> &gt;= 0.1 L <SUB>sun</SUB> and typical dipolar field components &lt;1 kG do not allow spin equilibrium solutions. Lower luminosity stars (L <SUB>UV</SUB> Lt 0.1 L <SUB>sun</SUB>) are compatible with a zero-torque condition, but the corresponding stellar winds are still very demanding in terms of mass and energy flux. We therefore conclude that accretion powered stellar winds are unlikely to be the sole mechanism to provide an efficient spin-down torque for accreting CTTS.

2765357

2011A&A...527A..49I

1012

1012.5335_arXiv.txt

At the end of the 20th century, observations of type Ia supernovae (SNIa) revealed that the Universe expansion is accelerating \citep{riess98,perlmutter99}. Since these publications, several efforts have been made to explain these observations (\citet{cunha09, frieman08, linder08, linder05, samsing10, freaza02, ishida05,ishida08} and references therein). In a standard analysis, dark-energy models are characterized by a small set of parameters. These are placed into the cosmic expansion rate by means of the Friedman equations, in substitution for the conventional cosmological-constant term. This approach assumes a specific dependence of the dark-energy equation of state ($w$) on redshift and provides some insight into the probable values of the parameters involved. However, the results remain restricted to that particular parametrization. An interesting question to attempt to answer is what can be inferred about the cosmic expansion rate from observations without any reference to a specific model for the energy content of the Universe? To perform an independent analysis, we used principal component analysis (PCA). In simple terms, PCA identifies the directions of data points clustering in the phase space defined by the parameters of a given model. Consequently, it allows a dimensionality reduction with as minimum an information loss as possible \citep{tegmark97}. The importance of a model-independent reconstruction of the cosmic expansion rate has already been investigated in the literature \citep{Huterer99,Huterer00,Tegmark02,Wang05,mignone08}. In this context, PCA has been used to reconstruct the dark-energy equation of state \citep{huterer03,Crittenden09,simpson06} and the deceleration parameter \citep{shapiro06} as a function of redshift. The use of PCA was also proposed in the interpretation of future experiments results by \citet{albretch09}. In the face of growing interest in the application of PCA to cosmology, \citet{kitching09} recall that some care must be taken in choosing the basic expansion functions and the interpretation assigned to the components. The main goal of this work is to apply PCA to reconstruct directly the \emph{Hubble} parameter redshift dependence without any reference to a specific cosmological model. In this context, the eigenvectors and eigenvalues of the Fisher matrix form a new basis in which the \emph{Hubble} parameter is expanded. For the first time, we show that it is possible to derive analytical expressions for the Fisher matrix if we focus on the \emph{Hubble} parameter ($H(z)$) as a sum of step functions. The reader should realize throughout this work that our procedure is mostly driven by the data, although there is a weak dependence of the components on our starting choices of parameter values. In other words, the functional form of each eigenvector is not of primary importance, we are more interested in how they are linearly combined. This approach allows us to avoid many interpretation problems pointed out by \citet{kitching09}. Our only assumption is that the Universe is spatially homogeneous and isotropic and can be described by Friedmann-Robertson-Walker (FRW) metric. The paper is organized as follows. In section \ref{sec:PCA}, we briefly review our knowledge of PCA and demonstrate how it can be applied to a \emph{Hubble} parameter analysis using type Ia supernova observations. Section \ref{sec:application} shows the results obtained with a simulated supernova data set, following the standard procedure for dealing with the linear combination coefficients. We demonstrate that the quality of our results derived from the simulated data are greatly improved if we consider the \emph{Hubble} parameter value in the upper redshift bound as a free parameter. We apply the same procedure to real type Ia supernova data compiled by the Sloan Digital Sky Survey team \citep{kessler09}. The results are shown in section \ref{sec:current_data}. Finally, in section \ref{sec:conclusions}, we present our conclusions.

\label{sec:conclusions} We have presented an alternative procedure for extracting cosmological parameters from type Ia supernova data. Our analysis is concentrated in the \emph{Hubble} parameter, although we emphasize that the same procedure can be applied to other quantities of interest. Our goal has been to be as general as possible, so we have tried to avoid parametric forms or specific cosmological models by using PCA. Writing $H(z)$ according to equation (\ref{eq:expanso em H}) and considering type Ia supernova observations, we have shown that it is possible to obtain analytical expressions for the Fisher matrix. We used a mock sample formed by 34 redshift bins of width $\Delta z=0.05$, with errors calculated following the prescription proposed by \cite{kim04}. This mock sample represents a simplification of future data sets, such as the JDEM, and is not a realistic representation of current data. Our goal in using it was to check the consistentency of our procedure. Our first attempt in reconstructing the \emph{Hubble} parameter as a linear combination of the eigenvectors of \textsf{\textbf{F}} was unsuccessful. In trying to fit high-redshift data with PCs that go asymptotically to zero, the most oscillatory modes propagate their behavior to the reconstructed $H(z)$ in the whole redshift range. As a consequence, the final result barely resembles our fiducial model. To suppress the influence of the high-redshift behavior present in all PCs of interest, we considered the value of the \emph{Hubble} parameter at high redshift as an extra free parameter in our analysis. This simple modification provided reliable results when used with simulated and real supernova data. Beyond that, our results are corroborated with measurements of red-envelope galaxies from \citet{stern10}. As a final remark, we emphasize that PCA provides a viable way of avoiding phenomenological parameterizations. It represents one of the few statistical methods that allow us to obtain the behavior of a chosen quantity directly from the data. It has its own assumptions, such as Gaussianity, independence of data points and in the specific case analyzed here, cosmologies that obey a FRW metric. In the final reconstruction phase, it also exhibits a bias in the upper redshift bound. On the other hand, the procedure proposed here can drastically suppress the influence of this bias. Beyond that, we show that in the context of this work, the Fisher matrix can be analytically obtained. This avoids all uncertainties related to numerical derivations of step functions and might be a good alternative to standard statistical analyses applied to cosmological data.

1012.5335

<BR /> Aims: A model-independent reconstruction of the cosmic expansion rate is essential to a robust analysis of cosmological observations. Our goal is to demonstrate that current data are able to provide reasonable constraints on the behavior of the Hubble parameter with redshift, independently of any cosmological model or underlying gravity theory. <BR /> Methods: Using type Ia supernova data, we show that it is possible to analytically calculate the Fisher matrix components in a Hubble parameter analysis without assumptions about the energy content of the Universe. We used a principal component analysis to reconstruct the Hubble parameter as a linear combination of the Fisher matrix eigenvectors (principal components). To suppress the bias introduced by the high redshift behavior of the components, we considered the value of the Hubble parameter at high redshift as a free parameter. We first tested our procedure using a mock sample of type Ia supernova observations, we then applied it to the real data compiled by the Sloan Digital Sky Survey (SDSS) group. <BR /> Results: In the mock sample analysis, we demonstrate that it is possible to drastically suppress the bias introduced by the high redshift behavior of the principal components. Applying our procedure to the real data, we show that it allows us to determine the behavior of the Hubble parameter with reasonable uncertainty, without introducing any ad-hoc parameterizations. Beyond that, our reconstruction agrees with completely independent measurements of the Hubble parameter obtained from red-envelope galaxies.

12147247

2010nuco.confE.184B

1012

1012.0456_arXiv.txt

It is commonly believed that the $s$- and $r$-processes derive from separate astrophysical sites \cite{burbidge57}. The nucleosynthesis of the $s$-process occurs in stars of low mass (1.3 $\leq$ $M/M_\odot$ $\leq$ 8) during their thermally pulsing asymptotic giant branch (TP-AGB) phase. % The main neutron source is the $^{13}$C($\alpha$, n)$^{16}$O, which burns radiatively at $T$ $\sim$ 0.9 $\times$ 10$^8$ K during the interpulse period in the region between the H- and He-shell (He-intershell). % A second neutron source, $^{22}$Ne($\alpha$, n)$^{25}$Mg, is marginally activated at the bottom of the recurrent convective thermal instability (thermal pulse, TP) in the He-intershell, mainly affecting the abundance at the branching points that are sensitive to temperature and neutron density. The $s$-process elements are mixed with the surface during the third dredge-up (TDU) episodes, in which the convective envelope engulfs part of the He-intershell, after the quenching of a TP. We refer to the reviews by \cite{busso99,kaeppeler10} for major details on the AGB nucleosynthesis. \\ Instead, the physical environment for the $r$-process is still unknown, although SNII are the most promising candidates. For elements from Ba to Bi, observations of very metal-poor stars with high $r$-enhancement (e.g., CS 22892--052 \cite{sneden03}) show an abundance distribution well reproduced by a scaled solar $r$-process residual contribution \cite{arlandini99}. Instead, lighter neutron capture elements with Z $\leq$ 47 show values lower than the scaled solar-system $r$-process \cite{wasserburg96,SCG08}. The nucleosynthesis site(s) and the exact contributions from different primary processes to Sr, Y, Zr is highly debated \cite{travaglio04,farouqi10,pignatari10,qian08,montes07}, although also related to massive stars. A large spread is observed for [Eu/Fe] and for [Sr,Y,Zr/Fe] in unevolved halo stars. For [Fe/H] $<$ $-$2, different ranges are observed for Eu and Sr, Y, Zr: $-$1 $\leq$ [Eu/Fe] $\leq$ 2 with an average around 0.5 dex, while $-$1 $\leq$ [Sr,Y,Zr/Fe] $\leq$ 0.5 with an average around 0 dex \cite{travaglio04,montes07}. This may be interpreted as a signature of incomplete mixing in the gas cloud from which these stars have formed \cite{ishimaru99,travaglio01}, as well as an indication of different and uncorrelated primary process contributions. In the last years, a quite large number of % carbon and $s$-process enhanced metal-poor (CEMP-$s$) stars have been detected. CEMP-$s$ are main-sequence/turnoff or giants of low mass ($M$ $<$ 0.9 $M_\odot$). The most plausible explanation for their peculiar high $s$-element abundances is mass transfer by stellar winds from the most massive AGB companion (now a white dwarf). About half of these CEMP-$s$ stars are also highly enhanced in $r$-process elements (CEMP-$s/r$). The observed $r$-enhancement in these stars reflects the observations of unevolved Galactic stars at low metallicity. CEMP-$s/r$ stars show abundance patterns incompatible with a pure $s$-process nucleosynthesis. While a pure $s$-process predicts [La/Eu] $\sim$ 0.8 -- 1.1 (where La and Eu are typical $s$- and $r$-process elements, respectively), CEMP-$s/r$ stars show 0.0 $\leq$ [La/Eu] $\leq$ 0.4, with [La/Fe] and [Eu/Fe] up to $\sim$ 2 dex. Different scenarios have been proposed in the literature to explain the origin of CEMP-$s/r$ (e.g., \cite{jonsell06,cohen03}). \\ We suggest that the molecular cloud from which the binary system formed was already enriched in $r$-process elements by local pollution of SNII ejecta \cite{SCG08,bisterzo09pasa}. This hypothesis is supported by numerical simulations by \cite{vanhala98}, who found that SNII explosion in a molecular cloud may trigger the formation of binary systems. These simulations may explain the very high fraction of CEMP-$s/r$ ($\sim$ 50\%) among the CEMP-$s$. We present here a preliminary analysis of a comparison between AGB theoretical predictions and spectroscopic observations of CEMP-$s$ and CEMP-$s/r$ stars. A detailed discussion will be presented in Bisterzo et al. (in preparation).

To explain the origin of CEMP-$s/r$, we hypothesised that the molecular cloud from which the binary system formed was already enriched in $r$-process elements. Subsequently, the $s$-process elements synthesised by the AGB companion are transferred by stellar winds on to the observed star. The $s$-process nucleosynthesis is not affected by the initial $r$-enhancement of the molecular cloud. However, for high $r$-process enrichment ([r/Fe]$^{\rm ini}$ = 2), one should account for the $r$-process contribution to solar La, Nd and Sm (30\%, 40\%, 70\%). In agreement with the [Y,Zr/Fe] observed in unevolved halo stars, we adopt solar scaled initial Y and Zr values. This increases [hs/ls] by $\sim$ 0.3 dex. This is sustained by observations in CEMP-$s/r$ stars, which show an [hs/ls] ratio in average higher than that observed in CEMP-$s$. Note that the [hs/ls] observed in CEMP-$s/r$ stars may be in agreement with pure $s$-process predictions within the errorbars. A deeper analysis will be given in Bisterzo et al., in preparation.

1012.0456

About half of carbon and s-process enhanced metal-poor stars (CEMP-s) show a high r-process enrichment (CEMP-s/r), incompatible with a pure s-process contribution. CEMP-s stars are of low mass (M &lt; 0.9 Msun) and belong to binary systems. The C and s-process enrichment results from mass transfer by the winds of the primary AGB companion (now a white dwarf). The nucleosynthesis of the r-process, instead, is believed to occur in massive stars exploding as Supernovae of Type II. The most representative r-process element is Eu (95% of solar Eu). We suggest that the r-process enrichment was already present by local SNII pollution in the molecular cloud from which the binary system formed. The initial r-enrichment does not affect the s-process nucleosynthesis. However, the s-process indicators [hs/ls] (where ls is defined as the average of Y and Zr; hs as the average of La, Nd, Sm) and [Pb/hs] may depend on the initial r-enhancement. For instance, the hs-peak has to account of an r-process contribution estimated to be 30% for solar La, 40% for solar Nd, and 70% for solar Sm. A large spread of [Eu/Fe] is observed in unevolved halo stars up to [Eu/Fe] ~ 2. In presence of a very high initial r-enrichment of the molecular cloud, the maximum [hs/Fe] predicted in CEMP-s/r stars may increase up to 0.3 dex. Instead, the spread of [Y,Zr/Fe] observed in unevolved halo stars reaches a maximum of only ~ 0.5 dex, not affecting much the predicted [ls/Fe]. This is in agreement with observations of CEMP-s/r stars that show an observed [hs/ls] in average higher than that observed in CEMP-s. Preliminary results are presented.

12137134

2010evn..confE..81G

1012

1012.2665_arXiv.txt

1012.2665

Following the high detection rate achieved by EVN observations of the central regions of local Seyfert galaxies (Giroletti &amp; Panessa 2009, ApJL 706, 260), we have targeted a few additional sources from a complete sample. We have detected three more sources (NGC 3982, NGC 3227, and NGC 4138) at both 1.6 and 5 GHz and present preliminary results. Moreover, the declination of the sources was suitable to include Arecibo in the EVN observations, which provides important clues on the compactness of the emission region.

12132531

2010arXiv1012.2386P

1012

1012.2386_arXiv.txt

Very little is known about the magnetic fields of hot, massive OB stars, due at least in part to the challenges of measurement. Even as a member of the elusive class of magnetic massive stars, the B0.2 V star $\tau$\,Sco is recognised to be a peculiar and outstanding object. The magnetic field of $\tau$\,Sco is unique because it is structurally far more complex than the mostly-dipolar fields ($l=1$) usually observed in magnetic OB stars, with significant power in spherical-harmonic modes up to $l=5$ with a mean surface field strength of $\sim300$\,G (Donati et al. 2006) $\tau$\,Sco also stands out from the crowd of early-B stars because of its stellar wind anomalies, as diagnosed through its odd UV spectrum (see Fig. 1). The UV wind line morphology of normal, early B stars conforms to a 2-D spectral grid (Walborn et al. 1995). Typically, C\,\textsc{iv} strengthens with increasing temperature and luminosity. N\,\textsc{v} is at most a trace on the main sequence at spectral type B0\,V but strengthens with temperature and luminosity for more luminous stars. For stars with fixed C\,\textsc{iv} and N\,\textsc{v}, Si\,\textsc{iv} is strictly luminosity dependent and breaks the degeneracy. $\tau$\,Sco does not fit into this grid. It has strong N\,\textsc{v} indicating that it should be well above the main sequence. However, its C\,\textsc{iv} lines are only slightly stronger than typical and not distinctly wind-like, suggesting a near main sequence luminosity. Finally, its Si\,\textsc{iv} profiles are unique, a bit stronger than typical class V stars, but unlike normal, early giants. As a result, this stars lie outside of the normal classification grid, which suggests a more highly ionized outflow than typical. The hard X-ray emission of $\tau$\,Sco also suggests hot plasma, in excess of 10\,MK (Cassinelli et al. 1994). \begin{figure} \begin{center} \includegraphics[width=2.1in]{Petit_fig3.eps} \includegraphics[width=2.1in]{Petit_fig2.eps} \includegraphics[width=2.1in]{Petit_fig1.eps} \caption{IUE spectra of $\tau$\,Sco and its analogues HD\,66665 and HD\,63425 (second to fifth spectra). For comparison, a typical spectrum for a B0 dwarf and a B0 giant is shown (bottom and top respectively). The dashed lines indicate the wind lines N\,\textsc{v}\,$\lambda \lambda 1239, 1243$, Si\,\textsc{iv}\,$\lambda\lambda 1393, 1403$ and C\,\textsc{iv}\,$\lambda \lambda 1548, 1550$. } \label{fig1} \end{center} \end{figure} Interestingly, the wind lines of $\tau$\,Sco vary periodically with the star's 41\,d rotation period (Donati et al. 2006). Clearly the magnetic field exerts an important influence on the wind dynamics. What is not clear is whether the wind-line anomalies described above are a consequence of the unusual complexity of $\tau$\,Sco's magnetic field, a general consequence of wind confinement in this class of star, or perhaps even unrelated to the presence of a magnetic field. Because such wind anomalies have never been observed in any other star, magnetic or not, this issue has remained unresolved. The identification and analysis of additional stars with wind properties similar to $\tau$\,Sco would therefore represent an important step toward understanding the origin of these peculiarities.

We have presented the characteristics of two stars - HD\,66665 and HD\,63425 - which we believe are analogues to the magnetic massive star $\tau$\,Sco. The UV spectra of these stars are similar to the once-unique spectrum of $\tau$\,Sco. We have shown that these three stars have similar fundamental properties, although the mass-loss rate values we estimate for HD\,66665 and HD\,63425 are lower than the value generally assumed for $\tau$\,Sco. We have shown that these two stars host a magnetic field. Our modelling of the LSD Stokes V profiles by an inclined dipole model results in field strengths that are comparable with, or maybe slightly larger than, the mean surface field strength of $\tau$\,Sco. The current observations can be acceptably reproduced by the dipole model, although more phase-resolved observations are required in order to assess the potential complexity of their magnetic field, and verify if the wind anomalies are linked to the field complexity.

1012.2386

The B0.2 V magnetic star tau Sco stands out from the larger population of massive magnetic OB stars due to its high X-ray activity and remarkable wind, apparently related to its peculiar magnetic field - a field which is far more complex than the mostly-dipolar fields usually observed in magnetic OB stars. tau Sco is therefore a puzzling outlier in the larger picture of stellar magnetism - a star that still defies interpretation in terms of a physically coherent model. Recently, two early B-type stars were discovered as tau Sco analogues, identified by the striking similarity of their UV spectra to that of tau Sco, which was - until now - unique among OB stars. We present the recent detection of their magnetic fields by the MiMeS collaboration, reinforcing the connection between the presence of a magnetic field and wind anomalies (Petit et al. 2010). We will also present ongoing observational efforts undertaken to establish the precise magnetic topology, in order to provide additional constrains for existing models attempting to reproduce the unique wind structure of tau Sco-like stars.

12131881

2010arXiv1012.1499I

1012

1012.1499_arXiv.txt

The photometric decomposition of disk galaxy images is a long-standing classical topic in studies of galaxy structure. However, the Sloan Digital Sky Survey (SDSS) has provided new possibilities due to its depth and large sky area covered by the observations. The statistics of structure characteristics obatined with the SDSS data for nearby late-type galaxies has been already analysed \cite{pt2006}. Here we present some results for a sample of early-type disk galaxies.

