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JSALT2024 Evaluating LLMs for Research Astronomy 7

12217139

2011PASA...28..128C

1012

1012.0566_arXiv.txt

One problem frequently encountered in astronomical research is that of estimating a confidence interval (CI) on the value of an unknown population proportion based on the observed number of success counts in a given sample. The unknown population proportion may be, for instance, the intrinsic fraction of barred disk galaxies at a specific epoch to be inferred from the observed number of barred disks in a volume-limited sample (e.g.\ \citealt{elm90,van02,cam10,nai10}), with the corresponding binomial CI used to evaluate the hypothesis that the bar fraction changes with redshift relative to a local benchmark (e.g.\ \citealt{cam10}). Experiments to investigate the role of mass and environment in quenching star-formation via measurement of the galaxy red sequence fraction (e.g.\ \citealt{bal06,hes10,ilb10}), or to investigate whether or not major mergers were more frequent at high redshift via measurement of the close-pair/asymmetric fraction (e.g.\ \citealt{dep05,con08,lop10}), also routinely present this class of problem. However, the two most commonly used methods for estimating CIs on binomial population proportions---the `normal approximation' and the \citet{clo34} approach---exhibit significant flaws under routine sampling conditions (cf.\ \citealt{vol93,san98,bro01,bro02}). In particular, the `normal approximation' (also called the `Poisson error') may systematically under-estimate the CI width necessary to provide coverage at the desired level, especially for small samples, but even for rather large samples when the true population proportion is either very low or very high. If used na\"ively the `normal approximation' has the potential to mislead one into over-stating the significance of one's inferences concerning the physical system under study formulated on the basis of the observed data. Astronomers aware of these flaws in the `normal approximation' often adopt the alternative \citet{clo34} approach to CI estimation by way of reference to the CI tables in \citet{geh86}. Unfortunately, the \citet{clo34} approach suffers from the opposite problem to that of the `normal approximation'---namely, a systematic over-estimation of the CI width required to provide the desired coverage \citep{clo34,ney35,geh86,agr98}. In scientific research this over-estimation of the statistical measurement uncertainties may mislead one into placing insufficient confidence in the experimental outcomes, resulting in an inefficient use of the measured data (and, hence, the resources expended in obtaining that data). Indeed, it has been well argued by \citet{agr98} that in many practical applications even the `normal approximation', despite its flaws, is preferable to the \citet{clo34} approach. Fortunately, there exist a multitude of alternative methods for generating CIs on binomial population proportions, many of which exhibit far more satisfactory behaviour than either the `normal approximation' or the \citet{clo34} approach---see \citet{agr98} and \citet{bro01} for various examples. Here I review both the theory and application of one of these methods---use of the beta distribution quantiles---deriving from a simple Bayesian analysis in which a uniform (`non-informative') prior is adopted for the true population proportion (e.g.\ \citealt{gel03}). As I will demonstrate, the beta distribution generator for binomial CIs is both theoretically well-motivated and easily applied in practice using widely available mathematical software packages (e.g.\ \textsc{R}, \textsc{matlab}, \textsc{mathematica}, \textsc{IDL}, \textsc{python}). Ultimately, I advocate strongly that this strategy for estimating binomial CIs be adopted in future studies aiming to constrain fundamental population proportions in astronomical research (e.g.\ the galaxy bar fraction, red sequence fraction, or merger fraction)---especially for samples intrinsically of small-to-intermediate size, or when the subdivision of larger samples for analytical purposes produces sparsely populated data bins.

I have reviewed the performance of three alternative methods for estimating confidence intervals on binomial population proportions; namely, the beta distribution quantile technique, the `normal approximation', and the \citet{clo34} approach (cf.\ \citealt{geh86}). Despite their current popularity in astronomical research, the latter two CI generators are demonstrated to perform poorly under sampling conditions routinely encountered in observational studies---with the `normal approximation' failing to provide CIs of sufficient width to achieve coverage at the nominal confidence level, and the \citet{clo34} approach producing CIs far wider than necessary to achieve the nominal coverage. In contrast, the (Bayesian) beta distribution quantile technique, is revealed to be a well-motivated alternative, consistently providing a mean level of coverage close to the nominal level, even for small-to-intermediate sample sizes. Given that the beta distribution generator for binomial CIs may be easily implemented using modern mathematical software packages, I advocate strongly that this technique be adopted in future studies aiming to constrain the true values of astronomical propulation proportions (e.g.\ the galaxy bar fraction, red sequence fraction, or merger fraction). \appendix

1012.0566

I present a critical review of techniques for estimating confidence intervals on binomial population proportions inferred from success counts in small to intermediate samples. Population proportions arise frequently as quantities of interest in astronomical research; for instance, in studies aiming to constrain the bar fraction, active galactic nucleus fraction, supermassive black hole fraction, merger fraction, or red sequence fraction from counts of galaxies exhibiting distinct morphological features or stellar populations. However, two of the most widely-used techniques for estimating binomial confidence intervals - the `normal approximation' and the Clopper & Pearson approach - are liable to misrepresent the degree of statistical uncertainty present under sampling conditions routinely encountered in astronomical surveys, leading to an ineffective use of the experimental data (and, worse, an inefficient use of the resources expended in obtaining that data). Hence, I provide here an overview of the fundamentals of binomial statistics with two principal aims: (i) to reveal the ease with which (Bayesian) binomial confidence intervals with more satisfactory behaviour may be estimated from the quantiles of the beta distribution using modern mathematical software packages (e.g. r, matlab, mathematica, idl, python); and (ii) to demonstrate convincingly the major flaws of both the `normal approximation' and the Clopper & Pearson approach for error estimation.

12167527

2011ApJ...728..101T

1012

1012.1176_arXiv.txt

Stars form in dense ($>$ 10$^{5}$ cm$^{-3}$) molecular-cloud cores. The IRAS survey of dense cores has revealed that approximately half of dense cores are associated with bright ($>$ 0.1 $L_{\odot}$) infrared point sources, and the dense cores associated with the IRAS sources have been considered as dense cores with the onset of the star formation \cite{bei86,mye87,lee99}. Recent high-sensitivity mid- and far-infrared observations with the Spitzer Space Telescope ($\sim$ 25 times higher sensitivity than with IRAS), however, have found infrared point sources in dense cores previously thought to be starless. The first such ``starless'' core is L1014, where Young et al. (2004) have discovered a Spitzer source (L1014-IRS) with a faint internal luminosity ($L_{int}$ $\sim$ 0.09 $L_{\odot}$). The internal (star+disk) luminosity, $L_{int}$, is defined as $L_{bol}$ -- $L_{isrf}$, where $L_{isrf}$, typically $\sim$ 0.2 $L_{\odot}$, is the luminosity due to the interstellar radiation field. Molecular-line observations have found plenty of dense molecular gas associated with L1014-IRS, which could accrete onto and further raise the mass of the central object \cite{cr05b}. A compact ($\sim$ 500 AU) molecular outflow has also been found with the SMA in L1014-IRS \cite{bou05}. Subsequent surveys of low-mass star-forming regions with Spitzer have discovered more infrared sources in dense cores previously thought to be starless \cite{bou06,dif06,dun06,dun08,ter09}. Those Spitzer sources with rising or flat SEDs between the longest detected Spitzer IRAC wavelength and MIPS 24 $\micron$ as well as between 24 and 70 $\micron$, and $L_{int}$ below or equal to 0.1 $L_{\odot}$, are dubbed ``Very Low Luminosity Objects'' or VeLLOs \cite{dun08}. So far, $\sim$ 15 VeLLOs have been identified out of 67 ``starless cores'', including L1014-IRS, mentioned above, and hence $\sim$ 80 $\%$ of these ``starless'' cores remain starless down to the sensitivity limit of $L_{int}$ $>$ 4 $\times$ 10$^{-3}$ ($d$/140 pc)$^2$ \cite{dun08}. Because of the characteristics of their SEDs and the very low internal luminosity, those VeLLOs could represent the very early phase of star formation with a lower central stellar mass and/or mass accretion rate than those of typical Class 0 sources. The detection of faint $K_{s}$ emission coincident with L1014-IRS, however, suggests that L1014-IRS may not be a young protostellar source, although it is difficult to tell whether part of the $K_{s}$ emission directly arises from the central star or entirely from the reflection nebula tracing the envelope cavity presumably evacuated by an outflow \cite{hua06}. L1521F is one of ``starless'' cores in Taurus, where the Spitzer survey has revealed the presence of a VeLLO, L1521F-IRS ($L_{int}$ $\sim$ 0.06 $L_{\odot}$) \cite{bou06,ter09}. L1521F-IRS was directly detected at mid-infrared wavelengths (24 and 70 $\mu$m), while at shorter infrared wavelengths ($<$ 5 $\mu$m) only the scattered light was detected. The scattered light shows a bipolar nebulosity oriented along the east to west direction, which probably traces an outflow cavity. The 1.3 mm and 850 $\mu$m continuum maps in L1521F show a centrally-peaked dusty condensation with a mass of 1.5 $M_{\odot}$ around L1521F-IRS \cite{cra04,shi04}. The position of the Spitzer source is approximately (within $\sim$ 5$\arcsec$) consistent with the peak position of the dust-continuum emission \cite{bou06}. In the central 20$\arcsec$ region the molecular hydrogen density ($n_{\rm H_{2}}$) is estimated to be $\sim$ 10$^{6}$ cm$^{-3}$, and CO molecules are depleted by a factor of 15 from the nominal CO abundance of 9.5 $\times$ 10$^{-5}$ \cite{cra04}. BIMA observations of L1521F in the CCS ($J_N$ = 3$_{2}$--2$_{1}$) and N$_{2}$H$^{+}$ (1--0) lines have shown that the N$_{2}$H$^{+}$ emission correlates well with the dust-continuum emission while the CCS emission traces the outer region surrounding the N$_{2}$H$^{+}$ emission \cite{shi04}. Onishi et al. (1999) have studied L1521F (MC 27 in their paper) in millimeter and submillimeter HCO$^{+}$ ($J$ = 4--3, 3--2) and H$^{13}$CO$^{+}$ ($J$ = 4--3, 3--2, 1--0) lines and have found a highly condensed gas structure and a signature of infalling motion with an infalling velocity of 0.2 - 0.3 km s$^{-1}$ at 2000 - 3000 AU. They have also found that only the central position of the HCO$^{+}$ ($J$=3--2) line shows winglike components (5.0 - 5.8 km s$^{-1}$ and 7.4 - 8.6 km s$^{-1}$), suggesting the compactness of the associated molecular outflow. From these results, they suggested, before the discovery of the VeLLO, that L1521F is on the verge of star formation. IRAM 04191+1522 (hereafter IRAM 04191) has previously been identified as a Class 0 protostar with its very low T$_{bol}$ of $\sim$ 18 K, located in the southern part of the Taurus molecular cloud \cite{an99a,an99b}. Spitzer observations have revealed that IRAM 04191 is another VeLLO with $L_{int}$ $\sim$ 0.08 $L_{\odot}$ \cite{dun06}. IRAM 04191 was directly detected as an infrared point source at all six wavelengths observed by Spitzer (3.6, 4.5, 5.8, 8.0, 24, and 70 $\mu$m), associated with a nebulosity at wavelengths shorter than 24 $\mu$m extending to the south. The Spitzer source position agrees within 0.2$\arcsec$ with the source position given by Belloche et al. (2002) based on the 227 GHz continuum emission detected by the IRAM PdBI interferometer. IRAM 04191 is associated with a well-developed ($>$ 10000 AU) molecular outflow extending from northeast (red) to southwest (blue), which appears to be similar to outflows from more luminous protostars \cite{an99a,lee02,dun06}. The southern infrared nebulosity appears to be consistent with the blueshifted molecular outflow \cite{dun06}. Millimeter molecular-line observations of IRAM 04191 show a 10000-AU scale rotating molecular envelope around IRAM 04191 \cite{bel02,tak03,lee05}. Comparison between the 1.3 mm dust-continuum and N$_{2}$H$^{+}$ (1--0) observations of IRAM 04191 shows that N$_{2}$H$^{+}$ depletion by a factor of $\sim$ 4 is present at the center of the envelope ($r$ $<$ 1600 AU, $n_{\rm H_{2}}$ $>$ 5 $\times$ 10$^{5}$ cm$^{-3}$), whereas at the outer part the N$_{2}$H$^{+}$ emission correlates with the dust-continuum emission \cite{bel04}. It is still not clear, however, how to track the evolutionary stage and the differences among these VeLLOs. The chemical status of dense gas around the VeLLOs could be a key, since chemistry in dense cores is non-equilibrium and time-dependent (e.g. Leung et al. 1984; Herbst \& Leung 1989). Carbon-chain molecules are of particular interest. Since carbon-chain molecules are produced through hydrogenation of C$^{+}$ and C, and carbon in the form of C$^{+}$ and C is only present at the early ($\leq$ 10$^{5}$ yr) evolutionary stage of the chemical evolution, they are likely to trace the early evolutionary stage of star formation. Observations of organic molecules such as CH$_{3}$OH and CH$_{3}$CN towards VeLLOs are also intriguing. At the ``real'' starless phase the low temperatures at the central regions result in the depletion of most molecules onto dust grains. Once a protostellar source forms at the center of the dense core, heating from the protostar desorbs molecules from grain surfaces, and the desorption region extends as the protostellar evolution proceeds \cite{sch04,aik08}. Since organic molecules are thought to be formed on grain surfaces \cite{aik08} (for example CH$_{3}$OH is formed through hydrogenation of CO on grain surfaces), the amount of organic molecules in the gas phase towards VeLLOs could be an independent measure of the evolutionary stage. In order to study the chemical status of the molecular material surrounding VeLLOs and to investigate the evolution of the earliest protostellar formation, we have conducted observations of L1521F and IRAM 04191 in several carbon-chain and organic molecular lines with the Nobeyama 45-m telescope. Our observations of the two representative VeLLOs in carbon-chain and organic molecular lines should shed light on the chemical evolutional status of dense gas surrounding these VeLLOs and the initial stage of star formation.

\subsection{Chemical Evolutionary Difference between L1521F and IRAM 04191} As shown in the previous section, our observations of the carbon-chain molecules show clear chemical differentiation between L1521F and IRAM 04191. The chemical differentiation could reflect the evolutionary difference between the embedded VeLLOs of L1521F-IRS and IRAM 04191. Here, we will discuss the link between the chemical differences and the evolutionary difference in these VeLLOs. Chemistry in dense molecular-cloud cores is known to be non-equilibrium and time-dependent, and molecular abundances in dense cores change as a function of the evolutionary time \cite{leu84,her89,ber97}. In the gas phase, the time dependences of molecular abundances are mainly controlled by the status of carbon. At an early evolutionary stage ($\leq$ 10$^{5}$ yr) carbon is mainly in the form of C$^{+}$ and C, and at late evolutionary stages ($>$ 10$^{5}$ yr) almost all carbon is in the form of CO. Carbon-chain molecules are produced through hydrogenation of C$^{+}$ and C, and hence are more abundant in the earlier evolutionary stage. On the other hand, HCO$^{+}$ is formed via CO + H$_{3}^{+}$, and HC$^{18}$O$^{+}$ is more abundant at later evolutionary phases. Thus, the higher abundances of the carbon-chain molecules in L1521F than in IRAM 04191, as well as the sightly higher HC$^{18}$O$^{+}$ abundance in IRAM 04191 than in L1521F, are likely to indicate that IRAM 04191 is chemically more evolved than L1521F. This chemical evolutionary difference is also consistent with the different outflow activity between L1521F-IRS and IRAM 04191; there is a well-developed CO molecular outflow in IRAM 04191, as seen in Figure \ref{i04mom0}, while no clear CO molecular outflow is observed in L1521F-IRS, suggesting that the star-forming activity is more advanced in IRAM 04191 than L1521F-IRS. Aikawa et al. (2001, 2003, 2005) theoretically studied the evolution of molecular distributions in collapsing dense cores. Their model shows that the radial distribution of molecules varies as a function of the evolutionary time, and that the distribution of carbon-chain molecules such as CCS first shows a central hole after $\sim$ 10$^{6}$ yr (with $\alpha$ = 1.1, where $\alpha$ denotes the internal gravity-to-pressure ratio in the collapsing dense core.). This is because at the center the gas density is higher than that at the outer part, and hence the chemical evolution proceeds faster than that at the outer part. On the other hand ``late-type'' molecules such as CO, HCO$^{+}$, and N$_{2}$H$^{+}$ show centrally-peaked column density profiles in the early evolutionary stage, and then CO and HCO$^{+}$ start to show central holes due to the depletion onto grain surfaces ($>$ 10$^{6}$ yr). Eventually, N$_{2}$H$^{+}$ also becomes depleted \cite{ber02,bel04}. In L1521F, the CH$_{3}$CCH, C$_{4}$H, and the CCS emissions do not show a centrally-peaked spatial distribution but the N$_{2}$H$^{+}$ emission shows a centrally-peaked distribution (see Figure \ref{hiroko}). On the other hand, in IRAM 04191 the CH$_{3}$CCH and C$_{4}$H emissions are as weak as our detection limit, and the depletion of N$_{2}$H$^{+}$ in the central $\sim$ 1600 AU has been reported \cite{bel04}. Comparison between the theoretical model and these observational results implies that in L1521F carbon-chain molecules are still abundant but at the center, they become less abundant due to the faster gas-phase chemical evolution, while in IRAM 04191 these carbon-chain molecules become undetectable and the depletion of N$_{2}$H$^{+}$ proceeds at the center. \subsection{Comparison with Other Dense Cores} In $\S$4.1., we have discussed the possibility that L1521F-IRS is in an earlier evolutionary stage than IRAM 04191. It is interesting to then compare the evolutionary stage of the surrounding dense gas around L1521F-IRS and IRAM 04191 to that of starless dense cores in Taurus. In Table 4, we compiled several evolutionary indicators of three starless cores (L1521B, L1498, and L1544) and dense cores around the two VeLLOs (L1521F and IRAM 04191). The evolutionary indicators include the ratio between the N$_{2}$H$^{+}$ and CCS column densities, the deuterium fractionation measured from the N$_{2}$D$^{+}$ / N$_{2}$H$^{+}$ column density ratio, the CO depletion factor defined by Crapsi et al. (2004), the central molecular-gas density, and the CH$_{3}$CCH and C$_{4}$H column densities. As already discussed in $\S$4.1., the ratio between the N$_{2}$H$^{+}$ and CCS column densities is an excellent indicator of the gas-phase chemical evolution in dense cores, and the ratio increases as the dense core evolves. The deuterium fractionation and the CO depletion factor are measures of chemical evolution of cold dense gas. Under the cold gas ($\leq$ 10 K) condition the deuterium fractionation proceeds due to preferable condition for the reaction of H$_{3}^{+}$ + HD $\rightarrow$ H$_{2}$D$^{+}$ + H$_{2}$, and CO molecules keep depleting onto grain surfaces \cite{cra04,cr05a}. The central gas density is a direct measure of the physical evolutionary stage of dense cores. Table 4 shows that the collected dense cores can be sorted following these evolutional criteria. In L1521B and L1498 the N$_{2}$H$^{+}$ to CCS ratio, the deuterium fractionation, the CO depletion factor, and the central gas density are lower than those in L1521F and IRAM 04191. The CO depletion factor and the central gas density in L1521F are similar to those in L1544, another starless core in Taurus \cite{oha99,cr05a}, but the deuterium fractionation is twice as high as in L1544 than in L1521F. These results imply that L1544 is in a similar evolutionary stage as L1521F. In fact, towards L1544 an infalling gas motion has been observed \cite{wil99,oha99}, suggesting the core is in the process of protostellar formation, though inspection of the Spitzer data in L1544 does not reveal any protostellar candidate \cite{bou06}. IRAM 04191 shows a higher N$_{2}$H$^{+}$ to CCS ratio, and CO depletion factor, and lower carbon-chain abundances than L1521F, and IRAM 04191 is probably more evolved than L1521F as already discussed. \placetable{tbl-4} \subsection{Molecular Desorption} Once a protostar forms at the center of the dense core, heating from the protostar desorbs molecules from grain surfaces and alters the chemical conditions. Aikawa et al. (2008) showed that the warm molecular region heated by the protostar expands during protostellar evolution, and that molecules such as CO are desorbed from grain surfaces and show abundance enhancement within the radius where the gas temperature exceeds their sublimation temperatures. Organic molecules such as CH$_{3}$OH and CH$_{3}$CN formed on grain surfaces are also desorbed and enhanced in the gas phase by orders of magnitude within the radius where the gas temperature is above $\sim$ 100 K ($\equiv$ $R_{100}$). The model by Aikawa et al. (2008) predicts that $R_{100}$ reaches 100 (AU) after $\sim$ 10$^{5}$ years from protostellar formation, and that the mean CH$_{3}$OH abundance ($\equiv$ $X_{\rm CH_3OH}$) and the mean gas kinetic temperature ($\equiv$ $T_{\rm K}$) averaged over the 45-m beam ($\sim$ 18$\arcsec$ $\sim$ 2500 AU) are $\sim$ 8.0 $\times$ 10$^{-9}$ and $\sim$ 52 K, respectively. On the other hand, N$_{2}$H$^{+}$ and 1.3 mm dust-continuum observations of L1521F \cite{cra04} and IRAM 04191 \cite{bel04} show that the gas density towards the central $\sim$ 20$\arcsec$ region ($\equiv$ $n_{\rm H_2}$) is $\sim$ 10$^{6}$ cm$^{-3}$ and that the velocity width over the core size ($\equiv$ $dv/dr$) is $\sim$ 10 km s$^{-1}$ pc$^{-1}$. With input values of $n_{\rm H_2}$ = 10$^{6}$ cm$^{-3}$, $T_{\rm K}$ = 52 K, and $X_{\rm CH_3OH} / dv/dr$ = 8.0 $\times$ 10$^{-10}$ km$^{-1}$ s pc, our LVG calculation of CH$_{3}$OH \cite{tak98,tak00} predicts the expected CH$_{3}$OH (6$_{-2}$--7$_{-1}$ E) line intensity to be $\sim$ 0.63 K, which is significantly above the 3$\sigma$ upper limit ($\sim$ 0.2 K) of our observations. This LVG calculation suggests that $R_{100}$ in L1521F and IRAM 04191 has not yet reached 100 AU, and that the evolutionary stage of L1521F-IRS and IRAM 04191 is younger than 10$^{5}$ years from protostellar formation. Similarly, Aikawa et al. (2008) showed that at $R_{100}$ = 100 (AU) the beam-averaged CH$_{3}$CN abundance reaches on the order of 10$^{-10}$, while the observed upper limit of the CH$_{3}$CN abundance is $\sim$ 4 $\times$ 10$^{-13}$. We consider that the non-detection of these organic molecular lines is a sign of the youthfulness of these VeLLOs. Although there is no other observation in the 6$_{-2}$--7$_{-1}$ E transition of CH$_{3}$OH towards low-mass protostars reported, interferometric observations of Class 0 protostars have found compact ($\leq$ 500 AU) millimeter and submillimeter CH$_{3}$OH emission associated with the central protostars in L1157 \cite{gol99,vel02}, NGC1333 IRAS 2A \cite{jo05b}, and in IRAS 16293-2422 \cite{kua04,cha05}. Similar compact components are also seen in the CH$_{3}$CN (6--5) \cite{bot04} and CH$_{3}$CN (12--11) emission \cite{bis08} towards IRAS 16293-2422. The CH$_{3}$OH abundances in those compact components were estimated to be $\sim$ 10$^{-8}$ towards L1157 \cite{gol99,vel02}, $\sim$ 3 $\times$ 10$^{-8}$ towards NGC1333 IRAS 2A \cite{jo05b}, and $\sim$ 9.4 $\times$ 10$^{-8}$ towards IRAS 16293-2422 \cite{cha05}, which are one order of magnitude higher than the CH$_{3}$OH abundance in cold dark clouds ($\sim$ 2 $\times$ 10$^{-9}$) \cite{fri88,tak98,tak00}. Maret et al. (2005) have conducted a CH$_{3}$OH (5$_{K}$--4$_{K}$; 7$_{K}$--6$_{K}$) survey of seven Class 0 sources with the IRAM 30 m and JCMT telescopes. From their radiative transfer and ``jump'' models, they have found that four sources (IRAS 16293-2422, NGC1333 IRAS 2, IRAS4B, and L1448-MM) show CH$_{3}$OH abundance enhancements up to 1-7 $\times$ 10$^{-7}$ at the innermost part of the envelopes. Similar single-dish studies of the CH$_{3}$CN lines have revealed more than 2 orders of magnitude of abundance enhancements of CH$_{3}$CN in NGC 1333 IRAS 2 (7 $\times$ 10$^{-9}$) \cite{jo05a} and IRAS 16293-2422 (10$^{-8}$) \cite{caz03}. These results show that towards several Class 0 protostars the CH$_{3}$OH and CH$_{3}$CN abundance enhancements take place. On the other hand, non-detection of the CH$_{3}$OH (6$_{-2}$--7$_{-1}$ E) and CH$_{3}$CN (6$_{K}$--5$_{K}$) lines towards L1521F-IRS and IRAM 04191 does not imply the absence of the desorption region. If $R_{100}$ = 10 AU for example, $R_{20}$, the radius of the CO desorption, reaches 250 AU \cite{mas98,mas00}. Our recent observations of the C$^{18}$O (2--1) line towards L1521F-IRS and IRAM 04191 with the SMA have revealed compact ($\sim$ 500 AU) CO emission associated with these VeLLOs, suggesting the possible desorption of CO. We suggest that the VeLLOs of L1521F-IRS and IRAM 04191 are in the ongoing stage of the desorption processes, and that the abundance enhancement of organic molecules of CH$_{3}$OH and CH$_{3}$CN can be used as a chemical evolutionary indicator from the VeLLO to Class 0 stage. \subsection{Different Origin of L1521F-IRS and IRAM 04191 ?} In the above discussion we implicitly assume that both L1521F-IRS and IRAM 04191 are protostars in a common evolutionary sequence and will form a star with a similar final mass. On the other hand, our recent SMA observations of L1521F-IRS have revealed compact ($\sim$ 1500 AU) molecular outflows \citep{tak10}, and the outflow properties are similar to those around very low-mass stars or brown dwarfs \citep{nat04,pha08}. It is therefore possible that L1521F-IRS is a very low-mass, more evolved star than protostars. In this final subsection, we will discuss this alternative possibility. The estimated outflow mass around L1521F-IRS is $\sim$ 3.6 $\times$ 10$^{-5}$ $M_{\odot}$, which is orders of magnitude lower than that of IRAM 04191 ($\sim$ 3 $\times$ 10$^{-2}$ $M_{\odot}$; Lee et al. 2002). There are several possible explanations for the difference in the outflow masses. One is that both L1521F-IRS and IRAM 04191 are protostars and the mass ejection event of L1521F-IRS has started more recently than IRAM 04191, and hence the amount of the molecular material entrained by L1521F-IRS is lower than that by IRAM 04191. This interpretation is consistent with the arguments in $\S$4.1 - $\S$4.3. The second interpretation is that the mass accretion rate towards L1521F-IRS is intrinsically smaller than that towards IRAM 04191, presumably due to the different effective sound speed between the L1521F and IRAM 04191 region. The lower mass accretion rate towards L1521F-IRS could yield the lower total outflow mass and the lower internal luminosity as compared to those of IRAM 04191. The other possible interpretation is that L1521F-IRS is a very low-mass, more evolved star close to the end of the mass accretion and ejection phase, and that the amount of the total mass of the molecular outflow is proportional to the amount of the total accreted material, and hence the central stellar mass. In this case, there may be a substantial difference in the central stellar mass between L1521F-IRS and IRAM 04191, and the most of the luminosity of L1521F-IRS may originate from nuclear burning whereas the luminosity of IRAM 04191 originate mostly from mass accretion \cite{nat04}. In this case, how can we explain the observed chemical difference of the surrounding dense gas between L1521F and IRAM 04191 ? As discussed in $\S$3.1, L1521F is located in the main ridge of the Taurus Molecular Cloud complex \cite{dam01,gol08}, and there is ample molecular material surrounding the L1521F region. On the other hand, IRAM 04191 is not located in the main Taurus Complex, but at the south-western edge of the Taurus-Auriga Complex away from the Galactic Plane ($b$ $\sim$ -23.5$\degr$) \cite{dam01}. The surrounding ambient gas can be a source of the ``chemically-fresh'' molecular gas and maintain the higher carbon-chain abundances around L1521F, whereas around IRAM 04191 there is less such molecular gas. The observed chemical difference between L1521F and IRAM 04191 may be due to the different surrounding environment.

1012.1176

We have observed dense gas around the Very Low Luminosity Objects (VeLLOs) L1521F-IRS and IRAM 04191+1522 in carbon-chain and organic molecular lines with the Nobeyama 45 m telescope. Toward L1521F-IRS, carbon-chain lines of CH<SUB>3</SUB>CCH (5<SUB>0</SUB>-4<SUB>0</SUB>), C<SUB>4</SUB>H ({17/2}-{15/2}), and C<SUB>3</SUB>H<SUB>2</SUB> (2<SUB>12</SUB>-1<SUB>01</SUB>) are 1.5-3.5 times stronger than those toward IRAM 04191+1522, and the abundances of the carbon-chain molecules toward L1521F-IRS are 2-5 times higher than those toward IRAM 04191+1522. Mapping observations of these carbon-chain molecular lines show that in L1521F the peak positions of these carbon-chain molecular lines are different from each other and there is no emission peak toward the VeLLO position, while in IRAM 04191+1522 these carbon-chain lines are as weak as the detection limits, except for the C<SUB>3</SUB>H<SUB>2</SUB> line. The observed chemical differentiation between L1521F and IRAM 04191+1522 suggests that the evolutionary stage of L1521F-IRS is younger than that of IRAM 04191+1522, consistent with the extent of the associated outflows seen in the <SUP>13</SUP>CO (1-0) line. The non-detection of the organic molecular lines of CH<SUB>3</SUB>OH (6<SUB>-2</SUB>-7<SUB>-1</SUB> E) and CH<SUB>3</SUB>CN (6<SUB>0</SUB>-5<SUB>0</SUB>) implies that the warm (~100 K) molecular-desorbing region heated by the central protostar is smaller than ~100 AU toward L1521F-IRS and IRAM 04191+1522, suggesting the young age of these VeLLOs. We propose that the chemical status of surrounding dense gas can be used to trace the evolutionary stages of VeLLOs.

12213700

2011MNRAS.415..944B

1012

1012.4565_arXiv.txt

\label{intro} Usually, neutron stars (NS) have magnetic field $B \sim 10^{12}$G and rotate with a period of fraction of a second, the Crab pulsar being a typical example. However, we now know that the ``zoo'' of NS is much more diverse and they can have both much weaker and much stronger magnetic field and rotate with both much longer and much shorter periods. Around 10\% of NS have surface dipolar magnetic field $B\sim 10^{14}-10^{15}$G \citep{kd98}. These ``magnetars'' are believed to be born in core collapse explosions of rapidly rotating stars. During the collapse, the proto-neutron star naturally develops strong differential rotation, which allows for generation of magnetic field via $\alpha$-$\Omega$-dynamo \citep{DT92,TD93}. The strength of saturated magnetic field strongly depends on the rotational period, with shorter periods leading to stronger magnetic field and more rapid release of rotational energy. In order to generate the magnetar strength magnetic field the rotation period must be around few milliseconds \citep{DT92}. Such a strong field allows to release up to $10^{52}\,$erg of magnetar's rotational energy in very short period of time. This is sufficient to drive extremely powerful supernova explosions, on the hypernova scale, and to produce powerful Gamma Ray Bursts \citep[GRB, e.g.][]{U92,tcq04}. Turbulence required for the magnetic dynamo action can also be generated via the magneto-rotational instability \citep[MRI,][]{BH91}. Calculations based on the linear theory show that in the supernova context strong saturation field can be reached very quickly, on the time scale of only several tens of rotational periods \citep{AWML03,ocma09}. Millisecond pulsars are found in low mass binaries, and it is generally thought that they have been spun up via disk accretion \citep{acrs82,ar09}. This origin implies mass increase by about 0.2$M_\odot$ compared to normal radio pulsars, whose masses are narrowly distributed around $1.35M_\odot$ \citep{ts99}. It is rather difficult to measure the mass of a millisecond pulsar, but the few available results agree with this prediction of the accretion model \citep{ktr94,jh05,dp10}. The most massive pulsar found to date, almost 2$M_\odot$, is a millisecond pulsar \citep{dp10}. The magnetic field of these millisecond pulsars is very low, down to $10^9$~Gauss. Most likely, their initial magnetic field was of similar strength to normal pulsars, but now it is buried under the layers of accreted matter \citep{bkk74,acrs82}. The reason why these pulsars could not generate magnetar-strength magnetic field in the same way as in the core-collapse scenario is the very long time scale of spinning up compared to the viscous time-scale. As the result, the differential rotation remains weak and there is not enough energy for effective magnetic dynamo \citep{s99}. The rotational frequency of the fastest X-ray pulsar is $\sim760$~Hz \citep{cmm03}, whereas the fastest known radio pulsar has the frequency of $641$~Hz \citep{bkh82}. Most likely, it is the gravitational radiation losses what places the upper limit on the rotation rates because at high spin the r-mode oscillations become excited \citep{ST83,l99}. \citet{s99} argued that this instability may also result in magnetic explosion. The idea is that the heating of NS, associated with these oscillations, reduces its viscosity and decouples its interior from the outer layers. Being most disturbed, the outer layers rapidly loose some of their angular momentum via gravitational radiation. This leads to strong differential rotation and generation of magnetar strength magnetic field in the NS interior. This field becomes unstable to buoyancy, emerges on the surface, and magnetically driven pulsar wind rapidly extracts the rotational energy of the NS. \citet{s99} proposed this as an alternative scenario for long GRBs. It is unlikely that a supernova-like event can accompany a GRB in this scenario. Although the wind energetics is sufficient, only a small fraction of this energy can be deposited into the companion star, simply because of its small geometrical cross-section. Moreover, recent results suggest that the amplitude of r-modes may saturate at a much lower level due to nonlinear interaction with other modes \citep{afm03,btw05,btw07}. A similar recycling of NS may occur during the common envelop (CE) phase, after the primary becomes a red super giant \citep[RSG; ][]{bks71,P76,ty79,PY06}. Due to the dynamic friction, the NS then spirals inside the RSG, accreting on its way. Now one can imagine two interesting outcomes of such process. First, the neutron star may accumulate too much mass and collapse into a black hole (BH). This BH is likely to be rapidly rotating and drive a stellar explosion in the collapsar fashion \citep{fw98,zf01,BK10}. Second, the NS may first spin up to a millisecond period and drive a magnetic explosion of the type proposed by \citet{s99} but now inside the common envelope. The magnetar wind can keep energising such supernovae, producing a similar effect to radioactive decay \citep{w10,kb10}. High accretion rates may modify the way the NS is recycled. The accreted gas can form a massive rapidly rotating layer above the NS crust \citep{is99,is10}. The strong differential rotation between the layer and the NS core may result in development of the Kelvin-Helmholtz instability when the NS crust melts down under the weight of the layer. This may lead to turbulence and strong amplification of the NS magnetic field. The accretion onto neutron stars during the in-spiral has been studied by \citet{ch96}, who concluded that the high angular momentum of the gas gravitationally captured by NS prevents it from effective neutrino cooling and keeps the mass accreting rate well below the rate of the Bondi-Hoyle capture. As the result, the NS accumulates very little mass while still inside the common envelope. On the other hand, he suggested that during the merger with the companion core the mass accretion rate rises and the NS collapses into a black hole. While carefully analysing various effects of rotation, \citet{ch96} did not use any particular model for the primary. Moreover, there is a great deal of uncertainty with the regard to the specific angular momentum of the gravitationally captured gas. In our study we come back to this problem, consider realistic models of RSG stars and allow for the uncertainty. The paper is organised as follows. In Section 2 we consider the accretion and recycling of NSs during the in-spiral. We conclude that the outcome is very sensitive to the assumed specific angular momentum of the gravitationally captured gas. Given the current uncertainty with the regard to the angular momentum, it seems possible that NSs begin to accrete with Bondi-Hoyle rate while still inside common envelopes. In this case, they rapidly spin up to millisecond periods. In Section 3 we speculate on the possible mechanisms of generating magnetar-strength magnetic field and study the magnetic interaction of millisecond magnetars with accretion flows typical for the in-spiral problem. We find that if the magnetar forms inside the common envelope and its magnetic field is about $10^{15}\,$G, the magnetospheric pressure can overcome the gravity and drive an outflow, with eventual release of up to $10^{52}$erg of the magnetar rotational energy. If the magnetar forms only after the merger with the core then higher magnetic field, $\sim10^{16}$G, is required to drive stellar explosions. In Section 4 we discuss the properties of such explosions and their observational signatures. Because the RSG envelope is rich in hydrogen the supernova will be classified as type-II. Because of the very high energy release and relatively small mass of the RSG, the speed of the ejecta is expected to be very high, $\sim 10^9$cm/s, and the luminosity at the plateau phase $\sim 10^{43}$erg/s. However, due to the small amount of $^{56}$Ni generated in the explosion and the rapid spin-down of the magnetar, the plateau is followed by a sharp drop in brightness and a steep light-curve tail. We show that the termination shock of the magnetar wind produces a high energy synchrotron and inverse Compton emission which may energise the tail, but the supernova ejecta soon becomes transparent to this emission, limiting its potential to mimic the effect of radioactive decay. The flux of gamma-ray emission is expected to be rather low and difficult to observe, unless the explosion occurs in the Local Group of galaxies. In the case of off-center explosion, the remnant is a very close binary system consisting of a WR star and a magnetar, but the strong magnetic field prevents magnetar from accreting plasma of the WR wind. Our results are summarized in Section 5. In this paper, the dimensional estimates are presented using the following notation: the time $t\sub{x,n}$ is measured in $10^n$s, the distance $R\sub{x,n}$ in $10^n\,$cm, the speed $V\sub{x,n}$ in $10^n\,\mbox{cm}/\mbox{s}$, the mass $M\sub{x,n}$ in $10^nM_\odot$, mass accretion rate $\dot{M}\sub{x,n}$ in $10^n M_\odot/\mbox{yr}$, and the magnetic field $B\sub{x,n}$ in $10^n$G.