After having decomposed the SDSS photometric images of 85 early-type disk galaxies lacking large-scale bars, we conclude: \begin{itemize} \item{all early-type disk galaxies in our sample lacking large-scale bars demonstrate double-tier large-scale stellar disks; } \item{the outer and inner stellar disks have probably different nature: the central surface brightness distribution is strongly bimodal, with the inner and outer disks having well separated peaks, and at the diagram `$\mu _0$ vs $r_0$' the outer and inner disks represent two quite separate point clouds; } \item{the galaxies with nuclear bars and those lacking them do not separate at any diagram analyzed here; we conclude that the presence of nuclear bars does not relate with any structure re-building at large scales. } \end{itemize} The work is supported by the grant of Russian Foundation for Basic Researches no. 10-02-00062a.

1012.1499

The results of photometric decomposition of surface brightness distributions in 85 early-type unbarred galaxies are presented. The SDSS r-images are analysed. Double-tiered exponential disks are found in all galaxies which are studied; the statistics of the disk parameters is derived.

12202004

2011ASPC..448.1173L

1012

1012.1720_arXiv.txt

Star-planet interactions (SPI) in stellar chromospheres could occur as recurring enhanced Ca \textsc{ii} H \& K flux following the periodicity of the planet orbiting its host star. So far signs for SPI were found only in few systems and the empirical evidence is still heavyly debated \citep{popp, shkol1, scharf}. The possible physical scenarios for SPI are gravitational and/or magnetic interactions. Gravitational interaction in the system could lead to tidal bulges on the surface of the star, changing the local geometry. It is perhaps conceivable that this may in turn favor magnetic reconnection thereby enhancing the observed stellar activity. Tidal interaction would lead to two enhancement peaks over the planetary phase. In the case of magnetic interaction, the reconnection of field lines from the planet and the star could lead to enhanced stellar activity peaking once per orbit \citep{cuntz}. Recent observations for HD 179949 even suggest an ``On/Off" behavior: after some years the periodicity of the Ca \textsc{ii} K emission changed from the planetary to the stellar rotation period \citep{shkol1}.

We find no flux variations in phase with the planetary orbit in the chromospheric Ca~\textsc{ii}~K and H$\alpha$ lines in the spectra of our sample stars that have close-in planets. The detected variability of $\upsilon$And is probably due to the rotation of the star and not bound to the planet. HD 41004 AB potentially is a very good candidate to observe interactions since the mass-ratio of the M dwarf and its brown dwarf companion is between that of interacting binaries and star-planet systems. More spectra will be obtained in the future to search for periodicities of the flux variations.

1012.1720

Massive planets in very close orbits around their central stars can induce so-called star-planet interactions (SPI), which may be of magnetic or gravitational nature. In both cases, SPI can potentially cause recurring chromospheric emission on the host star visible in Ca II H &amp; K and/or Hα. The emission would be bound to the planetary orbit, not to the rotation period of the star. We searched for SPI in a sample of 7 stars with massive close-in planets using high-resolution spectroscopic data taken at HRS (HET) and FEROS (La Silla). We find no periodically recurring emission in the planet-hosting stars. In the case of HD 41004 AB, a binary system consisting of a K dwarf and an M dwarf, where the M dwarf is orbited by a brown dwarf companion, we find signs of cyclic variation in the Ca II K and Hα emission lines that could be associated to interactions between the M dwarf and its companion. We present our first results of this interesting system that may become an important system for the understanding of SPI.

3396492

2011IAUS..272..638V

1012

1012.3403_arXiv.txt

\begin{figure}[t] \begin{center} \includegraphics[width=2.0in]{S6-20_Delia_Volpi_fig1.eps} \includegraphics[width=2.0in]{S6-20_Delia_Volpi_fig2.eps} \caption{Flux in the radio band at 6 cm for Cyg~OB2 No.~8A: on the left the observations from \cite[Blomme et al. (2010)]{Blomme_etal10}, on the right our simulated results. Periastron is at phase $\approx 0$} \label{fig1} \end{center} \end{figure} During recent years many OB stars have been discovered to be binary systems. Non-thermal radio emission is observed to be produced by some of these binary stars. The non-thermal emissivity is thought to be due to synchrotron emission radiated by relativistic electrons. The electrons are accelerated up to high energies by strong shocks produced by the collision between the two radiatively driven stellar winds (\cite[Eichler \& Usov 1993]{EichlerUsov93}). Several parameters of the system can be constrained by the synchrotron emission, among them the mass loss rates from the primary and the secondary. Investigating the synchrotron radiation is thus necessary. We model the non-thermal emission for two colliding wind systems, Cyg~OB2 No.~8A and Cyg~OB2 No.~9, and compare the obtained results with the observations.

1012.3403

Some OB stars show variable non-thermal radio emission. The non-thermal emission is due to synchrotron radiation that is emitted by electrons accelerated to high energies. The electron acceleration occurs at strong shocks created by the collision of radiatively-driven stellar winds in binary systems. Here we present results of our modelling of two colliding wind systems: Cyg OB2 No. 8A and Cyg OB2 No. 9.

12217826

2011PhLB..698....1P

1012

1012.0546_arXiv.txt

1012.0546

The tight-coupling approximation (TCA) used to describe the early dynamics of the baryons-photons system is systematically built to higher orders in the inverse of the interaction rate. This expansion can be either used to grasp the physical effects by deriving simple analytic solutions or to obtain a form of the system which is stable numerically at early times. In linear cosmological perturbations, we estimate numerically its precision, and we discuss the implications for the baryons acoustic oscillations. The TCA can be extended to the second order cosmological perturbations, and in particular we recover that vorticity is not generated at lowest order of this expansion.

12234596

2012APh....35..421H

1012

1012.5155_arXiv.txt

Cosmic neutrinos, as a probe of the universe to the highest energy regime, are indeed wonderful in many aspects. Due to their extremely small interaction cross section, they can penetrate through galactic infrared (IR) and cosmic microwave background (CMB) photons, while photons of energy above 10 TeV would be attenuated. Furthermore, being uncharged, they propagate along straight lines and are therefore able to point directly back to their sources, while protons or other charged particles would be deflected by the magnetic fields in the universe. Ultra-high energy cosmic rays (UHECRs) have been observed up to $\approx$ $10^{19.6}$ eV. The source of such amazingly energetic events remains a mystery. Above this energy scale, UHECRs interact with CMB photons through the Greisen-Zatsepin-Kuzmin(GZK) processes~\cite{gzk_process}. The GZK cut-off of the cosmic ray energy spectrum was first observed by the High Resolution Fly's Eye Experiment~\cite{cutoffHiRes} and later confirmed by the Pierre Auger Observatory~\cite{cutoff}. As such, the corresponding GZK neutrinos, a necessary by-product of the GZK process, are almost guaranteed to exist. Nevertheless, none of these have been observed so far. Detection of the GZK neutrinos would provide critical information for unraveling the mystery of the origin and evolution of cosmic accelerators, and will be one of the most exciting prospects in the coming decade~\cite{pisin_whitepaper}. A promising way of detecting UHE neutrinos is the radio approach. When an ultra-high energy cosmic neutrino interacts with ordinary matters on the Earth, it would lead to a hadronic debris, caused by either charged-current or neutral-current weak interaction. The former also produces a lepton with corresponding flavor. Both the high energy leptons and hadronic debris induce particle showers. As proposed by Askaryan in the 1960's~\cite{Askaryan}, the high energy particle shower developed in a dense medium would have net negative charges. This charge imbalance appears as a result of the knocked-off electrons being part of the shower, as well as positrons in the shower annihilating with electrons of the medium. The net charges of the showers, typically $20 \%$ of total shower particles, serve as a source which emits the Cherenkov radiations when they travel in the medium. The sizes of the showers are quite localized (tens of cm in radial and a few meters in longitudinal development) compared to those developed in the air (km scale), and therefore result in coherent radiations for wavelengths longer than the shower sizes. The corresponding coherent wavelength turns out to be in the radio band, from hundreds of MHz to a few GHz. This Askaryan effect has been confirmed in a series of experiments at the Stanford Linear Accelerator Center (SLAC), where different dense media such as silica sand, rock salt and ice were used~\cite{aska1, aska2, aska3, aska4}. Various experiments have been proposed based on the idea of Askaryan: the balloon-borne antenna array (e.g. ANITA~\cite{anita1, anita2}), the ground-based antenna array buried in salt (e.g. SalSA~\cite{salsa}) or ice (e.g. RICE~\cite{rice}), and the radio telescope searching for lunar signals (e.g. LUNASKA~\cite{lunaska}). As ground-based experiments have the advantage of low noise and low energy threshold, they play an important role in the next-generation experiments aimed at detecting some tens of GZK neutrinos per year. However, for the extremely high energy neutrino event, especially for the electron neutrino, it is very likely that the longitudinal size of the shower would become comparable to the distance between the antenna and the shower. Under this situation, the common assumption of far-field radiation does not apply and near-field effects become nonnegligible. In this paper, we study the impact of the near-field effect to the radiation pattern via a numerical method. The organization of this paper is as follows: Section II will discuss the underlying physics of the shower elongation. Section III will introduce the numerical method and our setup. Section IV will present the implementation of parallel computing to gain a satisfying efficiency of our numerical calculation. In Section V, we analyse the results in both time-domain and frequency-domain. We discuss the features of the near-field radiation pattern and compare the far-field pattern between our results and the theoretical ones for validation.

We have studied the near-field radiation in both time domain and frequency domain. We shown that even for a shower with symmetric longitudinal development (e.g. Gaussian distribution), the resulting near-field waveform would be asymmetric in time. The longitudinal development is as important as the lateral distribution even at the Cherenkov angle. Future work on the parameterization of near-field radiation should take this into account. As the radiation propagates, the waveform would gradually become symmetric. Moreover, this transition occurs at different $R$ for different $\theta$. To be specific, it occurs at shortest distance for the observer located at the Cherenkov angle. For a ground array neutrino detector, with the size of LPM-elongated showers becomes comparable with the typical detection distance, the near-field effect is an indispensable factor. The correct relation of distance dependence in near-field prevents underestimation of the signal strength in a Monte Carlo simulation. Furthermore, the correct angular spread in near-field is necessary in order not to underestimate the detection solid angle of a neutrino detector. The Fraunhofer approximation leads to an angular spread that is quite narrow for LPM-elongated showers. It is incorrect since a shower of hundred meters long would at least generate radiation which also spans hundred meters long in near-field. The overall detector sensitivity should be better than adopting traditional radiation formula in Fraunhofer limit in the Monte Carlo simulation. We plan to find the parameterization formula suitable in all cases in the future work. On the other hand, due to the complicated features of near-field radiation, new reconstruction methods are required. For example, the normal way to reconstruct the direction of incoming Cherenkov pulses is by the arrival time differences between antennas. This method treats the wavefront as a spherical one and is based on the far-field assumption which fails in near-field as we have shown. Charged-current interaction of electron neutrino is the main source of ultra-high energy electromagnetic showers. In a typical detection distance, the near-field condition should be easily satisfied. Identification of such radiation therefore implies electron neutrino events. This opens an opportunity for neutrino flavor identification, since muon and tau neutrinos only induce hadronic showers whose sizes are too compact to produce near-field radiation in a typical detection distance.

1012.5155

The radio approach based on the Askaryan effect for detecting the ultra-high energy cosmic neutrinos has become a mature experimental technique. So far the existing calculations of the Cherenkov radiation associated with the Askaryan effect has been mostly based on the far-field approximation, whose validity maybe challenged when the detector is close to the event. In this paper we present an alternative approach to calculate the Cherenkov pulse by a numerical code based on the finite difference time-domain (FDTD) method. This approach has the advantage of providing the solution everywhere in space, contrary to other methods that rely on the far-field approximation. We also present a one-dimensional theoretical model for the shower with analytical solution, which helps to elucidate our nonzero-width simulation results. We show that for a shower with symmetric longitudinal development, the resulting near-field waveform would be asymmetric in time. In addition, we demonstrate that for a shower elongated by the LPM (Landau-Pomeranchuk-Migdal) effect and thus with a multi-peak structure, a bipolar, asymmetric waveform is still preserved in the near-field regime irrespective of the specific variations of the multi-peak structure, which makes it a generic, distinctive feature. This should provide an important characteristic signature for the identification of ultra-high energy cosmogenic neutrinos.

12167936

2011ApJ...728..127D

1012

1012.5363_arXiv.txt

The neutral interstellar medium (ISM) of the Galactic disk is riddled with loops, shells and cavities \citep[e.g.][]{heiles79,mcclure02,ehlerova05}. The largest of these structures -- H{\sc i} supershells -- may reach hundreds of parsecs in diameter, with formation energies as high as % $\sim10^{53}$ ergs. Such objects strongly influence the structure and evolution of the Disk ISM. The dominant paradigm holds that supershells are formed through the cumulative action of multiple stellar winds and supernovae, which blow hot, overpressurized bubbles, and sweep up the surrounding medium into cool, dense % shells \citep{bruhweiler80,tomisaka81,mccray87}. The accumulation of the ISM in such superstructures is one means of generating the high densities and column densities required for the production of molecular gas, and supershells have long been suggested as drivers of molecular cloud formation \citep[e.g.][]{mccray87,mashchenko94,fukui99,hartmann01}. However, despite a substantial body of theoretical work \citep[e.g.][]{koyama00,bergin04,heitsch08,inoue09}, and an growing list of articles documenting supershell-associated molecular gas \citep[e.g.][]{handa86, jung96, fukui99, matsunaga01, yamaguchi01b, dawson08b}, conclusive observational evidence of this phase change occurring in the walls of shells has not yet been found. The degree to which large-scale stellar feedback is driving the production of molecular clouds remains unconstrained. Conversely, in a pre-structured and highly inhomogeneous ISM, an expanding supershell will undergo encounters with pre-existing dense clouds, which may compress, disrupt, fragment, or even completely destroy them \citep[e.g.][]{klein94,foster96,mellema02}. The relationship between supershells and the molecular ISM should therefore be governed by the interplay between this potentially destructive process and the creation of new molecular material. This paper presents parsec-scale resolution observations of atomic hydrogen and $^{12}$CO(J=1--0) in two Galactic shells, GSH 287+04--17 and GSH 277+00+36, with the aim of investigating the role played by supershells in the evolution of the molecular ISM. These are some of the highest resolution images of any supershell, and the first time a dedicated comparison between the atomic and molecular material in shell walls has been performed. In the remainder of this introduction we briefly summarize the basic properties of the two shells, before moving on to describe the observational and data reduction techniques in \S\ref{observations}. Sections \ref{results:obschar} and \ref{fitting} discuss the distribution and morphology of the two tracers, and examine the properties of the two phases quantitatively via Gaussian fitting to the shell spectra. Section \ref{coenhancement} investigates the degree of molecularization in the shell volumes, finding evidence of enhanced molecular fractions in both objects, and \S\ref{totalmass} gives their total H{\sc i} and H$_2$ masses. In \S\ref{insitu} and \S\ref{preexisting} we discuss H{\sc i}-CO structures in the shell walls in the context of different formation scenarios. Section \ref{stars} takes a preliminary look at star formation activity in the shell molecular clouds, and our conclusions are finally summarized in \S\ref{conclusions}. \subsection{GSH 287+04--17: Basic Properties} \label{gsh287:basics} GSH 287+04--17, also known as the `Carina Flare' supershell, is a medium-sized, gently expanding Galactic chimney, located in the Sagittarius-Carina Arm at a distance of $2.6\pm0.4$ kpc and Galactocentric radius of $\sim8$ kpc. It was originally discovered in $^{12}$CO(J=1--0) by \citet{fukui99}, who reported a scattering of molecular clouds extending in a wide swath above the Galactic Plane. These clouds showed telltale signs of global expansion, leading the authors to correctly identify them as parts of an expanding superstructure. \citet{dawson08a} later confirmed the existence of the counterpart atomic shell using low resolution $\sim16'$ H{\sc i} data from the Southern Galactic Plane Survey \citep[SGPS;][]{mcclure05}, revealing its global properties and large-scale morphology for the first time. The main body of the shell measures $\sim230\times360$ pc, and it has broken out of the disk at a height of $z\sim280$ pc, with a associated high-latitude emission seen up to heights of $\sim450$ pc above the midplane. The molecular clouds form co-moving parts of the shell, which is estimated to contain total H{\sc i} and H$_2$ masses of $M_{\mathrm{HI}}\sim7\pm3\times10^5~M_{\odot}$ and $M_{\mathrm{H}_2}\sim2.0\pm0.6\times10^5~M_{\odot}$, respectively. The expansion velocity of the shell is $\sim10$ km s$^{-1}$, and its age and formation energy are estimated to be $\sim1\times10^7$ yr and between $0.5-1\times10^{52}$ ergs, based on comparisons with analytical and numerical models. \subsection{GSH 277+00+36: Basic Properties} \label{gsh277:basics} GSH 277+00+36 is a large outer Galaxy chimney, located at the edge of the Sagittarius-Carina spiral arm at a kinematic distance of $6.5\pm0.9$ kpc and Galactocentric radius of $\sim10$ kpc. It was originally discovered in the low resolution portion of the SGPS \citep{mcclure00}, where it is seen as a prominent void in the bright H{\sc i} emission of the Plane. With a main body $\sim610$ pc in diameter and chimney extensions that reach more than 1 kpc above the midplane, this large and evolved supershell has swept up an estimated atomic mass of $M_{\mathrm{HI}}\sim3\pm1\times10^6~M_{\odot}$, and is expanding with a velocity of $\sim20$ km s$^{-1}$. The age and formation energy of the shell are estimated to be between $1-2\times10^7$ yr and $0.9-2.4\times10^{53}$ ergs, where the higher values in these ranges are taken directly from \citet{mcclure00} and the lower values are derived using the same analytical formula as \citet{dawson08a} in order to better illustrate the difference between the two shells. High-resolution ($\sim3'$) H{\sc i} observations by \citet{mcclure03} revealed narrow, well-defined walls, sharply delineated along their inner edges, with a great deal of complex substructure including knots, filaments and `drips'. No molecular observations have been previously reported for this shell.

\label{conclusions} We have presented parsec-scale resolution observations of H{\sc i} and $^{12}$CO(J=1--0) in two Galactic supershells, GSH 287+04--17 and GSH 277+00+36, with the aim of investigating the role played by supershells in the evolution of the molecular ISM. The main findings may be summarized as follows. 1. Both shells contain large quantities of associated molecular gas in the form of discrete co-moving CO clouds distributed throughout the atomic shell walls. These high resolution observations reveal rich substructure in both tracers. Molecular gas is seen elongated along the inner edges of atomic shell walls, embedded within H{\sc i} filaments and clouds, or taking the form of small CO clouds at the tips of tapering `fingers' of H{\sc i}. We note a similarity in features observed in both objects, despite differences in location and evolutionary stage. 2. The atomic shell walls are dominated by cold gas, showing narrow linewidths reaching as low as $\Delta v\sim2$ km s$^{-1}$. Mean temperatures and densities are estimated to be roughly $T_k\sim100$ K and $n_0\sim10$ cm$^{-3}$ in regions where the spectral signature from the shells is well determined. 3. An enhanced level of molecularization is observed over the volumes of both GSH 287+04--17 and GSH 277+00+36, providing the first direct observational evidence of increased molecular cloud production due to the influence of supershells. Our results imply that the amount of molecular matter in the volumes affected by these shells is enhanced by $\sim3$ times with respect to neighboring regions. If this is found to hold true on Galactic scales, it has powerful implications for our understanding of the role played by stellar feedback on the evolution of the molecular phase. Already, the confirmation of this phenomenon in two shells at quite different locations and evolutionary stages is very compelling, and we suggest that more followup work should be undertaken to repeat this analysis for other objects. 4. CO clouds embedded in the main atomic shell walls provide excellent candidates for the in-situ formation of molecular gas from the swept up medium. This scenario has been explored in detail for archetypal examples of embedded molecular clouds, demonstrating that the formation timescales implied by theory are consistent with shell ages and number densities, and that the requirements for the shielding of the CO molecule are met. We thus confirm on a local scale the viability of the triggered formation of molecular clouds due to the influence of the shells. 5. Small offset CO clouds located at the tips of tapering `fingers' of H{\sc i} may be the remnants of molecular gas present in the ISM prior to the formation of the shells. We have demonstrated the plausibility of this scenario for archetypal examples of such offset clouds, by comparing estimates of cloud destruction times with the survival times implied by the data. 6. A preliminary examination of YSO candidates in the shell regions confirms that active star formation is occurring in the molecular clouds in both GSH 287+00-17 and GSH 277+00+36. This includes massive star forming regions in embedded molecular cloud complexes, as well as less luminous YSO candidates associated with offset CO clouds. \\ We wish to thank the anonymous referee for comments that led to the improvement of this manuscript. We also thank Shu-ichiro Inutsuka and Tsuyoshi Inoue, whose helpful discussions have contributed to this work. We gratefully acknowledge the past staff and students of Nagoya University who made the CO observations utilized in this paper. The NANTEN project was based on a mutual agreement between Nagoya University and the Carnegie Institute of Washington, and its operation was made possible thanks to contributions from many companies and members of the Japanese public. The Australia Telescope Compact Array and Parkes Telescope are part of the Australia Telescope which is funded by the Commonwealth of Australia for operation as a National Facility managed by CSIRO. This research has made use of the NASA/IPAC Infrared Science Archive, which is operated by the Jet Propulsion Laboratory, California Institute of Technology, under contract with the National Aeronautics and Space Administration. We also gratefully acknowledge the Southern H-Alpha Sky Survey Atlas (SHASSA), which is supported by the National Science Foundation.