\label{conclusions} The main aim of this study was to investigate whether during the common envelope phase of a close binary, involving RSG and NS, the NS can spin up to a millisecond period and generate magnetar-strength magnetic field. If possible, this would result in magnetically driven stellar explosion, releasing up to $10^{52}\,$erg of magnetar's rotational energy inside the RSG. It turns out that the outcome is very sensitive to the specific angular momentum of the gas gravitationally captured by the NS during the in-spiral. Should it be as high as suggested in the early papers by \citet{is75} and \citet{sl76}, the accretion rate is low and neither mass nor spin of the NS will increase significantly until the NS settles into the center of RSG. This is because the accretion disk forms too far from the NS, its temperature remains rather low, and it is unable to cool effectively via neutrino emission, which is very sensitive to temperature. As the result, the accretion shock is pushed too far, beyond the sonic point of the Bondi-type flow, thus preventing the accretion from reaching the Bondi-Hoyle rate \citep{ch96}. The CE can either be ejected, leaving behind a close NS-WR binary or survive, leading to a merger of the NS with the RSG core \citep{Taams00}. It is not clear what exactly would occur following the merger, as we cannot treat this case using our approach. The common belief is that the NS will begin to accrete at the Bondi rate and collapse into a black hole. However, this conclusion is based on models with spherical symmetry, whereas the strong rotation developed by the system during the merger makes the assumption of spherical symmetry unsuitable. It may still be possible that the NS spins up to a millisecond period and becomes a magnetar before its mass reaches the limit of gravitational collapse. If, on the contrary, the angular momentum is much lower, more in line with the analysis of \citet{dp80}, the accretion can begin to proceed at the Bondi-Hoyle rate when the NS is still inside the common envelope. Further investigation of the Bondi-Hoyle accretion via 3D numerical simulations are required to clarify this issue. Once the NS becomes a millisecond magnetar, its emerging magnetosphere interacts with the accretion flow. If this occurs while the NS is still inside the common envelope and the magnetic field is as strong as $\sim10^{15}$G, the interaction is most likely to begin in the propeller regime and then quickly proceed to the ejector regime, given the small radius of the light cylinder. Strong magnetar wind blows away the common envelope and drives an explosion whose energy is likely to exceed the standard $10^{51}$erg of normal supernovae. However, the core of the companion survives and the explosion leaves behind a very close binary, consisting of a magnetar and a WR star. In spite of the very close separation and the very strong wind from the WR star, the NS is shielded from the wind by its magnetosphere and prevented from accreting. Later on, when the WR explodes this will be a {\it third} explosion produced in the system. Moreover, the rapid rotation and the compactness of this WR star show that this explosion can be accompanied by a gamma-ray burst \citep{BK10}. If the magnetar forms only after the merger with the core then a higher magnetic field, $\sim10^{16}$G, is required to explode the star. In this case, the remnant is a solitary magnetar. If, however, the magnetar-driven explosion fails and the NS collapses into a black hole, this will be a rapidly rotating black with a massive accretion disk. Although this is exactly the configuration proposed in the collapsar model of gamma-ray bursts, the relativistic jets will not be able to escape from the extended envelope of RSG and produce such a burst. Instead, they will deposit their energy inside the envelope and drive a type-II supernova. Given the very high rotational energy of a millisecond NS the supernova is expected to be very bright at the plateau phase and show very broad spectral lines. However, only a small amount of $^{56}$Ni is expected to be produced in the explosion due to the relatively low densities, in comparison to those reached during the normal core-collapse explosions. By the end of the plateau phase the power of the magnetar wind is also significantly reduced and the supernova brightness is expected to drop sharply. The combination of the rapidly declining power of magnetar wind and the increasing transparency of the supernova ejecta to the high energy emission from the wind termination shock result in steeper than normal light-curve tails. The unique property of supernovae produced in this way is the high energy synchrotron and inverse Compton emission from the magnetar wind nebula. Unfortunately, the supernova ejecta does not become transparent to this emission until a hundred days after the explosion, by which time the NS rotation is already very slow and its wind is no longer that powerful. The flux of gamma-ray emission is expected to be rather low and difficult to observable, unless the explosion occurs in the Local Group of galaxies.

1012.4565

In this paper, we propose a new plausible mechanism of supernova explosions specific to close binary systems. The starting point is the common envelope phase in the evolution of a binary consisting of a red supergiant and a neutron star. As the neutron star spirals towards the centre of its companion it spins up via disc accretion. Depending on the specific angular momentum of the gas captured by the neutron star via the Bondi-Hoyle mechanism, it may reach millisecond periods either when it is still inside the common envelope or after it has merged with the companion core. The high accretion rate may result in the strong differential rotation of the neutron star and generation of the magnetar-strength magnetic field. The magnetar wind can blow away the common envelope if its magnetic field is as strong as 10<SUP>15</SUP> G and can destroy the entire companion if it is as strong as 10<SUP>16</SUP> G. The total explosion energy can be comparable to the rotational energy of a millisecond pulsar and reach 10<SUP>52</SUP> erg. However, only a small amount of <SUP>56</SUP>Ni is expected to be produced this way. The result is an unusual Type II supernova with very high luminosity during the plateau phase, followed by a sharp drop in brightness and a steep light-curve tail. The remnant is either a solitary magnetar or a close binary involving a Wolf-Rayet star and a magnetar. When this Wolf-Rayet star explodes, it will be a third supernova explosion in the same binary.

12213713

2011MNRAS.413.1206B

1012

1012.3345_arXiv.txt

Studying AGN at X-ray wavelengths allows an assessment of the intrinsic source properties, such as the measurement of the continuum and the source power, due to the ability of X-rays to penetrate and measure large columns of obscuring material. This makes X-ray observations essential for testing AGN unification schemes \citep{awaki91} which invoke obscuration to explain the perceived difference between Seyfert 1s and Seyfert 2s \citep{antonucci93, urry95}. The \nh\ and $\Gamma$ distributions for AGN are also an important ingredients for synthesis models of the cosmic X-ray background \citep*[XRB,][]{gilli07}. Most of our knowledge of the X-ray properties of local AGN comes from X-ray selected samples, such as the \citet{piccinotti82} sample derived from {\it HEAO 1} observations \citep[e.g.][]{turner89, nandra94}. X-ray selection of AGN is in principle very effective as in this band the nucleus predominantly outshines the host galaxy, and can penetrate obscuration. The hard X-ray telescopes aboard {\it INTEGRAL} and {\it Swift} and their large sky coverage have allowed for the compilation of hard X-ray selected AGN catalogues. \citet{beckmann06} and \citet{beckmann09} have presented catalogues for {\it INTEGRAL} selected AGN, which total 199 after 5 years of observations, and \citet{tueller08} and \citet{tueller10} have presented 229 {\it Swift/BAT} selected Seyfert galaxies after 22 months of observations. However, the discovery from these surveys seems to be that the fraction of Compton thick AGN detected is smaller than expected from XRB studies \citep*{gilli07}. This points to the empirical fact that even the X-ray band is biased against the most obscured, Compton thick, AGN. X-ray selection at energies greater than 10 keV is less affected by Compton thick obscuration than 2-10 keV selection though. An alternative to X-ray AGN selection is mid-infrared (MIR) selection where the primary radiation of the AGN is re-emitted after having been reprocessed by hot dust. The extended {\it IRAS} 12 micron galaxy sample \citep*[12MGS -][RMS]{rush93} is a sample of 893 MIR selected local galaxies which contains a high fraction of AGN (13\% at the time of publication, RMS). The sample is taken from the {\it IRAS Faint Source Catalogue, version 2 (FSC-2)} and imposes a flux limit of 0.22 Jy, including only sources with a rising flux density from 12 to 60 microns (to exclude stars) and with a galactic latitude of $|b| \geq$25 deg. Being selected in the mid-infrared (MIR) this sample is also relatively unbiased against absorption, low luminosity systems and `hidden AGN'. \citet{spinoglio89} showed that a wide variety of AGN types emit a constant fraction of their bolometric luminosities at this wavelength, and furthermore shown to be true for star forming galaxies as well by \citet{spinoglio95}. The 12MGS should be therefore also representative of the true number of different active galaxy types in relation to each other. It is possible, though, that even at 12 \mic, the AGN emission may be suppressed in the most heavily obscured nuclei. \citet{pier92} present radiative transfer modelling for the infrared emission from dust tori for varying torus orientations. The most significant result of this study was that throughout the infrared, the emission in the edge-on direction is weaker with respect to the face-on direction, and at 12 \mic, the difference can be up to several orders of magnitude for optically thick tori. These models suggest that AGN selection at 12 microns is biased against heavily obscured nuclei. However, recently this was tested observationally by \citet{horst07}. They used high resolution 12.3 \mic\ imaging of a sample of local Seyfert nuclei, with X-ray data to explore the ratio of the MIR to intrinsic X-ray nuclear luminosities in order to detect any difference between the obscured nuclei (Seyfert 2s) and unobscured nuclei (Seyfert 1s). As the X-ray luminosities have been corrected for absorption, $L_{\rm X}$ should represent the intrinsic emission from the nucleus. For smooth tori modelled by \citet{pier92}, obscured systems should then present lower $L_{\rm MIR}/L_{\rm X}$ ratios than unobscured systems due to the suppressed MIR flux seen in these systems. A difference of an order of magnitude for a difference in \nh\ of between $\sim10^{20}$ and $\sim10^{24}$ \cmsq\ is predicted by the torus models, and should be measurable. However, this is not seen observationally, as \citet{horst07} find no significant difference between Seyfert 1s and Seyfert 2s, or even Compton thick nuclei. They argue from this that the torus may not be smooth, and instead clumpy, where the torus is seen to be optically thin, but the clumps are optically thick, though with a small volume filling factor. Most importantly for this work though, it suggests that 12 micron selection is in fact not heavily biased against obscuration. This therefore makes the 12MGS the ideal parent sample for study of obscuration and unification using X-ray observations and is as such the subject of this paper. The 12MGS has previously been studied at X-ray wavelengths by \citet{barcons95} who used hard X-ray observations by {\it HEAO 1}, and at soft X-rays by \citet{rush96x} with {\it ROSAT} data. \citet{barcons95} use the X-ray and IR data sets to show that most of the local X-ray emissivity originates from Seyfert galaxies and that most of the local 2-10 keV X-ray luminosity function between 10$^{42}$ and 10$^{46}$ \ergs\ can be accounted for by 12 micron emitting AGN. They also test the unification scheme and come to the conclusion that it must be modified to include the reduction of the covering fraction of the torus with increasing source luminosity, an observation first suggested by \citet{lawrence91} which has come to be known as the `receding torus model'. \citet{rush96x} conducted a spectral analysis of the 0.1-2.4 keV {\it ROSAT} data. For spectra with enough counts, they fit absorbed power-laws, finding a median soft X-ray spectral index of 2.3 for all Seyfert types. In this paper, we present an X-ray spectral analysis of the galaxies of the 12MGS with an observation by \xmm, which consists of a heterogeneous mix of Seyferts, LINERs and star-forming galaxies. We focus on the X-ray spectral properties of the sources without regard for their optical spectroscopic types and use X-ray luminosity to select unambiguous AGN. In a companion paper though, we continue our investigation into this sample combining the X-ray data from this paper with optical spectroscopic and infrared data. In Section \ref{transsec} we present work on Monte-Carlo simulations of X-ray radiative transport which we use in our spectral analysis. Section \ref{specfit} describes our spectral fitting methods and Section \ref{specfitresults} gives our spectral fitting results. Finally we discuss all of these results in Section \ref{discussion} and present our conclusions in Section \ref{conclusions}.

\label{conclusions} In summary, we have presented an X-ray spectral analysis of the 126 galaxies with \xmm\ coverage in the 12MGS, which is a relatively unbiased and representative selection of galaxies. We have done so with the help of new X-ray radiative transfer calculations, based on Monte-Carlo simulations which take into account photo-electric absorption, Compton scattering and includes fluorescent iron line emission. We have compiled new X-ray spectral models from these data, {\tt trans} and {\tt torus} which we have presented here. These table models are available publicly on the web at {\tt http://astro.ic.ac.uk/mbrightman/home} Our main results from these models have been: \begin{itemize} \item unobscured AGN with a presumed toroidal distribution of matter around it can achieve a maximum iron K$\alpha$ equivalent width of $\sim150$ eV, a useful limit to consider when assessing if a source is truly unobscured or not. From our calculations, we show that in order to exceed this EW, the line of sight must be obscured with \nh$>10^{23}$ \cmsq\ (e.g. NGC 3690). \item for \nh=$10^{25}$ \cmsq, the flux seen in the 10-40 keV, 20-100 keV ({\it INTEGRAL}), and 14-195 keV ({\it Swift/BAT}) bands is only 10\% of the intrinsic flux, which is important for considering the biases present against hard X-ray selected, heavily obscured AGN. At this \nh, 10\% or less of the intrinsic flux is seen in any X-ray band at all redshifts, which reveals the bias present in all X-ray surveys against such heavily obscured systems. \item using spectral models based on slab geometries such as {\tt pexrav} or {\tt pexmon} will underestimate the reflection fraction and intrinsic luminosities with respect to toroidal geometries by a factor of 2-3, leading to underestimation of the intrinsic luminosity of the source, as also found by \citet{murphy09}. \end{itemize} By combining the positive attributes of 12 micron selection with the penetrative power of X-rays, we have been able to determine the intrinsic X-ray properties of galaxies such as source power, power-law index, and in addition we have assessed the absorption occurring in the sources. We have improved upon previous works on this sample with better signal to noise, spectral resolution and larger band pass by using \xmm\ data. We have developed a detailed systematic approach to the fitting of the spectral data, with the aim of determining the intrinsic power, $\Gamma$ and \nh\ as accurately as possible. We have paid particular attention to the uncovering of Compton thick sources using our new model, {\tt trans} and the {\tt pexmon} model. The main conclusions from this study have been: \begin{itemize} \item we find the mean primary power-law index to be $<\Gamma>=1.90_{-0.07}^{+0.05}$ for our sample with an intrinsic spread of $\sigma = 0.31_{-0.05}^{+0.05}$. The index is consistent with previous works, though we find a larger intrinsic dispersion in our results than previously reported. \item we also find that the mean power-law index for obscured and unobscured sources is consistent with the two source populations being the same, supporting unified schemes. \item we find no dependence of $\Gamma$ on X-ray luminosity as has been previously reported for local AGN. \item we find an 18\% Compton thick fraction for the X-ray AGN defined here, which is higher than the hard X-ray selected samples, but roughly consistent with XRB predictions. \item the obscured fraction for our sample is a strong function of X-ray luminosity, peaking at a luminosity of $\sim10^{42-43}$ \ergs. The decline in the obscured fraction at high luminosities, where these sources are unambiguously AGN, is a confirmation of previous works. The decline at lower luminosities has also been suggested but needs further explanation accounting for optical spectroscopic classifications. \end{itemize}

1012.3345

We present an X-ray spectral analysis of 126 galaxies of the 12 μm galaxy sample. By studying this sample at X-ray wavelengths, we aim to determine the intrinsic power, continuum shape and obscuration level in these sources. We improve upon previous works by the use of superior data in the form of higher signal-to-noise ratio spectra, finer spectral resolution and a broader bandpass from XMM-Newton. We pay particular attention to Compton thick active galactic nucleus (AGN) with the help of new spectral fitting models that we have produced, which are based on Monte Carlo simulations of X-ray radiative transfer, using both a spherical and torus geometry, and taking into account Compton scattering and iron fluorescence. We use this data to show that with a torus geometry, unobscured sightlines can achieve a maximum equivalent width of the Fe Kα line of ∼150 eV, originally shown by Ghisellini et al. In order for this to be exceeded, the line of sight must be obscured with N<SUB>H</SUB> &gt; 10<SUP>23</SUP> cm<SUP>-2</SUP>, as we show for one case, NGC 3690. We also calculate flux suppression factors from the simulated data, the main conclusion from which is that for N<SUB>H</SUB>≥ 10<SUP>25</SUP> cm<SUP>-2</SUP>, the X-ray flux is suppressed by a factor of at least 10 in all X-ray bands and at all redshifts, revealing the biases present against these extremely heavily obscured systems inherent in all X-ray surveys. Furthermore, we confirm previous results from Murphy &amp; Yaqoob that show that the reflection fraction determined from slab geometries is underestimated with respect to toroidal geometries. For the 12 μm selected galaxies, we investigate the distribution of X-ray power-law indices, finding that the mean (&lt;Γ&gt;= 1.90<SUP>+0.05</SUP><SUB>-0.07</SUB> and σ<SUB>Γ</SUB>= 0.31<SUP>+0.05</SUP><SUB>-0.05</SUB>) is consistent with previous works, and that the distribution of Γ for obscured and unobscured sources is consistent with the source populations being the same, in general support of unification schemes. We determine a Compton thick fraction for the X-ray AGN in our sample to be 18 ± 5 per cent which is higher than the hard X-ray (&gt;10 keV) selected samples. Finally we find that the obscured fraction for our sample is a strong function of X-ray luminosity, peaking at a luminosity of ∼10<SUP>42-43</SUP> erg s<SUP>-1</SUP>.

12205221

2011IAUS..260E..16C

1012

1012.3035_arXiv.txt

We, astronomers and people around us now, tend to see in astronomy an activity that is mainly, if not solely devoted to the understanding of the Universe and the objects to be found within. This is an intellectual quest that is particularly fruitful since the beginning of the space age and the opening of the electro-magnetic spectrum from radio waves to gamma rays to the observations of the sky. This is certainly a valid point of view now. The quest for understanding as an activity for itself has, however, certainly not been the main activity of astronomers over history. Most of the astronomical workforce has been invested not so much in a cultural endeavour, but rather in a tedious quest for the measurement and keeping of time. Agriculture requires that seeds are planted well in advance of the season of growth. It is, therefore, necessary to know when to plant. This time cannot be estimated on the current weather, but requires advance planning for which the only useful information is based on the positions of Sun, Moon and stars. Since the Moon month and the year do not have a simple relation, the quest for timekeeping during the year is a complex one. Astronomers have, therefore, spent large efforts world wide to solve this problem. The same is true for other needs of society, navigation certainly requires seaworthy ships, it also requires the capacity to locate oneself on the surface of the Earth in unknown territories and on sea. Again, this is provided by a knowledge of the respective positions of Sun, Moon, planets and stars, together with a good mastering of time keeping. All of this knowledge is based on astronomy. The daily needs of human society also require some level of time keeping, be it only to be able to meet at given time and place. Here again the local timekeeping in many organised societies has relied on astronomical observations. The results being then relayed to the population by bells, canons and other signals. Astronomy has thus been a most practical endeavour for most of the human history. The present conference amply demonstrates this not only in the European culture, but also in other ones. The benefits of astronomy for society cannot be overestimated. Agriculture, navigation and the organisation of societies have depended crucially on astronomical observations for almost the whole of the development of human civilisations. The development of a "Weltbild" should be seen in this context as an, important, side benefit. The practical importance of astronomy has only rather recently ceased to dominate our activities. The observatory of Geneva has, as one example, been in charge of certifying chronometers for the local manufactures until the late 1960's. It is also interesting to learn that the observatory of Besan\c{c}on, in France, North of the Jura mountains, was founded in 1878 not so much to contribute to the great human quest for knowledge about the world than to help the local watch manufacturers who considered themselves at a disadvantage compared to their competitors of the South of the Jura mountains, in Neuch\^atel and Geneva, who had access to astronomical observatories. It is, therefore, only very recently that Astronomy as become a mainly cultural quest, the practical benefits of which are only minor and indirect, to be found in the knowledge gained in making complex instruments capable of observing in remote and hostile environments.

1012.3035

Astronomy has played a major part in the development of civilisations, not only through conceptual developments, but most importantly through the very practical gains obtained through the observation of Sun, Moon planets and stars. Space sciences, including astronomy, have also played a major rôle in the development of modern societies, as an engine for most subsequent space technology developments. Present trends tend to decrease the rôle of science in space development. This trend should be reversed to give modern ``societies'' their independence in space-related matters that permeate the lives of all inhabitants of the Earth.

2857948

2011PhRvD..83b3510A

1012

1012.3203_arXiv.txt

Currently, the most successful model of cosmology that we have is $\Lambda$CDM. It is constructed in order to match a wide variety of modern cosmological observations that have stunned the physicist community in the last decades. Among them, there are precision measurements of anisotropies in the cosmological microwave background radiation \cite{WMAP,WMAP7y}, baryon acoustic oscillation \cite{Eisenstein05,Percival07}, and Type Ia supernovae \cite{Riess98,Perlmutter99,Union2}. All these observations point out that our Universe at present is dominated by a cosmological constant -dark energy- and there are about 5 times of some unknown nonbaryonic, dark matter, over the baryonic matter which is well understood by the standard model of particles. Nonetheless, some objections to the $\Lambda$CDM model exist, both theoretical \cite{Weinberg,CopSamTsu06} and observational (see for example \cite{Peebles10}). Thus, alternative proposals have appeared in the literature giving rise to the idea that dark energy varies with time. Scalar fields have attracted special attention, mainly quintessence \cite{RatraPeeb88,Caldwell98,Copeland98}, among some other alternatives \cite{CopSamTsu06}. Recently there has been a lot of interest in studying couplings in the dark sector species \cite{Amendola99,Bean01,Amendola02,FarPeeb04,Koivisto05,Bean08, Corasaniti08,Caldera09,Costa09,Mota07}. This is in part motivated by the fact that until today we can only extract information of these components through gravitational interaction, a feature that has been dubbed {\it Dark Degeneracy} \cite{Hu99,Rubano02,Wasserman02,Kunz20091,Kunz20092}. Specifically, we can define the energy content of the dark sector, and in fact the dark sector itself, using the Einstein field equations \begin{equation} 8 \pi G T_{\mu\nu}^{dark} = G_{\mu\nu} - 8 \pi G T_{\mu\nu}^{obs}, \end{equation} where $G_{\mu\nu}\,$ comes from the observed geometry of the Universe and $T_{\mu\nu}^{obs}\,$ from its observed energy content. In this sense the dark sector reflects our lack of knowledge. The easiest way -mathematically and conceptually- to accommodate these ideas into the observed history of the Universe is decomposing $T_{\mu\nu}^{dark}$ in two species, dark energy and dark matter, \begin{equation} \sla{canondecomp} T_{\mu\nu}^{dark} = T_{\mu\nu}^{DE} + T_{\mu\nu}^{DM}. \end{equation} But this decomposition is not unique. In fact, the dark sector could be composed by a large zoo of particles and complicated interactions between them. Or, it could be even just one unknown field. To accomplish this last possibility we note that in the $\Lambda$CDM model the equation of state parameter of the total dark sector is given by \begin{equation} \sla{lcdmweff} w_{eff} \equiv \frac{\Sigma \rho_i w_i}{\Sigma \rho_i }\simeq - \frac{1\,\,}{1+ \,0.315\, a^{-3}}, \end{equation} where $i$ index the dark matter component and the cosmological constant, and in the last equality we have used $\Omega^{(0)}_{DM} / \Omega^{(0)}_{DE} \approx 0.315$ \cite{WMAP7y}. In order to mimic this model with just one dark field, we must have $ w_{dark} \simeq w_{eff}$. Any fluid with equation of state parameter equal to $ w_{eff}$ will produce the same expansion history of the Universe. On the other hand, some string theory-inspired models of dark energy share the peculiarity that scalar fields, like the dilaton, couple directly to matter with gravitational strength. To have cosmological influence at present, these fields must be nearly massless, leading to long-range fifth-forces and to large violations of the equivalence principle. Some mechanisms have been proposed in order to avoid this unacceptable behavior, as the Damour-Poliakov effect \cite{Damour90,Damour94}, in which the interaction is dynamically driven to zero by the expansion of the Universe, and as the chameleon mechanism \cite{Khoury041,Khoury042}, where the mass of the scalar field has an environment density dependence, becoming very huge in overall high density regions, such as those in which Einstein principle equivalence and fifth-force search experiments are performed. In this work we follow the both lines of thought outlined above. To this end we consider an interaction Lagrangian between a nearly massless scalar and the trace of the energy momentum tensor of the ordinary matter fields given by \begin{equation} \sla{intlag} \mathcal{L}_{int} = \sqrt{-g} A(\phi) T. \end{equation} This type of coupling has been investigated in a cosmological context in \cite{NarPad85,SinghPad88,SamiPad03}. This interaction has some attractive properties, the field does not couple to the electromagnetic field and in this sense is dark; also, it does not couple to relativistic matter and then do not affects the success of the early Universe cosmology, although it could couple to the inflaton field. As we shall see, the coupling could also give mass to the field. At a more fundamental level we can consider, in a first approximation, a fermionic matter free fields with energy momentum tensor $T_{\mu\nu}^{(f)} = -i \bar{\psi} \gamma_{\mu}\partial_{\nu} \psi\,$ and trace given by $T^{(f)} = -i \bar{\psi} \gamma^{\mu}\partial_{\mu} \psi = - m_{\psi} \bar{\psi}\psi$, where in the last equality we have used the Dirac equation. The coupling then becomes a Yukawa-like one and gives mass to the fermions. If we chose correctly the function $A(\phi)$ we can also interpret this result as that the interaction has given mass to the field $\phi$. This paper is organized as follows: in Section II, we present the details of the general theory; in Section III, we derive the background cosmology equations and we choose a specific model that mimics the $\Lambda$CDM model, then we numerically obtain the cosmological solutions; in Section IV, we work the theory of linear perturbations; finally, in Section V, we present our conclusions.

It is possible that future observations and experiments reveal that there exist some interactions between the dark sector and the known fields of the standard model; indeed, this is the hope for terrestrial experiments in detecting dark matter. The dark degeneracy allows us to study a wide variety of models that mimic the $\Lambda$CDM at today's observation accuracy but with a richer dynamics. In this paper we have presented a plausible mechanism in which an interaction between a quintessence-like field and baryonic matter gives rise to effects similar to those of dark matter. The interaction is given through the trace of the energy momentum tensor of the matter fields, and in the case that we have considered, dust matter, it is the same as an interaction of the type $\,f(\phi)\mathcal{L}$. The differences are manifest if we instead use a radiation field, {\it e.g.} photons, in this case the interaction vanishes -neglecting the trace anomaly \cite{BirrelDavies}- and in this sense is dark. Because of this coupling the radiation era is the same as in the standard big bang and the differences appear once matter dominates. We have described the background cosmology and by fitting the parameters of the model we can mimic as far as we want the late time $\Lambda$CDM background solutions, from decoupling until today. Specifically we have shown the differences for the cases of $\beta = 0.4$ ($\epsilon \approx G$) and $\beta=0.01$. For this last value of $\beta$, we made fits to the Union 2 supernovae data set and found that the best fit is obtained by choosing the other two parameters of the model equal to $C_1=0.305\,$ and $\,C_2(\phi_0) = 4.66$. We have worked out the first order perturbation theory and shown that these models are capable to form linear structure without dark matter. Initial perturbations of a free scalar field dark energy are erased because its sound speed is equal to one, and by the fact that dark energy accelerates the Universe, it tends to freeze any matter perturbation. One of the effects of the interaction, when the field is slow-rolling, is to decrease the sound speed. In our model this induces the scalar field perturbations to oscillate about a nonzero negative value and yield a repulsive gravitational force over the baryonic matter perturbations. Despite this effect, we showed that the baryonic matter density contrast could grow as fast as in the $\Lambda$CDM model. This is because the interaction enhances the gravitational strength of the baryons. This latter effect increases as $\beta$ does, but at the same time the increase of $\beta$ can break the slow-roll. Thus, suitable $\beta$ are found for $\beta \lesssim 0.04$.

1012.3203

We present a scenario in which a scalar field dark energy is coupled to the trace of the energy momentum tensor of the baryonic matter fields. In the slow-roll regime, this interaction could give rise to the cosmological features of dark matter. We work out the cosmological background solutions and fit the parameters of the model using the Union 2 supernovae data set. Then, we develop cosmological perturbations up to linear order, and we find that the perturbed variables have an acceptable behavior, in particular, the density contrast of baryonic matter grows similar to that in the ΛCDM model for a suitable choice of the strength parameter of the coupling.

12137726

2010idm..confE.114S

1012

1012.1206_arXiv.txt

The Large Area Telescope (LAT) on board the Fermi Gamma-ray Space Telescope (\emph{Fermi})~\cite{Atwood:2009ez} is observing the high-energy sky with unprecedented precision and sensitivity. In addition to numerous individual sources, the LAT observes a substantial level of diffuse gamma-ray emission. This emission includes a Galactic component resulting from the interactions of cosmic rays with the interstellar gas and radiation fields, as well as a component that appears isotropic on large angular scales, which is often assumed to originate from unresolved members of cosmological source populations. Many astrophysical sources are guaranteed to contribute to the large-scale isotropic gamma-ray background (IGRB), including cosmological populations such as blazars~\cite{Stecker:1996ma, Narumoto:2006qg} and star-forming galaxies~\cite{Thompson:2006np, Fields:2010bw}, as well as Galactic source classes whose sky distributions extend to high Galactic latitudes, e.g., millisecond pulsars~\cite{FaucherGiguere:2009df}. Proposed but unconfirmed sources of gamma-ray emission, such as the annihilation or decay of dark matter particles in Galactic or extragalactic structures, may also contribute to the IGRB~\cite{Ullio:2002pj, Zavala:2009zr, Abdo:2010dk}. Interestingly, the \emph{Fermi}-measured IGRB energy spectrum~\cite{Abdo:2010nz} is consistent with a single power law over a large range of energies, and hence lacks any spectral features which could aid in identifying contributions from individual source classes. Moreover, recent \emph{Fermi} results~\cite{Collaboration:2010gqa} indicate that the majority of the IGRB does not originate from members of the source classes already detected by \emph{Fermi}, leaving the origin of this emission a mystery. To complement searches based on spectral features, recent work has considered the possibility of using angular anisotropy information in the IGRB to help to reveal its contributors~\cite{Zavala:2009zr, Ando:2005xg, Ando:2006cr, SiegalGaskins:2008ge, Fornasa:2009qh, Ando:2009fp, Ando:2009nk}. If the IGRB is composed of emission from unresolved members of gamma-ray source populations, characteristic small-scale fluctuations are expected to be present due to the variation in the number density of sources along the line-of-sight in different sky directions. In this study we searched for angular anisotropies in the IGRB measured by the \emph{Fermi}-LAT\@. The angular power spectrum of the emission, after masking low Galactic latitudes and known sources, was calculated in several energy bins. The results from the data were compared with those from a simulated model of the gamma-ray sky in order to identify any significant differences in anisotropy properties. Preliminary results from this angular power spectrum analysis of the IGRB are presented.

1012.1206

The contribution of unresolved sources to the diffuse gamma-ray background could produce anisotropies in this emission on small angular scales. Recent studies have considered the angular power spectrum and other anisotropy metrics as tools for identifying contributions to diffuse emission from unresolved source classes, such as extragalactic and Galactic dark matter as well as various astrophysical gamma-ray source populations. We present preliminary results of an anisotropy analysis of the diffuse emission measured by the Fermi-LAT.