1012.5363

We present parsec-scale resolution observations of the atomic and molecular interstellar medium in two Galactic supershells, GSH 287+04-17 and GSH 277+00+36. H I synthesis images from the Australia Telescope Compact Array are combined with <SUP>12</SUP>CO(J = 1-0) data from the NANTEN telescope to reveal substantial quantities of molecular gas closely associated with both shells. These data allow us to confirm an enhanced level of molecularization over the volumes of both objects, providing the first direct observational evidence of increased molecular cloud production due to the influence of supershells. We find that the atomic shell walls are dominated by cold gas with estimated temperatures and densities of T ~ 100 K and n <SUB>0</SUB> ~ 10 cm<SUP>-3</SUP>, respectively. Locally, the shells show rich substructure in both tracers, with molecular gas seen elongated along the inner edges of the atomic walls, embedded within H I filaments and clouds, or taking the form of small CO clouds at the tips of tapering atomic "fingers." We discuss these structures in the context of different formation scenarios, suggesting that molecular gas embedded within shell walls is well explained by in situ formation from the swept-up medium, whereas CO seen at the ends of fingers of H I may trace remnants of molecular clouds that pre-date the shells. A preliminary assessment of star formation activity within the shells confirms ongoing star formation in the molecular gas of both GSH 287+04-17 and GSH 277+00+36.

12167920

2011ApJ...728..125A

1012

1012.3365_arXiv.txt

Precise timing of exoplanet transits over many years is a powerful technique to learn more about a planetary system. Transit timing variations, or TTVs, can arise from many potential scenarios, such as orbital precession, orbital decay, and perturbations by additional planets, small stellar companions, large moons or rings; these effects have been modeled in numerous papers over the last decade \citep{MiraldaEscude2002, Sasselov2003, Holman2005, Agol2005, Heyl2007, Ford2007, Simon2007, Kipping2009a, Kipping2009b, Levrard2009}. Recently, \citet{Holman2010} announced the first incontrovertible evidence of transit timing variations in the Kepler-9 system, where a pair of Saturn-mass planets are trapped in a 2:1 mean-motion resonance, producing strong, complementary deviations of over an hour in the midtimes of both transits. Other observational evidence is less clear cut, with recent claims of timing variations in WASP-3b \citep{Maciejewski2010a} and WASP-10b \citep{Maciejewski2010b} announced but unconfirmed, while a claimed detection of TTVs for OGLE-TR-111b has not held up \citep{Adams2010a, Diaz2008}. Additionally, there have been claims of a tentative detection of variations in the transit duration of GJ436b, increasing by roughly 3 minutes per year \citep{Coughlin2008}, consistent with the $<8~M_{\oplus}$ limit for a second planet in this system placed by transit timing \citep{Bean2008}. Finally, previous work by our group has found very tentative evidence that the orbital period of OGLE-TR-113b is decreasing at a rate of $-60\pm15$ ms yr$^{-1}$, though additional observations are needed to confirm this result \citep{Adams2010b}. Though perhaps less exciting, the $non$-detection of transit timing variations is both important and expected. It has been suggested that systems with hot Jupiters are statistically less likely to be part of multi-planet systems \citep{Wright2009}, and transit timing limits can place tighter mass limits on potential nearby companions than other detection methods to date, with upper limits of a few Earth masses possible for companions in strong mean-motion resonances. With observations spanning many years, upper limits can also be placed on the rate of orbital decay of the known hot Jupiter, providing valuable direct observational constraints on theoretical models of tidal dissipation in stars. Currently tidal dissipation in stars is only estimated indirectly, through arguments about circularization timescales of binary stars in clusters \citep[see e.g.][]{Meibom2005, Ogilvie2007}. In this work we present the case of the hot Jupiter OGLE-TR-132b, with seven new light curves observed and no timing variations larger than $-108 \pm 49$ s detected during observations spanning seven years. In Section~\ref{section:obs} we describe the observations, data analysis and light curve fitting. Section~\ref{section:fitting} describes the modeling and improvements to the system parameters, and places limits on parameter variability between transits. Section~\ref{section:timing} discusses the timing constraints placed on the planetary system. We conclude in Section~\ref{section:conclusions}.

\label{section:conclusions} We have measured seven new transits of the hot Jupiter OGLE-TR-132b, which triples the number of transits available for this planet to date. The new transits have timing precisions between 47 and 87 seconds and photometric precision from 1-2 mmag in two minute bins. The updated observational dataset of OGLE-TR-132b spans seven years (2002--2009) and a total of ten transits after combining the new transits presented in this paper with three previously published transit epochs. We find no evidence of transit timing variations with amplitudes larger than $-108 \pm 49$ s, which allows us to place significant constraints to the presence of additional close-in planets in the system. In particular, near both the internal and external 2:1 and 3:2 mean motion resonances, objects as small as 5-10 Earth masses would have been detectable if they existed. We also tested for orbital period decay, with the best fit rate of period decrease, $\dot{P}_{132b}=-125\pm80$ ms yr$^{-1}$, consistent with no change. The predicted orbital period decay rate for this planet is 20 ms yr$^{-1}$ if $Q_{*} = 10^5$, and 140 ms yr$^{-1}$ if $Q_{*} = 1.6 \times 10^4$ as suggested for the star OGLE-TR-113 \citep{Adams2010b}, so additional transits with better timing precision ($<30$ s) or a significantly longer time baseline are needed to search for slower rates of period change. Finally, we have measured a revised radius value for the planet, $R_p = 1.23 \pm 0.07$ $R_J$, which is fully consistent with the radius derived by \citet{Southworth2008} from a reanalysis of the \citet{Gillon2007} and \citet{Moutou2004} data. We used simple aperture photometry to extract all transit light curves, so the depth of our transits should not be affected by the potential normalization problems found when using image differencing techniques, as reported for this same planet by \citet{Gillon2007}. Our planet radius provides an independent confirmation, within errors, of the transit depth obtained by \citet{Gillon2007} using image deconvolution techniques.

1012.3365

We report the results of the first transit timing variation analysis of the very hot Jupiter OGLE-TR-132b, using 10 transits collected over a seven-year period. Our analysis combines three previously published transit light curves with seven new transits, which were observed between 2008 February and 2009 May with the new MagIC-e2V instrument on the Magellan Telescopes in Chile. We provide a revised planetary radius of R<SUB>p</SUB> = 1.23 ± 0.07R<SUB>J</SUB> , which is slightly larger, but consistent within the errors, than that given by previously published results. Analysis of the planet-to-star radius ratio, orbital separation, inclination, and transit duration reveals no apparent variation in any of those parameters during the time span observed. We also find no sign of transit timing variations larger than -108 ± 49 s, with most residuals very close to zero. This allows us to place an upper limit of 5-10 M <SUB>⊕</SUB> for a coplanar, low-eccentricity perturber in either the 2:1 or 3:2 mean-motion resonance with OGLE-TR-132b. We similarly find that the data are entirely consistent with a constant orbital period and there is no evidence for orbital decay within the limits of precision of our data. <P />This paper includes data gathered with the 6.5 m Magellan Telescopes located at Las Campanas Observatory, Chile.

12164667

2011AJ....141...68Z

1012

1012.1145.txt

Dwarf galaxies play an important role in our understanding of the formation and evolution of galaxies. In the hierarchical model of galaxy formation, they are proposed to be the building-blocks from which larger systems have been created by merging (Kauffmann et al. 1993). However, these building-block galaxies are too small and faint to be studied at high redshifts. Therefore, of special interest are the studies aimed at the search for some local example galaxies which might have much in common with the galaxies at high redshifts. Blue compact dwarf galaxies (BCDs) are such local example galaxies, whose characteristics are believed to have been common among unevolved low-mass galaxies at intermediate to high redshift. BCDs are small, gas-rich (H\,{\sc i} mass fraction typically higher than 30\%; e.g. Thuan \& Martin 1981; Salzer et al. 2002) and metal-poor ($1/50 \leq Z \leq 1/3\ Z_\odot$; Hunter \& Hoffman 1999) extragalactic objects. They have dramatically different properties compared to normal dwarf galaxies (Zwicky 1966; Gil de Paz et al. 2003, hereafter G03), and are spectroscopically characterized by a faint, blue optical continuum accompanied, in most cases, by strong narrow emission lines, due to the intense star formation activity in one or several star-forming regions (Cair\'os et al. 2001). Stars are formed at high rates in BCDs (Fanelli et al. 1988), exhausting their gas content with a timescale of $10^9$ yr (e.g. Lee et al. 2002), which is much shorter than the age of the Universe. This fact, combined with the low metal abundances, led Sargent \& Searle (1970) to suggest that BCDs are primeval galaxies undergoing star formation for the first time. The question as to whether BCDs are truly young galaxies or they are old galaxies exhibiting a strong starburst has been widely discussed for many years (e.g. Sargent \& Searle 1970; Searle et al. 1973; Schulte-Ladbeck et al. 2000, 2001a; Izotov \& Thuan 1999, 2004; Aloisi et al. 2005). Most BCDs are known to have a red low surface brightness background of presumably older stars (e.g., Loose \& Thuan 1986; Gil de Paz \& Madore 2005; Caon et al. 2005). Moreover, some individual red giant stars (RGBs) in the nearest BCDs can be resolved (e.g., Lynds et al. 1998; Drozdovsky et al. 2001; Corbin et al. 2008). This evidence ruled out the possibility that the majority of BCDs began to form stars within the last billion years. An alternative view proposed by Searle et al. (1973) is that BCDs are chemically primitive objects which experience an episodic star formation history (SFH). The star formation in BCDs occurs in intense bursts which are interleaved by long quiescent periods. This traditional picture also seems to be problematic. Several nearby BCDs with resolved stars do not show evidence of long gaps ($>1$ Gyr) in their recent star formation history (Schulte-Ladbeck et al. 2001a; Crone et al. 2002; Annibali et al. 2003; McQuinn et al. 2009). S\'anchez Almeida et al. (2008) identified a large sample of quiescent counterparts to BCDs (QBCD) in SDSS Data Release 6 (DR6; Adelman-McCarthy et al. 2008), which might support this recursive BCD phase scenario. They also argued that the quiescent phase could last 30 times longer than the starburst phase, and thus the quiescent period of BCDs might be several hundred Myrs under the assumption of a single 10 Myr long starburst per BCD phase. However, S\'anchez Almeida et al. (2008) showed that only $\sim 15\%$ of their QBCD candidates exhibit H$\alpha$ absorptions and therefore are lack of star-forming activities. Combining the results described above with the fact that very few late-type and H\,{\sc i} selected galaxies have no observable H$\alpha$ emissions (Meurer et al. 2006; James et al. 2008; Kennicutt et al. 2008), the SFR of the quiescent phase is probably not zero, but rather may be characterized by a continuous, low level of activity as discussed in Lee et al. (2009). The detailed studies of SFHs of BCDs are rather difficult. For the most nearby systems, where individual stars can be resolved, the Color-Magnitude Diagrams (CMDs) synthesis method could be used to reconstruct the SFHs (e.g. I Zw 18, Aloisi et al. 1999, 2007; Mrk 178, Schulte-Ladbeck et al. 2000; I Zw 36, Schulte-Ladbeck et al. 2001a; NGC 1705, Annibali et al. 2003). However, this method can not be applied to more distant objects, which are the great majority of BCDs. The only way to study their stellar populations is to compare their integrated properties with the predictions of evolutionary synthesis models. There have been several works using (evolutionary) synthesis models to study the observed spectral energy distribution of BCDs (e.g. Mas-Hess \& Kunth 1999; Guseva et al. 2001; Corbin et al. 2006). However, these studies only include one or a few galaxies. In this paper, we present a detailed study on the stellar populations of a BCD sample containing 31 galaxies, using a Simple Stellar Population (SSP) synthesis method. This is capable of yielding the various stellar components, the internal reddening (both stellar continuum and nebular line emission) and the SFH. The pure emission-line spectra allow us to measure the fluxes of emission lines more accurately. We can also study the star formation properties of these 31 BCDs by the using the continuum-subtracted, narrow band \Ha\ imaging data (G03) combined with the results presented in this work. The paper is organized as follows: Section 2 describes the BCD sample and our data reductions. Our results and analysis are given in Section 3. We discuss the uncertainties of the stellar population synthesis and the aperture effects of the fiber spectra in Section 4, and summarize our results in the last section. Where required we adopt a Hubble constant of $H_0 = 70\ {\rm km s}^{-1}$ Mpc$^{-1}$, $\Omega_{\rm M} = 0.3$ and $\Omega_\Lambda=0.7$.

In this section we mainly discuss the uncertainties of the spectral synthesis and the aperture effects. \subsection{Uncertainties of the Spectral Synthesis} The SEAGal Group has used the {\sc starlight} to analyze several large samples of SDSS galaxies as shown in their series of works (C05; Cid Fernandes et al. 2007; Mateus et al. 2006; Asari et al. 2007). As mentioned in Section 3.1, they have tested the uncertainties of the resulted stellar populations. In the study of C05, they found that the individual components of {\boldmath $x$} are very uncertain, whereas the binned vectors of {\boldmath $x$}, i.e. the young, inter-mediate and old populations, have uncertainties less than 0.05, 0.1 and 0.1, respectively, for S/N$\geq$10. For each fitting, {\sc starlight} provides the last-chain-values of the contributed light (mass) fraction of each SSP, $\chi^2$ and mass, for seven Markov chains. For the results of our sample, the median values of the rms of these adopted values are 1.8\%, 17.4\%, 8.8\%, 8.1\%, 21.0\%, 1.9\%, 0.3\%, 0.5\%, 1.8\%, 0.06\%, 5.1\% for $x_{\rm Y}$, $x_{\rm I}$, $x_{\rm O}$, $\mu_{\rm Y}$, $\mu_{\rm I}$, $\mu_{\rm O}$, $\left<\log t_\star\right>_L$, $\left<\log t_\star\right>_M$, $A^\star_V$, $\chi^2$ and $M_\star$, respectively. Therefore, the uncertainties of the resulting stellar populations and stellar mass will not much affect our conclusions. Most of the age sensitivity comes from the 4000 \AA\ break and continuum shape in the fitting, thus degeneracy with extinction should be discussed carefully. Stellar mass-related parameters (e.g. $\left<\log t_\star\right>_M$ and $M_\star$) will suffer more severely from this degeneracy because of the non-constant mass-to-light ratio of stars. As described in Section 3.3, the dust extinction is considered as a variable parameter in {\sc starlight}, and it is overestimated (comparing to $A_{V,\,{\rm neb}}$) for at least half of our sample, which might result in a large underestimation of $\left<\log t_\star\right>_M$. For example, out of the 11 galaxies whose fitted $\left<\log t_\star\right>_M$ are less than 1 Gyr based on the fiber spectra (see Table 2), seven have their $A^\star_V$ overestimated. To check to what extent this overestimation will affect the fitted results, we re-run {\sc starlight} using the same parameters except that we limit $A^\star_V$ to the value of $A_{V,\,{\rm{neb}}}$ for ten galaxies whose $A^\star_V/A_{V,\,{\rm neb}}>1.5$. We find that $\Delta\left<\log t_\star\right>_L=0.24\pm0.10$, $\Delta\left<\log t_\star\right>_M=1.15\pm1.03$, and $\Delta \log M_\star=0.41\pm0.27$. Among these ten galaxies, six have $\left<\log t_\star\right>_M < 1$ Gyr (as listed in Table 2), and we obtain $\Delta\left<\log t_\star\right>_L=0.25\pm0.12$, $\Delta\left<\log t_\star\right>_M=1.74\pm0.92$, and $\Delta \log M_\star=0.56\pm0.24$. Therefore, the uncertainties of our fitting results are dominated by the age-extinction degeneracy. \subsection{Aperture effects} The SDSS is a fiber-based survey, and thus we should consider the biases introduced by the use of small apertures to measure the galaxy spectra. Kewley et al. (2005) have examined this problem in detail and they conclude that 20\% of the galaxy light is required to minimize the aperture effects. \begin{figure}[pthb] \centering \includegraphics[width=0.8\textwidth]{fig9.ps} \caption{ The mean light- (upper panels) and mass-weighted (bottom panels) stellar age as a function of the $(u-g)$ and $(g-r)$ colors within the $3''$-diameter region. Open and solid circles show the spectra with $f_r < 20\%$ and $f_r > 20\%$, respectively.} \label{Fig9} \end{figure} As described earlier, the SDSS spectra for our sample galaxies only cover a small fraction of the galaxy light, and the median $f_r$ is 14.4\%. The fiber spectra are bluer than the integrated SEDs for most galaxies. Based on the fiber spectra, in general, $x_{\rm Y}$ is overestimated and $x_{\rm O}$ is underestimated (see Figure 6), resulting in an underestimation of both $\left<\log t_\star\right>_L$ and $\left<\log t_\star\right>_M$. Therefore we want to check to what extent our estimates of the stellar properties may be biased on the basis of the fiber spectra. In Figure 9 we plot the dependence of the mean stellar ages on the fiber $(u-g)$ and $(g-r)$ colors. We find that $\left<\log t_\star\right>_L$ correlates both with $(u-g)$ and $(g-r)$ (see Figures 9$a$ and 9$b$). Using these two correlations, we can make a coarse estimation of the bias in $\left<\log t_\star\right>_L$. For our sample, the $3''$-regions are generally 0.15 and 0.17 mag (median value) bluer than the whole galaxies in $(u-g)$ and $(g-r)$, respectively, which result in 0.30 and 0.42 dex underestimations of $\left<\log t_\star\right>_L$ respectively. However, no dependence of $\left<\log t_\star\right>_M$ on these two colors can be established. This is because that $\left<\log t_\star\right>_M$ has much less direct relation with the observed spectrum and is much more affected by the age-extinction degeneracy than $\left<\log t_\star\right>_L$ due to the non-constant mass-to-light ratio of stars. As shown in Section 3.1, the aperture effects generally cause an underestimation of $\left<\log t_\star\right>_M$ for our sample galaxies, and thus, combining with the results derived with the integrated spectra, they might not much affect our conclusion that BCDs are old galaxies experiencing starbursting activities.