12167631

2011ApJ...728...30T

1012

1012.0346_arXiv.txt

\label{s:intro} Understanding how solar and stellar coronae are heated to high temperatures is one of the most important open issues in astrophysics. Coronal heating is clearly related to the strong magnetic fields that fill the atmospheres of the Sun and solar-like stars, but the mechanism that converts magnetic energy into thermal energy remains unknown. Theoretical models based on steady heating failed to explain the physical properties observed in coronal structures. A promising and widely studied framework for understanding coronal heating is the nanoflare model proposed by \citet{Parker72}. In this model, convective motions in the photosphere lead to the twisting and braiding of coronal magnetic field lines. This topological complexity ultimately leads to the formation of current sheets, where the magnetic field can be rearranged through the process of magnetic reconnection. This model has been further refined and adapted to a scenario where nanoflares occur in unresolved strands in coronal loops, below the spatial resolution of available instrumentation \citep{Parker88,Cargill94,Cargill97,Klimchuk06,Warren03,Parenti06}. Nanoflare models predict the presence of very hot plasma with temperatures in excess of 3~MK in non flaring solar regions; depending on the energy of the nanoflare, temperatures may reach or even exceed 10~MK. Unambiguous detection of such extreme temperatures in non-flaring solar regions can provide convincing evidence for the presence of nanoflares. However, such detection is not easy, as the amount of very hot plasma produced by nanoflares is expected to be very small. Recent studies have provided some evidence of the presence of hot plasma in the non-flaring Sun. \citet{McTiernan09} carried out an analysis of the temperature and emission measure determined from quiescent plasma during the 2002-2006 decay phase of solar cycle 23, as derived from GOES and RHESSI observations. He found a persistent faint plasma component with temperatures approximately constant during the entire 2002-2006 interval, and approximately between 5 and 10~MK. However, GOES and RHESSI provided different values of such temperature, and their results were not necessarily well correlated; also, this analysis relied on the isothermal plasma assumption. Other studies tried to identify the hot plasma through the X-ray emission observed by the X-ray Telescope (\xrt; \citealp{Golub07}) onboard \hinode\ \citep{Kosugi07}. \citeauthor{Reale09} carried out a temperature analysis of a non-flaring active region, first with only \xrt\ multi-filter data \citep{Reale09} and then combining \xrt\ and RHESSI observations \citep{Reale09b}. Their findings point to the presence of small amounts of very hot plasmas, with temperatures of $\simeq 5-10$~MK, and emission measure of the order of few percent of the dominant cool component. These characteristics of the emission measure distribution are compatible with the predictions of nanoflare models. However, they spelled out and discussed the main limitations of their study and similar analyses. First, \xrt\ is also sensitive to plasma at normal active region temperatures (2-3~MK) and thus contamination from the colder active region plasma is a considerable obstacle to the detection of the much smaller amounts of hot plasma; also, the limited temperature resolution of \xrt\ prevents a detailed study of the cold component. Second, RHESSI sensitivity makes it very hard to even detect the quiescent active region plasma. Third, instrumental calibration is an issue for both instruments. \citet{Schmelz09a} also detected a faint hot temperature tail to the emission measure distribution of active region plasma, and determined its temperature to be around 30~MK. Subsequent analyses by \citet{Schmelz09b} included RHESSI data and, while confirming the presence of such hot material, could not reconcile the \xrt\ and RHESSI observations using the standard calibration of both instruments. A self-consistent solution was only found if a series of instrumental parameters and the plasma element abundances were adjusted, and the temperature of the hot plasma decreased. Ample efforts have been devoted to the accurate determination of the thermal structuring of coronal plasma to derive robust observational constraints to the mechanism(s) of coronal heating. The plasma temperature distribution of the quiet corona and of active regions has been investigated through imaging data and spectroscopic observations \citep[e.g.,][]{Brosius96,Landi98, Aschwanden00,Testa02,DZM03,Reale07,Landi09,Shestov10,Sylwester10}. Several recent works have focused on EUV spectra obtained with the \hinode\ Extreme Ultraviolet Imaging Spectrometer (\eis; \citealt{Culhane07} ) which provides good temperature diagnostic capability, together with higher spatial resolution and temporal cadence than previously available \citep[e.g.,][]{Watanabe07,Warren08loops,Patsourakos09, Brooks09,Warren09}. In the present work, we address the issue of determining the temperature distribution of coronal plasma from a different perspective: we investigate thermal properties of coronal plasma in non-flaring active regions using simultaneous \hinode\ observations with \xrt\ and with \eis, which provide complementary diagnostics for the X-ray emitting plasma. The multi-filter \xrt\ dataset together with \eis\ spectra, including its entire wavelength range, allow to accurately determine the thermal structure of the active region plasma, and to explore the presence of hot plasma in non-flaring regions. We use spectroscopic observations from the \hinode/\eis\ instrument of a quiescent active region to constrain the emission measure distribution of the bulk of the active region plasma with the spectral lines observed by \eis\ in the 171-212\AA\ and 245-291\AA\ spectral ranges (see also e.g., \citealt{Young07,Doschek07}). Since \eis\ is most sensitive to plasma with temperatures of 0.6-2~MK, \eis\ spectra allow us to accurately determine the emission measure distribution of the quiescent active region plasma, to evaluate the fraction of the observed \xrt\ count rates that it emits, and thus investigate the true amount of emission from the nanoflaring plasma. Thus, the combination of \xrt\ and \eis\ observations of the same active region allows us to characterize the plasma temperature distribution with better detail than in previous studies. In fact, while some previous studies have made use of data from both imaging and spectroscopic data to constrain the properties of the emitting plasma \citep[e.g.,][]{Warren10,Landi10,ODwyer10}, to our knowledge no previous work has carried out a determination of the temperature distribution by combining \xrt\ and \eis\ data, nor a quantitative comparison of the independent analysis from the different instruments, as we do here. Independent temperature analysis from the two datasets provide a cross-check of the different temperature diagnostics techniques applicable to spectral and broad-band data respectively, and insights into cross-calibration of the two instruments. The observations are described in Section~\ref{s:obs}. The data analysis and results of the determination of the plasma temperature distribution are presented in Section~\ref{ss:results}. Our findings are discussed in Section~\ref{s:discuss} and summarized in Section~\ref{s:conclusions}.

1012.0346

We analyze coordinated Hinode X-ray Telescope (XRT) and Extreme Ultraviolet Imaging Spectrometer (EIS) observations of a non-flaring active region to investigate the thermal properties of coronal plasma taking advantage of the complementary diagnostics provided by the two instruments. In particular, we want to explore the presence of hot plasma in non-flaring regions. Independent temperature analyses from the XRT multi-filter data set, and the EIS spectra, including the instrument entire wavelength range, provide a cross-check of the different temperature diagnostics techniques applicable to broadband and spectral data, respectively, and insights into cross-calibration of the two instruments. The emission measure distributions, (EM(T)), we derive from the two data sets have similar width and peak temperature, but show a systematic shift of the absolute values, the EIS (EM(T)) being smaller than the XRT (EM(T)) by approximately a factor two. We explore possible causes of this discrepancy, and we discuss the influence of the assumptions for the plasma element abundances. Specifically, we find that the disagreement between the results from the two instruments is significantly mitigated by assuming chemical composition closer to the solar photospheric composition rather than the often adopted "coronal" composition. We find that the data do not provide conclusive evidence on the high temperature (log T(K) &gt;~ 6.5) tail of the plasma temperature distribution, however, suggesting its presence to a level in agreement with recent findings for other non-flaring regions.

1047101

2011ApJ...729....2A

1012

1012.2200.txt

Blazars, a subclass of active galactic nuclei (AGN), are the dominant extragalactic source class in $\gamma$-rays. They have been observed to show rapid variability and non-thermal spectra, presenting a broad continuum across nearly the entire electromagnetic spectrum. This implies that the observed photons originate within highly relativistic jets oriented very close to the observer's line of sight \citep{Urry}. This orientation results in Doppler beaming that boosts the intensity and frequency of the observed jet emission, often overwhelming all other emission from the source. Therefore, blazars make excellent laboratories for studying the physical processes within the jets of AGN. They were among the first sources to be detected in the VHE band, and as of this writing there are 34 known VHE $\gamma$-ray blazars\footnote{http://tevcat.uchicago.edu/}. Markarian 501 (Mrk 501; 1H1652+398), at a redshift of $z=0.034$, was the second blazar to be detected at VHE \citep{501discovery}. The spectral energy distribution (SED) of Mrk 501 characteristically shows a double-peaked profile. These peaks occur at keV and TeV energies when the SED is plotted in the $\nu$F$_\nu$ vs $\nu$ representation. This structure is common among all VHE $\gamma$-ray blazars, and several models have been developed to account for the double-peaked structure. These models uniformly attribute the peak at keV energies to synchrotron radiation from relativistic electrons and positrons within the blazar jets, but they differ in accounting for the source of the VHE peak. The models are generally divided into two classes: leptonic and hadronic, named for their attributed source for the VHE peak. The leptonic models advocate inverse-Compton scattering to VHE of either the synchrotron photons from within the jet or an external photon field \citep[e.g.,][]{Marscher,Maraschi,Dermer,Sikora}. The hadronic models, however, account for the VHE emission by $\pi^0$ decay or by $\pi^{\pm}$ decay with subsequent synchrotron and/or Compton emission from decay products, or by synchrotron radiation from ultra-relativistic protons \citep[e.g.,][]{Mannheim, Aharonian, Pohl}. Observationally, Mrk 501 has been known to undergo both major outbursts on long time scales and rapid flares on short time scales, most prominently in the keV and VHE range \citep[e.g.,][]{Catanese, Pian, xue501, albert}. During these outbursts, both of the SED peaks have been observed to shift towards higher energies. During the most extreme cases, the keV peak has been observed above 200 keV, well above typical values below 1 keV. Historically, the SED has been measured in the VHE band primarily during outbursts, due to the lower sensitivity of previous generations of instruments. A previous study examined the quiet state, but it was performed before the launch of the Fermi Large Area Telescope (\FermiLATc), which provides coverage in the range between keV and VHE energies \citep{magic501}. This work attempts to provide state-of-the-art short-term multiwavelength measurements of the quiescent state of Mrk 501, with broad spectral coverage in the critical keV and VHE bands as well as coverage with the \FermiLATc. These measurements are then compared to observations by BeppoSAX in X-rays and by the Whipple 10m and CAT Cerenkov telescopes in VHE $\gamma$-rays during the 1997 extreme outburst. This outburst has been well studied using multiple instruments \citep[e.g.,][]{hegra501, Catanese, hegra501, Pian} and provides a good comparison to the quiet state observed in 2009.

This data set provides a high-quality sampling of the broadband SED of Mrk 501 in the quiescent state, and it allows comparisons to be made over a broad energy range with the extreme outburst observed in 1997. In choosing the values for the SED model parameters, we attempted to be consistent with previous work in choosing $\gamma_{min}$, $\gamma_{max}$, $B$, $R$, and $D$ while still matching the data (see Table \ref{tab:sscmodels}). In addition, when applying the models to the data from the two states, we attempted to limit the differences to the electron spectral indices and break energy ($\gamma_{b}$), similar to \cite{Pian} and \cite{magic501}, with a small shift in $B$ included to provide a better match to the data. The SSC model for the 1997 outburst was matched to the data from \citet{Pian} and then the shape compared to the data from \citet{cat}. The resulting SSC models can reproduce the measured spectra from the keV to the VHE range for both the quiescent state and the 1997 outburst. The successful match of the SED models to the data from both the outburst and the quiescent state primarily by a modification of the injected electron spectrum implies that this may be the primary explanation for the dramatic shift in the peak frequency. Specifically, $\gamma_{b}$ and the power-law spectral index above the break appear to drive the changes. In addition, there is a difference in electron densities, with a density of $2.7\times10^4$ cm$^{-2}$ for the 1997 data and $2.2\times10^3$ cm$^{-2}$ for the 2009 data. The SSC model matches indicate that during the high-flux state the break energy is shifted an order of magnitude higher and that past the break energy the spectrum is significantly harder than in the low-flux state. It should be noted that the match of these models to the data is not unique due to the number of free parameters. Although we did not fully explore the multi-dimensional parameter space of the model, which is not the main objective of this work, the general conclusion on the shift of the SED peaks seems to be robust. The X-ray data do not allow us to measure the location of the low-energy SED peaks directly, but they can be used to place limits on the peak energies. From inspection of the X-ray data points alone, the keV peak should be near 230 keV ($5.5 \times 10^{19}$ Hz) during the 1997 outburst and near 0.6 keV ($1.5 \times 10^{17}$ Hz) during the quiescent state. This means that with a change in flux of around one order of magnitude between the quiescent state and the 1997 outburst, the keV peak of the SED shifts in frequency by more than two orders of magnitude. In contrast to this large shift, the VHE peak of the SED does not seem to shift as dramatically in location. The SED peak at VHE is better constrained with the 2009 data than with the 1997 data, but the SED models applied to the data imply that even with a dramatic shift in flux at VHE, the peak energy in the VHE range is stable compared to that in the keV range. This may be due to Klein-Nishina (KN) effects which reduce the cross-section for scattering when dealing with electrons of high energy, as also noted by \citet{Pian}. To examine the possible KN effects, which become important above $h \nu \sim m_ec^2$ in the electron rest frame, we examined the energies of typical electrons from the model distributions. From the 1997 SED model, the peak attributed to synchrotron radiation is located at $3.6\times 10^{19}$ Hz. The peak energy of the emission in the jet frame, taking the Doppler factor of $D=20$ into account, is $E_{peak} \approx 7.4$ keV. From \cite{rybicki}, it can be shown that for synchrotron radiation: \begin{equation} \gamma_{e}^2 \approx \frac{8 E_{peak} m_e c}{3 \pi q B \hbar}\;, \label{eqn:gammae} \end{equation} \noindent where $q$ is charge, $\gamma_e$ is the characteristic Lorentz factor of the electrons contributing to the bulk of the X-ray emission and $E_{peak}$ is the peak energy of the synchrotron model. This yields a value of $\gamma_e = 1.5\times10^6(B/0.23\mbox{ G})^{-1/2}(D/20)^{-1/2}$ for the 1997 outburst. For the 2009 model, the synchrotron emission model peaks at $2.24\times10^{17}$ Hz, giving a peak energy in the jet frame of $E_{peak} \approx 46$ eV. This yields a value of $\gamma_e = 1.0\times10^5(B/0.34\mbox{ G})^{-1/2}(D/20)^{-1/2}$ for the 2009 measurements. Using these electron energies, we can examine the importance of KN effects during the course of the two measurements. For the 1997 outburst, $\gamma_{e}h \nu_{peak} \approx 1.1\times10^4$ MeV, well above the electron rest mass energy of 0.511 MeV, placing this model scenario in the extreme KN range. For the 2009 quiescent state model, $\gamma_{e}h \nu_{peak} \approx 4.6$ MeV, moderately above the KN limit. The extreme energies involved during the 1997 outburst may indicate that the emission is piling up in the VHE range due to the reduced KN cross-section. The SED models also allow us to examine the characteristic cooling timescales of the synchrotron and inverse-Compton processes that are assumed to be the sources of the keV and VHE emission, respectively. The synchrotron cooling time is shown by \cite{rybicki} to be: \begin{equation} \tau_{syn} \approx \frac{6\pi m_e c}{\sigma_{T} \gamma_{e} B^2}\;, \label{eqn:syn} \end{equation} \noindent where $\sigma_T$ is the Thompson cross-section and $\gamma_e$ is the characteristic Lorentz factor of the electrons contributing to the bulk of the X-ray emission. The inverse-Compton cooling timescale was calculated for the models, taking into account KN effects using the method discussed in \citet{kncalc}. As discussed above, models for both sets of observations indicate that KN effects are important, so these effects were carefully included in the calculation. Using the parameters for the 1997 model, the synchrotron cooling timescale is comparable to the SSC cooling timescale ($\tau_{syn} \approx 9.8 \times 10^3$ s, $\tau_{SSC} \approx 3.7\times 10^3$ s) for electrons with $\gamma_e = 1.5\times10^6(B/0.23\mbox{ G})^{-1/2}(D/20)^{-1/2}$. However, for the quiescent state model from 2009, the timescale for inverse-Compton cooling is two orders of magnitude shorter than the synchrotron cooling timescale ($\tau_{syn} \approx 6.7 \times 10^4$ s, $\tau_{SSC} \approx 4.6 \times 10^2$ s) for electrons with $\gamma_e = 1.0\times10^5(B/0.34\mbox{ G})^{-1/2}(D/20)^{-1/2}$, indicating that in this model scenario, radiative cooling is dominated by the SSC process for the values of $\gamma_e$ calculated above. This seems to indicate that the SSC peak energy is determined by the transition into the KN regime rather than the maximum or peak electron energies. Beyond this transition energy, the ratio of inverse-Compton to synchrotron cooling timescales should decrease dramatically. An estimate for the KN transition can be found using the peak of the synchrotron spectrum. For example, from the 2009 data, this peak resides at 46 eV in the jet frame. This corresponds to a value of $\epsilon'= E'/(m_e c^2)=9\times10^{-5}$. The transition to the KN regime is then at the observed energy of $E_{KN} \sim m_e c^2 D (1/\epsilon') \sim 110$ GeV, very near the observed SSC peak. In addition, all of the timescales calculated were comparable to or less than the light-crossing time for the modeled emission regions, indicating that light travel time effects may dominate the light curve profiles. Further intense multiwavelength observations of the source, in conjunction with long-term monitoring campaigns, will continue to be important to understand the broadband behavior of Markarian 501 and to shed light on the blazar phenomenon in general. % If you have acknowledgments, this puts in the proper section head. \bigskip % extra skip inserted

1012.2200

The very high energy (VHE; E &gt; 100 GeV) blazar Markarian 501 (Mrk 501) has a well-studied history of extreme spectral variability and is an excellent laboratory for studying the physical processes within the jets of active galactic nuclei. However, there are few detailed multiwavelength studies of Mrk 501 during its quiescent state, due to its low luminosity. A short-term multiwavelength study of Mrk 501 was coordinated in 2009 March, focusing around a multi-day observation with the Suzaku X-ray satellite and including γ-ray data from VERITAS, MAGIC, and the Fermi Gamma-ray Space Telescope with the goal of providing a well-sampled multiwavelength baseline measurement of Mrk 501 in the quiescent state. The results of these quiescent-state observations are compared to the historically extreme outburst of 1997 April 16, with the goal of examining variability of the spectral energy distribution (SED) between the two states. The derived broadband SED shows the characteristic double-peaked profile. We find that the X-ray peak shifts by over two orders of magnitude in photon energy between the two flux states while the VHE peak varies little. The limited shift in the VHE peak can be explained by the transition to the Klein-Nishina (KN) regime. Synchrotron self-Compton models are matched to the data and the implied KN effects are explored.

12213038

2011mast.conf..473G

1012

1012.2357_arXiv.txt

The presence of strong, globally-organized magnetic fields in hot, massive stars is rare. To date, only a handful of massive O-type stars are known to host magnetic fields. In 2009, Grunhut et al. reported the discovery of a strong magnetic field in the weak-wind O9IV star HD\,57682 from the presence of Zeeman signatures in mean Least-Squares Deconvolved (LSD) Stokes~$V$ profiles (see Fig.~\ref{fig1}). Grunhut et al. also used $IUE$ ultraviolet and ESPaDOnS optical spectroscopy to determine the following physical parameters using CMFGEN: $T_{\rm eff}=34.5$\,kK, $\log(g)=4.0\pm0.2$, $R=7.0^{+2.4}_{-1.8}$\,$R_{\odot}$, $M=17^{+19}_{-9}$\,$M_{\odot}$, and $\log({\dot{M}})=-8.85\pm0.5$\,$M_{\odot}$\,yr$^{-1}$. Of particular interest, we highlight the low mass-loss rate derived from UV wind diagnostic lines, which show variability characteristic of other magnetic OB stars (e.g. Schnerr et al. 2008), as shown in Fig.~\ref{fig1}. With only 7 observations of this star at our disposal at that time we employed a Bayesian inference method to determine the best-fitting dipole parameters to characterize the magnetic field (Petit et al. in prep). We concluded that the dipolar field of HD\,57682 is characterized by a polar strength of $\sim$1.7 kG, and a magnetic axis aligned within 10 to 50$^{\circ}$ of the rotation axis, depending on the inclination of the rotation axis to our line of sight. Since our original report of this star in 2009, we have obtained an additional 10 high-resolution spectropolarimetric observations with the ESPaDOnS instrument at CFHT. \begin{figure} \centering \includegraphics[width=2.5in]{grunhut_hd57682_fig1.eps} \includegraphics[width=3.25in]{grunhut_hd57682_fig1b.eps} \caption{{\bf Left:} Observed mean LSD Stokes $V$ (top), diagnostic null (middle), and unpolarized Stokes $I$ (bottom) profiles of HD\,57682 from 2008-12-06. {\bf Right:} Overplot of 3 $IUE$ UV spectra of the Si\,{\sc IV} line profiles (top). The significance of the variability is displayed at the bottom.} \label{fig1} \end{figure}

1012.2357

We will review our recent analysis of magnetic properties of the O9IV star HD 57682, using spectropolarimetric observations obtained with the ESPaDOnS at the Canada-France-Hawaii telescope within the context of the Magnetism in Massive Stars (MiMeS) Large Program. We discuss our most recent determination of the rotational period from longitudinal magnetic field measurements and Hα variability, -- the latter obtained from over a decade's worth of professional and amateur spectroscopic observations. Lastly, we report on our investigation of the magnetic field geometry and the effects of the field on the circumstellar environment.

12165890

2011AN....332..191S

1012

1012.5741_arXiv.txt

The major components of disk galaxies are bulges and disks, basically different in their support against gravitational collapse (\cite{J3_96}). Various correlations involving bulge and disk parameter have been established, e.g., bulge vs. disk scale lengths (MacArthur,~Courteau~\& Holtzman \cite{MCH_03}; \cite{M_04}; \cite{AEC_05}), bulge effective surface brightness (SB) vs. Hubble type (\cite{J3_96};\linebreak \cite{M_04}), bulge effective colour index (CI) vs. disk central CI (\cite{J4_96}). Furthermore, bulge-to-disk ratio underlies morphological classification of disk galaxies. Correlations between bulge parameters and black hole mass (\cite{FF_05}) evidence the coevolution of the\linebreak black hole and its host galaxy. Considering active galaxies, the problems related to the origin and angular momentum reduction mechanisms of the fuel have given rise to much discussion (e.g., Jogee \cite{J_06}). In this regard, comparative analysis of the morphology and local environment of matched active and inactive galaxy samples have been performed (e.g., \cite{MR_97}; De~Robertis, Yee \& Hayhoe \cite{RYH2_98}; Virani, De~Robertis \& VanDalfsen \cite{VRV_00};\linebreak \cite{SSF_07}). Thus, studies on the fueling me\-chanisms of active galactic nuclei and correlations among galaxy parameters, as well as the precise morphological classification, all demand morphological characterization, i.e., disclosure of the features present. We analysed the evidence of non-axisymmetric perturbation of the potential in a sample of 35 Seyfert galaxies and in a matched sample of inactive galaxies, based on a detailed morphological characterization and study of the local environment. The results are presented in Slavcheva-Mihova \& Mihov (\cite{SM_11}, hereafter Paper\,I). Here we present a homogeneous set of global, isophotal, and bar parameters of the Seyfert galaxies\footnote{Three more galaxies~-- Mrk\,1040, NGC\,5506, and Mrk\,507, were added.}. Global and isophotal parameters are involved in inclination corrections and various galactic structure studies. While databases provide global and isophotal parameters for the bulk of the galaxies, they are generally based on photographic data of less photometric accuracy than CCD data. The classical bar signature on the profiles is an ellipticity maximum, accompanied by a position angle (PA) plateau and a SB bump (Wozniak \& Pierce \cite{WP_91}; Wozniak et~al. \cite{WFM_95}); more detailed bar criteria were introduced later on (Knapen, Shlosman \& Peletier \cite{KSP2_00}; Men\'endez-Delmestre et~al. \cite{MSS_07}; Marinova \& Jogee \cite{MJ_07}; Aguerri, M\'endez-Abreu \& Corsini \cite{AMC_09}). There are various methods for bar length estimation, based on: visual inspection of images, analysis of the SB profile along the bar major axis, isophote fitting with ellipses, Fourier analysis, etc. (see the reviews of Athanassoula \& Misiriotis \cite{AM_02}; Erwin \cite{E_05}; Michel-Dansac \& Wozniak \cite{MW_06}). The semi-major axis (SMA) corresponding to the ellipticity maximum ($\ell_{\rm max}$, see Wozniak \& Pierce \cite{WP_91}) proved to be the most robust, objective, and reproducible among the bar length estimates. It, however, underestimates bar length (Wozniak et~al. \cite{WFM_95}) and is not related to any of the bar dynamical characteristics (e.g., Michel-Dansac \& Wozniak \cite{MW_06}). Moreover, bar strength can be defined as the maximum tangential force in terms of the mean radial force after \cite{CS_81}. Thus, it generally depends on the bar ellipticity, bar mass, and central force field. The tight correlation found between bar strength and deprojected bar ellipticity shows that the latter is a good measure of bar strength (Laurikainen, Salo \& Rautiainen \cite{LSR_02}). Values of the deprojected ellipticities below 0.15 are among the signatures of ovals and lenses (\cite{KK_04}). Details about the sample selection, observations, data reduction, and Johnson-Cousins $BVR_{\rm \scriptstyle C}I_{\rm \scriptstyle C}$ surface photometry, as well as contour maps and profiles of the SB, CI, ellipticity ($\epsilon$), and PA could be found in Paper\,I. The extra added galaxies were reduced as the rest of sample; observational details, contour maps, and profiles are presented in Appendix\,\ref{AppendixA}. The paper is structured as follows. In Sect.~\ref{glopar} we present the global values of the ellipticity, PA, inclination, and total magnitude. The isophotal parameters~-- SMA and CIs at 24 $V$ mag arcsec$^{-2}$, are discussed in Sect.~\ref{isopar}. The bar parameters are outlined in Sect.~\ref{barpar}. A summary of our results is given in Sect.\,\ref{summary}. A set of contour maps and profiles of the extra added galaxies is presented in Appendix\,\ref{AppendixA}. Global ellipticities and deprojected bar ellipticities of the matched inactive galaxies are given in Appendix\,\ref{AppendixB}. Throughout the paper the linear sizes in kpc have been calculated using the cosmology-corrected scale given in \linebreak NED\footnote{NASA/IPAC Extragalactic Database.} ($H_0$$=$73 km s$^{-1}$ Mpc$^{-1}$, $\Omega_{\,\rm M}$$=$0.27, $\Omega_{\,\Lambda}$$=$0.73,\linebreak Spergel et~al. \cite{SBD_07}). \begin{table}[ht] \caption{Global values of the ellipticity, PA, and inclination, estimated over all available passbands.} \label{T_glopar} \begin{center} \begin{tabular}{@{}l@{\hspace{0.2cm}}l@{\hspace{0.2cm}}l@{\hspace{0.3cm}}r@{\hspace{0.3cm}}l@{}} \hline \noalign{\smallskip} ~~Galaxy & Region$^{\rm a}$ & ~~Ellipticity & PA~~~~~~~ & Inclination \\ & (arcsec) & & (degree)~~~ & ~~(degree) \\ \noalign{\smallskip}\hline \noalign{\smallskip} Mrk\,335 & $ ~~10 $ & $0.09 \pm 0.02$ & $109.2 \pm 13.7$ & $25.0 \pm 3.0$ \\ III\,Zw\,2 & $ ~~12 $ & $0.17 \pm 0.02$ & $ 14.7 \pm~~4.7$ & $34.7 \pm 2.2$ \\ Mrk\,348 & $ ~~20\,(45)$ & $0.09 \pm 0.02$ & $ 90.1 \pm 41.3$ & $25.0 \pm 3.0$ \\ I\,Zw\,1 & $ ~~12 $ & $0.11 \pm 0.02$ & $139.0 \pm 10.9$ & $27.7 \pm 2.7$ \\ Mrk\,352 & $ ~~~~7 $ & $0.21 \pm 0.01$ & $ 84.6 \pm~~1.1$ & $38.7 \pm 1.0$ \\ Mrk\,573 & $ ~~30 $ & $0.13 \pm 0.02$ & $ 76.2 \pm~~6.8$ & $30.2 \pm 2.5$ \\ Mrk\,590 & $ ~~15 $ & $0.06 \pm 0.02$ & $127.4 \pm 14.5$ & $20.4 \pm 3.8$ \\ Mrk\,1040 & $ ~~20 $ & $0.73 \pm 0.03$ & $ 74.2 \pm~~1.1$ & $79.3 \pm 2.8$ \\ Mrk\,595 & $ ~~16 $ & $0.32 \pm 0.01$ & $ 92.3 \pm~~1.2$ & $48.4 \pm 0.8$ \\ 3C\,120 & $ ~~24 $ & $0.27 \pm 0.02$ & $115.0 \pm~~1.5$ & $44.2 \pm 1.8$ \\ Ark\,120 & $ ~~10 $ & $0.13 \pm 0.01$ & $ 8.5 \pm~~2.3$ & $30.2 \pm 1.2$ \\ Mrk\,376 & $ ~~11 $ & $0.29 \pm 0.02$ & $164.8 \pm~~2.5$ & $45.9 \pm 1.8$ \\ Mrk\,79 & $ ~~47 $ & $0.14 \pm 0.02$ & $148.6 \pm~~2.5$ & $31.4 \pm 2.4$ \\ Mrk\,382 & $ ~~17 $ & $0.16 \pm 0.02$ & $155.1 \pm~~6.8$ & $33.6 \pm 2.3$ \\ NGC\,3227 & $ ~~80 $ & $0.47 \pm 0.01$ & $150.1 \pm~~3.1$ & $59.9 \pm 0.8$ \\ NGC\,3516 & $ ~~35 $ & $0.19 \pm 0.01$ & $ 44.3 \pm~~4.5$ & $36.8 \pm 1.0$ \\ NGC\,4051 & $ 110 $ & $0.39 \pm 0.03$ & $109.3 \pm~~6.3$ & $54.0 \pm 2.4$ \\ NGC\,4151$^{\rm b}$&$~\ldots$&$0.07 \pm 0.03$& $ 26.0 \pm~~3.5$ & $21.0 \pm 5.0$ \\ Mrk\,766 & $ ~~24 $ & $0.19 \pm 0.03$ & $ 66.9 \pm~~4.1$ & $36.8 \pm 3.2$ \\ Mrk\,771 & $ ~~11 $ & $0.13 \pm 0.01$ & $ 80.2 \pm 11.7$ & $30.2 \pm 1.2$ \\ NGC\,4593 & $ ~~75 $ & $0.35 \pm 0.02$ & $ 70.6 \pm~~6.5$ & $50.9 \pm 1.6$ \\ Mrk\,279 & $ ~~18 $ & $0.33 \pm 0.01$ & $ 30.8 \pm~~2.3$ & $49.3 \pm 0.8$ \\ NGC\,5506 & $ ~~40 $ & $0.75 \pm 0.03$ & $ 87.8 \pm~~0.7$ & $81.2 \pm 2.8$ \\ NGC\,5548 & $ ~~35\,(60)$ & $0.19 \pm 0.01$ & $ 96.2 \pm~~6.0$ & $36.8 \pm 1.0$ \\ Ark\,479 & $ ~~10 $ & $0.31 \pm 0.02$ & $117.2 \pm~~1.5$ & $47.6 \pm 1.7$ \\ Mrk\,506 & $ ~~12 $ & $0.27 \pm 0.01$ & $116.3 \pm~~2.3$ & $44.2 \pm 0.9$ \\ Mrk\,507 & $ ~~~~5 $ & $0.30 \pm 0.03$ & $ 9.9 \pm~~ 5.8$ & $46.8 \pm 2.6$ \\ 3C\,382 & $ ~~15 $ & $0.24 \pm 0.03$ & $ 90.0 \pm~~3.3$ & $41.6 \pm 2.9$ \\ 3C\,390.3 & $ ~~~~8 $ & $0.12 \pm 0.02$ & $112.7 \pm~~9.9$ & $29.0 \pm 2.6$ \\ NGC\,6814 & $ ~~60 $ & $0.08 \pm 0.01$ & $ 92.2 \pm 17.8$ & $23.6 \pm 1.6$ \\ Mrk\,509 & $ ~~~~8 $ & $0.16 \pm 0.02$ & $ 72.4 \pm~~3.6$ & $33.6 \pm 2.3$ \\ Mrk\,1513 & $ ~~13 $ & $0.54 \pm 0.01$ & $ 57.9 \pm~~1.2$ & $65.0 \pm 0.7$ \\ Mrk\,304 & $ ~~~~6 $ & $0.03 \pm 0.01$ & $ 54.5 \pm 26.5$ & $14.4 \pm 2.6$ \\ Ark\,564 & $ ~~16 $ & $0.22 \pm 0.01$ & $112.2 \pm~~1.9$ & $39.7 \pm 1.0$ \\ NGC\,7469 & $ ~~50 $ & $0.26 \pm 0.01$ & $124.4 \pm~~1.5$ & $43.4 \pm 0.9$ \\ Mrk\,315 & $ ~~13 $ & $0.14 \pm 0.01$ & $ 33.2 \pm~~5.4$ & $31.4 \pm 1.2$ \\ NGC\,7603 & $ ~~18 $ & $0.32 \pm 0.01$ & $166.8 \pm~~2.9$ & $48.4 \pm 0.8$ \\ Mrk\,541 & $ ~~20 $ & $0.36 \pm 0.01$ & $172.3 \pm~~1.3$ & $51.6 \pm 0.8$ \\ \hline \end{tabular} \end{center} $^{\rm a}$ Start SMA, defining the region, over which the global values of the ellipticity and PA were estimated. The region generally extends to the profile end with two exceptions, for which the end SMA is given in parentheses.\\ $^{\rm b}$ The global parameters were taken from Simkin (\cite{S_75}). \end{table} \begin{table*}[t] \caption{Total magnitudes.} \label{t_totmag} \begin{center} \begin{tabular}{@{}llrrrr@{}} \hline \noalign{\smallskip} ~Galaxy & ~Civil Date & $B_{\rm tot}$~~~~~~ & $V_{\rm tot}$~~~~~~~~ & $R_{\rm C,tot}$~~~~~~\, & $I_{\rm C,tot}$~~~~~~~~ \\ & (yyyy\,mm\,dd) & (mag)~~~~~\, & (mag)~~~~~\, & (mag)~~~~~\, & (mag)~~~~~\, \\ \noalign{\smallskip} \hline \noalign{\smallskip} Mrk\,335 & 1998 08 22 & $14.19 \pm 0.03$ & $13.93 \pm 0.03$ & $13.47 \pm 0.03$ & $13.08 \pm 0.05$ \\ & 2007 08 20 & $ \ldots~~~~~~$ & $14.02 \pm 0.02$ & $13.59 \pm 0.02$ & $13.30 \pm 0.04$ \\ III\,Zw\,2& 1997 09 09 & $15.85 \pm 0.05$ & $14.86 \pm 0.04$ & $14.39 \pm 0.02$ & $13.79 \pm 0.05$ \\ Mrk\,348 & 1997 09 07 & $13.79 \pm 0.05$ & $12.90 \pm 0.04$ & $12.38 \pm 0.04$ & $11.86 \pm 0.04$ \\ I\,Zw\,1 & 1998 08 20 & $14.48 \pm 0.05$ & $14.06 \pm 0.05$ & $13.64 \pm 0.05$ & $13.15 \pm 0.06$ \\ Mrk\,352 & 2007 08 21 & $ \ldots~~~~~~$ & $ \ldots~~~~~~$ & $13.82 \pm 0.01$ & $ \ldots~~~~~~$ \\ & 2008 01 30 & $15.38 \pm 0.01$ & $14.57 \pm 0.02$ & $13.99 \pm 0.02$ & $13.42 \pm 0.03$ \\ Mrk\,573 & 1997 09 07 & $13.97 \pm 0.06$ & $13.17 \pm 0.04$ & $12.67 \pm 0.04$ & $12.07 \pm 0.05$ \\ Mrk\,590 & 1997 09 06 & $13.23 \pm 0.03$ & $12.52 \pm 0.03$ & $11.98 \pm 0.03$ & $11.28 \pm 0.04$ \\ Mrk\,1040 & 1997 09 06 & $13.65 \pm 0.03$ & $12.80 \pm 0.03$ & $12.19 \pm 0.03$ & $11.73 \pm 0.04$ \\ Mrk\,595 & 1997 09 09 & $15.12 \pm 0.05$ & $14.29 \pm 0.03$ & $13.61 \pm 0.03$ & $12.92 \pm 0.05$ \\ 3C\,120 & 1997 09 09 & $14.55 \pm 0.04$ & $13.86 \pm 0.03$ & $13.19 \pm 0.03$ & $12.51 \pm 0.05$ \\ & 2008 02 01 & $14.34 \pm 0.04$ & $13.75 \pm 0.03$ & $13.24 \pm 0.03$ & $12.69 \pm 0.02$ \\ Ark\,120 & 1994 09 29 & $13.97 \pm 0.07$ & $ \ldots~~~~~~$ & $12.89 \pm 0.03$ & $ \ldots~~~~~~$ \\ & 1991 12 08 & $ \ldots~~~~~~$ & $13.45 \pm 0.05$ & $ \ldots~~~~~~$ & $ \ldots~~~~~~$ \\ Mrk\,376 & 2008 02 03 & $15.21 \pm 0.05$ & $14.62 \pm 0.07$ & $13.86 \pm 0.03$ & $13.33 \pm 0.04$ \\ Mrk\,79 & 1999 02 16 & $13.72 \pm 0.03$ & $13.08 \pm 0.03$ & $12.48 \pm 0.03$ & $11.91 \pm 0.05$ \\ & 2008 02 01 & $13.85 \pm 0.03$ & $13.14 \pm 0.03$ & $12.65 \pm 0.04$ & $12.12 \pm 0.04$ \\ Mrk\,382 & 1998 02 27 & $15.30 \pm 0.03$ & $14.63 \pm 0.03$ & $14.09 \pm 0.04$ & $13.56 \pm 0.05$ \\ & 2008 02 02 & $15.21 \pm 0.03$ & $14.63 \pm 0.03$ & $14.11 \pm 0.03$ & $13.58 \pm 0.03$ \\ NGC\,3227 & 1999 04 17 & $11.62 \pm 0.07$ & $10.84 \pm 0.05$ & $10.23 \pm 0.06$ & $ 9.43 \pm 0.06$ \\ NGC\,3516 & 2008 01 08 & $ \ldots~~~~~~$ & $11.65 \pm 0.01$ & $11.13 \pm 0.01$ & $10.55 \pm 0.01$ \\ NGC\,4051 & 1995 05 06 & $11.09 \pm 0.01$ & $ \ldots~~~~~~$ & $ \ldots~~~~~~$ & $ \ldots~~~~~~$ \\ & 2000 03 30 & $ \ldots~~~~~~$ & $ \ldots~~~~~~$ & $ 9.78 \pm 0.02$ & $ \ldots~~~~~~$ \\ & 2001 04 09 & $ \ldots~~~~~~$ & $ \ldots~~~~~~$ & $ \ldots~~~~~~$ & $ 9.18 \pm 0.02$ \\ NGC\,4151 & 1999 03 10 & $11.13 \pm 0.03$ & $10.52 \pm 0.04$ & $10.00 \pm 0.04$ & $ 9.33 \pm 0.05$ \\ & 1999 04 19 & $11.25 \pm 0.05$ & $10.56 \pm 0.04$ & $10.05 \pm 0.04$ & $ 9.45 \pm 0.05$ \\ Mrk\,766 & 1999 02 15 & $13.83 \pm 0.07$ & $13.10 \pm 0.03$ & $12.58 \pm 0.04$ & $11.97 \pm 0.06$ \\ Mrk\,771 & 1990 06 23 & $ \ldots~~~~~~$ & $14.37 \pm 0.03$ & $ \ldots~~~~~~$ & $13.11 \pm 0.01$ \\ NGC\,4593 & 2008 01 08 & $ \ldots~~~~~~$ & $11.11 \pm 0.01$ & $10.59 \pm 0.01$ & $10.00 \pm 0.01$ \\ Mrk\,279 & 2008 02 02 & $14.43 \pm 0.04$ & $13.76 \pm 0.03$ & $13.23 \pm 0.04$ & $12.79 \pm 0.06$ \\ NGC\,5506 & 1999 04 17 & $12.59 \pm 0.07$ & $11.91 \pm 0.05$ & $11.16 \pm 0.06$ & $10.34 \pm 0.06$ \\ NGC\,5548 & 1999 04 19 & $12.76 \pm 0.05$ & $12.27 \pm 0.04$ & $11.79 \pm 0.05$ & $11.18 \pm 0.06$ \\ Ark\,479 & 2007 07 19 & $ \ldots~~~~~~$ & $14.62 \pm 0.06$ & $14.05 \pm 0.07$ & $13.48 \pm 0.08$ \\ Mrk\,506 & 1997 06 01 & $15.09 \pm 0.04$ & $14.29 \pm 0.04$ & $13.76 \pm 0.05$ & $13.01 \pm 0.05$ \\ & 1998 07 18 & $15.01 \pm 0.04$ & $14.20 \pm 0.02$ & $13.62 \pm 0.03$ & $12.99 \pm 0.05$ \\ & 2007 06 17 & $15.38 \pm 0.03$ & $14.33 \pm 0.02$ & $13.82 \pm 0.03$ & $ \ldots~~~~~~$ \\ Mrk\,507 & 1998 07 20 & $16.77 \pm 0.06$ & $15.85 \pm 0.05$ & $15.21 \pm 0.05$ & $14.23 \pm 0.06$ \\ 3C\,382 & 1998 08 23 & $14.18 \pm 0.04$ & $13.85 \pm 0.04$ & $13.31 \pm 0.04$ & $12.78 \pm 0.05$ \\ 3C\,390.3 & 1998 08 20 & $15.93 \pm 0.06$ & $15.05 \pm 0.04$ & $14.40 \pm 0.04$ & $13.93 \pm 0.05$ \\ NGC\,6814 & 1997 07 06 & $12.16 \pm 0.02$ & $11.30 \pm 0.02$ & $10.37 \pm 0.02$ & $ 9.62 \pm 0.03$ \\ & 1997 07 10 & $12.30 \pm 0.03$ & $11.36 \pm 0.02$ & $10.56 \pm 0.03$ & $ 9.69 \pm 0.04$ \\ & 1997 09 07 & $12.14 \pm 0.05$ & $11.10 \pm 0.03$ & $10.36 \pm 0.03$ & $ 9.49 \pm 0.04$ \\ & 1998 07 18 & $12.28 \pm 0.04$ & $11.26 \pm 0.02$ & $10.54 \pm 0.03$ & $ 9.79 \pm 0.05$ \\ Mrk\,509 & 1997 07 10 & $14.09 \pm 0.03$ & $13.59 \pm 0.02$ & $13.08 \pm 0.03$ & $12.65 \pm 0.04$ \\ & 1997 09 08 & $13.89 \pm 0.03$ & $13.48 \pm 0.02$ & $12.99 \pm 0.03$ & $12.69 \pm 0.04$ \\ & 1998 07 20 & $13.40 \pm 0.05$ & $13.16 \pm 0.04$ & $12.70 \pm 0.04$ & $12.30 \pm 0.05$ \\ Mrk\,1513 & 2007 08 20 & $ \ldots~~~~~~$ & $14.54 \pm 0.02$ & $14.08 \pm 0.02$ & $13.79 \pm 0.01$ \\ Mrk\,304 & 1998 07 19 & $14.61 \pm 0.04$ & $14.31 \pm 0.03$ & $13.82 \pm 0.03$ & $13.50 \pm 0.04$ \\ Ark\,564 & 1998 07 18 & $14.22 \pm 0.04$ & $13.71 \pm 0.02$ & $13.27 \pm 0.03$ & $12.97 \pm 0.04$ \\ & 1998 08 20 & $14.34 \pm 0.05$ & $13.82 \pm 0.04$ & $13.37 \pm 0.04$ & $12.90 \pm 0.05$ \\ NGC\,7469 & 1997 09 06 & $12.73 \pm 0.03$ & $12.17 \pm 0.03$ & $11.67 \pm 0.02$ & $11.01 \pm 0.04$ \\ & 1998 07 19 & $12.70 \pm 0.04$ & $12.18 \pm 0.03$ & $11.77 \pm 0.03$ & $11.11 \pm 0.04$ \\ & 1998 08 23 & $12.87 \pm 0.04$ & $12.40 \pm 0.04$ & $11.80 \pm 0.04$ & $11.05 \pm 0.05$ \\ & 2003 07 28 & $13.33 \pm 0.03$ & $12.42 \pm 0.03$ & $11.75 \pm 0.02$ & $11.45 \pm 0.03$ \\ Mrk\,315 & 2007 08 22 & $ \ldots~~~~~~$ & $ \ldots~~~~~~$ & $13.89 \pm 0.03$ & $ \ldots~~~~~~$ \\ NGC\,7603 & 2007 07 19 & $ \ldots~~~~~~$ & $12.47 \pm 0.07$ & $ \ldots~~~~~~$ & $11.51 \pm 0.08$ \\ Mrk\,541 & 2007 07 19 & $ \ldots~~~~~~$ & $14.73 \pm 0.07$ & $14.12 \pm 0.07$ & $13.52 \pm 0.08$ \\ \hline \end{tabular} \end{center} \end{table*} \begin{figure*}[t] \centering \begin{minipage}[t]{5cm} \includegraphics[width=5cm]{AN_2201_fig1.eps} \caption{Comparison between the ellipticities listed in HyperLeda and ours. Named are the galaxies with $|\Delta\epsilon|\,$$>$$\,0.3$. The line of exact correspondence is plotted.} \label{E} \end{minipage} \hspace{0.5cm} \begin{minipage}[t]{5cm} \includegraphics[width=5cm]{AN_2201_fig2.eps} \caption{Comparison between the PAs listed in HyperLeda and ours. Named are the galaxies with $|\Delta\rm PA|\,$$>$$\,30\degr$. The line of exact correspondence is plotted.} \label{PA} \end{minipage} \end{figure*} \begin{figure*}[t] \centering \begin{minipage}[t]{5cm} \includegraphics[width=5cm]{AN_2201_fig3.eps} \caption{ Comparison between $a_{24}$ estimated by us and those published in Hunt et~al. (\cite{HMR2_99}). Named are the outliers. The line of exact correspondence is plotted.} \label{comp_isosiz} \end{minipage} \hspace{0.5cm} \begin{minipage}[t]{5cm} \includegraphics[width=5cm]{AN_2201_fig4.eps} \caption{Distribution of $(B$$-$$I_{\rm \scriptstyle C})_{24}^{(0)}$ (solid) and $(V$$-$$I_{\rm \scriptstyle C})_{24}^{(0)}\,$ (dashed), whose median values are denoted by the right and left arrow, respectively.} \label{col_hist} \end{minipage} \hspace{0.5cm} \begin{minipage}[t]{5cm} \includegraphics[width=5cm]{AN_2201_fig5.eps} \caption{Distribution of $(V$$-$$I_{\rm \scriptstyle C})_{24}^{(0)}$ for the galaxies with/without outer rings (solid/dashed) with the median values denoted by the left/right arrow.} \label{comp_isosize} \end{minipage} \end{figure*}

\label{summary} This paper is third in a series, studying the optical properties of a sample of Seyfert galaxies. The first paper addresses the evidence of non-axisymmetric perturbation of the potential in a sample of 35 Seyfert galaxies and in a matched inactive sample. A homogeneous set of global (ellipticity, PA, inclination, and total magnitude) and isophotal (SMA and CIs at 24 $V$ mag arcsec$^{-2}$) parameters of the Seyfert sample are reported in this study. Correction for galactic absorption, inclination and internal absorption, cosmological dimming, as well as $K$- and $E$-correction was applied to the isophotal parameters. We found the following median isophotal parameters: \begin{eqnarray} a_{24}^{\rm (0)}&=&13.9\,\rm kpc, \nonumber \\ (B-I_{\rm \scriptstyle C})_{24}&=&2.2\,\rm mag\,arcsec^{-2}, \nonumber \\ (V-I_{\rm \scriptstyle C})_{24}&=&1.3\,\rm mag\,arcsec^{-2}, \nonumber \\ (B-I_{\rm \scriptstyle C})_{24}^{(0)}&=&1.9\,\rm mag\,arcsec^{-2}, \nonumber \\ (V-I_{\rm \scriptstyle C})_{24}^{(0)}&=&1.1\,\rm mag\,arcsec^{-2}. \nonumber \end{eqnarray} The estimated parameters can be further used in various galactic structure studies. We presented a set of bar parameters~-- ellipticity, PA, SMA corresponding to the ellipticity maximum in the bar region, and length; deprojected values of the bar ellipticity, length, and relative length in terms of galaxy isophotal SMA are also given. As bar length we adopted the minimum of the SMAs corresponding to 15\% ellipticity decrease from its maximal value, both before and after the ellipticity maximum. The so obtained bar length and the most often used bar length estimate~-- the SMA, corresponding to the ellipticity maximum, show a tight correlation with a median ratio $\ell/\ell_{\rm max}$ of 1.22, which we further used to obtain the bar length in cases the above approach did not work. The median of the deprojected bar ellipticity, length, and relative length are 0.39, 5.44\,kpc, and 0.44, respectively. The deprojected bar length correlates with the corrected isophotal SMA at 24 $V$ mag arcsec$^{-2}$. Seventeen of the galaxies have large-scale bars, three of which are strong, based on the deprojected bar ellipticity as a rough estimate of bar strength. The deprojected relative bar length and bar ellipticity show no clear correlation. Global ellipticities and deprojected bar ellipticities of the matched inactive sample are also presented.

1012.5741

This paper is third in a series, studying the optical properties of a sample of Seyfert galaxies. Here we present a homogeneous set of global (ellipticity, position angle, inclination, and total magnitude) and isophotal (semi-major axis and colour indices at 24 V mag arcsec<SUP>-2</SUP>) parameters of the galaxy sample. We find the following median corrected isophotal colour indices: {(B-I_C)<SUB>24</SUB><SUP>(0)</SUP>=1.9} mag arcsec<SUP>-2</SUP> and {(V-I_C)<SUB>24</SUB><SUP>(0)</SUP>=1.1} mag arcsec<SUP>-2</SUP>. A set of bar parameters (ellipticity, position angle, semi-major axis corresponding to the ellipticity maximum in the bar region, and length) are also reported; deprojection has been applied to the bar ellipticity, length, and relative length in terms of galaxy isophotal semi-major axis. Regarding bar length estimation, we use a method, based on the relation between the behaviour of the profiles and orbit analysis. The so estimated bar length tightly correlates with the semi-major axis, corresponding to the ellipticity maximum with a median ratio of the former to the latter of 1.22. The median of the deprojected bar ellipticity, length, and relative length are 0.39, 5.44 kpc, and 0.44, respectively. There is a correlation between the deprojected bar length and the corrected isophotal semi-major axis at 24 V mag arcsec<SUP>-2</SUP>. Three of the 17 large-scale bars appear strong, based on the deprojected bar ellipticity as a first-order approximation of bar strength. The deprojected relative bar length does not appear to correlate with the bar ellipticity. <P />Based on observations obtained with the 2-m telescope of the Institute of Astronomy and National Astronomical Observatory, Bulgarian Academy of Sciences.

12167346

2011ApJ...726L...3A

1012

1012.0022_arXiv.txt

Grain growth, crystallization, and settling occur in primordial circumstellar disks around both stellar and substellar objects. These processes produce several observational signatures. The mid-infrared (IR) silicate emission becomes weaker as the grains in the upper layers grow and settle towards the midplane while crystallization modifies the shape of the band. IR continuum emission is reduced at longer wavelengths with grain growth and settling because the intercepted stellar flux decreases as the disk flattens. In the substellar domain, these hallmarks of dust evolution have been reported previously for disks at ages $\tau\sim1-10$~Myr \citep{Apai04,Apai05,Scholz07,Morrow08,Pascucci09,Riaz09,RLG09}. For instance, disks around the youngest brown dwarfs exhibit evidence of small, crystalline silicate grains \citep[][]{Apai05,Riaz09} while older brown dwarf disks have weaker silicate emission that indicates the presence of large grains \citep{Scholz07,Morrow08,RLG09}. The evolution of grains in the inner annuli of these disks, e.g., the zone where the 10~\micron\ silicate band emerges \citep[within $10^{-3}-0.1$~AU,][]{KS07}, appears to occur more rapidly than in disks around solar-mass stars. One possible explanation for this phenomena is that the low levels of turbulence for brown dwarf disks produce lower mass accretion rates and reduce the replenishment of the inner disk with primordial grains from outer annuli \citep{SA07}. Also, the collisional rate between large size grains would be small, preventing the fragmentation to small size grains. During a mid-IR spectroscopic survey of young brown dwarfs with the {\it Spitzer Space Telescope} \citep{wer04}, we have identified a brown dwarf disk that shows evidence of advanced grain evolution at a relatively young age of $\sim1$~Myr. This brown dwarf, \object{2MASS J04414489+2301513} \citep[][henceforth 2M~J04414489]{Luhman06}, lies in the Taurus star-forming region and was recently shown to harbor a companion with a projected separation of $0\farcs105$ ($\sim15$~AU) and a mass of 5--10~$M_{\rm Jup}$ \citep{TLM10}. In this Letter, we use our mid-IR spectrum of this system and our models of settled, irradiated accretion disks to constrain the properties of the circumstellar disk that resides around the primary.

We have presented a mid-IR spectrum of the binary brown dwarf 2M~J04414489 in the Taurus star-forming region. The spectrum exhibits excess emission that we attribute to a circumprimary disk, but silicate emission at 10~\micron\ is not present, indicating that significant grain growth has occurred, leading to a disk on which small grains no longer exist ($5\lesssim a_{max}<1000$~\micron). This is one of $\sim2$ brown dwarf disks at such a young age ($\sim1$~Myr) that have been found to lack silicate emission. Our models of the SED of 2M~J04414489 suggest that the outer radius of its disk is rather small ($R_d\sim0.2$--0.3~AU), which may reflect truncation by the binary companion. The absence of an outer disk with a reservoir of primordial grains, combined with weak turbulence, could explain the advanced grain growth in the inner disk of 2M~J04414489.

1012.0022

Using the Spitzer Infrared Spectrograph, we have performed mid-infrared spectroscopy on the young binary brown dwarf 2MASS J04414489+2301513 (15 AU) in the Taurus star-forming region. The spectrum exhibits excess continuum emission that likely arises from a circumstellar disk around the primary. Silicate emission is not detected in these data, indicating the presence of significant grain growth. This is one of the few brown dwarf disks at such a young age (~1 Myr) that has been found to lack silicate emission. To quantitatively constrain the properties of the disk, we have compared the spectral energy distribution of 2MASS J04414489+2301513 to the predictions of our vertical structure codes for irradiated accretion disks. Our models suggest that the remaining atmospheric grains of moderately depleted layers may have grown to a size of gsim5 μm. In addition, our model fits indicate an outer radius of 0.2-0.3 AU for the disk. The small size of this circumprimary disk could be due to truncation by the secondary. The absence of an outer disk containing a reservoir of small, primordial grains, combined with a weak turbulent mechanism, may be responsible for the advanced grain growth in this disk. <P />Based on observations made with the Spitzer Space Telescope, which is operated by the Jet Propulsion Laboratory, California Institute of Technology under a contract with NASA.

12167991

2011ApJ...727L..17Z

1012

1012.2211_arXiv.txt

As the early surveys revealed \citep{po_87}, Wolf-Rayet (WR) stars are sources of X-ray emission and massive WR$+$OB binaries are the brightest amongst them. Pointed X-ray observations with modern observatories ({\it Chandra, XMM-Newton}) confirmed this and provided us with grating spectra of WR$+$OB binaries (\citealt{sk_01}; \citealt{raa_03}; \citealt{schi_04}; \citealt{po_05}; \citealt{zhp_10b}) that are rich in spectral lines from different ionic species which indicates a range of X-ray plasma temperatures. These findings suggest that the enhanced X-ray emission is produced in colliding-stellar-wind (CSW) shocks, as first proposed by \citet{pril_76} and \citet{cherep_76}. The available pointed {\it Chandra {\rm and} XMM-Newton} observations suggest that the X-ray properties of presumably single WR stars might correlate with their subtype. Namely, the WN (nitrogen-rich) objects have X-ray spectra with prominent emission lines originating from an admixture of cool (kT~$< 1$~keV) and hot (kT~$> 2$~keV) plasma \citep{sk_10}. Only one of the four WO (oxygen-rich) stars in the Galaxy, WR 142, has been observed so far and it is a weak but extremely hard X-ray source (\citealt{os_09}; \citealt{sokal_10}). In contrast, the single WC (carbon-rich) stars are conspicuously X-ray quiet: all the {\it Chandra} and {\it XMM-Newton} pointed observations of WCs so far have resulted in non-detections (\citealt{os_03}; \citealt{sk_06}; \citealt{sk_09}). The X-ray production mechanism in single WR stars is still not well-understood, but the ubiquitous presence of a high-temperature X-ray component in WR spectra is a clear indication that other processes besides embedded, radiative wind shocks are contributing (\citealt{sk_10} and references therein). Overall, binary systems, especially wide binaries, are more X-ray luminous than comparable single WR stars. For single WR stars $L_X$ is typically in the range $10^{31} - 10^{33}$~ergs s$^{-1}$ \citep{sk_10}, while for the known wide WR$+$OB binaries it is about an order of magnitude higher, with WR~140 having the largest known luminosity of $(1-2)\times10^{34}$~ergs s$^{-1}$ \citep{po_05}. As a rule, if a WR star is detected with $L_{\rm X} > 10^{33}$~ergs s$^{-1}$, it is most likely a wide CSW binary system. In this {\it Letter}, we report the first X-ray detection of the Wolf-Rayet star \WR, which establishes it as a very luminous X-ray source and as such, a likely multiple system.

\label{sec:discussion} The two most important results from the model fits to the {\it XMM-Newton} spectra of \WR are that it is a very luminous X-ray source and thermal plasmas with high temperature dominate its emission (Table~\ref{tab:fits} and Fig.~\ref{fig:dem}). Using the unabsorbed flux values from Table~\ref{tab:fits}, the X-ray luminosity of \WR is L$_X (0.5 - 10~\mbox{keV}) = (0.5 - 2.1)\times10^{34} d^2_{kpc}$ ergs s$^{-1}$, where $d_{kpc}$ is the distance in units of kpc. The upper limit is from the NEI shock model which gives a better quality of fit. The distance to \WR is not tightly constrained and a range of values is available in the literature: $ d_{kpc} = 1.21 - 4$ (\S~\ref{sec:thestar}). But, based on its proximity to the open clusters Danks 1 and 2 and the similar interstellar extinction (e.g., \citealt{danks_83}; \citealt{clark_04}; \citealt{baume_09}), the distance of 4 kpc seems more realistic, which puts the X-ray luminosity in the range L$_X = (0.8 - 3.4)\times10^{35}$ ergs s$^{-1}$. In this case and excluding the most X-ray luminous WR star, Cyg X-3, which is a WR binary with a compact companion (a neutron star or a black hole), \WR is the most luminous WR star in the Galaxy among those observed so far. For example, the X-ray luminosity of the brightest CSW binary, WR 140, is: $(0.5 - 3)\times10^{34}$ ergs s$^{-1}$ (\citealt{zhsk_00}; \citealt{po_05}). Adopting the bolometric luminosity from \citet{clark_04}, we note the very high value of $\lg L_X/L_{bol} = [-4.3, -3.7]$ for \WRE. All this raises the interesting question about the X-ray production mechanism in \WRE. {\it Colliding Stellar Winds.} As already mentioned, \WR is an episodic dust-maker \citep{williams_95} which suggests that CSWs in a wide binary system with an eccentric orbit might provide most of its X-ray emission, as is the case for the prototype episodic dust-maker amongst WR stars, WR 140 \citep{williams_90}. Unfortunately, the lack of information about the stellar wind parameters of \WR did not allow us to make a comparison between the theoretical predictions based on hydrodynamic CSW modeling and observations. On the other hand, some simple (e.g. {\it qualitative}) considerations are possible. The temperature of the hot plasma component deduced from the global spectral fits (\S \ref{sec:global}) can provide an estimate of the stellar wind velocity (see \S 5.2 in \citealt{zh_07} for discussion of CSW models versus discrete-temperature models ). For the case of typical WC abundances, the postshock plasma temperature is kT~$ = 3.09 ~V_{1000}^2$~keV, where $V_{1000}$ is the shock velocity in units of 1000\kms. From Table~\ref{tab:fits} and Fig.~\ref{fig:dem}, we see that the stellar wind of the WC star in \WR must have a velocity V$_{wind} \geq 1000$\kms. Such high velocities would be expected for WC8-9 stars since they have typical average wind speeds 1400\kms \citep{prinja_90}, 1300\kms \citep{eenens_94}. We note that the derived values for ionization age of the shocks in the two-temperature model (Table~\ref{tab:fits}) and in the model with a distribution of NEI shocks (Fig.~\ref{fig:dem}) are also qualitatively consistent with the CSW picture. Namely, the higher density and higher temperature plasma is located near the axis of symmetry (the line-of-centers between the two stars) while the less dense and cooler plasma is found downstream from that axis. Thus, it is natural in the CSW picture that the higher temperature plasma will have a larger ionization age. We can use the scaling law for the CSW X-ray luminosity with the mass-loss rate ($\dot{M}$), wind velocity ($v$) and binary separation ($D$): $L_X \propto \dot{M}^2 v^{-3} D^{-1}$ (\citealt{luo_90}; \citealt{mzh_93}) for a comparison between \WR and the classical CSW binary WR 140. Adopting the WR 140 binary parameters (\citealt{williams_90}; \citealt{po_05}), a wind velocity of $1300-1400$ \kms in \WR and assuming that its mass loss is equal to that of WR 140, we see that the binary separation in \WR could be similar to that in WR 140 even when the \WR X-ray luminosity is about an order of magnitude larger than that of WR 140 (see above). Therefore, the CSW picture in a wide binary system is at least qualitatively consistent with the observational data for \WRE. {\it Magnetically Confined Wind Shocks (MCWS).} This mechanism is capable of producing hard X-ray emission. It was proposed to explain the high plasma temperature found in X-ray emission from young massive stars and requires the presence of a relatively strong magnetic field (\citealt{sch_03}, \citealt{gagne_05} and references therein). Then, could it be that the WC star in \WR is a magnetized object or has a close magnetized companion? We note that the X-ray luminosity of the prototype MCWS object, $\theta^1$~ Ori C, is L$_X (0.5 - 10~\mbox{keV}) \approx 10^{33}$ ergs s$^{-1}$ \citep{gagne_05} which is one to two orders of magnitude below that of \WRE. Thus, we can likely rule out the case of a close magnetized companion. On the other hand, the MCWS mechanism does not seem promising if adopted directly to the WC star in \WRE. First, in order to be efficient this mechanism requires a relatively strong global magnetic field to confine the massive WR wind and no such fields have been reported at present for these stars. Second, because of the expected decay of the magnetic field strength with the age of a massive star, the MCWS mechanism is associated only with young massive stars (age $\leq 1$~Myr; \citealt{sch_03}). Thus, the likely association of \WR with the open clusters Danks 1, 2 and the recent estimate of $\sim 5$~Myr for their age \citep{baume_09} present serious difficulties for the MCWS model in the case of \WRE. {\it Wind Accretion Shocks.} The presence of a close degenerate companion (e.g. a neutron star) seems very intriguing and accretion onto such an object can in general provide high X-ray luminosities (e.g., \citealt{do_73}). Details depend on the actual binary parameters and wind properties of the main stellar component which are yet unknown for \WRE. But, this case may not be unlikely if \WR is indeed a member of the 5-Myr-old open clusters Danks 1, 2 and if it once had a more massive close companion which has already evolved and exploded as a supernova. We have listed above some physical mechanisms that might be able to explain the observed X-ray properties of \WRE. One could continue this list by considering other mechanisms or even more complex combinations of those mentioned above. But, all such mechanisms and the corresponding physical picture remain quite speculative because of the scarcity of detailed observational data for \WR. Thus, we will end our discussion by simply mentioning the cases that look least speculative to us and briefly discuss what observational information may help distinguish between them. For the moment, the following cases are the most likely explanation for the X-ray properties of \WRE: (i) CSWs in a wide binary system; or (ii) accretion wind shocks in a close binary: WR$+$compact companion (a neutron star). Variability is a key parameter for distinguishing between these cases. Namely, if long-term (months, years) X-ray variability is established from future observations, then case (i) would be strongly favored. Furthermore, case (i) would be strengthened if variable non-thermal radio (NTR) emission from \WR is eventually discovered and even more so if the X-ray emission is found to modulate at a similar period. In general, NTR WRs are commonly associated with wide binary systems (\citealt{do_00}). Alternatively, if short-term (a few days) variability is established or a sudden change of the X-ray luminosity is detected, case (ii) would be favored. The {\it XMM-Newton} data show no significant variability on timescales $\le 1$~day; the Kolmogorov-Smirnov test rules out variability at the 95\% confidence level. On the other hand, based on the infrared variability of \WRE, \citet{williams_03} proposed that it is a triple stellar system. If so, we may expect some very long-period variability due to the X-rays from CSWs in a wide binary with eccentric orbit and some with a much shorter period that could result from the stellar wind shocking onto a non-degenerate companion. In this case, WR 48a would be very similar to the CSW binary WR 147 which was recently resolved into a double X-ray source with {\em Chandra}, of which one component is variable and possibly an unresolved binary (\citealt{zhp_10a},b). However, WR 48a has a much higher X-ray luminosity. Finally, we hope that the unusually high L$_{X}$ of \WR will motivate deeper follow-up observations across the entire spectrum (radio, infrared, optical, UV, X-rays). Such observations would provide us with valuable information about the star and its wind properties that could be very helpful for constraining the physical picture in this remarkable but understudied Wolf-Rayet system.

1012.2211

We present XMM-Newton observations of the dusty Wolf-Rayet (W-R) star WR 48a. This is the first detection of this object in X-rays. The XMM-Newton EPIC spectra are heavily absorbed and the presence of numerous strong emission lines indicates a thermal origin of the WR 48a X-ray emission, with dominant temperature components at kT <SUB>cool</SUB> ≈ 1 keV and kT <SUB>hot</SUB> ≈ 3 keV, the hotter component dominating the observed flux. No significant X-ray variability was detected on timescales &lt;=1 day. Although the distance to WR 48a is uncertain, if it is physically associated with the open clusters Danks 1 and 2 at d ~4 kpc, then the resultant X-ray luminosity L <SUB> X </SUB>~ 10<SUP>35</SUP> erg s<SUP>-1</SUP> makes it the most X-ray luminous W-R star in the Galaxy detected so far, after the black hole candidate Cyg X-3. We assume the following scenarios as the most likely explanation for the X-ray properties of WR 48a: (1) colliding stellar winds in a wide WR+O binary system, or in a hierarchical triple system with non-degenerate stellar components and (2) accretion shocks from the WR 48a wind onto a close companion (possibly a neutron star). More specific information about WR 48a and its wind properties will be needed to distinguish between the above possibilities.

12224503

2011PhRvD..83c2010W

1012

1012.3021_arXiv.txt

\noindent Since the early days of solar neutrino flux measurements, efforts have been made to discover a variation of the neutrino rate over time. The most prominent effect is the experimentally confirmed annual variation of the flux by 7\,\% that is induced by the eccentricity of the Earth's orbit (e.\,g.\,\cite{ran07a,sk08sol}). But a variety of alternative sources of flux modulation is conceivable: The survival probability of solar electron neutrinos might be influenced by fluctuations of the solar matter density \cite{mir01,akh02,cha05} or by traversing terrestrial matter \cite{car86,bah97,deg01} before reaching the detector on Earth. Solar neutrino production rates might even change in the course of the solar cycle of about 11 years \cite{kra90}, or might be subject to short-term variations correlated to the oscillation of the solar core temperature induced by helioseismic waves \cite{sno09}. The next-generation neutrino observatory LENA will be a low-background liquid-scintillator experiment \cite{mar06,mar08,wur10det}. In recent years, the potential of this detector technique has been demonstrated by the first real-time measurement of {$^7$Be} neutrinos in the solar neutrino experiment Borexino \cite{bx07be7}. Opposed to the comparatively low rate of $\sim$50 counts per day (cpd) in Borexino, about $10^4$\,cpd of {$^7$Be} neutrino events are expected in LENA due to the considerably larger target mass (Sect.\,\ref{SecNeuRat}). Given this large statistics, it is clear that LENA will be much more sensitive to temporal flux fluctuations as any of the preceding solar neutrino experiments, including the Super-Kamiokande detector \cite{sk05sm}. In the following analysis based on the Lomb-Scargle method \cite{lom76,sca82} (Sect.\,\ref{SecLomSca}), we demonstrate that LENA will be sensitive to fluctuations on a sub-percent level over a wide range of frequencies: Studies of the sensitivity cover modulation periods from several minutes up to tens of years (Sect.\,\ref{SecAnaRes}). The resulting discovery potential for a number of plausible modulation sources will be presented in Sect.\,\ref{SecDiscus}.

The next generation of large-volume, low-energy neutrino detectors will enhance the experimental sensitivity towards modulations in the solar neutrino flux \cite{tur10}. A liquid-scintillator detector seems especially attractive due to the possibility to observe the {$^7$Be} neutrino line at 866\,keV: As the expected event rate is of the order of $10^4$ counts per day, the high statistics will allow to search for modulations on a sub-percent level, by far surpassing the sensitivity of currently running experiments. The sensitivity of the search stretches from very short time scales of the order of tens of minutes to tens or even hundreds of years. This will allow to probe the frequency regions of the helioseismic observation for g-mode oscillations in the solar center \cite{app09}, the day-night effect \cite{bah97,deg01}, and to search for variations of the fusion rate with the solar cycle.

1012.3021

A next-generation liquid-scintillator detector will be able to perform high-statistics measurements of the solar neutrino flux. In LENA, solar Be7 neutrinos are expected to cause 1.7×10<SUP>4</SUP> electron recoil events per day in a fiducial volume of 35 kilotons. Based on this signal, a search for periodic modulations on a subpercent level can be conducted, surpassing the sensitivity of current detectors by at least a factor of 20. The range of accessible periods reaches from several minutes, corresponding to modulations induced by helioseismic g-modes, to tens of years, allowing to study long-term changes in solar fusion rates.

1314489

2011ASPC..448..625G

1012

1012.1354_arXiv.txt

Star formation occurs pre-dominantly in dense and cool molecular clouds, which collapse under their own gravity and fragment into smaller pieces. In our galaxy we observe those clouds over a range of masses, but it now seems that the majority of stars form in clusters with more than a hundred members \citep{2000AJ....120.3139C,2003ARA&A..41...57L,2003AJ....126.1916P}. Thus, studying star formation in those clusters is an important step to understand the history of the stars we currently observe. In the early stages the proto-star is still hidden by its parent molecular cloud and thus can only be observed as far-IR emission from the warm, dusty envelope (class 0). Due to the conservation of angular momentum the matter does not collapse radially onto the star, but forms an accretion disk in class I objects. They often drive powerful outflows and carve holes in their envelopes so radiation can escape. Eventually, the envelopes disperse and the stars become visible as classical T~Tauri stars (CTTS), or class II sources in IR classification. Their disks still cause an IR excess, and on the stars accretion shocks and coronal activity lead to X-ray emission. Later, the IR excess vanishes, because the disk mass decreases. In this stage (class III or weak-line T~Tauri stars = WTTS) cluster members cannot be distinguished from main-sequence (MS) stars by IR observations; one method to identify them is through their X-ray luminosity, which is far higher than for older stars \citep[for a review see][]{1999ARA&A..37..363F}. In dense clusters, stars may influence one another during their evolution. Most notably strong winds or strong radiation fields of early-type stars can evaporate the disks of their late-type neighbors. In the Orion nebula this process can be directly observed in distorted disks, called proplyds \citep[e.g.][]{1994ApJ...436..194O}. Despite significant observational progress many details of the star formation process in clusters are still under debate. The variety of competing processes makes it difficult to quantify the different contributions. For example, it is unclear if early or late-type stars form first in an undisturbed cluster. The theoretical assumption of a cluster evolving with minimal influence from the rest of the galaxy may be far from realistic. For e.g. $\rho$~Oph, \citet{1992A&A...262..258D} suggested that the expanding shell of the upper Sco star forming region traveled through the cloud and triggered star formation. If it turns out that this is the rule rather than the exception, it makes it even more difficult to disentangle the competing processes that influence star formation. In this situation, we decided to observe IRAS~20050+2720, which is a star forming region located at a distance of about 700~pc \citep{1989ApJ...345..257W}. Apparently, no massive star formed in IRAS~20050+2720, thus the intensity of the ambient radiation should be small and we can study the evolution of late-type stars in the absence of external irradiation on the disk. Observationally, this region is especially suited for our study, because the diffuse background at 24~$\mu$m is weak, thus \emph{Spitzer} observations should provide a more complete source list than in other clusters. IRAS~20050+2720 was initially discovered as a point source by the \emph{IRAS} satellite. Molecular line emission indicates a mass infall on the regions \citep{1997ApJ...484..256G,1999ApJS..122..519C}. The luminosity of this region is estimated to 388~$L_{\sun}$ from the \emph{IRAS} fluxes \citep{1996A&A...308..573M}. \citet{2001A&A...369..155C} obtained mm-maps and they estimate the total gas mass within 65\arcsec=0.2~pc radius as 200~$M_{\sun}$. The first attempt to establish an IR classification lead to about 100 cluster members, about half of which show an IR excess \citep{1997ApJ...475..163C}. They find four regions of enhanced stellar density and label the most significant ones as A, B and C. Subcluster A contains several deeply embedded YSOs (young stellar objects). While \citet{1997ApJ...475..163C} estimate an average cluster age of 1~Myr, they interpret the presence of radio lobes and the stronger reddening as signatures for a more recent star formation event in subcluster A in the last 0.1~Myr. \citet{2005ApJ...632..397G} confirmed these ideas, noting that subcluster~B lacks the 850~$\mu$m emission which is present in A and C. It seems, that subcluster B recently cleared the gas of the parental molecular cloud. With \emph{Spitzer} we can get very accurate photometry of star forming regions and classify individual objects according to their IR SED. \citet{2009ApJS..184...18G} identified 177 YSOs in IRAS~20050+2720, which could be grouped in two distinct cores. We extend the spatial coverage with additional \emph{Spitzer} observations and add information from a deep \emph{Chandra} observation and, surprisingly, the first dedicated optical photometry.