1012.1145

We study stellar populations, star formation histories (SFHs), and star formation properties for a sample of blue compact dwarf galaxies (BCDs) selected by cross-correlating the Gil de Paz et al. sample with the Sloan Digital Sky Survey Data Release 6. The sample includes 31 BCDs, which span a large range of galactic parameters. Using a stellar population synthesis method, we derive stellar populations and reconstruct SFHs for these BCDs. Our studies confirm that BCDs are not young systems experiencing their first star formation, but old systems undergoing a starburst activity. The stellar mass-weighted ages can be up to 10 Gyr, while the luminosity-weighted ages might be up to approximately three orders of magnitude younger (~10 Myr) for most galaxies. Based on multiwavelength data, we also study the integrated star formation properties. The star formation rate (SFR) for our sample galaxies spans nearly three orders of magnitude, from a few 10<SUP>-3</SUP> to ~1 M <SUB>sun</SUB> yr<SUP>-1</SUP>, with a median value of ~0.1 M <SUB>sun</SUB> yr<SUP>-1</SUP>. We find that about 90% of BCDs in our sample have their birthrate parameter (the ratio of the current SFR to the averaged past SFR) b&gt;2-3. We further discuss correlations of the current SFR with the integrated galactic stellar mass and explore the connection between SFR and metallicity.

12213964

2011MNRAS.416.1717K

1012

1012.3479_arXiv.txt

Cosmic shear has been identified as being a particularly sensitive tool in understanding dark energy (Albrecht et al., 2006; Peacock et al., 2006), dark matter (Massey, Kitching, Richards, 2010), neutrino physics (Hannestaad, 2010) and potential depatures from general relativity. There are a number of on-going (CFHTLenS, Pan-STARRS) and planned experiments (KIDS, DES, LSST, Euclid) whose primary scientific goals are to use cosmic shear to constrain cosmological parameters. However, as we show in this article, the majority of dark energy information comes from small scales, that are expected to be influenced by poorly understood non-linear effects such as baryonic processes (see White, 2004; Zhan \& Knox, 2004, Mead et al., 2010; Rudd et al., 2008; Guillet et al., 2010 for example); which may be very difficult to model to sufficient accuracy. Indeed we may not even know the sign of contribution from baryonic effects. If dissipation is the main effect, we can imagine baryonic collapse will lead to an enhancement of matter clustering, or Baryonic Compression. But nonlinear feedback from star or AGN production could also blow out baryons, reducing the overall mass and allowing the dark matter to disperse. The formation of stars or AGN in simulations is usually governed by sub-grid semi-analytic prescriptions, which makes the understanding of such processes uncertain. As a step forward it would be very useful to even quantify our uncertainty on the physics involved on different scales. In addition to the uncertain physical effects there is the practical consideration of the accuracy of fitting formulae. The matter power spectrum as a function of redshift can be computed using linear perturbation theory of the underlying initial density field, but by comparison with N-body simulations it is apparent that at scales where the matter overdensity $\delta$ becomes greater than $0.1$ the linear theory predictions cannot be used as non-linear effects in the growth of structure become dominant. The most widely used corrections are Smith et al. (2003) {\tt halofit}, Peacock \& Dodds (1996) and the {\tt Coyote} formula Heitmann, et al, (2010); {\tt halofit} and {\tt Coyote} are accurate to approximately $5-10\%$. In addition to this uncertainty the formula are proposed by only sampling a small discrete number of points in parameter space, and currently do not include baryons. Finally, to extract cosmological information from the statistical properties of the density field means we require a covariance matrix. Nonlinear growth of dark matter structure will already correlate estimates of the matter power spectrum on different scales (e.g. Kiessling, et al. 2010), and including baryonic physics we can expect even stronger covariances between scales due to feedback processes. The cosmic shear power spectrum depends on the matter power spectrum through an integral over the line of sight distance, with a geometric lensing kernel. The lensing power spectrum contains cosmological information, through the lensing kernel, even on small scales. Rather than throw this information away, we propose that one can include even very uncertain non-linear scales in a lensing analsysis, if one correctly marginalizes over the uncertainty. Using the path-integral marginalization techniques, presented in Taylor \& Kitching (2010) and Kitching \& Taylor (2010), we derive an expression, in the self-calibration regime, for the cosmic shear covariance that includes the geometric lensing kernel information from all scales. If an informative upper bound on the functional behaviour of the non-linear power spectrum can be determined (from simulations for example) the residual uncertainty in the functional form of the power spectrum can be marginalized over simultaneously with the cosmological parameters of interest. We place a requirement on the external fuctional prior such that the cosmological constraints from future all-sky cosmic shear experiments are not degraded below a level needed to determine dark energy physics. This article is organised as follows, in Section \ref{Method} we present the marginalized likelihoods and Fisher matrices for the non-linear tomographic cosmic shear power spectrum. In Section \ref{Results} we present results, and in Section \ref{Conclusion} we present our conclusions.

\label{Conclusion} To conclude we find that the non-linear and baryon-dominated part of the matter power spectrum contains, above wavenumbers of $k\gs 0.5h$Mpc$^{-1}$, half of the information content on dark energy parameters, parameterised through the Figure of Merit. However the lack of knowledge about this regime, and the complex simulated modelling needed to correctly constrain its behaviour as a function of environment, scale and cosmology means that the uncertainty on the non-linear power spectrum must be correctly accounted for in cosmic shear surveys. This article has some resonance with previous work, that parameterise the uncertainty in the non-linear power spectrum and marginalize over those parameters (Zhang et al., 2009; Rudd et al., 2008; Zenter et al., 2008; Huterer et al., 2006; Jing et al., 2000) or attempt to modify the data to minimise the effect (Huterer \& White, 2005), however all these assume parameterized models, that may not be able to reflect the real effect of baryons. We have derived likelihood expressions for the tomographic cosmic shear power spectrum in the cases that the functional matter power spectrum is self-calibrated from the data itself, and in the case that an external prior on the functional variation of the matter power spectrum is available. We summarise our results in Figure \ref{vbias3}. \begin{itemize} \item With no external priors, a Euclid-like cosmic shear survey, with a $k_{\rm max}=50h$Mpc$^{-1}$ could achieve a FoM $\approx 220$ from cosmic shear tomography alone.\\ \item If the non-linear matterpower spectrum is completely removed using a hard cut in $k$-modes of $k_{\rm max}=0.5h$Mpc$^{-1}$ then the FoM is reduced by a factor of $50\%$. \\ \item In the functional self-calibration regime the cosmic shear survey can recover the FoM, with only a $10\%$ reduction in the FoM. \\ \item By including an informative prior, from simulations for example, the orginal FoM can be recovered if the functional variation of the non-linear matter power spectrum is known to $\sim 1\%$ to $k=50h$Mpc$^{-1}$, or a physical scale of $\sim 120$Kpc$/h$. \end{itemize} Finally we note that in the self-calibration regime, the information used to constrain the cosmological information, through the lensing kernel, has some similarities with the shear-ratio method, where cluster scale weak lensing is isolated. Constraining the non-linear power spectrum to 1\% functional accuracy down to $120$kpc$/h$ as a function of scale, redshift and cosmology is a significant theoretical and observational challenge. On the theoretical side modeling the baryons on the scale of galaxy clusters is already a challenge. Extending this to group and individual galaxy haloes will require a much deeper understanding of the baryonic processes on these scales. On the observational side, weak lensing itself, and galaxy-galaxy lensing, can provide much empirical information about the mass distribution which can be compared with stellar and gaseous components. These are challenges that must be realised if we are to fully exploit the potential of tomographic cosmic shear experiments. \begin{figure} \includegraphics[angle=0,clip=,width=\columnwidth]{fig3.ps} \caption{Summary of the main results. If the non-linear scales are removed using a hard cut in $k$-modes then the FoM can be reduced by $50\%$. In the functional self-calibration regime the reduction is less severe with a relative reduction of $10\%$, finally if a $1\%$ external prior on the functional variation of the non-linear power can be applied then the FoM is recovered. All $k$-modes are in units of $[h$Mpc$^{-1}]$. The parameters $a_0$ and $a_1$ refer the shape of the functional variance as a function of scale, given by equation (\ref{44}).} \label{vbias3} \end{figure} \\ \noindent{\em Acknowledgements:} TDK is supported by STFC rolling grant number RA0888. We thank the eScience institute Edinburgh for hosting a workshop on n-body simulations in cosmology. We thank Alan Heavens, Catherine Heymans, Fergus Simpson for interesting discussions on this topic.

1012.3479

We present a new method that deals with the uncertainty in matter clustering in cosmic shear power spectrum analysis that arises mainly due to poorly understood non-linear baryonic processes on small scales. We show that the majority of information about new physics contained in the shear power comes from these small scales. Removing these non-linear scales from a cosmic shear analysis results in 50 per cent cut in the accuracy of measurements of dark energy parameters, marginalizing over all other parameters. In this paper we propose a method to recover the information on small scales by allowing cosmic shear surveys to measure the non-linear matter power spectrum themselves and marginalize over all possible power spectra using path integrals. Information is still recoverable in these non-linear regimes from the geometric part of weak lensing. In this self-calibration regime we recover 90 per cent of the information on dark energy. Including an informative prior, we find that the non-linear matter power spectrum needs to be accurately known to 1 per cent down to k= 50 h<SUP>-1</SUP> Mpc, or a scale of 120 kpc, to recover 99 per cent of the dark energy information. This presents a significant theoretical challenge to understand baryonic effects on the scale of galaxy haloes.

3773534

2011A&A...527A..88V

1012

1012.0961_arXiv.txt

Massive stars play an important role in determining physical, chemical, and morphological properties of galaxies. The last decade has seen considerable progress in the understanding of massive star formation \citep{2007ARA&A..45..481Z}. Since their detection by \citet{1996A&A...315L.165P} and \citet{1998ApJ...494L.199E}, cold and dense infrared dark clouds (IRDCs) appear to be the ideal sites to investigate the initial conditions for the process of massive star formation. IRDCs have typical sizes between 1 and 10 pc, masses from several hundreds to several thousands solar masses and H$_2$ column densities between 2 and 10$\times 10^{23}$~cm$^{-2}$ \citep[e.g.][]{2006ApJ...641..389R,2009A&A...499..149V,2009ApJ...698..324R}. Apart from continuum observations, molecular line data have been used to characterize the properties of IRDCs. The first molecular line data were obtained by \citet{1998ApJ...508..721C}, who detected H$_2$CO in 10 clouds, thus confirming the presence of dense gas. Using LVG modeling % they estimated H$_2$CO abundances of $\sim 10^{-10}$. That is a factor of 50 lower in comparison with low-density clouds and can be explained by accretion of gas-phase metals onto dust grains in the cold and dense IRDCs. Ammonia observations by \citet{2005ApJ...634L..57S} and \citet{2006A&A...450..569P} allowed temperature determination of IRDCs in a range from 10 to 20 K. While earlier studies \citep{2006ApJ...639..227S} suggested the association of most IRDCs with the so-called Galactic molecular ring (galactocentric distance of 5 kpc), the \citet{2008ApJ...680..349J} study gave new evidence that the IRDC distribution in the first and fourth galactic quadrant more closely follows a galactic spiral arm (the Scutum-Centaurus arm). Since in normal spiral galaxies, OB stars seem to form primarily in spiral arms, the association of IRDCs with a Milky Way spiral arm supports the idea that IRDCs are related to high-mass star formation. \citet{2008ApJ...678.1049S} observed N$_2$H$^+$(1-0), HC$_3$N(5-4), CCS(4$_3$-3$_2$), NH$_3$(1,1), (2,2), (3,3) and CH$_3$OH(7-6) lines toward the massive clumps associated with IRDCs, to study the chemical conditions in IRDCs. Analysing the CCS and N$_2$H$^+$ abundance ratio, they conclude that infrared dark clouds are chemicaly more evolved than low-mass pre-stellar cores. An estimation of the chemical evolutionary status of IRDCs was performed also by \citet{2009ApJ...705..123G}. Using C$^{18}$O, CS and N$_2$H$^+$ abundances, and a chemical evolution code, they showed that cores where all three lines are detected appear to be chemically young (10$^{4.5} < $ t $<$ 10$^{5.5}$~years). Cores where no N$_2$H$^+$ emission is detected are suspected to be especially young (t $<$ 10$^2$~years). This suggests that these regions may not have yet formed massive protostars. A molecular survey by \citet{2010ApJ...714.1658S} toward 20 massive clumps, including mid-IR bright and mid-IR dark sources, showed outflow activities and the possible presence of protostars in some clouds. Their line width analysis allowed them to reconstruct the distribution of the molecular species at the different evolutionary stages of massive clumps. While many characteristics of infrared dark clouds were determined during the last ten years, some of their properties are still not well known. Among the open questions is their chemical state. What is the chemistry in IRDCs? Is it really different from the chemistry in low-mass dark clouds? To enlarge the sample of well characterized IRDCs, we started a program to study gas and dust properties of southern infrared dark clouds \citep[][Paper I hereafter]{2009A&A...499..149V}. A set of 15 clouds has been selected in the pre--Spitzer era by visual examination of the extinction contrast of the MSX 8.3~$\mu$m images. In the meantime, the Spitzer satellite has since succeeded MSX and provided a far higher spatial resolution and sensitivity. GLIMPSE mid-infrared images of our regions were retrieved from the Spitzer Archive. Continuum 1.2~mm maps were obtained with the SIMBA bolometer array at the SEST telescope. A 1.2~mm and 8~$\mu$m study of these southern IRDCs showed that these objects are not just distant Taurus-like clouds, but a distinct type of clouds with a clear trend to higher H$_2$ column densities. It was found that the true peak column densities, extracted from millimeter data, exceed 3 $\times 10^{23}$~cm$^{-2}$ (or 1~g~cm$^{-2}$), which has been proposed as a threshold for high-mass star-forming clouds \citep{2008Natur.451.1082K}. This paper is our next step toward the understanding of the nature of the IRDCs. Here we present our investigations of the chemical composion of southern clouds. We perform molecular line observations in the 3 mm band with the Mopra single-dish radio telescope. Combining molecular line data and H$_2$ column densities from the previous study we estimate molecular abundances and compare them with results for low-mass pre-stellar cores. Analysis of molecular lines provides not only information about chemistry, but also about physical processes in molecular clouds. For instance, the presence of SiO emission and HCO$^+$ extended wings are evidences for outflow activity in a cloud. Specific line shapes can indicate infall motion. Detection or non-detection of some species help to determine the evolutionary status of our targets. The paper is organized in the following way. In Sect. 2 we describe our target selection, the selected molecular lines, observational and technical details. In Sect. 3 we present the results of the qualitative analysis of obtained molecular line spectra, line parameters and abundance estimates. We discuss the obtained results in Sect. 4 and conclude in Sect. 5.

In this paper, we present 3~mm molecular line observations with the 22-m Mopra radio telescope. In total 13 molecular lines were observed for all IRDCs. The results of our study can be summarized as follows: \begin{list}{}{} \item[1.] Using H$_2$ column densities from the previous investigation, we estimate molecular abundances of all species. We show that there is a tendency for the IRDCs to have molecular abundances similar to the low-mass pre-stellar core rather than to the HMPOs abundances. However, the derived abundances come with uncertainties of around one order of magnitude. Furthermore, also the comparison abundances for low--mass cores and HMPOs can be affected by considerable uncertainties, especially, when abundances have been computed in the literature by combining heterogeneous data sets (regarding beam sizes etc.). To make more solid statements about the evolutionary status of IRDCs therefore calls for subsequent systematical studies. \item[2.] According to the classification of \citet{2009ApJS..181..360C}, we subdivided our clouds to "quiescent" and "active" and added "middle" class to them. We have found a trend for more evolved regions to have higher line widths and integrated intensities. However, we do not detect clear evidence of different chemistry in these three groups. \item[3.] Comparison of the line width and integrated intensities of the IRDCs and low-mass dark clouds show several times higher values for IRDCs. Broader and more intense lines mean that in IRDCs we have more turbulent conditions in comparison with low-mass clouds. \item[4.] We detect the SiO emission in some clouds and complicated shapes of the HCO$^+$ emission line profile in all IRDCs, which indicates the presence of infall and outflow motions and the beginning of star formation activity, at least in some parts of the IRDCs. \item[5.] The analysis of the two available CH$_3$C$_2$H excitation diagrams and detection of the very weak CH$_3$CN (5$_0$-4$_0$) and CH$_3$CN (5$_1$-4$_1$) lines on the "average spectra" indicate the presence of a warm gas component in some IRDCs. However, these warm regions are compact and cannot be resolved with single-dish observations. \end{list}

1012.0961

Context. Massive stars play an important role in shaping the structure of galaxies. Infrared dark clouds (IRDCs), with their low temperatures and high densities, have been identified as the potential birthplaces of massive stars. In order to understand the formation processes of massive stars, the physical and chemical conditions in infrared dark clouds have to be characterized. <BR /> Aims: The goal of this paper is to investigate the chemical composition of a sample of southern infrared dark clouds. One important aspect of the observations is to check, whether the molecular abundances in IRDCs are similar to the low-mass pre-stellar cores, or if they show signatures of more evolved evolutionary stages. <BR /> Methods: We performed observations toward 15 IRDCs in the frequency range between 86 and 93 GHz using the 22-m Mopra radio telescope. In total, 13 molecular species comprising N<SUB>2</SUB>H<SUP>+</SUP>, <SUP>13</SUP>CS, CH<SUB>3</SUB>CN, HC<SUB>3</SUB>N, HNC, HCO<SUP>+</SUP>, HCN, HNCO, C<SUB>2</SUB>H, SiO, H<SUP>13</SUP>CO<SUP>+</SUP>, H<SUP>13</SUP>CN, and CH<SUB>3</SUB>C<SUB>2</SUB>H were observed for all targets. Hence, we included in general species appropriate for elevated densities, where some of them trace the more quiescent gas, while others are sensitive to more dynamical processes. <BR /> Results: We detect HNC, HCO<SUP>+</SUP>, and HNC emission in all clouds and N<SUB>2</SUB>H<SUP>+</SUP> in all IRDCs except one. In some clouds we detect SiO emission. Complicated shapes of the HCO<SUP>+</SUP> emission line profile are found in all IRDCs. Both signatures indicate infall and outflow motions and the beginning of star-formation activity, at least in some parts of the IRDCs. Where possible, we calculate molecular abundances and make a comparison with previously obtained values for low-mass pre-stellar cores and high-mass protostellar objects (HMPOs). We show a tendency for IRDCs to have molecular abundances similar to low-mass pre-stellar cores rather than to HMPOs abundances on the scale of our single-dish observations. <P />Appendices A-C are only available in electronic form at <A href="http://www.aanda.org">http://www.aanda.org</A>

1990014

2011A&A...526L...6R

1012

1012.2649_arXiv.txt

1012.2649

We present the first observations of a probable brown dwarf, obtained with the new spectrograph X-shooter mounted on the UT2@VLT. The target (2MASS J053825.4-024241) is a 0.06 M_⊙ object in the star-formation region σ Ori. The X-shooter spectrum covers simultaneously the whole range from UV to NIR (300-2500 nm). The J053825.4-024241 spectrum is rich in emission lines that are typical of accreting young object and clearly shows the Balmer jump. Moreover, many photospheric atomic and molecular absorption lines yield the spectral type and confirm that the object is young. We compute the mass accretion rate from all available observed accretion diagnostics. We find that there is a large spread in the dot M_acc values (up to a factor 40) that is not caused by variability; some of this spread may be intrinsic, i.e., owing to different physical conditions of the emitting region for the same dot M_acc. However, within the large error bars all dot M_acc measurements agree, and the mean value is log dot M_acc -9.86 ± 0.45 M_⊙/y. The hydrogen Balmer lines are clearly detected up to n = 25. Their ratios suggest that the emitting region is cold (T 2000-3000 K), dense and in thermal equilibrium (LTE), and that the lines are optically thick up to n 21. We briefly discuss the implications of this result for magnetospheric accretion models. <P />Based on observations collected at the European Southern Observatory, Chile. Program 084.C-0269(A).