We presented data on the young stellar cluster IRAS~20050+2720, which includes X-ray data from \emph{Chandra}, optical data and IR data from 2MASS and \emph{Spitzer}. The optical data is used to separate out foreground X-ray sources. Class~I and class~II sources are identified from their IR spectral energy distribution (SED). We treat all X-ray sources which are neither foreground objects nor classified as class~I or class~II according to their SED as potential class~III objects. The spatial distribution shows a strong subclustering with four subclusters. Subcluster A and B have the highest density of young stellar objects with members of all three classes. Subcluster C is so young that no class~III objects are present. Subcluster~D is separated from the other three by about 10\arcmin(=2~pc) and contains almost exclusively class~II objects. We identified some objects with optically thick disks.

1012.1354

We present early results of our multiwavelength study of the star forming region IRAS 20050+2720. While we use X-rays and IR to classify young stellar objects, the optical data can be used to exclude foreground objects. The dataset set contains 57 class I sources, 183 class II sources and 183 X-ray sources, of which 140 are class III candidates. Within IRAS 20050+2720 four subclusters are found. Subcluster A and B are the densest regions, which contain stars of all evolutionary stages. Subcluster C is much younger than the other two. It has not formed any class III objects yet. We newly identify a fourth subcluster, which consists mostly of class II objects and is located about 10' from the center of the cloud.

2358026

2011MNRAS.413.1764P

1012

1012.4092_arXiv.txt

Early-type dwarf galaxies (dEs, M$_{\it B}$ $>$ -18) are the numerically dominant population in the present-day Universe \citep{Sandage85, Binggeli87, Ferguson94}. They also exhibit strong clustering, being found predominantly in the close vicinity of giant galaxies, either as satellites of individual giants, or as members of galaxy clusters \citep{Ferguson89}. Although the dEs are characterized by their smooth appearance, having no recent or ongoing star formation and apparently no gas or dust content, the understanding of their origin and evolution remain major challenges for extragalactic astronomy. Stellar population studies show that dEs exhibit on average younger ages as compared to their giant counterparts, and also a lower metal content according to the correlation of metallicity and luminosity \citep{Michielsen08}. However, past studies provided a wide range of ages (e.g., \citealt{Poggianti01}; \citealt{Rakos01}; \citealt{Caldwell03}; \citealt{Geha03}; \citealt{van04}), from as old as being primordial objects to dEs with recently formed young stellar populations. It appears that dEs themselves are not a homogeneous class of objects. Sub-structures such as stellar disks, faint spiral arms or bars are quite frequent among the brighter dEs \citep{Lisker06a,Lisker07}. Many dEs were found to contain a central surface brightness enhancement consistent with a point source on top of the galactic main body \citep[e.g.][]{Binggeli91, Binggeli93}, referred to as so-called nucleated dEs. The studies from the HST/ACS Virgo cluster survey \citep{Cote06}, with their high angular resolution, not only verified the presence of such a distinct nucleus but also showed that nuclei are ubiquitous in bright dEs, covering a range in nucleus brightness. Interestingly, dEs with comparably faint nuclei that had not been identified before \citet{Cote06} show several systematically different properties as compared to dEs with bright nuclei \citep{Lisker07,Lisker08}. Different studies of dE nuclei from different data sets found several contradictory properties for the nuclei \citep{Grant05,Lotz04b,Cote06}. Particularly, the ground and space based data sets yielded different results. \citet{Grant05} found that the nuclei are on average redder than their surrounding galactic main body. On the other hand, studies using HST observations \citep{Cote06,Lotz04b} measured the dE nuclei to be slightly bluer than the galactic part. Furthermore, \citet{Cote06}, who used high quality data sets from the ACS Virgo Cluster Survey, proposed that the nuclei rather closely match the nuclear clusters of late type spiral galaxies in terms of size, luminosity and overall frequency. Another related scenario is also emerging: the recently discovered new (candidate) type of extremely small dwarf galaxies, the UCDs (Ultra Compact Dwarfs) with typical magnitudes of $-13 < M_b < -11$ \citep{Hilker99,Phillipps01}, might be the remnant nuclei of tidally stripped dwarf galaxies \citep{Bekki03,Drinkwater03,Goerdt08}. The formation mechanisms of the nuclei of dEs are poorly understood and various possibilities have been proposed, also depending on the evolution and formation of dEs as a whole. As the nucleated dEs are preferentially rounder in shape, \citet{Bergh86} proposed that the nuclei of dEs could have formed from the gas that sank to the centre of the more slowly rotating objects. Since they predominantly appear in highly dense environments, like the centre of a cluster of galaxies, the pressure from the surrounding inter-galactic medium may allow dwarf galaxies to retain their gas during star formation and produce multiple generation of stars \citep{Silk87,Babul92}, forming nuclei in the process. In both proposed scenarios the nuclei are formed along with the evolution of the galaxy itself, i.e., continuous star formation activity occurs at the dE centre as time passes. Unlike that, \citet{Oh00} suggested that dE nuclei might have formed in a different way, namely through subsequent migration or orbital decay of several globular clusters towards the centre of their host dE. It is difficult to provide a definitive observational test of these different scenarios for nucleus formation. Nevertheless, we can gain some insight by comparing the different observational properties, in particular relative ages and chemical enrichment characteristics, of the nuclei with their galactic main bodies, as well as with UCDs as their possible descendants. However, we need to bear mind that there may be a mixture of different formation scenarios. Our previous study based on this dataset \citep[hereafter Paper I]{Paudel10} has focused on the analysis of the inner stellar populations of dEs as a whole, without separating nuclei and galactic main bodies. Instead, our intention was to see the variation of the inner stellar population properties with different morphological subclasses of dEs \citep[cf.][]{Lisker07}, using a much larger sample of Virgo dEs than in previous Lick index studies. We showed that dEs with different substructure properties (with/without disk features, \citealt{Lisker06a}) have significantly different stellar populations: dEs with disk features are younger and more metal rich than dEs without disks. Therefore we concluded that these dEs probably do not have the same origin, as they also differ in their distribution with local environmental density in which they reside. By selection, all dEs in our sample contain a central nucleus, therefore it seems important to see the nature of the stellar populations of the nuclei and the surrounding galactic main bodies separately. And since there are different possibilities for the processes that form nuclei and also dEs themselves, we ask: can the nuclei thus tell us something about the formation history of dEs? This paper is organized as follows. In Section 2, we describe the sample of Virgo cluster dEs, observation and data reduction in brief. In section 3, we describe the measurement of line-strength indices in the Lick/IDS system. Our main results from the stellar population parameters are given in Section 4 and are discussed in Section 5. Finally, we summarize our findings in Section 6.

We have investigated the stellar population properties of the central nucleus and the surrounding galactic main body for a sample of 26 dEs in the Virgo cluster and compared the SSP-equivalent stellar population parameters of the dE nuclei with the ones of a small sample of UCDs. In addition to this, we have derived the radial profiles for age, metallicity and [$\alpha$/Fe] abundance for 13 dEs. Our main findings can be summarized as follows: \begin{itemize} \item We find that for most of the dEs the nuclei are significantly younger ($\sim$3.5 Gyr) and more metal rich ($\sim$0.07 dex) as compared to the galactic main body of the galaxies. Only five dEs have significantly older nuclei than their galactic main bodies, and dEs with old and metal poor nuclei are more likely to be distributed in the dense region of the cluster than the dEs with young and metal-enhanced nuclei. \item The metallicity of dE nuclei correlates with the total luminosity of dEs, and the observed metallicities of the nuclei have a fairly large range (+0.18 to -1.22 dex). All galactic main bodies of the dEs have sub-solar metallicity. \item While we see two distinct behaviours of SSP profiles (with and without a break) the overall trend of increasing age and decreasing metallicity with the radius is consistent with earlier studies. The $\alpha$-abundance as function of radius is consistent with no gradient. \item These observed properties suggest that the merging of globular clusters might not be the appropriate scenario for the formation of nuclei in dEs, at least not for the brighter dEs. The younger and comparably metal-rich nuclei support the idea that the central stellar populations of dEs were governed by continuous infall/accretion of gas in the centre of the potential well, building the nuclei. \item The heterogeneous nature of the stellar population characteristics of dEs hints at different formation scenarios of dEs, similar to the conclusion of our previous study \citep{Paudel10}. Our results suggest that the old, faint and metal-poor dEs are more likely to have a primordial origin, while those with relatively young ages and a higher metallicity and luminosity may have formed through morphological transformation. \end{itemize}

1012.4092

We present a comprehensive analysis of the spatially resolved stellar population properties of 26 early-type dwarf galaxies (dEs) in the Virgo cluster. Using Lick/IDS absorption line indices we derive simple stellar population (SSP) equivalent age, metallicity and [α/Fe] abundance ratio. In particular, we focus on the comparison of the stellar populations between the central nucleus and the surrounding galactic main body. The stellar populations of the nuclei are, for most dEs, significantly younger than those of the respective galactic main bodies, with an average difference of 3.5 Gyr. We find only five dEs with significantly older nuclei than their galactic main bodies. Furthermore, we observe most dE nuclei to be more metal rich compared to their host galaxies. These age and metallicity behaviours are shown by almost all dEs brighter than M<SUB>r</SUB>=-17 mag. <P />The metallicity of both nuclei and galactic main bodies correlates with the total luminosity of the dEs. However, the metallicity of the nuclei covers a larger range (+0.18 to -1.22 dex) than that of the galactic main bodies, which all have subsolar metallicity. The ages of dE nuclei show a statistically significant correlation with the local projected galaxy density within the cluster, such that younger ages are predominantly observed outside of the high-density central cluster region. The α-element abundance ratios are consistent with solar for both nuclei and galactic main bodies. <P />We also examine the presence of radial gradients in the SSP parameters for a subset of 13 dEs (up to 1.2 kpc or 15 arcsec radius). We notice two different types of gradients, namely smooth profiles that include the nucleus, and profiles where a break occurs between the nucleus and the rest of the galaxy. Nevertheless, an overall trend of increasing age and decreasing metallicity with radius exists, consistent with earlier studies. The α-abundance ratio as function of radius is consistent with no gradient. <P />Possible formation scenarios for the nuclei of dEs are discussed. The young and metal-enhanced population of nuclei suggests that these might have formed at later epochs, or the termination of star formation activity in the nuclei might have occurred relatively late, perhaps due to continuous infall of gas into the central potential well. Our stellar population analysis suggests that the merging of globular clusters is not an appropriate scenario for the formation of most dE nuclei, at least not for the brighter dEs. We speculate that there might be different formation processes which are responsible for the formation of dEs and their nuclei depending on their luminosity. Based on observations collected at the European Organization for Astronomical Research in the Southern hemisphere, Chile (programme 078.B-0178).

12164014

2011AIPC.1331...56P

1012

1012.0572_arXiv.txt

Currently, over 400 extra-Solar planetary systems have been found; most of hem around main sequence stars, with a few tens found in wide binary systems. Typically, such planets are thought to form in a protoplanetary disk left over following the central star formation in a protostellar disk \citep[e.g.][for a recent review]{arm07}. Several studies explored the later effects of stellar evolution on the survival and dynamics of planets around an evolving star \citep{deb+02,vil+07} or the possible formation of planets around neutron stars (NSs; see \citet{phi+94,pod95} for reviews). Others studied the formation and stability of planets in binary systems \citep[see][for a review]{hag09}. Here we focus on the implications of stellar evolution\emph{ in binaries} on the formation and growth of planets. More detailed discussion of these issues could be found in \citep{per10,per+10a,per+10b}. One of the most likely outcomes of stellar evolution in binaries is mass transfer from an evolved donor star {[}which later becomes a compact object; a white dwarf (WD), NS or a black hole (BH); here we mostly focus on low mass stars which evolve to become WDs{]} to its binary companion. If the binary separation is not too large, this process could result in the formation of an accretion disk containing a non-negligible amount of mass around the companion. Such a disk could resemble in many ways a protoplanetary disk, and could possibly produce a second generation of planets and/or debris disks around old stars. In addition, the renewed supply of material to a pre-existing ('first generation') planetary system (if such exists after surviving the post-MS evolution of the host star), is likely to have major effects on this system, possibly leading to the regrowth/rejuvenation of the planets and planetesimals in the system as well as possibly introducing a second epoch of planetary migration. The later evolution of evolved binaries could, in some cases, even lead to a third generation of planet formation in the system. When the previously accreting star (the lower mass star in the pre-evolved system) goes through its stellar evolution phase, it too can expand to become a mass donor to its now compact object companion. A new disk of material then forms around the compact object and planet formation may occur again, this time around the compact object (see also \citealp{bee+04} which tries to explain the formation the pulsar planet system PSR B1620-26). \begin{figure} \includegraphics[scale=0.4]{fig1d} \caption{Second (and third) generation planet formation. The various stages of second generation planet formation are shown schematically. (a) The initial configuration: a binary MS system, possibly having a first generation (I) circumstellar planet around the lower mass on an allowed (stable) orbit. (b) The higher mass stars evolves to the AGB phase, and sheds material which is accreted on the secondary, and forms a protoplanetary disk. The binary orbit expands, and the allowed stability region expands with it. The existing first generation planet may or may not survive this stage (see text). (c) Second generation debris disk and planets are formed, in regions previously forbidden for planet formation (in the pre-evolved system; panel a). (d) The secondary evolves off the MS, and sheds material to its now WD companion. A protoplanetary disk forms. The binary orbit and the planetary allowed region (around the WD) further expand. Second generation planets may or may not survive this stage (see text). (e) Third generation debris disk and planets are formed around the WD, in regions previously forbidden for planet formation (in the pre-evolved system; panel a).} \end{figure} In the following we discuss the role of binary stellar evolution in the formation of a second and third generation of planets in evolving binary systems (a schematic overview of this scenario is given in Fig. 1) and the interaction of accretion disks with pre-existing planets. We begin by discussing the conditions for the formation of a second generation protoplanetary disk and its properties (\S 2). We then explore the role of such disks in the formation of new planets (\S 3), their effects on pre-exiting planetary systems (\S 4), and on the formation of planets around compact objects (\S 5) . We then review the observational expectations for such second generation planetary systems as suggested from our discussion (\S 6). Finally we suggest several planetary systems as being candidate second generation planetary systems (\S 7) and then conclude (\S 8).

In this paper we discussed the implications of binary stellar evolution on planetary systems formation and evolution. We raised the possibility for the formation of second generation planetary systems in mass transferring binaries as a possible route for planet formation in old evolved systems and around compact objects in double compact object binaries. We presented possible implications for this process and the planetary systems it could produce, and detailed the possible observational signatures of second generation planetary systems. We also pointed out a few currently observed planetary systems with properties suggestive of a second generation origin. In addition we discussed the orbital evolution of pre-exiting planets in evolving binary systems due to mass loss from the system and the possible interaction with the formed accretion disk. The possibility of second generation planets and the study of planets in evolved binary systems may open new horizons and suggest new approaches and targets for planetary searches and research. It suggests that stellar evolution processes and stellar deaths may serve as the cradle for the birth and/or rejuvenation of a new generation of planets, rather than being the death throes or hostile hosts for pre-existing planets. In particular such processes could provide new routes for the formation of habitable planets, opening the possibilities for their existence and discovery even in (the previously thought) less likely places to find them. The environments of old stars and more so of compact objects could be very different from that of young stars. Such different environments can strongly affect the formation of second generation planets and possibly introduce unique processes involved in their formation and evolution. The discovery and study of second generation planets could therefore shed new light on our understanding of both planet formation and binary evolution, and drive further research on the wide range of novel processes opened up by this possibility. {\bf Acknowledgments} The author is a CfA, Rothschild, FIRST and Fulbright-Ilan Ramon fellow.

1012.0572

Exo-planets are typically thought to form in protoplanetary disks left over from protostellar disk of their newly formed host star. However, additional planetary formation and evolution routes may exist in old evolved binary systems. Here we discuss the implications of binary stellar evolution on planetary systems in such environments. In these binary systems stellar evolution could lead to the formation of symbiotic stars, where mass is lost from one star and could be transferred to its binary companion, and may form an accretion disk around it. This raises the possibility that such a disk could provide the necessary environment for the formation of a new, second generation of planets in both circumstellar or circumbinary configurations. Pre-existing first generation planets surviving the post-MS evolution of such systems would be dynamically effected by the mass loss in the systems and may also interact with the newly formed disk. Such planets and/or planetesimals may also serve as seeds for the formation of the second generation planets, and/or interact with them, possibly forming atypical planetary systems. Second generation planetary systems should be typically found in white dwarf binary systems, and may show various observational signatures. Most notably, second generation planets could form in environment which are inaccessible, or less favorable, for first generation planets. The orbital phase space available for the second generation planets could be forbidden (in terms of the system stability) to first generation planets in the pre-evolved progenitor binaries. In addition planets could form in metal poor environments such as globular clusters and/or in double compact object binaries. Observations of exo-planets in such forbidden or unfavorable regions could possibly serve to uniquely identify their second generation character. Finally, we point out a few observed candidate second generation planetary systems, including Gl 86, HD 27442 and all of the currently observed circumbinary planet candidates. A second generation origin for these systems could naturally explain their unique configurations.

12205369

2011IAUS..272..519L

1012

1012.5211_arXiv.txt

The HERMES instrument on the 1.2m Mercator telescope at La Palma (\cite[Raskin \& Van Winckel 2008]{Raski08}) is a new high-efficiency fiber-fed bench-mounted cross-dispersed echelle spectrograph that observes the complete wavelength range from 420 nm to 900 nm in a single exposure with $R$=80,000. TAC reviewed HERMES observation programs of the contributing research institutions started mid 2009. We present first results of a long-term high-resolution spectroscopic monitoring program of 3 LBVs and 2 LBV candidates (up to $V$=$11^{\rm m}.0$). The HERMES monitoring program will provide invaluable new clues about the structure and dynamics of LBV atmospheres, the physics of their extended winds, and the strong line broadening mechanisms in these rare massive hot stars near the Eddington luminosity limit. The monitoring program will be crucial for documenting the enigmatic LBV outburst events, for detecting new long-period LBV binaries and the reliable determination of the orbital parameters.

Long-term monitoring with Mercator-HERMES of the optical spectrum of the prototypical LBV P Cyg reveals variability at the base of its supersonic wind that can be linked to moderate $V$-changes (of $0^{\rm m}.1$ to $0^{\rm m}.2$) over a period of $\sim$4 months. We find strong indications for the binary nature of MWC~314 from large radial velocity changes observed in photospheric absorption lines during less than one week. We observe prominent P Cygni profiles in the optical spectrum of LBV candidate MWC~930, signaling the presence of a central massive hot star. The optical spectral lines of LBV candidate HD~168625 are less variable, although we also observe clear signatures of expanding H$\alpha$ wind variability on short time-scales of $\sim$1 month. The optical spectrum of HD~168607 reveals large line profile changes over the past 12 years confirming its LBV designation.

1012.5211

We present results of a long-term spectroscopic monitoring program (since mid 2009) of Luminous Blue Variables with the new HERMES echelle spectrograph on the 1.2m Mercator telescope at La Palma (Spain). We investigate high-resolution (R = 80,000) optical spectra of two LBVs, P Cyg and HD 168607, the LBV candidates MWC 930 and HD 168625, and the LBV binary MWC 314. In P Cyg we observe flux changes in the violet wings of the Balmer Hα, Hβ, and He i lines between May and Sep 2009. The changes around 200 to 300 km s<SUP>-1</SUP> are caused by variable opacity at the base of the supersonic wind from the blue supergiant. <P />We observe in MWC 314 broad double-peaked metal emission lines with invariable radial velocities over time. On the other hand, we measure in the photospheric S ii λ5647 absorption line, with lower excitation energy of ~14 eV, an increase of the heliocentric radial velocity centroid from 37 km s<SUP>-1</SUP> to 70 km s<SUP>-1</SUP> between 5 and 10 Sep 2009 (and 43 km s<SUP>-1</SUP> on 6 Apr 2010). The increase of radial velocity of ~33 km s<SUP>-1</SUP> in only 5 days can confirm the binary nature of this LBV close to the Eddington luminosity limit. <P />A comparison with VLT-UVES and Keck-Hires spectra observed over the past 13 years reveals strong flux variability in the violet wing of the Hα emission line of HD 168625 and in the absorption portion of the Hβ line of HD 168607. In HD 168625 we observe Hα wind absorption at velocities exceeding 200 km s<SUP>-1</SUP> which develops between Apr and June 2010.

2058795

2011A&A...526A.113F

1012

1012.0328_arXiv.txt

Gamma-ray burst (GRB) afterglows are commonly interpreted in the framework of the standard synchrotron shock model, in which an ultra-relativistic shock is expanding into the ambient medium swept up by the blast wave (M\'esz\'aros \cite{meszaros2}; Zhang \& M\'esz\'aros \cite{zhang3}; Piran \cite{piran}). For the simplified assumption that the shock front is spherical and homogeneous, a smooth afterglow light curve is expected. This smooth power-law decay with time was a common phenomenon in most of the pre-\emph{Swift} GRBs (Laursen \& Stanek \cite{laursen}), because the afterglow observations typically began $\sim$1 day after the burst compared to now when we can be on-target within minutes. The \emph{Swift} satellite (Gehrels et al. \cite{gehrels}) allows studies of the early afterglow phase thanks to its rapid slew, a precise localization of GRBs with its Burst Alert Telescope (BAT, Barthelmy et al. \cite{barthelmy}), and the early follow-up with two telescopes sensitive at X-ray (XRT, Burrows et al. \cite{burrows}) and ultraviolet/optical (UVOT, Roming et al. \cite{roming}) wavelengths. Since its launch in 2004, \emph{Swift}, together with ground-based follow-up telescopes, has provided many early and well-sampled afterglow light curves deviating from the smooth power-law decay (Panaitescu et al. \cite{panaitescu3}; Nousek et al. \cite{nousek}; Zhang et al. \cite{zhang2}; Panaitescu et al. \cite{panaitescu4}). Such variability can shed light on the central engine and its surroundings. Several major scenarios have been proposed for afterglow variability. The reverse shock emission might add to the emission from the forward shock (see \textsection 4.1, Sari \& Piran \cite{sari2}; M\'esz\'aros \& Rees \cite{meszaros}; Zhang et al. \cite{zhang}; Kobayashi \& Zhang \cite{kobayashi}), the shock might be refreshed by slower shells catching up with the decelerating front shells (see \textsection 4.2, Rees \& M\'esz\'aros \cite{rees}; Panaitescu \cite{panaitescu2}; Sari \& M\'esz\'aros \cite{sari3}; Panaitescu et al.\cite{panaitescu}; Granot et al. \cite{granot2}; Kumar \& Piran \cite{kumar}), the ambient density profile into which the blast wave expands might not be homogeneous (see \textsection 4.3, Lazzati et al. \cite{lazzati}; Nakar et al. \cite{nakar2}; Zhang et al. \cite{zhang2}; Nakar \& Piran \cite{nakar3}; Ioka et al. \cite{ioka}; Wang \& Loeb \cite{wang}; Dai \& Lu \cite{dai},;Nakar \& Granot \cite{nakar}), or the jet may have an angular structure different from a top hat (see \textsection 4.4, Peng et al. \cite{peng}; Granot et al. \cite{granot}; Berger et al. \cite{berger}; Racusin et al. \cite{racusin2}). Here we provide details of the \emph{Swift}, GROND, and REM observations of the afterglow of GRB 080413B and test the above alternative scenarios for consistency with these data. Throughout the paper, we adopt the convention that the flux density of the GRB afterglow can be described as $F_\nu (t) \propto t^{\alpha} \nu^{-\beta}$.

In this paper we study the optical/NIR light curve produced by the afterglow of GRB 080413B. The possibility that the jet of this GRB might have a narrow ultra-relativistic core and a wider, mildly relativistic outer component has been indicated by the observation of the afterglow emission. An on-axis coaxial two-component jet model provides a consistent description of the properties of GRB 080413B, and can additionally explain the wide range of light curve evolutions, the difference between optical/NIR and X-ray light curves, and the chromatic evolution of the optical light curve itself. The comparison with the two most prominent light curves modeled by the two-component jet to date - GRB 050315 and GRB 080319B - reveal consistency with the GRB 080413B afterglow light curve. The X-ray light curve of the afterglow of GRB 050315 (Granot et al. \cite{granot}; Nousek et al. \cite{nousek}) shows a remarkable resemblance to the optical/NIR light curve evolution of the afterglow of GRB 080413B. If we neglect the very steep tail of the prompt GRB emission, the initial XRT light curve of GRB 050315 is dominated by the narrow jet, followed by a slight rebrightening at around 1.5~ks caused by the wide jet in its pre-deceleration phase. After the peak, the light curve decay is dominated by the emission from the wide jet. Times of jet breaks of narrow ($\sim9$~ks) and wide ($\sim200$~ks) components, as well as their opening angles $\theta_w = 2 \theta_n = 3.2^\circ$ (Granot et al. \cite{granot}), are within an order comparable with those of GRB 080413B. The X-ray light curve of the naked-eye GRB 080319B (Racusin et al. \cite{racusin2}) shows similar evolution. The narrow jet dominates the first $\sim40$~ks of the afterglow. After the narrow jet decays, the wide jet dominates the late afterglow. There is no rising part of the wide jet and thus no sharp rebrightening or plateau, so the wide jet merely makes the decay flatter. The optical light curve is missing the emission from the narrow jet, suggesting that the optical flux from the wide jet must be much stronger than that of the narrow jet. The jet break of the narrow jet at $\sim2.8$~ks, which corresponds to an extremely narrow opening angle of $0.2^\circ$, is the earliest of these three bursts. The jet break of the wide component with opening angle $\sim4^\circ$ is, on the other hand, the latest at roughly 1~Ms. In general, the X-ray light curves of GRBs 050315, 080319B and the optical light curve of GRB 080413B are very similar. However, the afterglow of GRB 080413B is the only one showing both components in the optical/NIR wavelengths, while the emission from the wide jet in the X-rays is negligible. The X-ray flux from the wide jet must then be much less prominent than for the narrow jet. Following this line of reasoning the relative fluxes in optical/NIR and X-ray of the narrow and wide jets can be explained in the following way. The SED of the narrow jet (intervals I and II) shows a break, while that of the wider jet does not. For both jets we have argued that we cover the slow cooling regime. The spectral slope of the wide component implies that the cooling break is at frequencies below the near-infrared bands. Both the cooling frequency and the maximum power depend on the product of Lorentz factor $\Gamma$ of the shocked fluid and the magnetic field strength. It is generally assumed that the narrow jet comes with a larger Lorentz factor than the wide one, and a similar assumption can be reasonably made about the (self-created) magnetic field. The SEDs of the two jets show us that the product $\Gamma *B$ of the wide jet, and consequently also the emission at X-ray energies, are at least a factor 100 less than for the narrow one. Therefore, the wide jet does not contribute to the X-ray emission in any significant way. The situation is different in the optical/NIR since cooling break of the narrow jet leads to a reduced flux by a factor of $\approx$ 10 relative to a spectrum with no cooling break between the optical/NIR and X-rays. Consequently, the optical/NIR emission of the wide jet is much more prominent than for the narrow jet. The values derived from the modeling of GRB 080413B afterglow are fairly consistent with the collapsar jet breakout model of Zhang et al. (\cite{zhang5}), where the numerical simulations predict $\theta_n = 3-5^\circ$, $\Gamma_n \gtrsim 100$ for the narrow component and $\theta_w \sim 10^\circ$, $\Gamma_w \sim 15$ for the wide component (Peng et al. \cite{peng}). The characteristic Lorentz factors are very similar to those of the hydromagnetically accelerated, initially neutron-rich jet model of Vlahakis et al. (\cite{vlahakis}), where $\Gamma_n \sim 200$ and $\Gamma_w \sim 15$. These two models are distinguished by the ratio of the kinetic energy injected into the two components. For values typical of the collapsar model ($E_w / E_n \sim 0.1$), Peng et al. (\cite{peng}) predict that the contribution of the narrow component dominates at all times. However, for $E_w \gtrsim 2 E_n$ (as in the neutron-rich hydromagnetic model), the narrow component dominates at early times but the contribution of the wide jet becomes dominant around the deceleration time of the wide jet. If $E_w > E_n$, the jet break of the narrow jet could be masked by the rise (and subsequent dominance) of the flux from the wide jet as the deceleration time of the wide component is approached. That the only visible jet break in the optical light curve is the one of the wide jet may lead to overestimating the emitted gamma-ray energy if the opening angle of the wide jet is used in converting the measured energy into the beaming-corrected energy (see Peng et al \cite{peng} for detailed discussion). Because the deceleration time of the wide component is much longer than for the narrow component, a bump is expected to show up in the decaying light curve of the narrow component owing the emission of the wide component at its deceleration time. These predictions are in perfect agreement with our data, suggesting that the two-component jet model can be placed among models that explain the variability in the early light curves of the GRB afterglows.

1012.0328

<BR /> Aims: The quick and precise localization of GRBs by the Swift telescope allows the early evolution of the afterglow light curve to be captured by ground-based telescopes. With GROND measurements we can investigate the optical/near-infrared light curve of the afterglow of gamma-ray burst 080413B in the context of late rebrightening. <BR /> Methods: Multi-wavelength follow-up observations were performed on the afterglow of GRB 080413B. X-ray emission was detected by the X-ray telescope onboard the Swift satellite and obtained from the public archive. Optical and near-infrared photometry was performed with the seven-channel imager GROND mounted at the MPG/ESO 2.2 m telescope and additionally with the REM telescope, both in La Silla, Chile. The light curve model was constructed using the obtained broad-band data. <BR /> Results: The broad-band light curve of the afterglow of GRB 080413B is well fitted with an on-axis two-component jet model. The narrow ultra-relativistic jet is responsible for the initial decay, while the rise of the moderately relativistic wider jet near its deceleration time is the cause of the rebrightening of the light curve. The later evolution of the optical/NIR light curve is then dominated by the wide component, the signature of which is almost negligible in the X-ray wavelengths. These components have opening angles of θ_n ~ 1.7° and θ_w ~ 9°, and Lorentz factors of Γ_n &gt; 188 and Γ_w ~ 18.5. We calculated the beaming-corrected energy release to be E<SUB>γ</SUB> = 7.9 × 10<SUP>48</SUP> erg. <P />Tables 1-3 are only available in electronic form at <A href="http://www.aanda.org">http://www.aanda.org</A>

12231110

2011RAA....11..434T

1012

1012.1604_arXiv.txt

% \label{sect:intro} Galactic black hole X-ray binaries exhibit two main X-ray spectral states -- the hard state and the soft state (see the review Remillard \& McClintock 2006). In the soft state, the X-ray energy spectrum is dominated by a thermal component with a weak power law tail. In the hard state, the energy is characterized by a power law with a break or an exponential cutoff at high energy. These two spectral states are shared by the atoll sources in the neutron star X-ray binaries (van der Klis 1994), roughly corresponding to the banana state and the island state, respectively (Hasinger \& van der Klis 1989), in terms of the X-ray spectral and timing properties. The transition between the hard state and the soft state is called as \textquotedblleft state transition \textquotedblright. In this paper, we focus on the hard-to-soft (hereafter H-S) transition. H-S transition can usually be seen during the rising phase of a bright outburst of a black hole or neutron star transient, but there are exceptions. Yu \& Dolence (2007) discovered one H-S transition during a luminosity decline in Aquila X-1. A number of transient sources have been seen to stay in the hard state throughout their outbursts, e.g., XTE J1550-564 showed spectral and timing features typical of an hard state during its short outburst in January 2002 (Belloni et al. 2002). A complete H-S transition usually takes on time scales from a few days to a few weeks, varied source by source. Up to now, the origin of the state transition has not been completely understood. The widely accepted understanding is based on stationary accretion (e.g., Esin et al. 1996). In this framework, mass accretion rate determines the state transition. However, the study of the hysteresis effect of spectral state transitions (Miyamoto et al. 1995; Nowak 1995; Maccarone \& Coppi 2003) and the large span of H-S transition luminosities, which could vary by one order of magnitude in a single source (Yu \& Dolence 2007), suggest that an additional parameter is needed to interpret state evolution. Yu \& Yan (2009) showed that the rate-of-change of the mass accretion rate rather than the mass accretion rate itself drives the state transition in most bright Galactic X-ray binaries primarily, suggesting that we need to rely on non-stationary accretion to interpret state transitions. It has been found that the luminosity corresponding to the start of the H-S transition positively correlates with the outburst peak luminosity not only in individual transient sources Aql X-1, XTE J1550-564, and GX 339-4, and individual persistent, transient-like neutron star low-mass X-ray binary 4U 1705-44 (Yu et al. 2004, 2007; Yu \& Dolence 2007), but also in transient and persistent sources as a whole (Yu \& Yan 2009). This correlation holds for bright outbursts as well as low-luminosity flares (Yu \& Dolence 2007; Yu \& Yan 2009), and in view of no saturation toward high luminosities, brighter hard states than the ones currently known is expected to be observed in transient sources during brighter outbursts. On the other hand, one can predict the outburst peak luminosity during the rising phase of an outburst when the H-S transition occurs (Yu et al. 2004, 2007; Yu \& Dolence 2007), and even earlier, during the early outburst rise using measurements of the rate-of-change of the X-ray flux (Yu \& Yan 2009). It is therefore very important to keep tracking spectral transitions in the bright X-ray binaries --- understanding luminosity regimes of the hard state and the soft state and predicting or alerting further spectral and flux evolution. We have systematically studied state transitions of all bright X-ray binaries, seen with X-ray monitoring observations with the All-Sky Monitor (ASM) on board the Rossi X-Ray Timing Explorer (RXTE) and the Burst Alert Telescope (BAT) on board the Swift in the 2-12 keV and the 15-50 keV energy ranges, respectively, in a period of five years following the previous study Yu \& Yan (2009). Specifically, we applied an automatic program to search for H-S transitions in all bright X-ray binaries according to the hardness ratios between the BAT flux and the ASM flux and to determine the transition flux and the peak flux of the following soft state. We obtained the updated correlation between the luminosity corresponding to the H-S transition and the peak luminosity of the following soft state in both transient and persistent sources. In total we found 59 new H-S transitions in 2008-2010, in addition to 69 H-S transitions identified in 2005-2008.