12206014

2011IJMPD..20.1079I

1012

1012.2822_arXiv.txt

Since their publication, the Ho\v{r}ava's seminal papers\cite{horava1,horava2} on a four-dimensional theory of gravity which is power-counting renormalizable and, hence, can be considered a candidate for the short-range (UV) completion of the General Theory of Relativity (GTR), have raised a lot of interest and stimulated an intense research effort on many topics of the theory and its modifications (a brief review can be found in the recent paper by Sotiriou\cite{sotiriou}). Ho\v{r}ava's theory admits Lifshitz's scale invariance: $\vec{r} \rightarrow b \vec{r}, \quad t \rightarrow b^{q} t$, and, after this, it is referred to as Ho\v{r}ava-Lifshitz (HL) gravity. Actually, it has anistropic scaling in the UV domain, since it is $q=3$, while relativistic scaling with $q=1$ is recovered at large distances (IR). Among the other issues, spherically symmetric solutions in HL gravity have been investigated in details\cite{spherical1,spherical2,spherical3,spherical4}, also in five dimensions\cite{spherical5d} (as for slowly rotating solutions, see Reg. \refcite{rot1,rot2}); in particular, Kehagias and Sfetsos\cite{KS} obtained an asymptotically flat and static spherically symmetric solution that can be considered the analog of Schwarzschild solution in GTR. An open issue in HL theory pertains the role of matter and its coupling to gravity, which has not been clarified in full details. As a consequence, the motion of free particles in HL gravity is not trivial: since the fact that particles move along geodesics is not granted, it is possible to expect deviations from geodesic motion\cite{geo1,geo2,geo3,capasso}. However, many authors (see e.g. Ref. \refcite{light} and references therein) assumed that free particles followed geodesics of KS metric to focus on many issues HL gravity, such as gravitational lensing, quasi-normal modes, accretion disks and so on (moreover, an accurate analysis of KS geodesics can be found in Ref. \refcite{geoKS}). In other words, they used KS solution as a toy model to study some fundamental aspects of the theory. In this paper, starting from the same assumption, we aim at constraining the dimensionsless parameter $\psi_0$ of KS metric. In previous papers\cite{ijmpa,TOAJ} we obtained bounds both from solar system and extra solar system observations of orbital motions; other constraints were derived by studying light deflection\cite{light} and analyzing the impact on the classical GTR tests\cite{tiberiu}. Here we show that tighter results can be obtained by considering the perturbation $\Delta\rho$ induced by the KS solution on the two-body range $\rho$ for a pair of test particles A and B orbiting the same central body of mass $M$. The plan of the work is as follows: in Section \ref{anale}, \textcolor{black}{after a brief introduction to the KS solution in the context of HL gravity,} we will analytically work out the perturbation $\Delta\rho$ induced by the KS solution on the two-body range $\rho$, which is a very accurate, direct and unambiguous observable in solar system planetary studies. We will try to make such a part as more self-consistent as possible, in view of a readership which may not be fully acquainted with the subtleties of celestial mechanics applied to fundamental physics, an endeavor requiring interdisciplinary knowledge across different fields. In this way the reader has the possibility of following autonomously the future developments of such kind of investigations, and to apply the present approach to other exotic long-range modified models of gravity as well. In Section \ref{osservazioni} we compare our theoretical predictions to the range residuals $\delta\rho$ constructed with the latest planetary ephemerides applied to recent, accurate data sets of some planets of the solar system. Section \ref{conclusioni} is devoted to summarizing our findings.

\lb{conclusioni} In view of the fact that ranging from the Earth to some major bodies of the solar system is one of the most direct and accurate way to determine their orbits, in this paper we analytically worked out the effects that the Kehagias-Sfetsos solution of the Ho\v{r}ava-Lifshitz modified gravity at long distances have on the two-body range $\rho$. We successfully checked our analytical results by simultaneously integrating the equations of motion of the bodies A and B considered by including such exotic dynamical effects as well. By comparing our predictions with the range-residuals $\delta\rho$ of Mercury, Mars and Saturn produced with the recent ephemerides INPOP10a we have been able to effectively constrain the dimensionless free parameter $\psi_0$ entering the Kehagias-Sfetsos equations in the case of the Sun. We obtained $\psi^{\odot}_0\geq 7.2\times 10^{-10}\ ({\rm Mercury}),\ \psi^{\odot}_0\geq 9\times 10^{-12}\ ({\rm Mars}),\ \psi^{\odot}_0\geq 1.7\times 10^{-12}\ ({\rm Saturn})$. Such constraints are orders of magnitude better than those existing in literature. In principle, one should re-process the entire planetary data set by explicitly modeling the Kehagias-Sfetsos dynamical effects in addition to the standard Newtonian-Einsteinian ones and solving for a dedicated parameter in order to avoid the possibility that they, if unmodeled, may be removed from the signal in the estimation of the initial conditions. However, such a task is very time-consuming and can effectively be implemented only by skilful and expert teams of specialists in astronomical data processing. On the other hand, the relatively simpler approach outlined here can be easily and quickly reproduced, and extended to other exotic long-range modified models of gravity as well in order to yield reasonable evaluations of the magnitude of the effects of interest.

1012.2822

We analytically work out the perturbation Δρ induced by the Kehagias-Sfetsos (KS) spacetime solution of the Hořava-Lifshitz (HL) modified gravity at long distances on the two-body range ρ for a pair of test particles A and B orbiting the same mass M. We apply our results to the most recently obtained range residuals δρ for some planets of the solar system (Mercury, Mars, Saturn) ranged from the Earth to effectively constrain the dimensionless KS parameter ψ<SUB>0</SUB> for the Sun. We obtain ψ <SUP>⊙</SUP><SUB>0&gt;=</SUB> 7.2× 10<SUP>-10</SUP> (Mercury), ψ <SUP>⊙</SUP><SUB>0&gt;=</SUB> 9× 10<SUP>-12</SUP> (Mars), and ψ <SUP>⊙</SUP><SUB>0&gt;=</SUB> 1.7× 10<SUP>-12</SUP> (Saturn). Such lower bounds are tighter than others existing in the literature by several orders of magnitude. We also preliminarily obtain ψ <SUB>0<SUP></SUP>&gt;=</SUB> 8× 10<SUP>-10</SUP> for the system constituted by the S2 star orbiting the supermassive black hole (SBH) in the center of the galaxy.

12167447

2011ApJ...728..159W

1012

1012.5319_arXiv.txt

\par A fundamental aspect of understanding the structure of the interstellar medium (ISM) involves the origin of the small-scale structure that is observed. A number of formulations provide a framework for analyzing such structure, and explaining its creation and distribution. These formulations include those of Houlahan \& Scalo (1992), who developed a hierarchical tree structure, describing clouds as a series of partitions. Falgarone et al.\ (1991) explored fractal structure and showed that this description is applicable to molecular clouds. Vogelaar \& Wakker (1994) used fractal structure in an attempt to describe the structure of high-velocity clouds. Fractal structure may arise naturally from the density statistics of interstellar turbulence. More recently, Lazarian \& Pogosyan (2000, 2004, 2006, 2008) developed techniques for comparing observations of velocity and density structure in the ISM with statistics derived from theories of turbulence. Kowal et al.\ (2007) used these ideas to construct 3-D MHD models of the ISM that predict the spectrum of density fluctuations, while Burkhart et al.\ (2009) studied how statistical measures can be used to connect these models to observations. In this paper, we look at the column density distribution of neutral hydrogen, trying to compare measurements made at different angular resolutions, which allows us to apply one of the measures discussed by Burkhart et al.\ (2009). \par Galactic \HI\ column densities can be measured using 21-cm radio observations or by using \Lya\ absorption-line spectra. 21-cm observations are made with single-dish or interferometer radiotelescopes. Single-dish telescopes have a large beam size, typically 36\arcmin\ for all-sky surveys, and 9\arcmin--21\arcmin\ for more targeted observations. Interferometers produce beams of less than 2\arcmin, although for Galactic gas it is necessary to combine this with single-dish observations to derive an accurate total column density. With \Lya\ observations of ultra-violet bright background targets, however, one measures \NHI\ in a very small (sub arcsecond) area centered on the background target. If the background target is a galactic disk star, the \Lya\ absorption is caused only by the \HI\ in front of the star, whereas the 21-cm observations measure \HI\ both in front of and behind the star. When observing AGNs, both \Lya\ and 21-cm observations mesaure all \HI\ in the line of sight. Some halo stars may also be sufficiently high above the galactic plane to lie above most of the \HI. Thus, valid comparisons between \NHILya\ and \NHITWcm\ can only be made in the directions of AGNs or distant halo stars. \par Precise measurements of \NHI\ are also necessary to derive abundances of heavy elements in the interstellar medium (ISM). Usually, 21-cm data are used to derive the \HI\ reference column density, which is necessary whenever absorption lines show different components. However, the metal-line absorption is produced only by the ions in the small beam toward the AGN, while the 21-cm data average \NHI\ across the radiotelescope beam. Thus, there is a question as to whether it is correct to use the 21-cm data to derive \NHI. Do \Lya\ measurements and 21-cm measurements in fact give consistent results? If they do not, what is the reason for this difference? \par Hobbs et al.\ (1982) were the first to directly compare \NHI\ measured using \Lya\ absorption lines (from Copernicus data) to \NHITWcm, measured using the 140-ft Green Bank telescope (21\arcmin\ beam). Although they did not present errors on their meaurements, the ratios they found varied from about 0.9 to 3.9, with the outliers for the stars closest to the plane. For the three stars above most of the Galactic \HI\ layer ($z$$>$2~kpc) the ratios were given as 0.9, 0.9 and 1.1. \par This was followed by a study by Lockman et al.\ (1986a), who used {\it International Ultraviolet Explorer} (IUE) spectra to measure \NHILya\ toward 45 stars, by fitting the profiles of the damping wings. The resulting errors in $N$(\HI) were estimated to be about 0.1 dex (25\%). These values were again compared to 140-ft Green Bank Telescope 21-cm data, which were corrected for stray radiation. They also used a spin temperature of 75\,K to correct the column densities for optical depth effects, though in most cases this correcton is small (a few percent). For the stars closest to the plane ($z$$<$1~kpc), the ratio \NHILya/\NHITWcm\ tends to be $<$1, because not all of the \HI\ in the sightline is seen in absorption. For the six stars at $z$$>$1.5 kpc the average ratio was about 1, to within the errors, although the typical error in each ratio was about 0.3. \par Savage et al.\ (2000) revisited this comparison in the directions of 14 QSOs, using data from the G130H grating in the Faint Object Spectrograph (\FOS) on the Hubble Space Telescope (\HST). Unlike the Copernicus and IUE data, the \FOS\ spectra had low resolution (230~\kms). Savage et al.\ (2000) compared these masurements to 21-cm data from the Green Bank 140-ft telescope. For ten of the QSO spectra it was possible to correct for geocoronal emission, and for these Savage et al.\ (2000) found that the ratio \NHILya/\NHITWcm\ had values in the range 0.62 to 0.91 with errors of about 0.1. \par Savage et al.\ (2000) suggested two possible origins for differences between the column densities measured from \Lya\ absorption and 21-cm emission. First, differences might arise from a combination of systematic and random errors in the \Lya\ and 21-cm observations. In the \Lya\ observations, systematic errors can be produced by uncertain continuum placement, geocoronal \HI\ emission removal, the spectrograph background and scattered light correction, interfering QSO and IGM absorption, and detector fixed pattern noise. For the \FOS\ dataset of Savage et al.\ (2000) all of these effects were present. Using data at higher spectral resolution removes all of the trouble associated with geocoronal \HI, IGM absorption and fixed pattern noise, while it greatly reduces the other problems. In the 21-cm data, systematic errors can be created by the absolute calibration of the radio telescope, baseline fitting and the stray-radiation correction. Alternatively, the differences in \Lya\ and 21-cm column densities could be caused by the structure of the ISM. If there are small bright spots embedded in a smoother background, the 21-cm data will include these, but a random sightline to an AGN has a high probability of missing the brighter spots. To study such effects requires a large sample of sightlines. \par In this paper we analyze 59 sightlines using higher-resolution and higher S/N \Lya\ spectra than used by Hobbs et al.\ (1982), Lockman et al.\ (1986b) and Savage et al.\ (2000). We compared all these measurements to the column densities found in the LAB survey (Kalberla et al.\ 2005), as well as to the Lockman \& Savage (1995) Green Bank 140-ft data that is available in many directions. In addition, we obtained new 21-cm data with the Green Bank Telescope (GBT) toward 35 AGNs. In Sect.~\Sobs\ we describe the \Lya\ and 21-cm datasets that we used, as well as the method used to derive \NHILya. During our analysis, we discovered the presence of a spurious, broad underlying component in the LAB and 140-ft spectra. We show this in Sect.~\Scorrect. After removing this component, we find the results that are presented in Sect.~\Sresults\ and discussed in Sect.~\Sdisc. \begin{figure}\plotfiddle{f1.eps}{0in}{270}{250}{480}{0}{-350}\figurenum{\Fmap} \caption{All-sky map of \NHI\ integrated between $-$100 and 100~\kms, using the Leiden-Argentina Bonn survey (Kalberla et al.\ 2005), with the column density scale on the bottom. The directions to our 59 targets are shown by the stars and labels.} \end{figure} \input table1.tex