\label{sect:conclusion} Using X-ray monitoring observations with the RXTE/ASM and the Swift/BAT, we have performed an automatic search for the H-S state transitions in bright persistent and transient X-ray binaries during a period of five years from 2005 to 2010. The identification of transitions and the correlation between the transition luminosity and the outburst/flare peak luminosity in the following soft state obtained with our automatic routine from the data in 2005-2008 is consistent with that reported in Yu \& Yan (2009). From the recent observations in 2008-2010, we have identified 59 more H-S state transitions in 28 Galactic X-ray binaries, which also show the same positive correlation. \subsection{H-S Transitions Newly Identified in Individual Sources} Let's take a closer look at the new sources other than those in Yu \& Yan (2009). In Aql X-1, at least 4 H-S transitions were seen before 2005 (Yu \& Dolence 2007), but none in 2005-2008, then a new one occurred at the end of 2009. IGR J17473-2721 was only observed a week outburst in 2005, characterized by a hard state spectrum. A H-S state transition was seen in 2008 June. 4U 1630-47 is one of the most active black hole transients, and it has produced strong hard X-ray emission during its 17 detected outbursts (Tomsick et al. 2005). The latest outburst started in 2009 December, underwent the H-S transition, and went back to the hard state until 2010 August (Tomsick \& Yamaoka 2010). SWIFT J1842.5-1124 was first detected in 2008 July (Krimm et al. 2008). Its X-ray spectral and timing properties are similar to the black hole in the hard state (Markwardt et al. 2008). A H-S transition occurred in 2008 September according to the BAT and ASM light curves. XTE J1752-223 was discovered during RXTE scanning the Galactic bulge region on 23 October 2009 (Markwardt et al. 2009) and there is strong evidence indicating that it is a black hole X-ray transient. Since the discovery of the source, it showed little spectral evolution and had been in the hard state. Until 2010 January, it underwent a H-S transition (Homan 2010). In addition, 4U 1820-30 remained in the soft state after the recent H-S transition in 2005 August. In 2009 April, it entered the hard state and after two months, another H-S transition took place. This phenomenon is atypical for 4U 1820-30 over the past 10-15 years, since the source usually stays in the hard state for one to two weeks (Krimm et al. 2009). Furthermore, by analyzing H-S transitions during 2005-2010 in single sources, we have found that the transition luminosities of 4U 1608-52 in recent two years were much lower than that in 2005-2008, different from the brightest H-S transition by an order of magnitude. Similar large variation of the transition luminosity has also been seen previously in Aql X-1 (Yu \& Dolence 2007). Besides, in the single source 4U 1636-53, which had experienced many spectral transitions during the studied time period, there exists the correlation between the luminosities in Eddington units well. Such a correlation has been found in individual sources XTE J1550-564, Aql X-1 and 4U 1705-44 (Yu et al. 2004). Detailed study of the 4U 1608-52 and 4U 1636-53 will be presented elsewhere. \subsection{Estimates of the Significance of the H-S Transitions Identified} Yu \& Yan (2009) studied the duration of the bright Galactic X-ray binaries staying at certain hardness ratios and found the hardness ratios of the hard state and the soft state are in the range above 0.6 and below 0.35, respectively, showing distinct bimodal distribution. Therefore we used the two thresholds to estimate the significance of the H-S transitions we identified. We used the data point immediately before the transition whose hardness ratio is above 1.0 and the preceding two data points to calculate the significance that at least one 2-day hardness ratio is greater than 0.6 --- the significance that the source was in the hard state. Similarly, we used the first data point following whose hardness ratio is below 0.2 and the following two data points to calculate the significance that at least one 2-day hardness ratio is smaller than 0.35 --- the significance that it is in the soft state. A H-S transition should have occurred in this interval if the soft state appears after the hard state. The result was that in these 128 transitions, 125 exceeds $2\sigma$ and 101 exceeds $3\sigma$. If we added one data point following the hard state data and one data point preceding the soft state data to calculate the significance of reaching the hard state and the soft state respectively, there are 126 out of these 128 transitions larger than $2\sigma$ and 105 larger than $3\sigma$. Thus, most of the identifications of state transitions we identified are reliable. \subsection{Future Work} After including 2008--2010 data compared with the data used in Yu \& Yan (2009), the H-S transition luminosities remain to spread over a luminosity range of about 2 orders of magnitude, showing no cutoff or saturation at either end of the correlation. The large span of the transition luminosities indicates that using mass accretion rate to interpret the state transition is not correct. According to the conclusion of Yu \& Yan (2009), the rate of increase of the mass accretion rate drives the H-S transition in most of the state transitions in X-ray binaries. 2S 0918-549 and 4U 0614+091, which have the lowest transition luminosities, can be explained as that their $\frac{d\dot{M}}{dt}$ is the least. And Cyg X-3, if it contains a neutron star, has the highest transition luminosities because of the highest $\frac{d\dot{M}}{dt}$ (see Figure 26. in Yu \& Yan 2009). In general, transient sources tend to have higher $\frac{d\dot{M}}{dt}$ than persistent sources and therefore higher H-S transition luminosities. Modification of the program will allow us to automatically identify H-S transitions using monitoring data other than the BAT and the ASM, for example, MAXI, unless the instruments cover the soft component and the non-thermal components. Noting that the hardness ratio thresholds for the hard state and the soft state should be determined first, just as Yu \& Yan (2009), for the hardness ratio range of the hard state and the soft state will change when we use different instruments or different energy bands to calculate the hardness ratio. Additionally, the program can be modified to search for the soft-to-hard transition in the outburst or flare in which the H-S transition was identified. Application to instant alert or predict state transitions and to study the hysteresis effect of spectral state transitions is under further investigation. \normalem

1012.1604

We have studied X-ray spectral state transitions that can be seen in the long-term monitoring light curves of bright X-ray binaries from the All-Sky Monitor (ASM) onboard the Rossi X-ray Timing Explorer (RXTE) and the Burst Alert Telescope (BAT) onboard Swift during a period of five years from 2005 to 2010. We have applied a program to automatically identify the hard-to-soft (H-S) spectral state transitions in the bright X-ray binaries monitored by the ASM and the BAT. In total, we identified 128 hard-to-soft transitions, of which 59 occurred after 2008. We also determined the transition fluxes and the peak fluxes of the following soft states, updated the measurements of the luminosity corresponding to the H-S transition and the peak luminosity of the following soft state in about 30 bright persistent and transient black hole and neutron star binaries following Yu &amp; Yan, and found the luminosity correlation and the luminosity range of spectral transitions in data between 2008-2010 are about the same as those derived from data before 2008. This further strengthens the idea that the luminosity at which the H-S spectral transition occurs in the Galactic X-ray binaries is determined by non-stationary accretion parameters such as the rate-of-change of the mass accretion rate rather than the mass accretion rate itself. The correlation is also found to hold in data of individual sources 4U 1608-52 and 4U 1636-53.

12168431

2011ApJ...733L..29W

1012

1012.1268_arXiv.txt

\label{sec:introduction} One of the main accomplishments in cosmology during the last decade is the establishment of the $\Lambda$CDM inflationary concordance model. According to this model, the universe consists of 5\% baryonic matter, 22\% dark matter and 73\% dark energy \citep{jarosik:2010}, and is filled with random and Gaussian fluctuations. These fluctuations were generated during a short period of exponential expansion called inflation \citep[e.g.][ and references therein]{liddle:2000}, during which the universe expanded by a factor of $\sim10^{26}$ in $\sim10^{-34}$ s. With only a handful of free parameters, this model is able to successfully fit thousands of observational data points. Nevertheless, the $\Lambda$CDM model must at the current stage be considered an effective model rather than a fundamental model. First, it relies on several quantities that have never been directly observed except through their gravitational impact, such as both dark matter and dark energy. Second, it postulates the existence of an unknown scalar field, the inflaton. It is therefore important to put the inflationary framework to stringent tests, probing its range of validity in different ways. One approach to do so is to construct alternative cosmological theories, making different observational predictions than $\Lambda$CDM, and then compare the two models using high-precision data. One example of such work has been demonstrated by the development of the ``Conformal Cyclic Cosmology'' (CCC) model by \citet{penrose:2008,penrose:2009,penrose:2010}. In this picture, the history of the universe is described in terms of a series of ``aeons'', each of which is defined as a finite period between two Big Bang events. We will not consider further details of the CCC model in this paper, except for one interesting feature: According to \citet{gurzadyan:2010}, the model postulates that super-massive black holes may collide in earlier aeons, releasing tremendous amounts of energy in the form of gravitational waves. These collisions may be observable in our current aeon in the form of concentric circles of low variance in the cosmic microwave background (CMB). \begin{figure*}[t] \mbox{\epsfig{figure=wmap_example.eps,width=\linewidth,clip=}} \caption{Examples of single standard deviation profiles (thick histograms), computed from WMAP (left) and a simulation (right). The thin solid line shows the mean of all profiles, and the shaded regions show the 1 and $2\sigma$ confidence regions. The WMAP profile is centered on galactic coordinates $(l,b) = (105.04^{\circ}, 37^{\circ})$, reproducing Figure 2 of \citet{gurzadyan:2010}. Note that low-variance rings are found also in the simulation.} \label{fig:examples} \end{figure*} Following up on this prediction, \citet{gurzadyan:2010} analyzed the 7-year WMAP temperature sky maps, searching for concentric circles of low variance. And quite surprisingly, they claim to find evidence for such rings at the $6\sigma$ confidence level, by comparing the results obtained from WMAP with $\Lambda$CDM simulations. This claim is sufficiently spectacular (involving both pre-Big Bang phenomena and super-massive black holes) to catch the interest of the general media, with numerous news stories and live media appearances following in the weeks after the release of the paper. Given these widespread reactions, the claims of Gurzadyan and Penrose deserve closer scrutiny through independent analysis; this paper presents one such independent analysis.

\label{sec:conclusions} In this paper we search for concentric low-variance rings in the 7-year WMAP temperature sky maps, seeking to reproduce the results recently presented by \citet{gurzadyan:2010}. While our two analyses do agree in terms of specific variance profiles, they clearly disagree in terms of statistical interpretation and significance. Specifically, we find a substantially larger variance in our $\Lambda$CDM simulations than \citet{gurzadyan:2010} do in theirs. When taking into account this larger variance, and accounting for statistical selection effects, the evidence for concentric circles in the WMAP data appears minimal. Rather, the WMAP data appears fully consistent with our simulations. The main difference between the two analyses must lie in the construction of the simulations. However, we have good reason to believe that our simulations are correct. First, we note that in our simulations the mean variance profile decreases towards small angular distances. This is bound to happen for two reasons: The CMB anisotropies constitute a correlated field with a characteristic scale of $1^{\circ}$, corresponding to the first peak in the CMB spectrum. Also, the instrumental beam of WMAP smooths out all small scale structure. It is therefore surprising that this effect is not seen in Figure 2 of \citet{gurzadyan:2010}. Second, the fact that our simulations agree with WMAP provides further confidence; it is difficult to imagine an error in the simulation pipeline that would cause the simulations to become \emph{more} similar to the observed data. Third, as noted by \citet{gurzadyan:2010}, there are both low and high peaks in the reported variance profiles. While Gurzadyan and Penrose assign no significance to the high variance peaks (stating that ``the peaks of high variance are of no importance, as these can result from numerous irrelevant effects''), they are clearly relevant in our interpretation, as they illustrate the substantial statistical variance present in these profiles. This large variance is observed both in the real data and the simulations. Thus, we conclude that there is no evidence for the CCC model in the current WMAP data. Of course, Planck \citep{tauber:2010} may provide new light on this issue, having higher sensitivity and resolution than WMAP, but one should probably not have too high expectations in this regard. Even if the CCC model should turn out to describe the real universe, it will quite likely be difficult to unambiguously identify such concentric circles due to the dominant background cosmic variance.

1012.1268

In this Letter, we search for concentric circles with low variance in cosmic microwave background sky maps. The detection of such circles would hint at new physics beyond the current cosmological concordance model, which states that the universe is isotropic and homogeneous, and filled with Gaussian fluctuations. We first describe a set of methods designed to detect such circles, based on matched filters and χ<SUP>2</SUP> statistics, and then apply these methods to the best current publicly available data, the 7 year Wilkinson Microwave Anisotropy Probe (WMAP) temperature sky maps. We compare the observations with an ensemble of 1000 Gaussian ΛCDM simulations. Based on these tests, we conclude that the WMAP sky maps are fully compatible with the Gaussian and isotropic hypothesis as measured by low-variance ring statistics.

1158376

2011ASSP...21..559C

1012

1012.4351_arXiv.txt

\label{sec:1} There are only three confirmed galactic High Mass X-ray Binaries (HMXBs) with persistent TeV emission (PSR B1259-63, LS 5039 and LSI +61 303, see~\cite{holder09}) and one recently proposed candidate (HESS J0632+067~\cite{hinton09}). The VHE emission in the 3 confirmed binaries is strongly modulated with the orbital period, suggesting that the emitter is rather compact and close to the massive star. However, the origin of the VHE emission remains unclear, with competing leptonic and hadronic scenarios which invoke either inverse Compton scattering of stellar photons or proton-proton collisions in a jet/pulsar wind scenario. Obviously, the nature of the compact star is an important ingredient in the different $\gamma$-ray production models. PSR B1259-63 contains a pulsar whereas a black hole is not ruled out in LS 5039 and LSI +61 303. In this context optical studies can bring new insights through constraining the compact object mass with dynamical studies.

1012.4351

We present new optical spectroscopy of the γ-ray binary LS 5039. Our data show evidence for sub-orbital modulation in the radial velocities with amplitude ∼ 7 km/s and period ∼ Porb/4. This short-term oscillation is stable over at least 7 years and it is likely triggered by non-radial oscillations of the O6.5V optical star. We also present the results of a spectroscopic campaign on MWC 148, the optical counterpart of the new γ-ray binary candidate HESS J0632+067. Long-term variations in the Hα and Hβ emission line parameters are clearly detected which, if modulated with the binary orbit, would imply a period &gt;200 days.

12286486

2012MSAIS..19..271A

1012

1012.3051_arXiv.txt

\begin{figure*} \centering \centerline{\psfig{figure=JKCS041malaspina.ps,width=9truecm,clip=}} \caption{$gz'K$-band image of JKCS\,041. Red, green and blue channel use Ks, $z'$ and g bands. The smooth, blue emission is the X-ray emission. North is up, East is to the left, the field of view is $5\times5$ arcmin. Reprinted with permission. } \label{fig:malaspina} \end{figure*} Clusters markately differ from filaments, walls and sheets, up to the highest redshifts at which we have good data to allow such a discrimination. These structures are very different environments in terms of (hot) gas temperature and density, galaxy density, velocity dispersion, morphological composition, depth of the potential well, etc. While someone might argue that the difference between them is just semantic, as a matter of fact, we don't known any astronomer calling ``cluster" the Great Wall. If one wants to study the evolution of galaxies in clusters, of the X-ray scaling relations of clusters, or any sensible {\it cluster} quantity, it seems not reasonable to build samples where clusters, filaments, walls and sheets enter (and exit) from the sample in an uncontrolled way. For example, a cluster sample contaminated at high redshift by the presence of filaments, walls and sheets is prone to confuse the effect of (look-back) time (i.e. evolution) with environment. Such a contamination is possible at high redshift if the studied sample includes systems that we actually don't known whether they are clusters or not. In order to discriminate high redshift clusters from other large scale strucures (redshift spikes, filaments, sheets, etc), a reliable probe of the structure size in the line of sight direction is the most difficult observation to acquire, because measurements are generally easier in the plane of the sky. Low quality velocity dispersions, with common small samples, are of little help unless the found velocity dispersion is large enough, and its error small, to discard low values typical of large scale structures or of galaxy pairs. X-ray is useful, because if an X-ray source is unambigously extended and spatially coincident with a galaxy overdensity, it testifies the presence of a potential well deep enough to hot and retain the intracluster medium, i.e. to reject non-cluster possibilities. However, a generic X-ray detection, such as such as those of some high redshift structure candidates from Chiaberge et al. (2009) or Henry et al. (2010), is insufficient in this respect because it could originate from the ICM but also from the numerous (at the cluster count rate) point sources (or a blend of them) of the X-ray sky. Distinguishing between a truely extended or a blend of point sources with just the few counts typical of faint X-ray detections requires a good PSF, like the one of Chandra. We emphasize that if a probability computation is used instead of Chandra data, it should be realistic, and allow, for example, blended point sources to have all possible fluxes (as in the real universe) and not only a single value as in Henry et al. (2010) computation. Equally dangerous is allowing to enter in an uncontrolled way structures which won't necessarily become clusters by $z=0$. Several proto-clusters are in this situation, basically because the 1 (or 2) $\sigma$ interval of the collapsing time includes values larger than the look-back time of the structure. For example, the ''proto-clusters" in Hatch et al. (2010) have a collapsing time of 6 Gyr, shorter than the look back time at the redshift of the systems (about 10.8 Gyr), but widely uncertain and larger that the latter at better than 1.5 sigma. Therefore, there is very little, if any, evidence that these systems will be clusters by today. Assuming that they will become clusters by $z=0$, when such evidence is lacking, is risky and prone to mix different environments at different redshifts. To summarize, having detected a structure at high redshift, it not enough to call it cluster. If the nature of the detected structure is unknown, we risk to compare apples (at very high redshift) to oranges (at lower redshift) and therefore call ``evolution" what is instead ``contamination" or ``environemnt". In that sense, JKCS\,041 (Andreon et al. 2009, see Fig 1) is a uniq very high redshift cluster, as it has an unambiguously extended X-ray emission (from Chandra data) testifying the presence of an intracluster medium and deep potential well, that makes it unambigously a cluster. However, it lacks of a spectroscopic redshift: its 68 \% photometric redshifts interval [1.89,2.12]. In this proceeding we first report the presence of a clear red sequence of passive galaxies in the region of JKCS\,041 using a filter pair sampling the Balmer break at $z>1.2$: the $z'-J$ colour. Second, by comparing the red sequence colour of the JKCS\,041 and IRC0218A, a spectroscopically confirmed cluster at $z\sim1.62$ (Papovich et al. 2010: Tanaka et al. 2010), we confirm that JKCS\,041 lies at much higher redshift, and we quantify its photometric redshift ($z\sim 2.2$). More details can be found in Andreon \& Huertas-Company (2010). Throughout the paper, we assume the following cosmological parameters: $H_0=70$ km s$^{-1}$ Mpc$^{-1}$, $\Omega_m=0.3$ and $\Omega_\Lambda=0.7$. Magnitudes are in the AB system. \begin{figure*} \centerline{\psfig{figure=JKCS041_CM_proceed.ps,width=9truecm,clip=}} \caption{Colour--magnitude plot in the direction of JKCS\,041 (top panel) and toward IRCS0218A (bottom panel). Solid lines shows the fitted colour-magnitude relations of JKCS\,041 (top panel) and IRC0218 (bottom panel) clusters, whereas shading marks the simpler colour range $1.85<z'-J<2.15$. Slanted dashed lines mark magnitudes where the S/N=15 in $z'$, and therefore delimit the region where catalogs should be complete. JKCS\,041 red sequence galaxies are 0.32 mag redder than the red sequence of IRC0218A, at $z=1.62$, indicating that JKCS\,041 is at $z\gg1.62$. } \label{fig:CMRs} \end{figure*}

We show that galaxies $0.32$ mag redder than the red sequence of the $z_{spec}=1.62$ cluster IRC0218A are spatially concentrated where the JKCS\,041 X-ray emission is located. Their red colour implies that the cluster JKCS\,041 is at $z=2.20\pm0.11$, where the uncertainty accounts for uncertainties in stellar synthesis population models, in photometric calibration and in the red sequence colour of both JKCS\,041 and IRC0218A clusters. We can hence confirm that JKCS\,041 is a cluster of galaxies with the photometric redshift $z_{red \ sequence}=2.20\pm0.11$, with a formed potential well, deep enough to hot and retain the intracluster medium, and with a well defined red sequence. Incoming X--ray survey telescopes or red-sequence based surveys will likely return hundreds of $z\sim2$ clusters candidates. Getting spectroscopic redshifts for all of them, or even a small part, is too time consuming with present telescopes. Therefore, photometric redshifts based on the red sequence color will necessarily become very popular in the next years and we need to get used to them. Sunayev-Zeldovich surveys (e.g. High et al. 2010 for SPT; Menanteau et al. 2010 for ACT) are already using galaxy colours to infer the cluster redshift for their small (dozen of clusters) samples at low ($z<1$) redshift.

1012.3051

New deep z'-J data readly show a narrow red sequence co-centered with, and similary concentrated to, the extended X-ray emission of the cluster of galaxies JKCS 041. The JKCS 041 red sequence is 0.32±0.06 mag redder in z'-J than the red sequence of the z<SUB>spec</SUB>=1.62 IRC0218A cluster, putting JKCS 041 at z≫1.62. The colour difference of the two red sequences gives a red-sequence based redshift of z=2.20±0.11 for JKCS 041, where the uncertainty accounts for uncertainties in stellar synthesis population models, in photometric calibration and in the red sequence colour of both JKCS 041 and IRC0218A clusters.

12167294

2011ApJ...726L..21T

1012

1012.1865_arXiv.txt

\label{intro} The ``unification model'' (UM) of Seyfert galaxies proposes that Seyfert 1 (S1) and Seyfert 2 (S2) galaxies are basically the same type of objects viewed from different directions \citep{a93}. In an S1 nucleus, our view of the central engine is relatively unobstructed, allowing a direct observation of the ionizing continuum and broad-line region (BLR). In an S2 nucleus, however, this view is blocked by some form of obscuration, most often thought of as an optically thick torus, so that any broad emission lines (BELs) and ionizing continuum are not directly visible. While this UM has enjoyed considerable success over the years, and is undoubtedly correct for many Seyfert galaxies, questions regarding its universal applicability remain \citep[see, e.g.,][]{wz07}. Of particular interest is if there exists a ``true'' type of S2 that lacks a BLR, and whose appearance is unchanged regardless of orientation. The existence of such true or ``naked'' S2s has been implicated both observationally \citep{t01,pb02,t03,b03,nmm03,haw04} and theoretically \citep{nic00,l03,cze04,es06,eh09,cao10}. It would be of great interest to show definitively that such objects do exist in nature. Recently, a growing number of active galactic nuclei (AGNs), including NGC 3147, have been suggested to be just such objects \citep{gsf07,bi08,pan09,shi10}. In this Letter, we explore three of the best candidates for such type of naked S2s: NGC 3147 \citep{pap01,tw03,bi08}, NGC 4698 \citep{pap01,gz03}, and 1ES 1927+654 \citep{b03}. X-ray observations with {\it ASCA}, {\it ROSAT}, {\it XMM-Newton}, and {\it Chandra} show that all three exhibit characteristics typical of a type-1 view of the active nucleus: little or no intrinsic X-ray absorption from spectral fitting, high hard X-ray to \oiii~ratios indicating low obscuration \citep{b99}, and in the case of 1ES 1927+654, rapid, persistent, and strong X-ray variability observed over a 12 yr timescale. In addition, they are all classified as Compton thin, consistent with little or no intrinsic absorption above Galactic column density ($\sim$10$^{20}-10^{21}$ cm$^{-2}$). Thus, all indications appear to show that we have an unobscured, direct view of these active nuclei. Yet, optically they are classified as low-luminosity S2s, with no discernible sign of broad Balmer emission lines in their spectra \citep{ho97,b03}. Table \ref{oxp} summarizes the observed and derived optical and X-ray properties of the galaxies. Given the type-1 view inferred from the X-ray observations, the lack of BELs in these objects is puzzling and in apparent contradiction with the UM of AGNs. We have therefore conducted a deep search for any weak or hidden BELs in these objects using optical spectropolarimetry and near-infrared (near-IR) spectroscopy. Deep spectropolarimetry was designed to target \ha~to look for any scattered, polarized broad-line component, or hidden broad-line region (HBLR), and near-IR spectroscopy was used to probe deeper through any obscuration to directly uncover any broad permitted lines, such as \pb~and \bg, as has been done by, e.g., \citet{vgh97}. This Letter reports the main results of this search. \begin{deluxetable*}{lccccccccc} \tablecolumns{10} \tabletypesize{\scriptsize} \tablewidth{0pt} \tablecaption{Optical and X-ray Characteristics$^{a}$ \label{oxp}} \tablehead{ \colhead{Object} & \colhead{Type} & \colhead{$z$} & \colhead{$m_{B}$} & \colhead{$f^{b}_{\rm [O~III]}$} & \ha/\hb & \colhead{\nh} & \colhead{$f_{2-10~keV}/f_{[\rm O~III]}$} & log(\mbh) & log(\lbol/\ledd) \\ \colhead{} & \colhead{} & \colhead{} & \colhead{} & \colhead{(10$^{-15}$ erg cm$^{-2}$ s$^{-1}$)} & \colhead{} & \colhead{(cm$^{-2}$)} & \colhead{} & \colhead{(\msun)} & \colhead{} } \startdata NGC 3147 & Sey 2 & 0.0094 & 11.4 & 17.2$^{c}$ & 5.23$^{c}$& 1.5 $\times~10^{21}$ & $\sim$40 & 8.64 & ~~$ -3.05$ \\ NGC 4698 & Sey 2 & 0.0034 & 11.5 & 11.2~ & 3.12~ & 5 $\times~10^{20}$ & 1--3 & 7.43 & ~~$ -3.39$ \\ 1ES 1927+654 & Sey 2 & 0.019~ & 15.4 & 3.16 & 4.15$^{d}$ & 7.3 $\times~10^{20}$ & $\approx 800^{e}$ & 7.34 & $\sim$$-2.23~$ \\ \enddata \tablenotetext{a}{X-ray data are from \citet{tw03}, \citet{gz03}, and \citet{b03}. Unless otherwise noted, optical data are from this study. A cosmology of $H_o$ = 70 \hubu, $q_o$ = 0.5, and $\Lambda$ = 0 is assumed.} \tablenotetext{b}{Observed \oiii~\wave 5007 flux uncorrected for extinction measured from galaxy-subtracted spectrum.} \tablenotetext{c}{From \citet{ho97}.} \tablenotetext{d}{Due to the strong absorption at \hb~in the nuclear spectrum, only an upper limit can be determined \citep{b03}. This value is measured from our extended spectrum just off the nucleus, free of \hb~absorption.} \tablenotetext{e}{Only a 0.3-7 keV flux is available \citep{b03}. Using the photon index $\Gamma = 2.7$, we estimate the hard X-ray (2-10 keV) flux, $f_{2-10~keV}$ $\sim$ 0.2$f_{0.3-7~keV}$.} \end{deluxetable*}

NGC 3147, NGC 4698, and 1ES 1927+654 are three S2s with an unusual combination of properties: X-ray spectra show variability and little absorption indicative of a type-1 (direct) view, but optical spectra show only narrow emission lines, typical of a type-2 (obscured) view of the nucleus. A deep search for hidden BLR using Keck LRIS spectropolarimetry and direct near-IR spectroscopy with NIRSPEC does not reveal any BELs, hidden or direct. If typical broad lines were present, their non-detections would indicate an extinction of \av~$\sim$ 11$-$26, inconsistent with the ``naked'' nature of these galaxies. While the obscuration may be due to different material for X-ray and optical light, it appears plausible that the BLRs in these objects are anemically small or absent, due to the weakness of their active central engines.

1012.1865

NGC 3147, NGC 4698, and 1ES 1927+654 are active galaxies that are classified as Seyfert 2s, based on the line ratios of strong narrow emission lines in their optical spectra. However, they exhibit rapid X-ray spectral variability and/or little indication of obscuration in X-ray spectral fitting, contrary to expectation from the active galactic nucleus (AGN) unification model. Using optical spectropolarimetry with LRIS and near-infrared spectroscopy with NIRSPEC at the W. M. Keck Observatory, we conducted a deep search for hidden polarized broad Hα and direct broad Paβ or Brγ emission lines in these objects. We found no evidence for any broad emission lines from the active nuclei of these galaxies, suggesting that they are unobscured, completely "naked" AGNs that intrinsically lack broad-line regions.

12207140

2011JCAP...03..042A

1012

1012.3784_arXiv.txt

\label{sc:introduction} One of the outstanding problems of standard cosmology is to explain the origin of the asymmetry between matter and antimatter. Without such an asymmetry, all the structures we observe today would have never formed and mankind would not exist. The asymmetry can be expressed as the \emph{baryon-to-photon ratio} \be \frac{n_B}{n_\gamma} = (6.21 \pm 0.16 )\cdot 10^{-10} \ee whose numerical value is obtained from a combined analysis of data for large-scale structure and the spectrum of the Cosmic Microwave Background~\cite{WMAP}. In order to obtain a net baryon asymmetry, only three conditions need to be met, as outlined by Sakharov in his seminal paper~\cite{Sakharov}. Yet, providing a model that can successfully explain the measured baryon-to-photon ratio remains a challenging task. Several different scenarios how to realize the Sakharov conditions have been devised~\cite{Baryogenesis}. In the last decade, \emph{leptogenesis}~\cite{Fukugita} has become very popular. The basic idea of most leptogenesis models is to enlarge the particle content of the Standard Model (SM) with \emph{heavy Majorana neutrinos}. In the simplest realization they interact with the SM particles via a Yukawa coupling to ordinary, left-handed leptons and the Higgs bosons as follows: \begin{align}\label{Lint} \mathcal{L}_{\rm int} = h_{ij} \overline{N}_i {\widetilde \varphi}^\dagger \ell_{Lj} + \mbox{ h.c. }, \end{align} where $N$ stands for the Majorana neutrinos, $\ell_L$ and $\varphi$ are the left-handed lepton doublet and the Higgs doublet, and $\widetilde \varphi \equiv \varepsilon \varphi ^ \ast $ with $ \varepsilon _ { \alpha \beta } = - \varepsilon _ { \beta \alpha }$, $ \varepsilon _ { 12 } = 1 $. Finally, the indices $i,j$ label the fermion families, and $h_{ij}$ is the Yukawa coupling matrix which need not be diagonal. The Majorana neutrinos are unstable and decay both into leptons and antileptons, $N \to \ell \varphi $, $ N \to \bar {\ell} \varphi^\dag$. The CP symmetry is violated and the corresponding decay rates are not equal. Therefore, an excess of antileptons over leptons can be generated. The resulting asymmetry is converted into an excess of baryons over antibaryons via the sphaleron transitions which conserve $B - L$ but violate $B + L$~\cite{Shaposhnikov}. In addition to providing a source for the measured baryon asymmetry, this scenario offers a framework to explain the smallness of the neutrino masses via the \emph{seesaw mechanism}~\cite{seesaw}. This twofold virtue is what makes the scenario of leptogenesis particularly appealing. Despite a substantial amount of work and progress~\cite{CP,Giudice,Kiessig,CTP} a complete theory of leptogenesis is still lacking \cite{anisimov}. In this paper we study a type of processes which so far has not been considered in this context. We show that they contribute to the production of heavy Majorana neutrinos at leading order in the coupling constants, and we find that their contribution is numerically large. Our results therefore constitute an important step towards a complete treatment of thermal leptogenesis. We compute the production rate of a Majorana neutrino in a hot electroweak plasma that is fully equilibrated, except for the Majorana neutrinos themselves. The production rate is part of the network of Boltzmann equations that is solved to obtain the baryon asymmetry. Since it sets the initial conditions for leptogenesis, it is also of practical interest to study the production rate by itself. We assume that the number density of Majorana neutrinos is small compared to the equilibrium density, so that the inverse processes which reduce their number density can be neglected. We focus on the lightest Majorana neutrinos $N_1 \equiv N$, which we assume to be the dominant source of lepton asymmetry. When the temperature $ T $ is sufficiently above the Majorana neutrino mass $ M _ N $, a peculiar type of production mechanism occurs. It was already considered in Refs.~\cite{Giudice,Kiessig}. Interactions with the hot plasma generate thermal masses, which are much bigger for SM particles than for the Majorana neutrinos. Therefore the SM particles can become ``heavier'' than the Majorana neutrino, and the decay of a Higgs boson into Majorana neutrino and SM lepton can become possible. Since thermal masses are parametrically small compared to the typical particle momentum, all momenta involved in this decay are nearly collinear. In this paper we show that there are additional nearly collinear processes, involving soft electroweak gauge interactions, which contribute to the leading order production rate. We focus on the leading order in the SU(2) and U(1) gauge couplings $g$ and $g'$, the top quark Yukawa coupling constant $h _ t$ and the Higgs self-coupling $\lambda$. We do not consider the production via $ 2 \leftrightarrow 2 $ scattering processes \footnote{In the literature one can find several calculations of $ 2 \leftrightarrow 2 $ scattering rates (see, e.g. \cite{Giudice,pilaftsis,pedestrians,davidson,hahn-woernle}). However, to the best of our knowledge, there is no calculation which consistently treats all leading order thermal effects.}. For the power counting we assume that all these couplings are of the same order and collectively refer to them by $g$. All other SM couplings are neglected. We perform the computation in the high-temperature regime where $M_N \ll T$. This allows us to formally treat the mass of the Majorana neutrino as being soft, $ M_N \sim g T $, and therefore parametrically of the same order as the thermal Higgs and lepton masses. We demonstrate that even at leading order the production cannot simply be understood in terms of scattering processes involving only a handful of particles. '$ N $-strahlung' and inverse decay processes involving multiple interactions mediated by soft electroweak gauge bosons are not suppressed despite the large number of vertices. The emission occurs almost collinearly, so that propagators are nearly on-shell and compensate the suppression. In position space the radiated particle and its source overlap over large distances, and the interference of different interactions cannot be neglected~\cite{landaumigdal}. This phenomenon has been studied in various contexts such as parton energy loss \cite{baier}, photon~\cite{AurencheLPM,arnoldPhoton,BesakBodeker} and gluon~\cite{arnoldGluon} production in a quark-gluon plasma (for a general discussion see~\cite{arnoldKinetic}). Recently~\cite{BesakBodeker}, we presented a new approach how to consistently include soft gauge interactions in the computation of a thermal particle production rate. It is formulated in a way that is largely independent of the type of particles whose production we want to study, and can therefore easily be adapted to the case at hand. The paper is organized as follows. In Sec.~\ref{sc:rate} we relate the production rate, which describes out-of-equilibrium physics, to a real-time correlation function in equilibrium. The latter can be calculated in thermal field theory, which is subject to the following sections. We describe the physics of collinear emission and outline the relevant momentum scales in Sects.~\ref{sc:processes} and \ref{sc:kinematics}. In Sec.~\ref{sc:strategy} we give a short and qualitative summary how we proceed to obtain the leading order production rate due to collinear emission processes. The rest of Sec.~\ref{sc:LPM} provides all the technical details that are needed to arrive at the final results, and the reader who is not interested in the details of their derivation may skip directly to Eq.~\eqref{ratepsif}. In Sec.~\ref{sc:Numerical} we present numerical results and we conclude in Sec.~\ref{sc:conclusions}. The appendix finally explains how to obtain the numerical solutions.

\label{sc:conclusions} In this paper we have studied the production of relativistic Majorana neutrinos, relevant for models of thermal leptogenesis, in a hot electroweak plasma. Based on our previous work~\cite{BesakBodeker}, we have obtained an equation which sums all leading order collinear production processes. It includes the tree-level processes--the decay of the Higgs boson as well as the inverse decay of the Majorana neutrino--and in addition processes involving multiple scattering mediated by soft gauge boson exchange. All these turn out to contribute at leading order in the coupling constants and have never been included in previous treatments of thermal leptogenesis. Numerically, we find a very pronounced increase in the thermal production rate when soft gauge interactions are included. At high temperatures, when the production via tree-level Higgs decay is allowed, the rate increases by about a factor 3. When the tree-level processes are forbidden, the rate in units of temperature drops only slightly, and the production of Majorana neutrinos remains effective. When the temperature is close to $M_N$ the rate is not significantly affected by soft gauge interactions. We further showed that the production rate is dominated by helicity changing processes and that helicity conserving (inverse) decays only play a significant role at low temperatures, $T \sim M_N$. In addition to the production rate, we have studied the momentum spectrum of the produced Majorana neutrinos and found that it is strongly peaked in the infrared. The spectrum is thus not even approximately thermal. The leading order production rate of heavy Majorana neutrinos also receives contributions from $ 2 \leftrightarrow 2 $ scattering processes. They have been considered previously in the context of thermal leptogenesis (see e.g.\ \cite{Giudice,pilaftsis,pedestrians,davidson,hahn-woernle}), but a complete and consistent leading order calculation has not been done so far. In \cite{Giudice} the Higgs decay was found to do\-minate over the $ 2 \leftrightarrow 2 $ processes at high temperature. In this paper we have shown that adding soft gauge interactions to the Higgs decay process leads to a strong enhancement of the production rate. Therefore it would be very interesting to see how the processes discussed here, together with the $ 2 \leftrightarrow 2 $ scattering processes affect various scenarios of leptogenesis. \vspace{.5cm} {\bf Acknowledgments} This work was supported in part through the DFG funded Graduate School GRK 881. \appendix \renewcommand{\theequation}{\thesection.\arabic{equation}} \setcounter{equation}0

1012.3784

The production rate of heavy Majorana neutrinos is relevant for models of thermal leptogenesis in the early Universe. In the high temperature limit the production can proceed via the 1leftrightarrow2 (inverse) decays which are allowed by the thermal masses. We consider new production mechanisms which are obtained by including additional soft gauge interactions with the plasma. We show that an arbitrary number of such interactions gives leading order contributions, and we sum all of them. The rate turns out to be smooth in the region where the 1leftrightarrow2 processes are kinematically forbidden. At higher temperature it is enhanced by a factor 3 compared to the 1leftrightarrow2 rate.