\par We derive the column density of neutral \HI\ from the \Lya\ line using data from the \STIS-spectrograph on \HST, and from the 21-cm line using the Leiden-Argentina-Bonn (LAB) survey, Green Bank 140-ft and Green Bank Telescope (GBT) observations. Using these column densities, we compare the ratio of \NHILya\ to \NHITWcm\ and \NHITWcm\ to \NHITWcm\ in order to analyze the structure of the ISM. Based on the results, we conclude the following: \par (1) For 59 AGNs surveyed, 36 yield reliable \Lya\ column densities. There are 163 Green Bank 140-ft and 35 GBT sightlines for which \NHITWcm\ is derived. For each of the unique sightlines, we also measured $N$(\HI) using LAB data. \par (2) We conclude that the published LAB data, as well as our old (from the late 1980s) Green Bank 140-ft data, suffer from a problem that results in an excess column density. of $\sim$1.5\tdex{19}\,\cmm2. This problem is revealed by extracting the spectral regions outside the range where signal appears to be present. There is no residual emission in the combined GBT spectrum, in contrast to the LAB and 140-ft data. The residual can be fitted by a gaussian, which for the LAB dataset has $v$=$-$22~\kms, $T$=0.048~K and FWHM=167~\kms. The parameters of the residual are different for the 1~\kms\ and 2~\kms\ channel spacing 140-ft spectra. \par (3) We conclude that the \HI\ spectra from the LAB survey need to be corrected for the presence of a broad underlying gaussian. Without such a correction, the LAB data conflict with measurements of \NHI\ made with the GBT and made using \Lya\ absorption, as well as with UV absorption-line studies, and with the properties of the recently-found population of small \HI\ clouds. With such a correction, all tension between measures of \NHI\ at different resolutions disappears. \par (4) Using data from the \FOS\ on \HST, Savage et al.\ (2000) had found that on average \NHILya/\NHITWcm\ was 0.81$\pm$0.09 for 12 sightlines. Using our new data, the same set of sightlines yields \Lya\ column densities that are on average $\sim$0.06~dex higher (compared to the typical error of 0.2--0.3~dex in the \FOS\ results). We also find that a correction is needed to the 140-ft data, which corresponds to 0.05~dex in these 12 sightlines. As a result of these corrections, the average ratio \NHILya/\NHITWcm\ for these sightlines is now found as 1.02$\pm$0.08. \par (5) After applying the corrections to the LAB and 140-ft observations, we find that the ratios between \NHILya\ and \NHITWcm\ are on average 0.96$\pm$0.11 (\Lya/LAB), 0.95$\pm$0.09 (\Lya/140-ft) and 1.00$\pm$0.07 (\Lya/GBT). A statistical test shows that these averages do not differ from 1 in a statistically significant way. \par (6) A hierarchical model for the ISM matches the observed column density distribution for \NHILya/\NHITWcm\ ratios adequately well, although it underpredicts the number of high ratios. \par (7) A log-normal model matches column density ratio distribution moderately well, with a different width parameter for different cases. \par (8) Using a 3-D MHD model from Kowal et al.\ (2007), we can match most features of the column density distributions when choosing cases with $M_S$=0.65--0.90. These distributions are similar to a simple log-normal distribution, as is expected for turbulence. \par (9) We conclude that by comparing \HI\ column densities observed at very different resolutions it becomes possible to characterize the small-scale structure of the ISM. Although the number of sightlines in our sample is small, the distribution of column density ratios approximately follows a log-normal distribution, which is also similar to the predictions of 3-D MHD modeling, using a sonic Mach number in the range 0.65-0.90. \bigskip \bigskip Acknowledgements \par BPW acknowledges support from NSF grant AST-0607154 and NASA-ADP grant NNX07AH42G. The Green Bank Telescope is part of the National Radio Astronomy Observatory, a facility of the NSF operated under cooperative agreement by Associated Universities, Inc. The \Lya\ data in this paper were obtained with the NASA ESA {\it Hubble Space Telescope}, at the Space Telescope Science Institute, which is operated by the Association of Universities for Research in Astronomy, Inc.\ under NASA contract NAS5-26555. Spectra were retieved from the Multimission Archive (MAST) at STScI. We thank UW graduate students Alex Hill and Blakesley Burkhart for extracting the column densities from the simulation of Kowal et al.\ (2007) and for providing additional models. JMB acknowledges Steve Schmitt and Erin Conrad for mathematical discussions. We thank the anonymous referee for insisting that we investigate possible errors in the 21-cm surveys. \bigskip \bigskip Appendix \par {\it 3C\,249.1} -- The quality factor of the column density determination is only 2, because the spectrum is relatively noisy and the 21-cm spectrum shows the presence of several IVCs. There is an unidentified line near 1210~\AA, which cannot be intergalactic \Lya, but which also does not fit into any known system of absorbers toward this sightline. \par {\it 3C\,273.0} -- Intrinsic \OVI\ emission at 1195.32~\AA\ and 1201.91~\AA\ creates a moderately broad peak in the continuum at wavelengths shorter than Galactic \Lya. Combined with the intermediate-velocity gas at $v$$\sim$25~\kms\, this leads us to assign quality 2 to this sightline. A third order polynomial fits the right side of the \OVI\ wing well and creates a good fit overall from 1197~\AA\ to 1247~\AA. \par {\it 3C\,351.0} -- The spectrum of this AGN is flat near Galactic \Lya, and the 21-cm spectrum is simple, resulting in quality factor 4 for this spectrum. \par {\it H\,1821+643} -- The continuum of this target is flat and does not contain features. However, there are weak \HI\ components at high negative velocities ($<$$-$100~\kms), lowering the confidence in the resulting value of \NHI\ and leading us to give quality 3 to this target. \par {\it HE\,0226$-$4110} -- The spectrum of this AGN is flat near Galactic \Lya, and the 21-cm spectrum is simple, resulting in quality factor 4 for this spectrum. Where we find a value of 20.09$\pm$0.04 for the \HI\ columm density, Savage et al.\ (2007) reported 20.12$\pm$0.03 based on an earlier analysis of the same data. \par {\it HE\,0340$-$2703} -- Uncertain continuum placement due to the presence of multiple strong IGM lines results in a low quality fit for this QSO continuum. In addition, on the short-wavelength side there appears to be an (unidentified) intrinsic emission line, making the continuum even more uncertain. We therefore assign Q=1 to this target. \par {\it HE\,1029$-$1401} -- The spectrum of this AGN is flat near Galactic \Lya, and the 21-cm spectrum is simple, resulting in quality factor 4 for this spectrum. \par {\it HE\,1228+0131} -- The noisy continuum of this spectrum makes it is impossible to obtain a good fit. This is further complicated by the lack of knowledge about the intrinsic continuum of the QSO, which is higher at $\lambda$$<$1215~\AA\ than at $\lambda$$>$1215~\AA. Intrinsic \SIV\ $\lambda$1062.664 (redshifted to 1186~\AA) and intrinsic \NII\ $\lambda$1083.994 (redshifted to 1211~\AA) emission may be present, and there is also a steep continuum drop across Galactic \Lya. These problems make the continuum too uncertain to fit, and we do not use this sightline in our analysis. \par {\it HS\,0624+6907} -- The spectrum of this AGN is flat near Galactic \Lya, but the 21-cm spectrum is not simple, resulting in quality factor 2 for this spectrum. \par {\it HS\,1543+5921} -- A complicated continuum combined with low S/N make a good fit for the continuum almost impossible. In particular, \Lya\ absorption in SBS\,1543+593 at 2800~\kms\ blends with Galactic \Lya. We do not use this spectrum in our analysis. \par {\it HS\,1700+6416} -- The continuum is too difficult to obtain a good fit due to multiple IGM lines and moderate to high noise. We do not use this spectrum in our analysis. \par {\it MCG\,+10-16-111} -- The 21-cm spectrum shows two components of similar strength. As explained in Sect.~\Sfitmethod, we fix each one in turn and fit the other, in order to determine a systematic error. As can be seen in Table~\Tres, the sum of the two stays more or less constant. As we cannot reliably compare the 21-cm and \Lya\ column densities, we assign Q=1 to this sightline. \par {\it MRC\,2251$-$178} -- Intrinsic \Lya\ emission peaks at 1296~\AA, creating a moderately strong emission wing for measuring the Galactic \Lya\ absorption. A third order polynomial provides a good fit from 1195~\AA\ to 1241~\AA. Since the 21-cm spectrum is simple and the upward slope is minor, we assign Q=3 to this sightline. \par {\it Mrk\,110} -- Intrinsic \Lya\ emission peaks at 1259~\AA, creating an upward slope in the continuum. A third order polynomial fits this wing adequately and provides an adequate fit for the rest of the continuum in the wavelength range plotted, resulting in Q=3. \par {\it Mrk\,205} -- The 21-cm spectrum contains multiple \HI\ components including absorption originating in NGC\,4319 at 1289~\kms\ and an HVC at $v$=$-$204~\kms. The HVC is relatively strong, resulting in a 0.04~dex systematic error in the value of \NHILya\ for the Galactic emission. Extra curvature is present in the continuum and a second order polynomial was used to handle this. On balance, however, the result is reliable and we assign Q=3. \par {\it Mrk\,279} -- Many factors contribute to a low quality (Q=1) fit. Multiple \HI\ components are present in the 21-cm spectrum, but only one component at $v$=$-$40~\kms\ is used for the fitting. Strong \Lya\ emission is also present at 1252.69~\AA\ creating a large wing in the continuum. A fourth order polynomial is used to handle these features and fits the continuum adequately well. \par {\it Mrk\,335} -- Intrinsic \Lya\ emission peaks at 1247.02~\AA, creating a large wing beginning around 1208~\AA\ and producing significant curvature in the continuum. A fourth order polynomial fits this continuum moderately well. \par {\it Mrk\,478} -- An order three polynomial is needed to handle the curvature in the continuum, but the fitting lines provide a good fit over the range of 1201~\AA\ to 1234~\AA. The curvature is caused by intrinsic \FeIII-1122 emission, centered at 1209.02~\AA. \par {\it Mrk\,509} -- The 21-cm spectrum contains multiple components, including an IVC at $v$=61~\kms. All components were included in the measurement of the 21-cm and \Lya\ \HI\ column density. Strong intrinsic \Lya\ emission at 1256~\AA\ creates a large wing. A fourth order polynomial fits this wing and the continuum well from 1180~\AA\ to 1230~\AA. \par {\it Mrk\,771} -- The continuum is flat across the Galactic \Lya\ line. Although the 21-cm spectrum show two strong components, their separation in velocity is low enough that the final fit to the \Lya\ absorption results in Q=4. \par {\it Mrk\,876} -- The 21-cm spectrum shows multiple \HI\ components, including an HVC at $v$=$-$130~\kms, thus a systematic error in the \NHI\ value is present due to the HVC's impact. Intrinsic \OVI\ emission at 1165~\AA\ and 1171~\AA\ also causes the continuum to slope downward at the short-wavelength side of \Lya, but an order two polynomial still provides a good fit. These complications lead us to assign quality factor 2 to this measurement. \par {\it Mrk\,926} -- A third order polynomial provides a good fit to the continuum. However, multiple factors contribute to a systematic error in the value of \NHI. One factor is strong intrinsic \Lya\ emission at 1273~\AA\ that creates a broad wing in the continuum. Another factor is the fact that the G140M appears to show an extra upturn to the continuum below wavelengths of 1200~\AA\ (not easily visible in Fig.~\Fspectra). Combined with a relatively noisy spectrum, we decided to assign Q=1 to this sightline. \par {\it Mrk\,1044} -- Strong instrinsic \Lya\ emission at 1235.86~\AA\ adds a large wing to the continuum. A fourth order polynomial fits this wing well and fits the continuum adequately from 1203~\AA\ to 1223~\AA. The strong curvature across \Lya\ results in Q=2 for this target. \par {\it Mrk\,1383} -- An order two polynomial is needed to handle the slight curvature in the continuum, but the fit is still good enough to derive a result with quality factor 4. \par {\it Mrk\,1513} -- The continuum is flat, except near 1197~\AA\ where it shows an upturn, although there are no known intrinsic emission features near this wavelength. The tail-end of intrinsic \Lya\ emission also causes the continuum to rise at wavelengths above 1235~\AA. The net result of both issues is that the continuum fit is not as reliable as it might seem, resulting in quality 3 for this measurement. \par {\it NGC\,985} -- Intrinsic \Lya\ emission peaks at 1268~\AA, creating a moderate upward slope near Galactic \Lya. A third order polynomial provides a good fit from 1195~\AA\ to 1240~\AA. \par {\it NGC\,1705} -- The continuum is too complicated to make a good fit due to intrinsic \NV\ emission at 1241.42~\AA\ and 1245.41~\AA\ as well as intrinsic \Lya\ emission from the galaxy, which has a redshift of only 628~\kms. Therefore, we do not derive \NHILya\ for this sightline. \par {\it NGC\,3516} -- The continuum is too difficult to obtain a good fit due to moderate to high noise and strong intrinic \Lya\ emission at 1225~\AA. Therefore, we do not derive \NHILya\ for this sightline. \par {\it NGC\,3783} -- The continuum is too difficult to obtain a good fit due to the strong \Lya\ emission at 1227.50~\AA, as well as the presence of multiple strong \HI\ components at intermediate and high velocity. Therefore, we do not derive $N$(\HI;\Lya) for this sightline. \par {\it NGC\,4051} -- Strong \Lya\ emission at 1218.51~\AA\ and \NV\ emission at 1241.71~\AA\ and 1245.71~\AA\ creates a continuum too difficult to obtain a good fit. \par {\it NGC\,4151} -- The continuum is too difficult to make a good fit due to strong \Lya\ emission at 1219.70~\AA\ and \NV\ emission at 1242.93~\AA\ and 1246.93~\AA. \par {\it NGC\,5548} -- Strong \Lya\ emission at 1236.55~\AA\ adds a large wing to the continuum. A fourth order polynomial provides an adequate fit. Since the 21-cm profile is simple, the measurement is given a final quality of 3. \par {\it NGC\,7469} -- \Lya\ emission at 1235.51~\AA\ creates a large rise in the continuum, but this is fit adequately well by a fourth order polynomial over the range of 1197~\AA\ to 1230~\AA. This results in a quality 3 measurement. \par {\it PG\,0804+761} -- The continuum is flat from 1201~\AA\ to 1234~\AA\ but is pushed above the fitting line at 1196~\AA\ and below at 1245~\AA. The final fit is given quality 3. \par {\it PG\,0953+414} -- Intrinsic \CIII\ emission is present at 1204.92~\AA\, which adds curvature that creates a hill in the continuum over a range from 1188~\AA\ to 1236~\AA. A fourth order polynomial fits this hill well over a range of 1205~\AA\ to 1248~\AA, but the curvature leads us to assign Q=3 to this measurement. \par {\it PG\,1001+291} -- The quality of this continuum is diminished by multiple factors, including a moderate S/N ratio and the continuum resting above the fitting line at 1183~\AA. \par {\it PG\,1004+130} -- A high level of noise in the continuum makes it difficult to obtain a good quality fit and a reliable value of \NHI\ for this QSO. \par {\it PG\,1049$-$005} -- The continuum is flat, but the high level of noise greatly reduces the quality of the fit. As a result, we assign Q=1 to this sightline. \par {\it PG\,1103$-$006} -- Low S/N makes obtaining a reliable fit impossible for this target. \par {\it PG\,1116+215} -- The intrinsic \OVI\ emission lines are redshifted to 1214.06~\AA\ and 1220.75~\AA, which results in a large bump in the continuum across the Galactic \Lya\ line. This can be modeled by using a fourth order polynomial, which fits the continuum adequately well from 1180~\AA\ to 1248~\AA. The 21-cm spectrum shows two components of similar strength at $-$39 and $-$5~\kms. These factors combine to yield Q=1 for the resulting \Lya\ column density. \par {\it PG\,1149$-$110} -- Intrinsic \Lya\ emission peaks at 1275.24~\AA\ and a low S/N ratio make a good fit for the continuum too difficult to obtain. \par {\it PG\,1211+143} -- Intrinsic \FeIII\ $\lambda$1122.52 emission is redshifted to 1212.77~\AA, which pushes the continuum slightly upward on the short-wavelength side of Galactic \Lya. This lowers the quality of the fit, although a fourth order polynomial is used and fits the continuum adequately well from 1180~\AA\ to 1248~\AA. The final fit quality is assigned a value of 3. \par {\it PG\,1259+593} -- Multiple \HI\ components are present in the continuum including an HVC at $v$=$-$127~\kms\ and an IVC at $v$=$-$52~\kms. A two-sided fit is used, as described in Sect.~\Sfitmethod. \par {\it PG\,1302$-$102} -- The continuum is flat, but the low S/N visibly diminishes the quality of the fit. \par {\it PG\,1341+258} -- The continuum is slightly above the fitting line from 1223~\AA\ to 1230~\AA, but a linear fit still provides a high quality (Q=4) fit. \par {\it PG\,1351+640} - The continuum contains multiple absorption features, lowering the quality of the fit. The 21-cm spectrum shows multiple \HI\ components including an HVC at $v$=$-$156~\kms. Thus, a systematic error is introduced into the value for \NHI\ due to the HVC's impact on the continuum. Intrinsic \FeIII\ emission is also present at 1221.53~\AA. Although a fourth order polynomial is used to deal with these features and provides an adequate fit for the continuum, the uncertainty associated with the HVC's column density is such that we assign a final quality factor of 1. \par {\it PG\,1444+407} -- A flat continuum gives an acceptable fit. However, it is a little above the fitting line from 1180~\AA\ to 1200~\AA\ and from 1233~\AA\ to 1237~\AA, thus reducing the quality of this fit. Combined with the multiple 21-cm components, the final column density value is quality 2. \par {\it PHL\,1811} -- A flat continuum is a good fit despite the presence of intrinsic \OVI\ emission at 1230.06~\AA\ and 1236.84~\AA, which pushes the continuum slightly above the fitting line from 1224~\AA\ to 1236~\AA. Combined with the simplicity of the 21-cm profile, the final quality for this sightline is 4. \par {\it PKS\,0312$-$77} -- The 21-cm spectrum contains multiple \HI\ components, including the Magellanic Bridge at $v$=191~\kms\ and $v$=166~\kms. A two-sided fit is used as described in Sect.~\Sfitmethod. Lehner et al.\ (2008) gives \Lya-derived column densities of 20.78$\pm$0.06 and 20.12$\pm$0.30 for components at 5 and 210~\kms, respectively. The combined value is 20.86. From the LAB data, we find a column densities of 20.83 and 20.23 for these two components. Varying the Magellanic Bridge component between 20.14 and 20.31 results in values for the low-velocity column density of 20.67$\pm$0.18 to 20.71$\pm$0.17. A combined fit yield values between 20.79$\pm$0.14 and 20.87$\pm$0.11, depending on the precise selections. Thus, we find a column density for the Magellanic Bridge component that is 0.1 dex higher than that of Lehner et al, but within their error. We also find that Lehner et al.\ underestimated the error for the low-velocity gas by a factor $\sim$3. \par {\it PKS\,0405$-$12} -- The continuum is flat but a few extra features reduce the quality of the fit. There may be \NeVIII\ emission at 1211.06~\AA\ and intrinsic \OIV\ emission at 1238.75~\AA. \NIII\ emission at 1200.42~\AA\ pushes the continuum up slightly, and a dip from 1243~\AA\ to 1255~\AA\ pulls the continuum down, creating a twist in the continuum that causes the fitting line to slope upward. In the end, we only assign Q=2 to this sightline. \par {\it PKS\,2005$-$489} -- This spectrum is flat across Galactic \Lya, and the sightline has a simple 21-cm profile. The derived \Lya\ column density is quality four. \par {\it PKS\,2155$-$304} -- This spectrum is flat across Galactic \Lya, and the sightline has a simple 21-cm profile. The derived \Lya\ column density is quality four. \par {\it RX\,J0100.4$-$5113} -- The continuum placement is too uncertain for this QSO to obtain a good fit or a reliable value for \NHI. The continuum is further complicated by the presence of an HVC at $v$=92~\kms. \par {\it RX\,J1830.3+7312} -- The moderate curvature of the continuum reduces the quality of the fit to Q=3, but a third order polynomial still provides a good fit for this QSO. \par {\it Ton\,S180} -- The fit is of lower quality due to significant curvature in the continuum. This results from the intrinsic \Lya\ line centered at 1291~\AA, and a rise toward the lower wavelength edge that is seen in many G140M spectra, which is probably due to a calibration problem. \par {\it Ton\,S210} -- The fit is of lower quality due to multiple factors. One is the moderate S/N. Another is the uncertainty of the continuum, which rises toward the lower wavelengths due to the wing of the intrinsic \OVI\ emission, centered at about 1155~\AA. \par {\it UGC\,12163} -- This spectrum is relatively noisy, and the wing of the intrinsic \Lya\ line, centered at 1245~\AA, extends above the Galactic \Lya\ absorption, making a determination of the continuum almost impossible. \par {\it VII\,Zw\,118} -- A flat continuum is present, but the slope combined with a continuum dip between 1195~\AA\ to 1205~\AA\ reduces the quality of the fit.

1012.5319

We present a study of the small-scale structure of the interstellar medium (ISM) in the Milky Way. We used HST STIS data to measure N(H I) in a pencil beam toward 59 active galactic nuclei and compared the results with the values seen at 9'-36' resolution in the same directions using radio telescopes (Green Bank Telescope, Green Bank 140-ft, and LAB survey). The distribution of ratios N(Lyα)/N(H I) has an average of 1 and a dispersion of about 10%. Our analysis also revealed that spectra from the Leiden-Argentina-Bonn (LAB) all-sky H I survey need to be corrected, taking out a broad Gaussian component (peak brightness temperature 0.048 K, FWHM 167 km s<SUP>-1</SUP>, and central velocity -22 km s<SUP>-1</SUP>). The column density ratios have a distribution showing similarities to simple descriptions of hierarchical structure in the neutral ISM as well as to a more sophisticated three-dimensional magnetohydrodynamic simulation. From the comparison with such models, we find that the sonic Mach number of the local ISM should lie between 0.6 and 0.9. However, none of the models yet matches the observed distribution in all details, but with many more sightlines (as will be provided by the Cosmic Origins Spectrograph) our approach can be used to constrain the properties of interstellar turbulence.

12225829

2011PhRvD..83k3011D

1012

1012.5627_arXiv.txt

The large mixing angle (LMA) MSW solution \cite{w1,ms} has been established as the solution of the solar neutrino problem \cite{hom}, \cite{Abdurashitov:2009tn}, \cite{Altmann:2005ix}, \cite{Hosaka:2005um}, \cite{Aharmim:2005gt}, \cite{Aharmim:2007nv}, \cite{Aharmim:2008kc}, \cite{Arpesella:2008mt}. In assumption of the CPT conservation KamLAND confirms this result \cite{:2008ee}, \cite{Gando:2010aa}. One of the main goals of further precision measurements of the solar neutrino fluxes is to search for possible deviations from the LMA predictions which would indicate new physics beyond the Standard Model with three mixed neutrinos. In particular, new physics can show up at the neutrino energies $E = (1 - 7)$ MeV, {\it i.e.} in the transition region between the matter dominated conversion and vacuum oscillations. Here direct measurements of the spectrum are absent or inprecise and possible deviations from the LMA predictions can be relatively large. Some time ago in attempt to explain the low (about $2\sigma$) rate in the Homestake experiment~\cite{hom} in comparison to the LMA expectation as well as the absence of clear low energy upturn of the spectra of events at SuperKamiokande and SNO we have proposed a scenario with light sterile neutrino, $\nu_s$, which mixes weakly with active neutrinos \cite{pedroS}. Conversion of $\nu_e$ to $\nu_s$ driven by the mass squared difference $\Delta m^2_{01} \sim (0.2 - 2) \cdot 10^{-5}$ eV$^2$ and mixing in the mass state $\nu_1$, $\sin^2 2\alpha \sim 10^{-3}$, leads to appearance of a dip in the $\nu_e - \nu_e$ survival probability in the range (0.5 - 7) MeV which explains the data. After publication \cite{pedroS} several new experimental results have appeared which further support our proposal: \begin{itemize} \item Measurements of the solar neutrino spectrum by SuperKamiokande-III \cite{Abe:2010hy} with lower threshold still do not show the upturn. \item The SNO LETA analysis \cite{Aharmim:2009gd} gives even turn down of the spectrum in the two lowest energy bins. \item The Borexino measurements of the boron neutrino spectrum also hint some tendency of the spectral turn down \cite{Bellini:2008mr}. \end{itemize} Although separately these results are not statistically significant, being combined they can be considered an evidence of some new sub-leading effect. At the same time, the cosmological observations indicate possible presence of additional radiation in the Universe in the epoch of last photon scattering. This is quantified by the effective number of neutrino species, $N_{eff}$, which is bigger than 3. Combined analysis of WMAP-7, measurements of BAO (Baryon Acoustic Oscillations) and new value of the Hubble constant $H_0$) gives $N_{eff} = 4.34^{+ 0.86}_{- 0.88}$ \cite{Komatsu:2010fb}. WMAP-7 and Atacama Cosmology Telescope data lead to $N_{eff} = 5.3 \pm 1.3$ (68 \% C.L.) \cite{Dunkley:2010ge}. In the independent analysis~\cite{Hamann:2010bk} of these data the number of very light sterile neutrinos $\Delta N_{eff} = (0.02 - 2.2)$ (68 \% C.L.) has been obtained. All this confirms the earlier finding based on the WMAP-3 data: $N_{eff} = 5.3^{+ 0.4}_{- 0.6}{}^{+ 2.1}_{- 1.7}{}^{+ 3.8}_{-2.5}$~\cite{Seljak:2006bg}. These results do not contradict the recent Big Bang Nucleosynthesis (BBN) bounds $N_{eff} = 3.68^{+ 0.80}_{- 0.70}$ \cite{Izotov:2010ca} (see discussion in \cite{BBN} and theoretical considerations in \cite{BBNint}). Hence an additional radiation can be produced before the BBN epoch. In this connection we revisit our proposal of very light sterile neutrinos. We show that mixing of this neutrino in mass states $\nu_1$ or/and $\nu_2$ can consistently improve description of the solar spectral data. We introduce mixing of this neutrino in the mass eigenstate $\nu_3$ which allows $\nu_s$ to be produced in the Early Universe with nearly equilibrium concentration, so that $\Delta N_{eff} \approx 1$. The paper is organized as follows. In sect. 2 we consider properties of the $\nu_e$ conversion in the presence of $\nu_s-$mixing in the Sun generalizing our analysis in \cite{pedroS}. New feature, wiggles'' in the survival probability, is described which appear for relatively large $\Delta m^2_{10}$ at the $E > 5$ MeV. In sect. 3 we obtain bounds on the $\nu_s$ parameters from the Borexino measurements of the $Be-$neutrino flux. Spectra of the solar neutrino events have been computed for different experiments and confronted with the data. In sect. 4. the mixing of $\nu_s$ in $\nu_3$ is introduced and phenomenological consequences of this mixing are studied, in particular, generation of $\nu_s$ in the Early Universe. The conclusion is given in sect. 5. In appendix we give some details of appearance of the wiggles in the survival probability.