12104333

2010ApJ...724..386K

1012

1012.1112_arXiv.txt

Supermassive black holes (SMBHs), found at centers of massive spheroids and active galaxies, are considered to play an important role in the formation and the evolution of galaxies. It is suggested that SMBHs regulate the star formation activities of galaxies through their enormous energy output during active phase, providing a feedback mechanism to reconcile the observed trend of downsizing of the galaxy evolution with hierarchical galaxy formation models \citep{juneau05,bundy06,schawinski06}. However, the coevolution of SMBHs and their host galaxies stand as an unsolved astrophysical problem, as such a process needs to eventually lead to a rather unexpected tight correlation between host galaxy properties and SMBH mass today \citep{ferrarese00,gebhardt00,tremaine02}. One of the great challenges toward the understanding of the evolutionary sequence of AGNs and the co-evolution of SMBHs and their host galaxies comes from the so-called red or dusty AGNs which are believed to occupy more than 50$\%$ of the AGN population \citep{comastri01,tozzi06,polletta08}. Red, dusty AGNs are expected, if the initial phase of AGN activity starts in a dust-enshrouded environment such as within massive starburst regions of luminous infrared galaxies. In such a scenario, AGNs become more visible after sweeping away cold gas and dust that are necessary to sustain the massive star formation in their host galaxies \citep{hopkins06}. On the other hand, dust obscuration can arise in a different way in the unified model of AGN when accretion disks and broad line regions (BLRs) around SMBHs are viewed through dust torus. In any case, red, dusty AGNs can shed light on the properties and the evolution of the AGN population in general. The phenomenological definition of ``red, dusty AGNs'' is rather broad. It can include AGNs selected in a variety of ways such as those selected by very red colors in optical through NIR and the radio detection (e.g., $R-K > 5$ mag and $J-K > 1.3$ mag of the sample in Urrutia et al. 2009; also see Cutri et al. 2001), red MIR colors (Lacy et al. 2004; Lee et al. 2008), and hard X-ray detections (e.g., Polleta et al. 2007). These different kinds of AGNs have a common characteristic where their SEDs are red, due to the obscuration of their light by the foreground gas and dust. Hence, they are considered to be the intermediate population between the dust-enshrouded star forming galaxies and the unobscured AGNs. The dust-obscuration does not necessarily exclude Type-1 AGNs (AGNs with broad emission lines; e.g., Alonso-Herrero et al. 2006), as red, dusty AGNs with broad emission lines are found quite often. Of 23 X-ray QSOs with IR power-law SED and spectroscopic redshift identification, 14 are classified as broad-line AGNs (Szokoly et al. 2004; Alonso-Herrero et al. 2006). More than 50\% of the red AGNs with the radio and red optical/NIR color selection are found to be type-1 AGN which can have the reddening parameter of 2 or more (Uruttia et al. 2009; Glikman et al. 2007). The existence of broad line AGNs among red, dusty AGNs opens a possibility of studying the properties of the dust-obscured quasars in more detail using traditional type-1 AGN diagnostics. However, even the measurement of the most basic AGN parameters such as $M_{\rm BH}$ and the Eddington ratio (the accretion rate) is a difficult task for the red, dusty AGNs. In general, $M_{\rm BH}$ are derived from the optical or the ultraviolet (UV) part of AGN spectra for which spectral diagnostics are well established. In the case of dusty AGNs, the dust obscures or significantly reduces the UV or the optical light coming from the region around SMBHs, making the popular AGN optical/UV spectroscopic diagnostics useless. For example, popular $M_{\rm BH}$ estimators are based upon a virial relation between two parameters - the velocity width of the broad, H$\alpha$ or H$\beta$ lines, and the size of BLR estimated from the continuum luminosity at 5100 \AA ${}$ \citep{kaspi00} or the luminosities of H$\alpha$ or H$\beta$ \citep{greene05}. If the light from a red, dusty AGN is extincted by a color excess of $E(B-V)$ =2 mag \citep{glikman07,urrutia09}, its H$\alpha$ and H$\beta$ line fluxes would be suppressed by a factor of 100 and 1000 respectively. One can try to estimate the amount of the dust extinction from a continuum fitting of the optical-UV spectrum or through the Balmer decrement, but such estimates are often inconsistent with each other and accompanied with uncertainties of order of $\delta E(B-V) \sim 0.5$ mag or more which are related to the dispersion in the intrinsic properties of AGNs. The dust obscuration can arise from the Galacitic extinction for AGNs at low galactic latitude (e.g., Im et al. 2007; Lee et al. 2008). The problem can be substantially alleviated if we can use NIR lines instead of optical or UV lines, since NIR Hydrogen lines such as the P$\alpha$ and P$\beta$ are much less affected by the dust extinction than the UV/optical light. For the red, dusty AGN with the color excess of $E(B-V)$ = 2 mag, the line fluxes of P$\alpha$ and P$\beta$ are suppressed by a factor of only 2.3 and 4.7 respectively. This is a significant improvement over the optical lines. Since $M_{\rm BH} \propto L^{0.5}$, where $L$ is a luminosity of the continuum or an emission line, the suppression in the Paschen line luminosities introduces uncertainties in $M_{\rm BH}$ estimates only at the level of a factor of 2 or less, even without correcting for the dust extinction with the Paschen decrement. In this paper, we will derive the $M_{\rm BH}$ estimators based on the NIR Hydrogen lines of Type-1 AGNs and investigate the line ratios of the Paschen lines, keeping in mind future applications of such relations to studies of dusty, red AGNs with broad emission lines.

We derived new $M_{\rm BH}$ estimators, using Hydrogen P$\alpha$ and P$\beta$ lines of Type-1 AGNs. The derived estimator allows the determination of $M_{\rm BH}$ at the accuracy of $\sim 0.2$ dex, and they will be useful for estimating $M_{\rm BH}$ of red, dusty AGNs. Our analysis of the Paschen lines with respect to H$\alpha$ and H$\beta$ lines shows that the luminosities and FWHMs of the broad components of the Paschen lines correlate well with those of the Balmer lines. The Hydrogen line ratios from H$\beta$ through P$\alpha$ are consistent with a Case B recombination with the parameters of $\alpha= -1.0$, $U= 10^{-1.5}$ and $n= 10^9~\mathrm{cm^{-3}}$. The mean line ratios are $\mathrm{\frac{H\beta}{P\alpha}} = 2.70$, $\mathrm{\frac{H\alpha}{P\alpha}} = 9.00$, and $\mathrm{\frac{P\beta}{P\alpha}} = 0.91$ which can be used to estimate for the amount of dust extinction present in red AGNs in future. Future applications of these results on red, dusty AGNs will enable us to better understand the nature of red, dusty AGNs whose properties are still hidden behind a wall of dusty gas.

1012.1112

More than 50% of active galactic nuclei (AGNs) are suspected to be red and affected by dust obscuration. Meanwhile, popular spectral diagnostics of AGNs are based on optical or ultraviolet light, making dust obscuration the primary concern for understanding the general nature of AGNs and the supermassive black holes residing in them. To provide a method for investigating properties of dusty AGNs, we derive new black hole (BH) mass estimators based on velocity widths and luminosities of near-infrared (NIR) hydrogen emission lines such as Pα and Pβ, and also investigate the line ratios of these hydrogen lines. To derive the BH mass (M <SUB>BH</SUB>) estimators, we used a sample of 37 unobscured type 1 AGNs with an M <SUB>BH</SUB> range of 10<SUP>6.8</SUP>-10<SUP>9.4</SUP> M <SUB>sun</SUB>, where M <SUB>BH</SUB> comes from either the reverberation mapping method or single-epoch measurement method using Balmer lines. Our work shows that M <SUB>BH</SUB> can be estimated from the Paschen line luminosities and velocity widths to an accuracy of 0.18-0.24 dex (rms scatter). We also show that the mean line ratios of the Paschen lines and the Balmer lines are {Hα}/{Pα}} ≃ 9.00, {Hβ}/{Pα}} ≃ 2.70, which are consistent with case B recombination under a typical AGN broad-line region (BLR) environment. These ratios can be used as reference points when estimating the amount of dust extinction over the BLR for red AGNs. We expect the future application of the new BH mass estimators on red, dusty AGNs to provide a fresh view of obscured AGNs.

12137740

2010idm..confE.124S

1012

1012.3926_arXiv.txt

1012.3926

The determination of dark matter constraints from liquid xenon direct detection experiments depends upon the amount of scintillation light produced by nuclear recoils in the detector, a quantity that is characterized by the scintillation efficiency factor Leff. We examine how uncertainties in the measurements of Leff and the extrapolated behavior of Leff at low recoil energies (where measurements do not exist) affect the constraints from experiments such as XENON10 and XENON100, particularly in the light WIMP regions of interest for the DAMA and CoGeNT experimental results.

12165908

2011AN....332...83K

1012

1012.3321_arXiv.txt

Arcturus is a single giant star of spectral type K2 with an effective temperature of 4300K (Griffin \& Lynas-Gray 1999). Gray \& Brown (2006) found a two year modulation in the velocity span of the bisector of the Fe I $\lambda$6252.57 line which they interpret as the rotation period of the star. The two year period has also been found in the Ca II emission of Arcturus which has been monitored by the Mount Wilson H+K project since 1984 (Brown et al.~2008). The data also shows a variation on a longer time scale that could be an activity cycle with a period $\le$ 14 years. The so observed rotation period varies with an amplitude of 70days, changing by 20 days/year. A similar time dependence is found in the solar Ca II H+K emission, where it is a consequence of the differential rotation. As the activity moves to lower latitudes during the cycle, the rotation period decreases. By analogy, the variation in the rotation period of Arcturus can be interpreted as a variation of the active latitude on a differentially rotating stellar surface. Stellar butterfly diagrams are ambiguous, however, as a combination of anti-solar rotation and a polewards drift of the active latitude will produce the same pattern as if dynamo and rotation were both solar-type. The solar butterfly diagram can be explained as the result of an $\alpha \Omega$ dynamo, i.e.~a combination of differential rotation and the $\alpha$ effect caused by the helicity of the convective gas motions. However, significant radial shear is only found in the subsurface layer and in the tachocline at the bottom of the convection zone. More recent dynamo models therefore explain the butterfly diagram with the advection of magnetic flux by the large-scale meridional flow (cf.~Ossendrijver 2003, Charbonneau 2005). Flux transport by the meridional flow is only effective if the associated magnetic Reynolds number is large, i.e.~ \begin{equation} {\rm Rm} = \frac{u^{\rm m} d}{\eta} \gg 1 \end{equation} where $u^{\rm m}$ is the meridional flow speed, $d$ a characteristic length scale (e.g.~the stellar radius) and $\eta$ the magnetic diffusivity coefficient (K\"uker et al.~ 2001). The solar differential rotation is the result of angular momentum transport by convection and the large-scale meridional flow in the convection zone. On the one hand stratification causes a non-diffusive contribution to the Reynolds stress in addition to the diffusive part known as turbulence viscosity, on the other hand the Coriolis force causes a deviation of the convective heat flux from the radial direction. The result is a small horizontal temperature gradient that drives a meridional circulation. Mean field models of the solar convection zone reproduce the observed rotation pattern and surface meridional flow very well and allow predictions for other stars (Kitchatinov \& R\"udiger 1995, 1999, K\"uker \& Stix 2001). To examine the conditions for dynamo action on Arcturus, we apply our stellar rotation model to its convective envelope. Like solar-type stars, giants have deep outer convection zones. One might therefore expect similar rotation and magnetic activity patterns. There are, however, significant differences. With an equatorial rotation period of 25 days the Sun rotates much faster than Arcturus and with an effective temperature of 5780 K it is substantially hotter. Interestingly, the radius of Arcturus is larger by about the same factor by which its rotation period is longer, yielding about the same equatorial rotation speed (1.8 km/s vs.~2km/s). Finally, the geometrical depth of the outer convection zone is much larger than that of the Sun, both in absolute values and relative to the stellar radius. In the model described below it reaches down to three percent of the stellar radius, making the star almost fully-convective.

Despite the great geometrical depth, the properties of the convection zone of Arcturus are closer to those of the solar granulation and supergranulation layers than the deeper parts of the solar convection zone because of the long rotation period. With two years the latter is an order of magnitude longer than the convective turnover time, which reaches a maximum value of 85 days at a fractional radius of 0.4. Except for the layer around that radius the Coriolis number is less than one, making Arcturus a slow rotator in this context. The rotation pattern shows a strong variation with depth. For fractional radii greater than 0.2 there is a moderate decrease of the rotation rate with increasing radius and "solar-type" horizontal shear, i.e. the rotation rate is larger at the equator than at the poles. While the latter difference is large in relative terms it is only about 1/10th of the solar surface differential rotation in absolute terms. The most striking feature is the steep decline of the rotation rate with increasing radius at the lower boundary. Such a rotation profile would imply a core that rotates much faster than the convective envelope. A rotation rate that decreases with increasing radius is the natural consequence of angular momentum transport by Reynolds stress. For slow rotation the radial transport dominates and creates a negative gradient in the rotation rate, as is observed in the outermost layers of the solar convection zone. The meridional flow is much different from that found for the Sun. Both observations and mean field models find a surface flow of about 20 m/s amplitude towards the poles. The return flow has not been observed yet but is predicted by theory to occur at the bottom of the convection zone with about half the amplitude of the surface flow. Our model for Arcturus shows two flow-cells per hemisphere rather than a solar-like one-cell pattern. There is a large cell of fast flow at high latitudes and a smaller cell with lower flow speeds at low latitudes. The high-latitude flow is anti-solar, i.e. directed towards the equator at the surface while the slower flow at low latitude is solar-type. The situation for giant stars such as Arcturus is much different from that for Main Sequence stars as for the latter the mass is fixed by the effective temperature. For single giants there is a much larger uncertainty. Our alternative 1 $M_\odot$ model shows, however, that the differential rotation does not depend very strongly on the stellar mass provided temperature and radius are kept constant. Our model does not include the stably stratified core. In the Sun, the surface differential rotation persists throughout the whole convection zone while the core rotates rigidly with the same rotation rate as the surface at mid-latitudes. The transition between the two patterns occurs in a shallow layer at the bottom of the convection zone, the tachocline. It has been the subject of much debate both about the mechanism that causes such a sharp transition and the role it may play in the solar dynamo. With Arcturus we find a different situation. The horizontal shear ($\partial \Omega/\partial \theta$) disappears in the lower part of the convective envelope and a strong radial gradient appears. As the rotation is (horizontally) uniform at the lower boundary, a solar-type tachocline can not exist. The sharp decrease of the rotation rate with radius raises the question whether this rotation profile is dynamically stable. We can exclude the Taylor-Couette instability as the rotation rate decreases roughly like $1/\sqrt{\Omega}$, implying hydrodynamic stability. However, MHD instability caused by a sufficiently strong toroidal field can not be ruled out without a detailed dynamo model. The rotation pattern we find for Arcturus looks similar to the one assumed by early models of the solar dynamo. The combination of negative radial shear and the positive (negative) $\alpha$ effect that is expected in the northern (southern) hemisphere of a stellar convection zone (R\"udiger \& Kichatinov 1992) naturally produces a solar type butterfly diagram through a classical $\alpha \Omega$ dynamo without the need for a meridional circulation (Yoshimura 1975). This prompts us to have a look at the conditions for dynamo action. As the star is almost fully convective and has no horizontal shear at the bottom of the convection zone, there is no tachocline. Moreover, the turbulent magnetic diffusivity computed from the mean field model is 2--3 orders of magnitude larger than the corresponding value for the solar convection zone, as shown in Fig.~\ref{etaturb}. With meridional flow speeds of the order 100 m/s and a length scale of $2 \times 10^{12}$cm we find Rm$ \approx 2$. For the Sun a similar estimate yields a value of about 20, assuming a magnetic diffusivity coefficient of $10^{13}$ cm$^2$/s. Hence, conditions for a flux-transport dynamo are probably less favorable on Arcturus than on the Sun. There is some uncertainty about the turbulent magnetic diffusivity in mean field dynamos. Current models of the solar dynamo use smaller values than predicted by the SOCA. Dikpati \& Charbonneau (1999) assume $10^{10}$--$10^{11}$ cm$^2$/s. Chatterjee et al.~(2004) use similar values for the toroidal field but larger values of the order $10^{12}$ cm$^2$/s for the poloidal field. To create similarly favorable conditions for the flux transport dynamo in Arcturus we have to lower the magnetic diffusivity by as much as three orders of magnitude to $10^{13}$ cm$^2$/s to reach magnetic Reynolds numbers of the order 1000, as required for this type of dynamo. \\[5mm] {\it Acknowledgements:} This work was supported by Deutsche Forschungsgemeinschaft.

1012.3321

The spectroscopic variability of Arcturus hints at cyclic activity cycle and differential rotation. This could provide a test of current theoretical models of solar and stellar dynamos. To examine the applicability of current models of the flux transport dynamo to Arcturus, we compute a mean-field model for its internal rotation, meridional flow, and convective heat transport in the convective envelope. We then compare the conditions for dynamo action with those on the Sun. We find solar-type surface rotation with about 1/10th of the shear found on the solar surface. The rotation rate increases monotonically with depth at all latitudes throughout the whole convection zone. In the lower part of the convection zone the horizontal shear vanishes and there is a strong radial gradient. The surface meridional flow has maximum speed of 170 m/s and is directed towards the equator at high and towards the poles at low latitudes. Turbulent magnetic diffusivity is of the order 10<SUP>15</SUP>-10<SUP>16</SUP> cm<SUP>2</SUP> /s. The conditions on Arcturus are not favorable for a circulation-dominated dynamo.

2775551

2011ApJ...728..119W

1012

1012.1324_arXiv.txt

\label{sec:intro} Interactions between galaxies are an important process in the evolution of galaxies with cosmic time, as they can affect various galaxy properties such as star formation rate, morphology, and gas fraction. Simulations have shown that major interactions and mergers between galaxies can produce disturbed morphologies and trigger starbursts \citep[e.g.][]{toomre1972,barnes1991,barnes1992,mihos1992,kauffmann1993,mihos1994,mihos1996,springel2000,tissera2002,bundy2005,cox2006,dimatteo2007,dimatteo2008,lotz2008}. Interacting and merging systems can also trigger intense infrared emission in gas-rich galaxies, resulting in the formation of luminous and ultra-luminous infrared galaxies \citep[LIRGs and ULIRGs;][and references therein]{sanders1996}. The increase in the incidence of LIRGs and ULIRGs at intermediate redshifts \citep{lefloch2005,rujopakarn2010} suggests that the decline in the global star formation rate since redshift $z \sim 1$ \citep{lilly1996,madau1996,chary2001,perezgonzalez2005} may be caused by a decrease in the amount of star formation triggered by interactions over time. Past studies have shown that star formation triggered by major mergers is not a significant fraction of the overall star formation rate at intermediate redshifts, however. Using data from the Galaxy Evolution from Morphology and SEDs survey \citep[GEMS;][]{rix2004} and Classifying Objects by Medium-Band Observations in 17 Filters \citep[COMBO-17;][]{wolf2001,wolf2004} survey, \citet{wolf2005} found that morphologically-identified merging galaxies contribute roughly 20\% of the ultraviolet luminosity density (and thus, star formation rate density) at $z \sim 0.7$. \citet{robaina2009} used data from GEMS and COMBO-17, along with data from the Space Telescope A901/2 Galaxy Evolution Survey \citep[STAGES;][]{gray2009}, to show that $\lesssim 10$\% of star formation at $0.4 \leq z \leq 0.8$ is triggered by these major interactions. The conclusions drawn from these observational studies are in general agreement with the results of simulations performed by \citet{hopkins2010}, who find that only $\sim5-10$\% of the star formation rate density out to $z \sim 6$ is the result of merger-induced starbursts. While star formation triggered by major interactions and mergers may not be a significant contributor to the evolution of the global star formation density, it is also important to examine the properties of the likely progenitor population $-$ close galaxy pairs \citep[e.g.][]{patton2000,depropris2007} $-$ to investigate whether more frequent tidal interactions in close pairs could play a role. Searches for close galaxy pairs have been used by many independent studies to derive the evolution of galaxy merger rates out to intermediate redshifts \citep[e.g.][]{burkey1994,carlberg1994,patton1997,patton2002,lefevre2000,lin2004,lin2008,bell2006,kartaltepe2007,deravel2009}. Star formation triggered by these tidal interactions may also consume much of the galaxies' cold gas, which could be a key factor in the transformation of blue, star-forming galaxies to red, quiescent galaxies. We expect it would be a bigger factor in dense environments where these interactions are more common, so tidally-triggered star formation may also be an important contributor to the redshift evolution of galaxies in different environments. The first indication that tidal interactions could affect galaxy properties was found by \citet{larson1978}, who observed that interacting systems identified by peculiar morphologies showed much greater scatter in their optical color distributions in comparison to morphologically-normal galaxies. This was interpreted to be the result of a recent burst of star formation triggered by the tidal interaction. Subsequent studies \citep[e.g.][]{condon1982,keel1985,kennicutt1987} found similar results indicating that rapid bursts of star formation were associated with tidal interactions between galaxies. Most studies of tidally-triggered star formation in close galaxy pairs have been performed at low redshift ($z \lesssim 0.1$) using H$\alpha$ emission as the main diagnostic of star formation rate. These studies have indicated that there is an overall enhancement in the star formation rate of close pairs of galaxies relative to a similar population of isolated galaxies, and that the strength of the enhancement is anticorrelated with the pair separation. One of these early studies by \citet{kennicutt1987} found an enhancement in both H$\alpha$ and far-infrared emission in interacting galaxies compared to isolated galaxies, accounting for $\sim 6$\% of the total massive-star formation in luminous star-forming galaxies. However, they found that the degree of enhancement had a large variation and that they were strongly biased toward unusually bright and active systems. Thus, they were unable to draw conclusions about more typical interacting galaxy pairs. In order to isolate the effects of star formation triggered by tidal interactions, it is important to take into account the local environments of the galaxies in question. Dense environments, such as galaxy groups and clusters, are the most likely locations of tidal galaxy-galaxy interactions. However, these dense environments have a higher fraction of red elliptical galaxies with little ongoing star formation \citep{dressler1980,postman1984,hermit1996,guzzo1997,giuricin2001,gomez2003,kauffmann2004,blanton2005}. It has been shown that tidally-triggered star formation in galaxy pairs is more pronounced at low densities, likely due to the fact that galaxies in dense environments have had their gas content exhausted by previous interactions \citep{solalonso2006}. \citet{barton2007} used simulations to find that a significant fraction of close pairs tended to lie in host dark matter halos that contained more galaxies than just the two in the pair. They established that in order to isolate triggered star formation, a sample of isolated pairs must be constructed and compared to a sample of isolated galaxies in the same sparse environments. \citet{perez2009a} investigated the effects of various sources of bias that could influence the results derived from a comparison of galaxy pairs to a control sample. Using semi-analytical models, they showed that the local density in the environment of close pairs and the control galaxies is one of the most significant sources of bias and must be accounted for when selecting an appropriate control sample. For observational studies, selecting a large sample of galaxy pairs in low-density regions is difficult due to the low frequency of close pairs that lie in these environments. Only large-scale redshift surveys such as the Sloan Digital Sky Survey \citep[SDSS;][]{york2000} and 2dF Galaxy Redshift Survey \citep[2dFGRS;][]{colless2001} can provide datasets with sufficient volume for statistical studies of galaxy pairs in a variety of environments. The spectroscopic data from these surveys can be used to create a sample of true isolated galaxy pairs by filtering out interloping apparent pairs that are close in projected separation but far apart in redshift, as well as deprojecting nearby galaxies to get a better handle on their local environments. More recent galaxy pair studies \citep{nikolic2004,woods2007,ellison2008,li2008} have taken advantage of these larger samples from the SDSS and found an anticorrelation between specific star formation rate (SSFR) and projected separation, $r_{p}$, between the pairs within $\sim 30~h^{-1} - 100~h^{-1}$ kpc. These works do not explicitly account for the local environments of the pair and control galaxies, allowing them to define large pair samples containing several thousand galaxies. \citet{li2008} did test the relative enhancements in isolated and non-isolated pairs and did not find a significant difference. However, they characterize the environment only in a very small region (projected within $100~h^{-1}$ kpc) around their galaxies, which the results of \citet{barton2007} suggests is not large enough to accurately quantify the local density. \citet{nikolic2004} and \citet{woods2007} both find a stronger enhancement in tidally-triggered star formation for blue star-forming galaxies. Interestingly, \citet{nikolic2004} and \citet{li2008} find little dependence of the enhancement on the relative mass or luminosity of the pair galaxies to their companions, whereas \citet{woods2007} and \citet{ellison2008} find that the relative magnitude of the pair is important in the effectiveness of the tidally-triggered interactions. A similar dependence of star formation rate on projected separation in close pairs has been found by \citet{barton2000} and \citet{woods2006} in the CfA2 redshift survey \citep{geller1989}, and by \citet{lambas2003} in the 2dFGRS. While these studies have made use of large-volume surveys at low redshifts to investigate a large sample of close pairs, comparatively few studies of tidal interactions have been performed at intermediate redshifts, where the effects may be more prominent due to the higher star formation rate density and gas fraction in the universe at those epochs. These intermediate-redshift studies could be valuable in determining whether a decline in star formation triggered by tidal interactions contributes to the decline in the global star formation rate. Only recently have intermediate-redshift surveys with sufficient volume been utilized to examine close pairs in this regime. \citet{lin2007} used data from the DEEP2 survey \citep{davis2003} along with infrared fluxes from Spitzer and found an enhanced SSFR in \rpfifty close pairs and morphologically-identified merger systems, but were unable to draw conclusions about the redshift evolution of the signal between redshifts 0.1 and 1.1. They also found that the anticorrelation between star formation rate and pair separation seen in low-redshift pair studies was seen at higher redshift. \citet{deravel2009} found an enhancement in [O II] luminosity in close pairs out to $z \sim 1$ in the VIMOS VLT Deep Survey \citep[VVDS;][]{lefevre2005}, although they did not control for the local environment of the pairs and focused primarily on the evolution of the merger rate. \citet{woods2010} studied the enhancement of SSFR in the Smithsonian Hectospec Lensing Survey \citep[SHELS;][]{geller2005} at redshifts $0.08 \leq z \leq 0.38$ and found similar trends to the previous low-redshift studies, but they also did not look for redshift evolution in their sample. In this paper, we investigate the relative SSFR of galaxies in isolated close pairs compared to a control sample of isolated galaxies in the Prism Multi-Object Survey \citep[PRIMUS\footnote{http://cass.ucsd.edu/$\sim$acoil/primus/};][Cool et al. in preparation]{coil2010}, a low-dispersion prism intermediate-redshift survey that provides the largest sample of faint galaxy redshifts yet to $z \sim 1$. We use existing UV and optical photometry to infer rest-frame $UV-r$ colors as a proxy for SSFR. One advantage of studying galaxies at intermediate redshifts is that the ultraviolet emission, which measures the young intermediate-to-high-mass ($\gtrsim 5~M_{\odot}$) stellar population of galaxies, is shifted redward, close to or into observed-frame optical wavelengths. This makes it an ideal tracer of star formation over a range of redshifts \citep[][and references therein]{kennicutt1998}, although the UV is more susceptible to attenuation by dust. We have deblended UV data from the Galaxy Evolution Explorer \citep[$GALEX$;][]{martin2005} that we use to determine rest-frame fluxes in both the far-UV ($FUV$; $\lambda_{eff} \sim 1530$ \AA) and near-UV ($NUV$; $\lambda_{eff} \sim 2270$ \AA) bands. Our goal is to study the relative effects of tidal interactions on SSFR at intermediate redshifts ($z \sim 0.5$), where relatively few studies of this nature have been performed in comparison to low redshift studies. We investigate galaxies in pairs with a projected separation of \rpfifty, which has been used in previous studies \citep[e.g.][]{barton2000,woods2010} as a typical distance within which to define interacting systems. We also select a subsample of galaxies in pairs with a projected separation of \rpthirty to look for an increased enhancement in SSFR with decreasing separation, as has been indicated by past studies. We focus on the SSFR rather than the absolute star formation rate so that we are not biased toward intrinsically more massive and luminous galaxies with increasing redshift. We examine how those effects might evolve over a range of redshifts ($0.25 \leq z \leq 0.75$) to investigate whether the overall downward trend of star formation rate in field galaxies is reflected in isolated pair galaxies undergoing tidal interactions. A large redshift survey like PRIMUS ($\sim$120,000 robust galaxy redshifts over 9.1 deg$^{2}$ out to $z \sim 1$) is needed for such a study in order to define a clean sample of isolated pair galaxies that is large enough to compare statistically to an unbiased control sample. This paper is organized as follows. In \S~\ref{sec:data}, we describe the PRIMUS dataset and the cuts we apply to create a clean sample from which to select pair and isolated galaxies. In \S~\ref{sec:color}, we describe our methodology for determining the $FUV$ and $NUV$-band fluxes for our sample. We describe our method for selecting isolated pair galaxies and a corresponding control sample in \S~\ref{sec:sample}. We present our results in \S~\ref{sec:results} and discuss them in \S~\ref{sec:discussion}. We summarize our main conclusions in \S~\ref{sec:conclusions}. Throughout this paper, we assume a $\Lambda$CDM cosmology with $\Omega_{m} = 0.3$, $\Omega_{\Lambda} = 0.7$, and $h = 0.71$. All magnitudes used in this paper are on the AB system. Rest-frame colors are denoted by a preceding superscript zero, e.g. $^{0}(u-r)$.

1012.1324

Tidal interactions between galaxies can trigger star formation, which contributes to the global star formation rate (SFR) density of the universe and could be a factor in the transformation of blue, star-forming galaxies to red, quiescent galaxies over cosmic time. We investigate tidally triggered star formation in isolated close galaxy pairs drawn from the Prism Multi-Object Survey (PRIMUS), a low-dispersion prism redshift survey that has measured ~120,000 robust galaxy redshifts over 9.1 deg<SUP>2</SUP> out to z ~ 1. We select a sample of galaxies in isolated galaxy pairs at redshifts 0.25 &lt;= z &lt;= 0.75, with no other objects within a projected separation of 300 h <SUP>-1</SUP> kpc and Δz/(1 + z) = 0.01, and compare them to a control sample of isolated galaxies to test for systematic differences in their rest-frame FUV - r and NUV - r colors as a proxy for relative specific star formation rates (SSFRs). We find that galaxies in r<SUB>p</SUB> &lt;= 50 h <SUP>-1</SUP> kpc pairs have bluer dust-corrected UV - r colors on average than the control galaxies by -0.134 ± 0.045 mag in FUV - r and -0.075 ± 0.038 mag in NUV - r, corresponding to an ~15%-20% increase in SSFR. This indicates an enhancement in SSFR due to tidal interactions. We also find that this relative enhancement is greater for a subset of r<SUB>p</SUB> &lt;= 30 h <SUP>-1</SUP> kpc pair galaxies, for which the average color offsets are -0.193 ± 0.065 mag in FUV - r and -0.159 ± 0.048 mag in NUV - r, corresponding to an ~25%-30% increase in SSFR. We test for evolution in the enhancement of tidally triggered star formation with redshift across our sample redshift range and find marginal evidence for a decrease in SSFR enhancement from 0.25 &lt;= z &lt;= 0.5 to 0.5 &lt;= z &lt;= 0.75. This indicates that a change in enhanced star formation triggered by tidal interactions in low-density environments is not a contributor to the decline in the global SFR density across this redshift range.

12164728

2011AJ....141...29J

1012

1012.0264_arXiv.txt

Planetary ephemerides are used for multiple purposes including dynamical mass determination for solar system bodies, pulsar timing, high-precision tests of general relativity, and inter-planetary spacecraft navigation. During the past several decades a series of increasingly accurate ephemerides have been developed at the Jet Propulsion Laboratory (JPL) by adding new data types such as Very Large Array (VLA) astrometry (\citet{Muhleman85}; \citet{Muhleman86}), spacecraft tracking \citep{Duxbury89}, and radar range measurements \citep{Campbell78} to historical and modern optical observations. Accurate ephemerides are one of the basic tools of observational astronomy, in the same sense as star catalogs and redshift surveys. They represent a community resource whose value is proportional to their accuracy, and whose accuracy requires regular observational support to maintain and improve. A specific example is the great improvement between timing distances and kinematic distances for pulsars when using the newer DE405 ephemeris \citep{Standish04} compared with the older DE200 ephemeris ({\it e.g.}, \citet{Verbiest08}). The orbits of the inner planets are very accurately tied together with the current data set. For example, the angular ephemeris errors for Mars with respect to Earth are typically 0.2 milli-arcsecond (mas) or 1 nrad \citep{IAU09}. However, the outer planets are not as well tied to the inner planets (or each other) because there have been fewer opportunities to supplement optical observations with high precision spacecraft radio tracking data. The Pioneer and Voyager missions provided essentially single data points during their flybys of the outer planets, and the Galileo mission to Jupiter was severely constrained by the loss of its high gain antenna. This restricted Galileo downlink signals to a relatively low frequency (2.3 GHz) and a low signal/noise ratio. As a result, VLBI observations of Galileo had accuracies of only $\sim$5-10 mas \citep{Jacobson99}. Thus, the Cassini mission to Saturn (http://saturn.jpl.nasa.gov/index.cfm) is our first opportunity to incorporate high-accuracy data from a spacecraft orbiting an outer planet for an extended period. For reference, an angle of 0.1 mas (0.5 nrad) corresponds to about 750 meters at the average distance of Saturn from Earth. Our goal is to improve the position of Saturn in the International Celestial Reference Frame (ICRF, see \citet{Ma98}) through phase-referenced VLBI observations of Cassini using the Very Long Baseline Array (VLBA)\footnote{The VLBA is operated by the National Radio Astronomy Observatory (NRAO).} at 8.4 GHz (X band) combined with Cassini orbit determinations. The Cassini orbit can be determined to about 2 km at apoapse and 0.1 km at periapse relative to the center of mass of Saturn with range and Doppler tracking by the Deep Space Network \citep{Antreasian08}. The future Juno mission to Jupiter should allow a similar application of phase-referenced VLBI to improve the Jupiter ephemeris. Combined with our new data for Saturn this will lead to a better model for the gravitational interactions, and the orbital evolution, among all the outer planets. A covariance analysis \citep{Standish06} at JPL shows that with only a few years of VLBA data, the ephemeris improvement for Saturn extends for decades.

We have demonstrated repeatable phase referenced astrometry of the Cassini spacecraft using the VLBA, and verified consistent position determinations from direct imaging and from total delay measurements. Future observations will increase the time span of accurate position measurements, leading to ever improving constraints on the planetary ephemeris. In addition, continuing improvement in the accuracy of phase reference source positions will allow a more accurate tie of our Cassini positions to the ICRF. The Cassini mission has recently been extended until 2017, with further extensions likely in the future. By extending our VLBI observations beyond 2012 we will have high accuracy measurements over more than a quarter of Saturn's orbital period. The error in determining the plane of Saturn's orbit (latitude) decreases rapidly as the time span of observations approaches 1/4 of the orbital period. The error in longitude decreases approximately linearly with time span. The next mission to the outer planets will be the JUNO mission to Jupiter. This orbiting mission will provide an opportunity to use the same phase referenced astrometry techniques with the VLBA, and thereby improve the ephemeris of Jupiter in a similar manner.