1. Recent measurements of the energy spectra of the solar neutrino events at SuperKamiokande, SNO, Borexino do not shown the expected (according to LMA) upturns at low energies. The absence of the upturn can be explained by mixing of very light sterile neutrino in the mass states $\nu_1$ or/and $\nu_2$ with $\Delta m^2_{01} \sim (0.7 - 2) \cdot 10^{-5} $ eV$^2$ ($R_\Delta = 0.07 - 0.25$) and mixing $\sin^2 2 \alpha = (1 - 5) \cdot 10^{-3}$. Such a mixing leads to appearance of the dip in the $\nu_e-$ survival probability in the energy range (1 - 7) MeV, thus removing the upturn of the spectra. For $\Delta m^2_{01} \sim 2 \cdot 10^{-5} $ eV$^2$ and $\sin^2 2 \alpha \sim 5 \cdot 10^{-3}$ the $\nu_e - \nu_s$ conversion can even produce a turn down of the spectra. Description of the existing solar neutrino data in the presence of mixing with sterile neutrino is apparently improved. Values of $\Delta m^2 < 0.6 \cdot 10^{-5}$ eV$^2$ (for mixing angle interval $\sin^2 2 \alpha = (1 - 5) \cdot 10^{-3}$) are excluded by the Borexino measurements of the $Be-$neutrino flux. The presence of the dip can be further tested in future precision measurements of the low energy part of the $B-$neutrino spectrum as well as the $pep-$ neutrino flux.\\ \noindent 2. Mixing of $\nu_s$ in the $\nu_3$ mass eigenstate with $|U_{s3}|^2 \sim 0.02 - 0.2$ leads to production of significant concentration of $\nu_s$ via oscillations in the Early Universe. For $|U_{s3}|^2 \sim 0.1 - 0.2$ nearly equilibrium concentration can be obtained both in neutrino and antineutrino channels thus generating additional effective number of neutrinos $\Delta N_{eff} \sim 1$ before the BBN epoch. This can explain recent cosmological observations. \\ \noindent 3. Mixing of $\nu_s$ in $\nu_3$ leads to a number of phenomenological consequences, in particular, it can affect the atmospheric and accelerators neutrino fluxes as well as fluxes of the SN neutrinos. The mixing leads to existence of the $\nu_s - \nu_\tau'$ resonance. For neutrinos crossing the Earth the resonance should appear at energies $E \sim 10 - 15$ GeV. This can be tested in future atmospheric neutrino studies with Megaton-scale detectors as well in the long baseline experiments with accelerator neutrino beams.

1012.5627

Recent results from the SNO, Super-Kamiokande, and Borexino experiments do not show the expected upturn of the energy spectrum of events (the ratio R≡N<SUB>obs</SUB>/N<SUB>SSM</SUB>) at low energies. At the same time, cosmological observations testify for the possible existence of additional relativistic degrees of freedom in the early Universe: ΔN<SUB>eff</SUB>=1-2. These facts strengthen the case of a very light sterile neutrino, ν<SUB>s</SUB>, with Δm<SUB>01</SUB><SUP>2</SUP>∼(0.7-2)×10<SUP>-5</SUP>eV<SUP>2</SUP>, which mixes weakly with the active neutrinos. The ν<SUB>s</SUB> mixing in the mass eigenstate ν<SUB>1</SUB> characterized by sin⁡<SUP>2</SUP>2α∼10<SUP>-3</SUP> can explain an absence of the upturn. The mixing of ν<SUB>s</SUB> in the eigenstate ν<SUB>3</SUB> with sin⁡<SUP>2</SUP>β∼0.1 leads to production of ν<SUB>s</SUB> via oscillations in the Universe and to additional contribution ΔN<SUB>eff</SUB>≈0.7-1 before the big bang nucleosynthesis and later. Such a mixing can be tested in forthcoming experiments with the atmospheric neutrinos, as well as in future accelerator long baseline experiments. It has substantial impact on conversion of the supernova neutrinos.

12168032

2011ApJ...731..121B

1012

1012.2007_arXiv.txt

The most general stationary and axisymmetric black-hole (BH) solution of Einstein's equations in a four-dimensional, asymptotically flat spacetime is given by the Kerr geometry~\citep{Kerr:1963ud}. Today there are at least two classes of astrophysical BH candidates: stellar-mass objects in X-ray binary systems (mass $M \sim 5 - 20$~$M_\odot$)~\citep{stmbhs1,stmbhs2} and super-massive objects at the center of most galaxies ($M \sim 10^5 - 10^{10}$~$M_\odot$)~\citep{sumbhs}. The existence of a third class of objects, intermediate-mass BHs with $M \sim 10^2 - 10^4$~$M_\odot$~\citep{imbhs}, is still controversial because their detections are indirect and definitive dynamical measurements of their masses are still lacking~\citep{imbhs}. All these objects are supposed to be Kerr BHs because they cannot be explained otherwise without introducing new physics. In particular, stellar-mass BH candidates in X-ray binary systems are too heavy to be neutron or quark stars for any reasonable matter equation of state~\citep{ruf,kal}. Observations of stellar orbits around the super-massive BH candidate Sgr A$^\star$ at the center of the Galaxy show that this object is too massive, compact, and old to be a cluster of non-luminous bodies~\citep{maoz} or a fermion ball~\citep{2002Natur.419..694S} (\textit{i.e.,} an object made of sterile neutrinos, gravitinos or axinos supported by degeneracy pressure~\citep{1998ApJ...500..591T}). Other exotic alternatives such as boson stars~\citep{boson} and gravastars~\citep{gravastar,gravastar2,gravastar3} seem to be disfavored by the near-infrared observations of Sgr A$^\star$~\citep{2006ApJ...638L..21B,2009ApJ...701.1357B}. In spite of this body of indirect evidence, a definitive proof that BH candidates are really described by the Kerr solution of General Relativity is still elusive. A framework within which to test the Kerr BH hypothesis was first put forward by~\citet{ryan_multipoles,ryan_3,ryan_4}, who considered a general stationary, axisymmetric, asymptotically flat, vacuum spacetime. Such a generic spacetime can be used to describe the gravitational field around a central object, whatever its nature, and its metric can be expressed in terms of the mass moments $M_\ell$ and current moments $S_\ell$~\citep{multipoles1,multipoles2}. Assuming reflection symmetry, the odd $M$-moments and even $S$-moments are identically zero, so that the non-vanishing moments are the mass $M_0=M$, the mass quadrupole $M_2=Q$ and the higher-order even terms $M_4, M_6, \ldots$, as well as the angular momentum $S_1=J$, the current octupole $S_3$ and the higher-order odd terms $S_5, S_7, \ldots$. In the case of a Kerr BH, all the moments $M_\ell$ and $S_\ell$ are locked to the mass and angular momentum by the following relation: \begin{equation}\label{kerrMultipoles} M_\ell+{\rm i}S_\ell=M\left({\rm i}\frac JM\right)^\ell\;. \end{equation} This is the celebrated ``no-hair'' theorem \citep{hair1,hair2,hair3}: an (uncharged) stationary BH is uniquely characterized by its mass and spin angular momentum. Therefore, a measurement of the mass, spin and higher moments (starting with the quadrupole moment $Q$) of BH candidates would permit testing Eq.~(\ref{kerrMultipoles}) and therefore the Kerr-nature of these objects. Ryan's idea was to use future gravitational-wave observations of extreme-mass ratio inspirals (EMRIs, \textit{i.e.,} systems consisting of a stellar-mass BH orbiting a super-massive BH in a galactic center) to perform this test. EMRIs will be a key source for the future-space based detector LISA: because the stellar-mass BH will orbit the super-massive BH $\sim 10^6$ times during LISA's lifetime, slowing spiralling in due to the emission of energy and angular momentum via gravitational waves, even a small deviation from the Kerr geometry will build up an observable dephasing in the gravitational waveforms, thus allowing one to map the spacetime of super-massive BHs with very high accuracy. Ryan's spacetime mapping idea originated a whole line of research aiming at using LISA's observations of EMRIs to test the Kerr nature of super-massive BHs~\citep{bumpy1,bumpy2,apostolatos1,apostolatos2,kostas,gair,kesden_boson,torus,hydrodrag,barack_cutler} and even General Relativity itself~\citep{nicoTestCS,comment_on_psaltis}. Another independent (and complementary) test of the no-hair theorem with LISA uses BH quasi-normal modes~\citep{emanuele1}. Because the frequencies of these modes encode the multipolar structure~(\ref{kerrMultipoles}) of the Kerr geometry, they can be used to test consistency with the Kerr solution and to distinguish it from boson stars~\citep{emanuele2} or gravastars~\citep{gravastar2}. Besides these tests based on gravitational waves, there are other proposals using electromagnetic radiation. Constraints on the quadrupole moment of the compact companion of radio pulsars can be obtained with timing measurements~\citep{wex}. Astrometric monitoring of stars orbiting at milliparsec distances from Sgr A$^\star$ may be used to test the no-hair theorem for the super-massive BH candidate at the center of the Galaxy~\citep{will1,will2}. A very promising way to measure deviations from the Kerr metric is represented by future observations of the ``shadow'' of super-massive BH candidates through very long baseline interferometry (VLBI) experiments \citep{naoki,psaltis1,psaltis2}. The study of quasi-periodic variability in BH spectra may also test the geometry of the spacetime around BH candidates~\citep{psaltis3}. Remarkably, \citet{psaltis4} also shows that the data for iron K$\alpha$ emission lines in thin accretion disks can \textit{already} constrain deviations from the Kerr geometry. Although these measurements yield much less accurate constraints than what will be achieved with LISA, and can be subject to critiques (see~\citet{tita}, who show that iron-K$\alpha$ lines with the same features as those attributed to BH candidates are observed also around white dwarfs), these data are available \textit{now}, which is not the case for all the other tests reviewed above. However, because of the controversial interpretation of the origin of these lines, and because \citet{psaltis4} finds a degeneracy between the spacetime's quadrupole and spin (\textit{i.e.,} similar shapes for the iron-K$\alpha$ lines can be obtained with a Kerr BH or with a non-Kerr object with spin and quadrupole slightly shifted from the Kerr values), it is important to explore other techniques to test the no-hair theorem with \textit{present} data. In this paper we propose using the continuum spectrum of BH candidates, which has been shown to be potentially a promising tool to tell Kerr BHs from specific alternatives such as gravastars~\citep{harko_gs}, BHs in Chern-Simons gravity~\citep{harko_cs}, BHs in Horava-Lifschitz gravity~\citep{harko_hl} or certain classes of naked singularities~\citep{harko_ws,rohta,harko_ns}. While these attempts highlighted some important differences between the spectra of these objects and those of Kerr BHs, they relied on specific models for the BH candidate, and did not investigate whether presently available data allow one to break the degeneracy mentioned above between the parameters of these objects and those of a Kerr BH (\textit{i.e.,} whether present X-ray data can tell the spectrum of a non-Kerr object from that of a Kerr BH with arbitrary $J$ and $M$). In this paper we address both issues, \textit{(i)} by considering a very general model for the BH candidate (\textit{i.e.,} one which allows its quadrupole moment to slighlty deviate from the Kerr value, thus approximately describing a variety of almost-Kerr objects), \textit{(ii)} by comparing our model to present X-ray data, although in a simplified way, and \textit{(iii)} by discussing the sources of systematic error that might affect the data and that must be properly understood before one can perform robust tests of the no-hair theorem. In the range 0.1~keV -- 1~MeV, the generic spectrum of a stellar-mass BH candidate is characterized by three components, even if their relative intensities vary with the object and, for a given object, with time: $i)$ a soft X-ray component (energies $< 10$~keV), $ii)$ a hard power law X-ray component with an exponential cutoff (energies in the range $10-200$~keV, photon spectral index in the range $1-2.5$), and $iii)$ a $\gamma$-ray component (energies $> 300$~keV). For a review, see e.g. \citet{liang}. The soft X-ray component is commonly interpreted as the thermal spectrum of a thin disk, while the exact origin of the other two components is not so clear. Geometrically thin and optically thick accretion disks can be described by the Novikov-Thorne model~\citep{n-t}. They are expected when the accretion flow is radiatively efficient, which requires a luminosity $L \lesssim 0.3$~$L_{Edd}$, where $L_{Edd}$ is the Eddington limit. The emission is blackbody-like. Assuming that the inner edge of the disk is at the innermost stable circular orbit (ISCO)\footnote{Such an assumption is supported either by observational facts~\citep{steiner} and numerical simulations~\citep{shafee,penna} (but see~\citet{krolik}).}, the disk luminosity of a Kerr BH is determined only by its mass, $M$, the mass accretion rate, $\dot{M}$, and the spin parameter, $a=J/M^2$. This fact can thus be exploited to estimate the spin of stellar-mass BH candidates~\citep{zhang}. This is the continuum fitting method and at present has been used to estimate the spin parameter of a few stellar-mass BH candidates~\citep{mcclintock}\footnote{For super-massive BHs, the disk temperature is lower (the effective temperature scales like $M^{-0.25}$) and this approach cannot be applied.}. Basically, knowing the mass of the object, its distance from us, and the inclination angle of the disk, it is possible to fit the soft X-ray component of the source and deduce $a$ and $\dot{M}$. In this paper, we compute the thermal spectrum of a geometrically thin and optically thick accretion disk around a generic compact object. We use a subclass of Manko-Novikov spacetimes~\citep{m-n}, which are stationary, axisymmetric, and asymptotically flat exact solutions of the vacuum Einstein equations. In addition to the mass and the spin of the massive object, here we have the anomalous quadrupole moment, $q$. The latter measures the deformation of the massive object with respect to a Kerr BH: when $q>0$, the object is more oblate than a Kerr BH, when $q<0$, it is more prolate, while, for $q=0$, we recover the Kerr metric. The value of $q$ determines the radius of the ISCO and changes the high frequency region of the spectrum of the disk. In general, this makes the spectrum of the disk almost degenerate in $a$ and $q$. However, only in the Kerr case the radius of the ISCO goes to $M$ as $a$ approaches 1. For $q\neq0$, even a small deviation from the Kerr metric makes the radius of the ISCO significantly larger than $M$. Since current X-ray observations suggest that there are objects with small ISCO radius, one can in principle obtain interesting constraints on the value of $q$. The purpose of this paper is therefore to present a preliminary investigation, showing that X-ray continuum spectra can potentially be used to constraint small quadrupole deviations away from the Kerr metric, once all the physical effects have been included in the model and all systematics have been understood. In particular, our computation of the disk's spectrum does not include the effect of light bending. This is a simplification of our model and there are no reasons for the light bending to be negligigle with respect to the other relativistic effects (Doppler boosting, gravitational redshift, and frame dragging). Another subtle issue is the computation of the spectral hardening factor (here not discussed), which is another weak point of our approach. Our study has to be taken as a preliminary investigation and significant work has still to be done before the continuum fitting method can be used to obtain reliable constraints on the Kerr geometry around stellar-mass BH candidates. The content of this paper is as follows. In Sec.~\ref{s-disk}, we review the basic properties of a geometrically thin and optically thick accretion disk and how to compute its spectrum. In Sec.~\ref{s-kerr} and \ref{s-mn}, we present the results of our calculations, respectively for a Kerr BH and for a generic object with $q \neq 0$. In Sec.~\ref{s-cf}, we show how current observations can be used to constrain $q$, while in Sec.~\ref{systematics} we discuss the possible systematic errors that could affect the continuum fitting technique and therefore out analysis. Lastly, in~\ref{s-concl} we report our conclusions. The Manko-Novikov spacetime is reviewed in Appendix~\ref{a-mn}, and the properties of its ISCO are discussed in Appendix~\ref{a-isco}. Throughout the paper we use units in which $G_{\rm N}=c=1$, unless stated otherwise.

\label{s-concl} If current astrophysical BH candidates are Kerr BHs, their spacetime should be completely specified by two parameters, namely their mass $M$ and spin $J$. This can be tested by measuring at least three multipole moments of the BH candidate. While there are several proposals to obtain such a measurement with future experiments, in this paper we have shown that current X-ray observations of stellar-mass BH candidates in binary systems can already be used to constrain possible deviations from the Kerr metric. We have computed the thermal spectrum of a geometrically thin and optically thick accretion disk around a compact object with mass $M$, spin parameter $|a| \le 1$ and arbitrary anomalous quadrupole moment $q$. For $q = 0$, we recover the Kerr metric. The exact value of $q$ determines the inner radius of the disk, changing the high frequency region of the spectrum. The effect is small for low spin parameters or for counterrotating disk, but it becomes important for higher values of $a$. In general, the sole analysis of the disk spectrum cannot completely determine $q$, because the spectrum is degenerate in $a$ and $q$ and therefore one would need an independent measure of $a$. However, for very fast-rotating Kerr BHs the ISCO radius and therefore the disk's inner radius becomes very small. Any deviation from $q=0$ (\textit{i.e.,} from the Kerr solution) makes the inner radius grow quickly. Since current observations suggest that the inner radius of the accretion disk of some stellar-mass BH candidates is close to the gravitational radius $R_g=G M/c^2$, one can constrain the anomalous quadrupole moment of these objects very efficiently. In this paper we have considered a specific example, the stellar-mass BH candidate M33~X-7, whose estimated spin is $a = 0.84 \pm 0.05$ if one assumes it is a Kerr BH~\citep{m33x7,m33x7e}. Since stronger constraints on $q$ can be obtained from objects with higher $a$, we could have considered GRS~$1915+105$, whose spin parameter has been estimated to be larger than 0.98 in~\citet{1915} under the Kerr-BH assumption. However, the measurements of the distance, mass and viewing angle of M33~X-7 are more reliable, thus making this object more suitable to obtain preliminary constraints on the $a-q$ plane. To move our analysis beyond the simplified and preliminary stage we achieved in this paper, it is of paramount importance to properly understand the systematic errors that might affect the continuum fitting method, and which could in principle blur the difference between the spectra of Kerr BHs and those of other objects, and affect the measurements of the spin even if one adopts the Kerr BH hypothesis. Moreover, we will have to amend our disk model by including the following ingredients: \begin{enumerate} \item {\it The effect of light bending}. A rigorous computation of the spectrum requires to trace the light rays from the surface of the accretion disk to the distant observer in the background metric. The effect of light bending is presumably no less important than the other relativistic effects and further alters the observed spectrum. \item {\it The spectral hardening factor}. In the inner part of the accretion disk, the temperature is high and non-blackbody effects cannot be neglected. We thus need an accurate model of the disk atmosphere for computing the spectral hardening factor~\citep{shimura,merloni,davis}. \item Additional effects to be considered in an accurate study are the ones of limb darkening and of returning radiation. \end{enumerate}

1012.2007

Black holes in general relativity are known as Kerr black holes and are characterized solely by two parameters, the mass M and the spin J. All the higher multipole moments of the gravitational field are functions of these two parameters. For instance, the quadrupole moment is Q = -J <SUP>2</SUP>/M, which implies that a measurement of M, J, and Q for black hole candidates would allow one to test whether these objects are really black holes as described by general relativity. While future gravitational-wave experiments will be able to test the Kerr nature of these objects with very high accuracy, in this paper we show that it is possible to put constraints on the quadrupole moment of stellar-mass black hole candidates by using presently available X-ray data of the thermal spectrum of their accretion disk.

12163389

2011A&A...526A..83K

1012

1012.0833_arXiv.txt

Recurrent novae (RNe) are semi-detached binary systems consisting of a massive white dwarf (WD; M$_{WD}$$\ge$1.2M$_{\sun}$) accreting material from a companion, with mass transfer rates of $\ge$10$^{-7}$M$_{\sun}$/year (\cite{2000ApJ...534L.189H}). With total system mass exceeding the Chandrasekhar mass limit, they are prime candidates for SNeIa progenitors, therefore they are monitored diligently. Their ``recurrent'' designation reflects the fact that they exhibit multiple nova eruptions on time scales of several decades -- more frequent than classical novae whose eruptions have been observed only once. Therefore they are excellent natural test beds for the study and understanding of stellar explosions in accreting systems and, of course, the mechanism of mass accumulation on a near-Chandrasekhar limit WD. An excellent recent review on the properties of RNe can be found in \cite{2010ApJS..187..275S}. The object of this work, U Sco, is one of the most well-studied recurrent novae, not only for its short intra-outburst rate\footnote{The system has 10 recorded nova outbursts, starting in 1863, with an average inter-erruption interval of a decade; it also holds a record short inter-eruption interval of 7.88 years (1979 to 1987).} but also because it is an eclipsing binary with an orbital period of 1.23 days (Schaefer 1990, \cite{1995ApJ...447L..45S}). It consists of a massive WD (M$_{WD}$$\ge$1.2M$_{\sun}$; \cite{2000AJ....119.1359A}) and a subgiant mass-losing companion whose spectral type ranges between F8 (\cite{1992ApJ...396..267J}) and K2 (\cite{2000AJ....119.1359A}; also, Hanes 1985, Schaefer 1990). Because of its total mass exceeding the Chandrasekhar mass limit of 1.4 M$_{\sun}$, it is considered a favorable candidate for a SN Ia progenitor. In quiescence, its optical magnitude ranges from V$\sim$18 (outside eclipse) to V$\sim$19 in eclipse (Schaefer et al. 2010); the maximum brightness in outburst reaches V$\sim$8.0 mag. Its rapid decline by 3 magnitudes from maximum light occurs in 4 days (figure 1) making it a very fast recurrent nova. The principal mechanism leading to the eruption is believed to be a thermonuclear runaway on the mass-accumulating WD, enriching the ISM with CNO-enhanced material. This short communication presents one of the earliest U Sco 2010 outburst spectra, obtained at moderate spectral resolution within 24 hours of discovery and peak brightness and showing strong, wide Balmer lines with fast variability and blueshifted/redshifted velocities reaching 3000 km/sec.

This paper presents the earliest medium-resolution spectrum of U Sco one day after its 2010 outburst, covering the optical region of the spectrum (3000-8500 $\AA$). Wide, multicomponent Balmer emission lines are present, reaching velocities of $\sim$3000km/sec, indicative of an expanding shell resulting from the explosion. Broad emission components of Fe II, O III, He I, C III are identified in the blue part of the spectrum. The THEA complex is not present in U Sco, although the fact that this may be due to our viewing angle of the system can not be ruled out. Time-resolved spectra revealed intrinsic fast variations in the Balmer lines, similar to accretion disk flickering. \begin{table*} \caption{Kinematic characteristics of the main line profiles in the U Sco spectrum } \begin{tabular}{lcccccccc} \hline Line & rest wavelength & \multicolumn{3}{c}{Velocities of line components$^{\mathrm{a}}$} & EW$^{\mathrm{b}}$ & FWHM$^{\mathrm{c}}$ & em/abs & notes \\ & (\AA) & Blue & Red & center of line & (\AA) & (\AA)& & \\ & & (km/sec) & (km/sec) & (km/sec) & & & & \\ \hline H$\alpha$& 6563 & -2915 & 3283 & 150 & -309$\pm$5 & 109 & em &ejecta \\ NaD1& 5896 &-44 &-- &-- & 0.231$\pm$0.003 & 1.11 & abs & IS\\ &&-206 &-- &-- & 0.044$\pm$0.002 & 1.34 & abs & IS\\ NaD2& 5890 &-45&-- &-- & 0.313$\pm$0.003 & 1.15 & abs & IS\\ &&-174 & -- &-- &0.016$\pm$0.003 & 0.85 & abs & IS\\ H$\beta$& 4874 & -2966 & 3287 & -390 & -145$\pm$5 & 129 & em &ejecta \\ & &-3854 &-- &-- &-- &-- & abs & ejecta \\ & &-5864 &-- & -- & -- &-- & abs & ejecta\\ H$\gamma$ & 4342 & -- & -- & 584 & -138$\pm$6 & 99 & em &ejecta \\ & &-3806 & -- & -- & -- & --& abs & ejecta \\ & & -5864 & -- & -- & -- & -- & abs & ejecta \\ H$\delta$ & 4100 & -- & -- & 370 & -110$\pm$10 & 87 &em &ejecta \\ & &-3786 & -- & -- & -- & -- & abs & ejecta \\ & &-4304 & -- & -- & -- & -- & abs & ejecta \\ CaII~H & 3968 &-45 &-- &-- & 0.115$\pm$0.012 & 0.937& abs & IS\\ & & -187&-- &-- & 0.009$\pm$0.002 & 0.618 & abs &IS\\ & &-1547 & -- &-- & 2.340$\pm$0.050 & 14.92 &abs & ejecta\\ CaII~K & 3934 & -49 & -- &-- &0.181$\pm$0.023 & 0.902 & abs & IS\\ & &-187 &-- &-- & 0.018$\pm$0.005 & 0.617 & abs & IS\\ & & -1106$^{\mathrm{d}}$&-- &-- & 3.580$\pm$0.060 & 10.32 & abs & ejecta\\ \hline \end{tabular} \begin{list}{}{} \item[$^{\mathrm{a}}$]The blue/red velocities in the line profiles presented here were measured from the average spectrum with a Gaussian fit; the average error in the velocities is 2km/sec. The velocities of the central components of the Balmer lines weremeasured ith IRAF/splot's task ``e'', providing an average error of 5km/s. \item[$^{\mathrm{b}}$]Measured from the average spectrum. Following the conventional IRAF nomeclature, negative EW values correspond to emission lines. The errors in the measurements correspond to uncertainties in the determination of the line continuum. \item[$^{\mathrm{c}}$]Measured from the average spectrum; the error in FWHM measurements is less than 1$\AA$ in all cases. \item[$^{\mathrm{d}}$]Blend with H$\epsilon$ components. \end{list} \end{table*}

1012.0833

<BR /> Aims: We present optical spectra of the fast recurrent nova U Sco during its recent outburst, obtained within 24 h of maximum light. <BR /> Methods: We use medium resolution (R ~ 4000) spectra taken with the with the MagE spectrograph on the Magellan (Clay) 6.5 m telescope of the Las Campanas Observatories. <BR /> Results: The spectrum is notable for its lack of a low ionization transient heavy element absorption system that is visible in the large majority of novae near maximum light. We suggest that this may be due to the dominance of inner Lagrangian L1 mass transfer and the absence of a circumbinary gas reservoir in this object.

12213378

2011MNRAS.414..116B

1012

1012.0002_arXiv.txt

\label{i} The two pillars on which the presently accepted standard cosmological model is based, as witnessed by its acronym $\Lambda $CDM, are a cosmological constant $\Lambda $ which determines an acceleration of the expansion of the Universe, and a new type of non relativistic massive particles -- which go under the name of Cold Dark Matter (CDM) -- that source the gravitational potential wells in which cosmic structures can form. Far from being the ultimate description of our Universe, the $\Lambda $CDM model represents an efficient way to parametrize our ignorance about the fundamental constituents of roughly 95\% of its energy content. Although the combination of a cosmological constant $\Lambda $ and of a significant fraction of CDM particles provides a very good fit to most of the presently available observations \citep[see \eg][]{wmap5,wmap7,Percival_etal_2001,Cole_etal_2005,Reid_etal_2010,Riess_etal_1998,Perlmutter_etal_1999,SNLS,Kowalski_etal_2008,Percival_etal_2009}, the fundamental nature of the two main components of the Universe remains an open question. On one side, in fact, the observed value of the energy scale associated with the cosmological constant differs by tens of orders of magnitude from its theoretical predictions, thereby making of the cosmological constant an extremely fine-tuned parameter of the model. On the other side, all the proposed CDM candidates have so far evaded any attempt of direct detection and their properties can still be inferred only by cosmological and astrophysical observations. It is in this context that alternative models have been proposed, by taking the standard $\Lambda $CDM as an asymptotic state of more complex underlying scenarios. These generally involve, as a first step, the promotion of the cosmological constant $\Lambda $ to a dynamical quantity -- dubbed dark energy (DE) -- which evolves during the expansion history of the Universe just as all its other physical constituents. A convenient way to represent this scenario involves the dynamical evolution of a classical scalar field in a self interaction potential, which goes under the name of quintessence \citep{Wetterich_1988,Ratra_Peebles_1988} or k-essence \citep{kessence}. Other possible alternatives with respect to the standard cosmological scenario are given by modifications of the laws of gravity at large scales or at low spatial curvatures \citep[\eg][]{Starobinsky_1980,Hu_Sawicki_2007}, or by relaxing the assumptions of homogeneity and isotropy of the Universe \citep[see \eg][]{GarciaBellido_Haugboelle_2008}. Particular attention has been devoted, in recent years, to the idea that a dynamical DE scalar field might have direct interactions (besides gravity) with other cosmic fluids, by directly exchanging energy-momentum \citep{Wetterich_1995,Amendola_2000,Farrar2004,Amendola_Baldi_Wetterich_2008}. While such direct interactions with standard model particles would be heavily restricted by available observational constraints, the same does not happen for the case of a selective interaction with CDM particles only, as first suggested by \citet{Damour_Gibbons_Gundlach_1990}, for which observational bounds are much weaker. These models have been extensively studied in the recent past, with a particular focus on their possible distinctive effects on structure formation. In fact, all interacting DE models predict the existence of an additional attractive force -- mediated by the DE scalar field -- between massive particles which is expected to alter the growth of density perturbations. Besides this ``fifth-force", coupled massive particles also experience, in interacting DE models, a ``modified inertia" determined by the energy-momentum transfer between the DE and the CDM sectors. The combination of these new physical effects and of the modified cosmic expansion -- that arises as a consequence of the dynamical nature of the DE field -- determines a modified evolution of density perturbations thereby providing possible observational signatures of the models. The effects of interacting DE models on the evolution of cosmic structures have been studied both in the linear regime \citep{Amendola_2000,Amendola_2004,Pettorino_Baccigalupi_2008,DiPorto_Amendola_2008,Baldi_2010} and in the nonlinear regime by relying on simplified configurations, as for the case of spherical collapse \citep{Mainini_Bonometto_2006,Wintergerst_Pettorino_2010}, or in full generality by means of N-body simulations \citep[\eg by][]{Baldi_etal_2010,Baldi_2010,Li_Barrow_2010}. In this work we present an extension of previous analyses about how the different new physical effects featured by interacting DE models individually contribute to the overall modifications of the properties of linear and nonlinear cosmic structures, and we try to clarify some confusion that has been made in the literature concerning the relative importance of the different effects in these two different regimes. The paper is organized as follows: in Sec.~\ref{cde} we briefly review the main equations of interacting DE models both concerning the background evolution and the growth of linear perturbations, with a particular focus on the issue of model normalization which will be central for the rest of the present study; in Sec.~\ref{mog} we discuss the different effects that contribute to a modification of newtonian dynamics at small scales for CDM particles; in Sec.~\ref{sim} we present our set of simulations, describing which effects have been included in each simulation, and how the relevant quantities have been normalized in order to allow a meaningful comparison of the different outcomes; in Sec.~\ref{res} we present our results concerning linear and nonlinear properties of cosmic structures under different implementations of the interacting DE effects; finally, in Sec.~\ref{concl} we draw our conclusions.

\label{concl} Interacting DE models have become very popular in recent years as a viable alternative to the standard $\Lambda $CDM cosmological model. The investigation of these alternative cosmologies has shown significant progresses by moving from the study of the background evolution to the analysis of linear perturbations effects and ultimately to the impact of interacting DE models on nonlinear structure formation by means of specific modifications of N-body algorithms. The main effects through which an interaction between DE and CDM can affect the growth of density perturbations range from a modified background expansion, to a time variation of the CDM particle mass, to a long-range attractive ``fifth force" between CDM particles and a ``modified inertia" in the form of a new velocity-dependent acceleration. In this work we have presented a detailed study of the impact that interacting DE models have on linear and nonlinear structure formation processes. Our work significantly extends previous analyses and aims at a direct comparison between different numerical procedures previously adopted in the literature to investigate the relative importance of each of the above mentioned physical effects that characterize interacting DE cosmologies. By means of a suitable modification of the N-body code {\small GADGET-2}, we have performed a series of collisionless cosmological N-body simulations for a standard $\Lambda $CDM model and for an interacting DE model with coupling $\beta =0.24$. Such a large value for the coupling is already ruled out by several observational probes, however our aim here was not to describe an observationally viable scenario but rather to explore in detail how an interaction between DE and CDM affects structure formation, and a large coupling clearly makes this task more easily achievable by amplifying the effects under investigation. In addition to the $\Lambda $CDM and the interacting DE simulations, we have carried out other 8 simulations in which each of the specific effects of the DE-CDM interaction has been in turn artificially suppressed, in order to quantify its relative contribution to the different peculiar features of interacting DE cosmologies. In doing so, we have applied two different procedures previously adopted in the literature, namely the selective suppression of each individual effect only at the latest stages of structure formation when most of the nonlinear processes take place \citep[as first proposed in][]{Baldi_etal_2010}, and the switch off of the different effects during the whole simulations \citep[as more recently done by][]{Li_Barrow_2010b}. Our analysis therefore allows a direct comparison of these two methods and provides a direct way to test the relative importance of the different features of interacting DE models at high and low redshifts. \ \\ As a first test we have studied the impact of each individual effect of the DE-CDM interaction on the evolution of the matter power spectrum at different redshifts. The global effect of interacting DE on the matter power spectrum is to suppress power at small scales while the large scale amplitude has the same normalization as for $\Lambda $CDM at $z=0$. This first test already shows a significant difference between the two numerical procedures mentioned above. In fact, the suppression of any of the effects of the interaction right from the start of the simulations determines a different linear growth of density perturbations that produces a significant scatter in the large scale normalization of the power spectra at $z=0$. As a consequence, the final power spectrum of each of these simulations is modified at small scales by a superposition of the linear and nonlinear impact of the suppressed effect, making a direct test of the relative importance of each effect quite difficult. On the contrary, if the suppression is limited to low redshifts, the large scale normalization of the power spectrum remains consistent with the original $\Lambda $CDM and interacting DE models, and the small scale differences from simulation to simulation will be due only to the nonlinear impact of each specific effect. With this more suitable comparison procedure we have clearly shown that the velocity-dependent acceleration of CDM particles is the most important effect in suppressing small scale power, with the mass variation playing a minor but still significant role, while the fifth force has basically no effect in the nonlinear regime. \ \\ We have also tested how each of these effects alter the relative evolution of density perturbations in the uncoupled baryonic component as compared to CDM. In this case, the hierarchy of the different effects is reversed with respect to the impact on the power spectrum, and the fifth force shows the most important contribution to the faster growth of CDM perturbations with respect to baryonic perturbations both at linear and nonlinear scales. This different growth gives rise to the so called ``gravitational bias" between the two components which determines a significant baryon depletion of collapsed objects as compared to $\Lambda $CDM. Therefore, the fifth force is found to be the driving mechanism for the reduced baryon fraction of massive halos that characterize interacting DE cosmologies. \ \\ We have then investigated the impact on the halo mass function of each of the effects under study, where once again the two different numerical procedures used in our analysis show significant differences. Also in this case, suppressing individual effects from the beginning of the simulation determines a large scatter in the statistical distribution of collapsed objects at $z=0$ with respect to the original $\Lambda $CDM and full interacting DE simulations that by construction give very similar mass functions at the present time. This scatter is maximum at large masses and shows how, in the simulations run with this procedure, structures cannot be expected to form in the same locations and have similar masses and formation histories from simulation to simulation. This clearly makes any attempt to directly compare individual halos in different simulations particularly hard, since it will be difficult to safely identify objects in different runs as being the same structure, as we have directly shown in our analysis. On the contrary, again, we showed that suppressing individual effects of the DE-CDM interaction only at the latest stages of structure formation produces very similar statistical distributions of halos at the present time, with basically no scatter around the original mass functions of $\Lambda $CDM and of the full interacting DE model. This procedure therefore results clearly more suitable also in order to perform direct comparisons between individual halos in different numerical realizations, as the spatial locations of bound objects as well as their local environment and formation histories will be very similar from simulation to simulation. \ \\ We have finally studied the relative importance of each individual effect of the DE-CDM interaction on the radial density profiles of CDM halos. Consistently with the results found for the halo mass functions, we have shown that a suppression of any specific effect from high redshifts determines a large scatter of the characteristic overdensity of the halos at $z=0$, which have a very different total mass and show significant differences in the amplitude of the density profiles at all radii. This behavior again witnesses the superposition of linear and nonlinear effects when this numerical procedure is adopted. On the other hand, switching off the individual effects only at low redshifts preserves the characteristic overdensity of corresponding halos that show the same amplitude of their density profiles at large radii. As a consequence, the impact of each individual effect on the internal dynamics and on the distortion of the profile at small radii is much more clearly visible and can be safely disentangled from the linear normalization of the environment in which each halo is embedded. By adopting this more suitable procedure, we found once again that the most relevant effect in reducing the inner overdensity of halos in interacting DE models is given by the velocity-dependent acceleration of CDM particles, whereas the mass variation has a minor but not negligible impact, while the fifth force shows no influence whatsoever in this highly nonlinear context. These results are in full agreement with what previously found by \citet{Baldi_etal_2010} and show that the discrepancies claimed by \citet{Li_Barrow_2010b} are essentially due to the different procedure adopted in their numerical study, which determines a superposition of linear and nonlinear effects and is therefore less suitable to compare interacting DE models in the highly nonlinear regime of structure formation. \ \\ To conclude, we have performed a wide and detailed study of interacting DE cosmologies aimed at comparing the relative importance that each of the new physical effects that characterize these models has in affecting linear and nonlinear structure formation processes. Our study significantly extends previous works and allows a direct comparison between different numerical procedures recently adopted in the literature. Our outcomes fully confirm the early results of \citet{Baldi_etal_2010}, and show that the apparent discrepancies found by \citet{Li_Barrow_2010b} are actually due to the different numerical setup adopted by the latter work, which inevitably determines a superposition of linear and nonlinear effects and is therefore not particularly suitable to compare the different characteristic features of interacting DE models in the nonlinear regime.

1012.0002

We present a detailed numerical study of the impact that cosmological models featuring a direct interaction between the dark energy component that drives the accelerated expansion of the Universe and cold dark matter can have on the linear and non-linear stages of structure formation. By means of a series of collisionless N-body simulations, we study the influence that each of the different effects characterizing these cosmological models - which include among others a fifth force, a time variation of particle masses and a velocity-dependent acceleration - separately have on the growth of density perturbations and on a series of observable quantities related to linear and non-linear cosmic structures, as the matter power spectrum, the gravitational bias between baryons and cold dark matter, the halo mass function and the halo density profiles. We perform our analysis applying and comparing different numerical approaches previously adopted in the literature, and we address the partial discrepancies recently claimed in a similar study by Li &amp; Barrow with respect to the first outcomes of Baldi et al., which are found to be related to the specific numerical approach adopted in the former work. Our results fully confirm the conclusions of Baldi et al. and show that when linear and non-linear effects of the interaction between dark energy and cold dark matter are properly disentangled, the velocity-dependent acceleration is the leading effect acting at non-linear scales and in particular is the most important mechanism in lowering the concentration of cold dark matter haloes.

12203639

2011EL.....9319002R

1012

1012.5557_arXiv.txt

1012.5557

In this work we discuss if fermionic sources could be responsible for accelerated periods in a Friedmann-Robertson-Walker spatially flat universe, including a usual self-interaction potential of the Nambu-Jona-Lasinio type together with a fermion-scalar interaction potential of the Yukawa type. The results show that the combination of these potentials could promote an initially accelerated period, going through a middle decelerated era, with a final eternal accelerated period, where the self-interaction contribution dominates.