1012.0264

The planetary ephemeris is an essential tool for interplanetary spacecraft navigation, studies of solar system dynamics (including, for example, barycenter corrections for pulsar timing ephemerides), the prediction of occultations, and tests of general relativity. We are carrying out a series of astrometric very long baseline interferometry observations of the Cassini spacecraft currently in orbit around Saturn, using the Very Long Baseline Array (VLBA). These observations provide positions for the center of mass of Saturn in the International Celestial Reference Frame (ICRF) with accuracies ~0.3 mas (1.5 nrad) or about 2 km at the average distance of Saturn. This paper reports results from eight observing epochs between 2006 October and 2009 April. These data are combined with two VLBA observations by other investigators in 2004 and a Cassini-based gravitational deflection measurement by Fomalont et al. in 2009 to constrain a new ephemeris (DE 422). The DE 422 post-fit residuals for Saturn with respect to the VLBA data are generally 0.2 mas, but additional observations are needed to improve the positions of all of our phase reference sources to this level. Over time we expect to be able to improve the accuracy of all three coordinates in the Saturn ephemeris (latitude, longitude, and range) by a factor of at least three. This will represent a significant improvement not just in the Saturn ephemeris but also in the link between the inner and outer solar system ephemerides and in the link to the inertial ICRF.

1407911

2010ApJ...725.1192M

1012

1012.0787_arXiv.txt

One of the most enduring questions of astronomy is ``How do stars work?" Progress over the last century has led to a robust scientific framework that explains the physics of internal stellar structure and also how stars evolve in time. This framework strives to include stars of all types, e.g. low metallicity stars of the early universe, massive stars that will become supernovae, low-mass brown dwarfs as well as stars like our Sun. As a basic rule, our understanding is strongest and best proven for Sun-like stars and gets shakier and less robust for stars very much different from the Sun. This paper will highlight one critical aspect of stellar structure and evolution, but one that hardly affects our Sun -- that of stellar rotation. Stars in general are probably born with significant angular momentum, but most of them are low-mass (like the Sun) and are ``slowed" down by their magnetic winds in the early parts of their lives. On the other hand, most intermediate- and high-mass stars ($\simge$1.5 solar masses) have weak magnetic fields and do not live very long, thus are often observed to be rotating very quickly. ``Rapid rotation" is expected to change the star's luminosity and photospheric temperature distribution \citep{maeder2000} and also strongly modify the observed surface abundances of various elements \citep{pinsonneault1997}. For the most massive stars, rotation will partially determine when the star becomes supernovae, thus having a crucial impact on how heavy elements get dispersed into the interstellar medium. For all its importance, the effects of rapid rotation are only vaguely understood and cannot be confidently included in our stellar evolutionary codes. More observations are critically needed to support the recent renaissance in theoretical efforts using analytic calculations and 3-dimensional hydrodynamical computer simulations \citep[e.g.,][]{lee1997,townsend2003a,jackson2005,roxburgh2006,rieutord2006,reese2008}. Recent breakthroughs in making images of nearby rapidly-rotating stars using optical interferometry have coincided with new numerical efforts to simulate the internal stellar structure of these stars \citep[e.g.,][]{vanbelle2001, souza2003, monnierscience2007,zhao2009}. Many of the original assumptions that had been adopted are under renewed scrutiny and simple assumptions such as rigid-body rotation and early prescriptions for "gravity darkening" do not seem consistent with new data. Our group recognized the profound advances possible by combining the new images of nearby rapid rotators with the asteroseismic constraints \citep[see also][for general discussion of astereoseismology and interferometry]{cunha2007} that can be revealed by the precision photometer onboard the MOST satellite \citep{walker2003}. We have initially identified two rapidly rotating stars observable by MOST with extensive published interferometric datasets: Altair and $\alpha$~Oph (Rasalhague). Since there is no analytic theory that can predict the effects of rotation on the pulsational modes of a star spinning at nearly 90\% of breakup, we plan to use the known stellar and geometric parameters of our sample to constrain the asteroseismic models with unprecedented power. Essentially, we know everything about the external appearance of these stars (size, oblateness temperature gradients on the surface, viewing angle, rotation periods) from interferometry and thus we can use asteroseismic signal for deducing the internal structure without the normally disastrous degeneracies one confronts. Here we report on our first results of this project, namely the oscillation spectrum of $\alpha$~Oph using the MOST observations as part of the first year NASA Guest Observer program. $\alpha$~Oph is classified as an A5IV star \citep{gray2001} with an estimated mass of 2.18$\msun$ and is rotating at 88\% of breakup as judged by the extreme oblateness of the photosphere observed by interferometry \citep{zhao2009}. Interferometric imaging also revealed that our viewing angle is nearly exactly equator-on, a perspective that strongly suppresses the photometric amplitudes of $l-|m|$ odd-parity modes, simplifying mode identifications. An early K star orbits the primary star with a period of 7.9 years \citep[based on astrometric and speckle data, e.g.][]{gatewood2005}, which will allow the masses of the two stars to be determined eventually to high precision. Adaptive optics imaging has recently resolved (Hinkley et al., 2010, private communication) the two stars and finds the companion to be more than 100$\times$ fainter than the primary at wavelengths of the MOST instrument and so we have assumed all the pulsation signatures we see are from the primary. In this paper we present the full photometric dataset for Rasalhague along with an updated stellar model based on interferometry data. We find evidence for both p-mode and g-mode oscillations, along with a remarkable rotational modulation signal. It was beyond the scope of this paper to develop a mode analysis; however, we will pursue this in a future paper.

This first long stare by MOST at a rapidly rotating A star has led to number of a new results. We detect for the first time a rich p-mode spectrum consistent with low-amplitude $\delta$-Scuti pulsations, and measure a granulation spectrum below 26$\pm$2~c/d. In total, we have identified 57$\pm$1 distinct modes below 50 c/d including a complex set of low-frequency modes that we identify as rotationally-modulated g-modes with (co-rotating) frequencies $\sim$0.1 c/d. A mode analysis revealed linear relationships between the spacings of g-modes up to $m=7$, an unexpected result for a star rotating at $\sim$90\% of breakup. This periodicity can be explained as due to dispersion-free equatorial Kelvin waves (prograde $l=m$ modes) although some inconsistencies in our analysis demand follow-up study. Lastly, the long time-base has allowed us to study the mode lifetime, finding that most p-modes are stable while g-modes appear to live only a few times their intrinsic (co-rotating) periods. Understanding the effect of rapid rotation on stellar interiors is crucial to developing reliable models for massive star evolution in general. $\alpha$~Oph is emerging as a crucial prototype object for challenging our models and to spur observational and theoretical progress. Additional work is planned using adaptive optics to determine the mass of the star by measuring the 8-year visual orbit more precisely, using visible and infrared interferometry to strictly constrain possible differential rotation and gravity darkening laws, and using asteroseismology to follow-up on the new discoveries outlined here.

1012.0787

Despite a century of remarkable progress in understanding stellar interiors, we know surprisingly little about the inner workings of stars spinning near their critical limit. New interferometric imaging of these so-called rapid rotators combined with breakthroughs in asteroseismology promise to lift this veil and probe the strongly latitude-dependent photospheric characteristics and even reveal the internal angular momentum distribution of these luminous objects. Here, we report the first high-precision photometry on the low-amplitude δ Scuti variable star Rasalhague (α Oph, A5IV, 2.18 M<SUB>sun</SUB>, {ω}/{ω_c}∼ 0.88) based on 30 continuous days of monitoring using the MOST satellite. We have identified 57 ± 1 distinct pulsation modes above a stochastic granulation spectrum with a cutoff of ~26 cycles day<SUP>-1</SUP>. Remarkably, we have also discovered that the fast rotation period of 14.5 hr modulates low-frequency modes (1-10 day periods) that we identify as a rich family of g-modes (|m| up to 7). The spacing of the g-modes is surprisingly linear considering Coriolis forces are expected to strongly distort the mode spectrum, suggesting we are seeing prograde "equatorial Kelvin" waves (modes ell = m). We emphasize the unique aspects of Rasalhague motivating future detailed asteroseismic modeling—a source with a precisely measured parallax distance, photospheric oblateness, latitude temperature structure, and whose low-mass companion provides an astrometric orbit for precise mass determinations. <P />Based on data from the MOST satellite, a Canadian Space Agency mission operated by Dynacon, Inc., the University of Toronto Institute of Aerospace Studies, and the University of British Columbia, with assistance from the University of Vienna, Austria.

12101917

2010AIPC.1283..294M

1012

1012.5441_arXiv.txt

The distribution of objects in the Milky Way is important for understanding their formation and evolution, as well as the large scale structure of the Galaxy. Although Galactic coordinates are measured with respect to the location of the Sun, most analytical models for the locations of objects in the Galaxy use a coordinate system having an origin at the Galactic center (GC; Sartore et al. 2009; Lorimer et al. 2006). And although the locations are measured in a spherical coordinate system, the models are formed using cylindrical coordinates, with the height, azimuth, and equatorial radial distance of the objects taken as independent random variables (RVs). The distribution for azimuth is always uniform over $2\pi$, due to the symmetry provided by a galactocentric origin. The height of an object above the Galactic plane (GP) is almost always treated as an exponential RV (Lorimer et al. 2006; Bahcall 1986; Lyne et al. 1985; Gunn \& Ostriker 1970). Most models assume the equatorial radii of stellar coordinates follow gamma or exponential distributions (Lorimer et al. 2006; Bahcall 1986), while others invoke a normal distribution (Sartore et al. 2009). The parameterization of these models requires reasonably accurate estimates of the distances to the objects for a proper transformation of coordinates between the solarcentric and galactocentric coordinate systems. The distance estimates for pulsars (PSRs) are generally made from a model of the distribution of free electrons within the Galaxy (Taylor \& Cordes 1993; Cordes \& Lazio 2008) in combination with the observed value of a PSR's dispersion measure. The distance estimates can introduce a significant source of error in the model parameterization because actual measurements of PSR distances via parallax have shown some estimates to be in error by a factor of two or more (Cordes \& Lazio 2008; Brisken et al. 2002). The objective of this paper is to develop a solarcentric analytical model of spheroidal Galactic coordinates that can be parameterized without direct knowledge of the distances to objects in the Galaxy. The distribution of objects within the Galaxy is perhaps best illustrated by a Lambert equal area projection (LEAP), which is a polar plot of the Galactic coordinates of the objects in the two hemispheres of the celestial sphere (e.g. see Figs.~\ref{fig:psrleap} and~\ref{fig:snrleap}). The main advantage of a LEAP is it preserves the density of data points in the projection, unlike other projection methods (e.g. orthographic) which do not (Fisher et al. 1987). A LEAP consists of two sets of concentric circles centered on the location of the Sun. The left set of circles in a LEAP is the projection as viewed from above the GP, and the right set is the view from below the GP. The azimuth and radius of the points in the polar plots represent an object's Galactic coordinates. Galactic longitude, $\phi$, is equal to a point's azimuth in a LEAP and increases in the counterclockwise direction from the GC, which is located at the far right of each circle set. A point's radius in the polar plot is equal to $2\sin(\theta/2)$, where $\theta$ is the Galactic colatitude of the object. Objects located in the GP fall on the perimeter of the LEAP, and those located off the GP reside within the perimeter. The LEAPs of PSRs and supernova remnants (SNRs) are shown in Figures~\ref{fig:psrleap} and Figure~\ref{fig:snrleap}, respectively. Since PSRs are born in supernovae, one might expect their distributions on the sky to be similar. Both objects reside primarily in the GP in what is known as a girdle-type distribution in the statistical literature (Fisher et al. 1987). The LEAPs show the dispersion of PSRs perpendicular to the GP is much greater than that of their progenitor SNRs. The same may also be true of the dispersion along the GP. \begin{figure} \plotone{mckinnonLEAP_f1.eps} \caption{Lambert equal-area projection of pulsar Galactic coordinates. The perimeter of the figure corresponds to the Galactic plane. The concentric circles are lines of constant Galactic colatitude.} \label{fig:psrleap} \end{figure} \begin{figure} \plotone{mckinnonLEAP_f2.eps} \caption{Lambert equal-area projection of supernova Galactic coordinates. The numbers on the perimeter of the figure denote Galactic longitude. The numbers labelling the interior concentric circles denote Galactic colatitude.} \label{fig:snrleap} \end{figure}

An analytical model with a solarcentric, spherical coordinate system was developed for the distribution of objects in the Galaxy in an attempt to circumvent complications introduced by the galactocentric, cylindrical coordinate system used in other models. Unlike these previous models, the present model does not require distance estimates for its parameterization. The model reasonably describes histograms of measured coordinates for PSRs and SNRs and was used to quantify the differences in their distribution parameters. The dispersion of PSRs perpendicular to and along the GP is larger than that for SNRs. The difference is likely caused by the kick velocity acquired by a PSR in the explosion of its progenitor SNR. SNRs reside in a thin disk of the Galaxy, while PSRs are located in a thicker disk. When applied to globular clusters, the model shows the clusters reside in a spheroid centered on the GC. The model can be applied to other objects to quantify their spatial distribution and perhaps the large scale structure of the Galaxy.

1012.5441

The location of objects on the celestial sphere is a fundamental measurement in astronomy, and the distribution of these objects within the Milky Way is important for understanding their evolution as well as the large scale structure of the Galaxy. Here, physical concepts in Galactic astronomy are illustrated using straightforward mathematics and simplifying assumptions regarding the geometry of the Galaxy. Specifically, an analytical model for a smooth distribution of particles in an oblate ellipsoid is used to replicate the observed distributions of the Galactic coordinates for pulsars and supernova remnants. The distributions and the Lambert equal area projections (LEAPs) of the coordinates suggest that the dominant factors determining the general shape of the distributions are the heavy concentration of objects in the Galactic plane and the offset of the Galactic center from the coordinate system origin. The LEAPs and the distributions also show that the dispersion of pulsars about and along the plane are much larger than that for their progenitor supernovae. Additionally, the model can be used to derive an analytical expression for the dispersion measure along any line of sight within the Galaxy. The expression is used to create a hypothetical dispersion measure-distance map for pulsars in the Galaxy.

12137725

2010idm..confE.113A

1012

1012.0863_arXiv.txt

We have conducted a preliminary investigation of possible DM annihilation signal in the Milky Way Host Halo, and within the systematic uncertainties of our background model, find no significant detection. We model all known backgrounds, and use the Profile Likelihood method of marginalizing over the many nuisance parameters that come from uncertainty in the CR-induced galactic diffuse emission. Sampling the model parameter space adequately is a computationally difficult task, one we are addressing by intelligently pre-selecting our models with a classification tree. We expect the combination of increased sampling density of our maximum likelihood region and upcoming improvements to our models (anisotropic inverse compton \& inclusion of target uncertainty) to yield a robust limit that takes into account the large systematics inherent to a physical model of the galactic background.

1012.0863

Our Galaxy resides in the center of a vast "Halo" of Dark Matter (DM). This concentration produces, in many viable particle physics models, an indirect Weakly Interacting Massive Particle (WIMP) annihilation signal that peaks in the Fermi-LAT's energy range. Our knowledge of the diffuse background is essential to placing reasonable limits on the DM mass and cross-section. We incorporate a systematic variation of the GALPROP galactic diffuse background model, constrained by current cosmic-ray measurements, into a profile likelihood analysis and present preliminary upper limits on the DM annihilation cross-section using the Fermi-LAT data.

2371731

2011A&A...526A.141F

1012

1012.1168_arXiv.txt

Much recent theoretical work has gone into examining how planet formation depends on stellar mass, because stellar mass significantly changes the physical conditions which control the formation of planets. A comparison, for instance, of the planet populations around Sun-like stars on one hand, and around M dwarfs on the other hand, probes the sensitivity of the planetary formation process to several physical parameters: around lower mass stars gravity (hence disk rotation speed), temperature (which regulates the position of the ice line) are both lower, and, perhaps most importantly, disk mass scales approximately linearly with stellar mass \citep[e.g.][]{Scholz2006}. Within the ``core accretion'' paradigm, \citet[][]{Laughlin2004}, \citet[][]{Ida2005}, and \citet[][]{Kennedy2008} all predict that giant planet formation is inhibited around very-low-mass stars, while Neptune-mass planets should inversely be common. Within the same paradigm, but assuming that the properties of protoplanetary disks, contrary to observations, do not change with stellar mass, \citet[][]{Kornet2006} predict instead that Jupiter-mass planets become more frequent in inverse proportion to the stellar mass. Finally, \citet[][]{Boss2006} examines how planet formation depends on stellar mass for planets formed by disk instability, and concludes that the frequency of Jupiter-mass planet is largely independent of stellar mass, as long as disks are massive enough to become unstable. One needs to note, though, that proto-planetary disks of a realistic mass are likely to be gravitationally stable out to at least 10~AU. Planets can thus form through gravitational instability only beyond that distance, in a separation range only skimmed by radial velocity monitoring and probed mostly by microlensing searches and direct imaging. Massive planets formed by gravitational instability and found well within 5~AU must thus then have migrated inward. How giant planets migrating in the massive disks needed for gravitational instability can escape accreting enough mass to become a brown dwarf ($>13~\mathrm{M_{Jup}}$) is unclear \citet[e.g.][]{Stamatellos2009,Kratter2010} Observationally, just a dozen % of the close to 400 planetary systems currently known from radial velocity monitoring, % are centered around M dwarfs (M$<$0.6\Msol) \footnote{http://exoplanet.eu/catalog-RV.php}. This no doubt reflects in part a selection bias, since many more of the intrisically brighter solar-type stars than of the fainter M dwarfs have been searched for planets, but there is increasing statistical evidence \citep[e.g.][]{Bonfils2006,Endl2006,Johnson2007,Johnson2010} that M dwarfs also genuinely have fewer massive planets ($\sim$\Mjup) than the more massive solar-type stars. They may, on the other hand, and though no rigorous statistical analysis has yet been performed for that planet population, have a larger prevalence of the harder to detect Neptune-mass and super-Earth planets: a quarter of the $\sim$30~planets with M~sin(i)~$<$~0.1\Mjup % known to date orbit an M dwarf, when solar-type stars outnumber M dwarfs by an order of magnitude in planet-search samples. Conversely, the highest mass planets known around M dwarfs are the M~sin(i)~=~2\Mjup Gl~876b \citep{Delfosse1998,Marcy1998} and HIP79431b \citep{Apps2010}, and at a larger orbital separation of $\sim$3~AUs the M~=${\approx}$3.5\Mjup % OGLE-2005-BLG-071Lb microlensing planet \citep{Dong2009}, when over two dozen planets with masses over 10\Mjup are known around solar type stars. The statistical significance of that difference however remains modest, since the M dwarfs searched for planets only number in the few hundreds, when the apparent fraction of these very massive planets is under 1\% around solar type stars. We present here the detection of two giant planets around M0 dwarfs, a M~sin(i)~=~0.354\Mjup planet around HIP~12961, and a M~sin(i)~=~4.87\Mjup planet around Gl~676A.

As discussed above, HIP~12961 and Gl~676A are orbited by giant planets with minimum masses of approximately 0.5 and 5 Jupiter masses. The latter is twice the M~sin(i)~=~2\Mjup of Gl~876b \citep{Delfosse1998,Marcy1998} and HIP79431b \citep{Apps2010}, previously the highest mass planets found by radial velocity monitoring of M~dwarfs, and above the 3.8 or 3.4~\Mjup (from two degenerate solutions) of the OGLE-2005-BLG-071Lb \citep{Dong2009} microlensing planet. The M0V Gl~676A however is significantly more massive (0.71~\Msol, Table~\ref{table:stellar}) than the M4V Gl~876 \citep[0.33~\Msol][]{Correia2010} and the M3V HIP79431 \citep[0.49~\Msol][]{Apps2010}. The higher mass of its planet therefore remains in approximate line with the current upper envelope of the planetary versus stellar mass diagram. These most massive planets are rare at any stellar mass, with an occurence rate under 1\%, suggesting that they can form only under the most favorable conditions. They have been suggested to form through gravitational instability, with their lower mass counterparts forming by core accretion. Proto-planetary disks of any realistic mass, however, are expected be gravitationally stable out to beyond 10~AU. If Gl~676Ab formed through gravitational instability, it would therefore have undergone much inward migration, through a very massive disk. How it could escape accreting enough mass during this migration to become a brown dwarf is unclear. Gl~676A and HIP~12961 increase the sample of M dwarfs with giant planets (Saturn-mass and above) from 7 to 9, and therefore offer an opportunity to evaluate the trend \citep{Johnson2009,Schlaufman2010} for giant planets being more common around more metal-rich M~dwarfs. Adopting the very recent \cite{Schlaufman2010} metallicity calibration of the M$_{K_s}$ vs V-K$_S$ plane, which finds metallicities approximately half-way between those of the earlier \citet{Bonfils2005} and \citet{Johnson2009} calibrations, the metallicities of Gl~676A and HIP~12961 are $0.18$ and $-0.07$. Both values are above the [Fe/H]~=~-0.17 average metallicity for the solar neighborhood in the \cite{Schlaufman2010} metallicity scale, the latter very significantly so. The two new planets therefore clearly reinforce the incipient trend, and help suggest that more massive planets are found around more metal-rich M-dwarfs.

1012.1168

Fewer giants planets are found around M dwarfs than around more massive stars, and this dependence of planetary characteristics on the mass of the central star is an important observational diagnostic of planetary formation theories. In part to improve on those statistics, we are monitoring the radial velocities of nearby M dwarfs with the HARPS spectrograph on the ESO 3.6 m telescope. We present here the detection of giant planets around two nearby M0 dwarfs: planets, with minimum masses of respectively 5 Jupiter masses and 1 Saturn mass, orbit around Gl 676A and HIP 12961. The latter is, by over a factor of two, the most massive planet found by radial velocity monitoring of an M dwarf, but its being found around an early M-dwarf is in approximate line with the upper envelope of the planetary vs stellar mass diagram. HIP 12961 ([Fe/H] = -0.07) is slightly more metal-rich than the average solar neighborhood ([Fe/H] = -0.17), and Gl 676A ([Fe/H] = 0.18) significantly so. The two stars together therefore reinforce the growing trend for giant planets being more frequent around more metal-rich M dwarfs, and the 5 Jupiter mass Gl 676Ab being found around a metal-rich star is consistent with the expectation that the most massive planets preferentially form in disks with large condensate masses. <P />Based on observations made with the HARPS instrument on the ESO 3.6-m telescope at La Silla Observatory under program ID 072.C-0488Tables 3 and 4 are also 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/cgi-bin/qcat?J/A+A/526/A141">http://cdsarc.u-strasbg.fr/cgi-bin/qcat?J/A+A/526/A141</A>

12217153

2011PASJ...63S.437T

1012

1012.4860_arXiv.txt

Since recent observations in optical and near-infrared wavelength indicate that the volume-averaged star formation rate (SFR) increases from $z = 0$ to $z \sim 1$ and plateaus at $z \sim 2$ \citep{2004ApJ...615..209H,2006ApJ...651..142H}, it is likely that a large fraction of stars in galaxies at present-day formed at $z>1$. The AGN activity and the redshift distribution of submm galaxies (SMG) also peak at this epoch \citep{2005ApJ...622..772C}. Therefore the redshift range of $z$=2--3 is the epoch when galaxies have the most intensive evolution. It is absolutely imperative to build a statistical sample of star forming galaxies at $z$=2--3 in order to understand the cosmic star formation history and the early evolution of galaxies. A line-emitter survey with a narrow-band filter can provide a large sample of active galaxies efficiently in a limited range of redshift in contrast to color-color selections such as $UG\mathcal{R}$ \citep{2004ApJ...604..534S} or $BzK$ \citep{2004ApJ...617..746D}. Since red galaxies represent characteristic colors on the color-magnitude diagram as referred to as the ``red sequence'', we can relatively easily select such red galaxies located at a specific redshift \citep{1998A&A...334...99K}. The blue star-forming galaxies, however, are much more difficult to be identified at $z\sim2$ using a color-selection or a photometric redshift technique because their spectra are relatively flat and featureless in the optical-NIR regime. Moreover, in a narrow-band survey which captures emission lines from galaxies directly, we can sample star forming galaxies completely above a certain limiting flux and an equivalent width limit unless the line is attenuated by dust or stellar absorption. We are not biased by colors of galaxies, either. For these reasons, the narrow-band emitter survey is a very efficient and effective method for initially making a sample of star forming galaxies and obtaining their photometric properties at a particular redshift. We can then study them more in detail by spectroscopic follow-up in the near-infrared and also by radio observations targeting their molecular lines. One of the most important quantities characterizing a galaxy is SFR. There are many different SFR indicators, such as UV continuum, nebular emission lines (H$\alpha $ and [OII]) and mid-infrared, but the SFRs estimated by various measurements do not always provide consistent results due to selection biases and different amounts of dust extinction effects, and so on. For UV continuum radiated by hot OB type stars, we must correct for absorption by the surrounding dust. This effect is often corrected for by using the UV slope but this process may lead to a large uncertainty in the estimated SFRs. On the other hand, mid-infrared emission is not affected by dust ``extinction'' since it is dust ``emission''. However, we can observe only the galaxies that have relatively high SFRs so as to be detected by mid-infrared observations which are less sensitive to optical or near-infrared observations. Also, to derive SFRs from the mid-infrared luminosity one has to assume the dust temperature which is somewhat uncertain. The H$\alpha $ line, a hydrogen's Balmer series line emitted from ionized gas (i.e., HII region) around hot young stars, is one of the best SFR indicators. It has many great advantages; being less affected by dust extinction, providing a survey with high sensitivity, and having been well calibrated in the local Universe. For the rest-frame UV-selected galaxies, \cite{2006ApJ...647..128E} obtained 114 H$\alpha $ spectra of star forming galaxies at $z \sim 2$ and indicate that H$\alpha $ emission is attenuated by a typical factor of 1.7, which is about 1/3 compared to the UV attenuation. The narrow-band survey of H$\alpha $ line enables us to make a relatively-unbiased large sample of star forming galaxies at the same redshift as well as to obtain accurate estimates of SFR, unlike the rest-frame UV surveys which tend to miss a significant fraction of star-forming activities due to dust obscuration. However, the H$\alpha $ line is red-shifted into the near-infrared regime at $z > 0.5$. Because the near-infrared observation is severely affected by OH sky emission lines, only a fraction of redshift range can be surveyed with high sensitivities. At $z\sim $2.2, H$\alpha $ line falls just in between the OH line forest ($\lambda \sim 2.1 \mu $m). There are some pioneering previous works that searched for H$\alpha $ emitters at $z\sim $2 with narrow-band imaging \citep{1994AJ....107....1T,1995MNRAS.273..513B,1998ApJ...506..519T} or with slitless spectroscopy \citep{1999ApJ...519L..47Y,2010ApJ...723..104A}. In the past decade, some large H$\alpha $ surveys have been carried out at $z \sim 2.2$. \cite{2000A&A...362....9M} searched for H$\alpha $ emitting galaxies with ESO NTT telescope. The survey reached a limiting line flux of 5$\times 10^{-17}$ erg s$^{-1}$ cm$^{-2}$ and covered a 100~arcmin$^2$ area. The observation with the narrow-band filter yielded 10 candidates and 6 of them have been confirmed spectroscopically later. Then, \cite{2008MNRAS.388.1473G} conducted a narrow-band survey with a line flux limit of $\sim 10^{-16}$ erg s$^{-1}$ cm$^{-2}$ over a large area of 0.6 deg$^2$ with a 3.8m telescope (UKIRT) and identified 55 H$\alpha $ emitters. Although their survey provided a sample large enough to evaluate the H$\alpha $ luminosity function (LF) at the bright end, they pointed out that the faint end of H$\alpha $ LF is important in order to estimate the global star formation rate, and hence a deeper survey is required for this purpose. To date, the deepest survey of H$\alpha $ emitters at this high redshift is that of \cite{2010A&A...509L...5H}. They carried out a narrow-band survey with HAWK-I on ESO-VLT, and reached to a 3$\sigma $ flux limit of 4.1 $\times 10^{-18}$ erg s$^{-1}$ cm$^{-2}$ over 56~arcmin$^2$ area. Their survey suggested that the faint-end slope of the H$\alpha $ LF is steeper than the value obtained in the local universe. However, large scale structures at $z\sim 2$ may stretch the intrinsic slope of the LF and it is too early to conclude on the faint end of H$\alpha $ LF by only the data in a single field. The result is very susceptible to the effect of cosmic variance. Although the recent development of wide-field instruments in the near-infrared regime has enabled us to study distant galaxies at $z\sim 2$ systematically, there is an on-going debate as to how the ancient galaxies evolve into the ones at the present-day universe. This is attributed to the fact that the galaxy evolution is connected not only to time, but also to environment and mass \citep[e.g.][]{2007A&A...468...33E,2010A&A...518A..18T}. It is thought that there are two types of causes for the environmental dependence, inherent origin and acquired effects. In the former case, initial density fluctuation is originally a bit large in high density regions such as cluster cores, and galaxy formation take place earlier there compared to lower density regions. In the later case, some external factors, such as gas stripping by ram pressure and mergers between galaxies, influence galaxy properties during the course of their build-up. The spatial distribution of star forming galaxies helps us to tell whether the inherent origin or the acquired effects dominant the environmental dependence. In galaxy cluster RXJ1716 at $z=0.8$, H$\alpha $ emitters or mid-infrared sources, which are both tracers of active star formation, are preferentially found in the surrounding region of the cluster core \citep{2008MNRAS.391.1758K,2010MNRAS.403.1611K}. In contrast, star-forming galaxies in XCS2215 at $z=1.5$ are located in the cluster core as often as in the surrounding environments \citep{2010MNRAS.402.1980H,2010ApJ...719L.126T}. It seems that star formation activity is initially enhanced in massive galaxies in high density regions and is propagated to lower mass galaxies and to lower density regions with time. If the relative velocity between galaxies is too large, these galaxies would pass through each other without merging together. Therefore, at $z=0.8$, it is thought that mergers in mid-density regions would be more important than the inherent effect. However, at a higher redshift $z=1.5$, it seems that the inherent effect would surpass the acquired effects. To further look into and confirm this interesting result based on the two clusters, and to extend it further back in time, it is required to construct a larger sample of star forming galaxies over the critical era of $z=$1--3, that covers various environments from galaxy cluster core to blank field region in order to compare the properties between different environments. In this paper, we present a survey of H$\alpha $ emitters at $z=2.2$ in GOODS-N field in order to study star formation activities in the general field or low density regions. This paper is structured as follows. In \S~2, we describe our survey design, observations, and the data. In \S~3, our target selection of H$\alpha $ emitters and the spectroscopic follow-up observations are described. We show our results in \S~4 and \S~5, and discuss the star formation activities of galaxies in various environments and redshifts in \S~6. Finally, we summarize our study in \S~7. We assume the cosmological parameters of H$_0$=70 km s$^{-1}$ Mpc$^{-1}$, $\Omega_M=0.3$, and $\Omega_\Lambda=0.7$, and adopt AB magnitudes throughout this paper.

We have conducted a narrow-band imaging and spectroscopic surveys of H$\alpha $ emitters at $z=2.2$ in GOODS-North field, using MOIRCS on Subaru. Our survey has identified 11 H$\alpha $ emitters and one AGN over a 70 arcmin$^2$ area. We have confirmed probable H$\alpha $ emission lines for all of the seven targets by the spectroscopic follow-up observation, on top of the two already confirmed star forming galaxies at $z_{\mathrm{spec}}=2.2$. Therefore, our technique of searching for H$\alpha $ emitters at $z=2.2$ based on the excess fluxes in the narrow-band (NB209) and photometric redshifts, is proven to be robust and efficient. The results and conclusions of this study are summarized below: \begin{enumerate} \item The H$\alpha $ emitters have SFRs ranging from 12~M$_\odot \mathrm{yr}^{-1}$ to 60~M$_\odot \mathrm{yr}^{-1}$, with the mean SFR of $\langle SFR \rangle =27.8$M$_\odot \mathrm{yr}^{-1}$. Note that we have corrected for dust extinction by assuming the typical value of A(H$\alpha $)=1. The averaged stellar mass of the H$\alpha $ emitters is $4.0\times 10^{10}$~M$_\odot$, and we find the correlation between the stellar mass and the specific star formation rate in the sense that more massive galaxies tend to have lower specific star formation rates. \item The H$\alpha $ luminosity function is derived from our data by combining the data points reproduced from the previous works in the literature for the brighter magnitudes. The combined LF is represented by a Schechter function with log$L$ = 42.82, log$\phi = -2.78$ and $\alpha = -1.38$. Our result shows a moderate steepness of the faint-end slope. By integrating the luminosity function thus derived, we find that the cosmic star formation rate at $z=2.2$ is $\rho _{\mathrm{SFR}}=0.31$~M$_\odot$yr$^{-1}$Mpc$^{-3}$, which is consistent with other previous studies at $z\sim 2$. \item We have compared the properties of our emitters in the general field or low density environment, with those in the cluster environments to investigate the environmental dependence of galaxy evolution. There is a difference in the degree of time evolution of $\Sigma$SFR / $\Sigma M_{\mathrm{star}}$ between the two environments. This implies that the star formation activity is enhanced at $z>2$ in high density regions as a consequence of ``galaxy formation bias'' in the early universe. \end{enumerate} We must warn however that there is a possibility that these results are still severely affected by the cosmic variance, which may well be expected from the actual inhomogeneous spatial distribution of our H$\alpha $ emitters. It is necessary to extend the survey area to cover a representative volume of the universe and average over the cosmic variance in order to obtain more robust conclusions.

1012.4860

We present a pilot narrow-band survey of Hα emitters at z = 2.2 in the Great Observatories Origins Deep Survey North (GOODS-N) field with MOIRCS instrument on the Subaru Telescope. The survey reached a 3σ limiting magnitude of 23.6 (NB209), which corresponds to a 3σ limiting line flux of 2.5 × 10<SUP>-17</SUP>erg s<SUP>-1</SUP>cm<SUP>-2</SUP> over a 56 arcmin<SUP>2</SUP> contiguous area (excluding a shallower area). From this survey, we have identified eleven Hα emitters and one AGN at z = 2.2 on the basis of narrow-band excesses and photometric redshifts. We obtained spectra for seven out of the new objects, including one AGN; also, an emission line above 3σ was detected from all of them. We estimated star-formation rates (SFR) and stellar masses (M<SUB>star</SUB>) for individual galaxies. The average SFR and M<SUB>star</SUB> are 27.8M<SUB>odot</SUB>yr<SUP>-1</SUP> and 4.0 × 10<SUP>10</SUP>M<SUB>odot</SUB>, respectively. Their specific star-formation rates negatively correlate with their stellar masses. Fitting to a Schechter function yields the Hα luminosity function with logL = 42.82, logφ = -2.78, and α = -1.37. The average star-formation rate density in the survey volume is estimated to be 0.31M<SUB>odot</SUB>yr<SUP>-1</SUP>Mpc<SUP>-3</SUP> according to the Kennicutt relation between the Hα luminosity and the star-formation rate. We compared our Hα emitters at z = 2.2 in GOODS-N with narrow-band line emitters in other fields and clusters to see their time evolution and environmental dependence. We found that the star-formation activity is rapidly reduced from z = 2.5 to z = 0.8 in the cluster environment, while it only moderately changed in the field environment. This result suggests that the time scale of galaxy formation differs among different environments, and the star-forming activities in high density regions eventually overtake those in lower-density regions as a consequence of ``galaxy-formation bias'' at high redshifts.

1344499

2011A&A...527A..58B

1012

1012.3757.txt

\label{intro} Until just a few years ago, VV124=UGC4879\footnote{A09125+5303 in the nomenclature adopted by \citet{jansen_phot}.} was considered a unassuming isolated dwarf galaxy, classified as spheroidal/irregular, at a distance of $D\sim 10$~Mpc \citep{halfa}. Integrated multi-color optical photometry was obtained by \citet{jansen_phot} and \citet{taylor}, and J and K$_S$ images were obtained by \citet{nir}, within large surveys of nearby galaxies. From inspection of a low resolution integrated optical spectrum, \citet{jansen_spec} concluded that it was likely ``a young post-starburst galaxy''. From the $H_{\alpha}+N[II]$ equivalent width, \citet{jansen_spec} estimated a total star formation rate of $0.005~M_{\sun}$yr$^{-1}$ (assuming distance of 10 Mpc). The generally adopted distance of $D\sim 10$~Mpc was based entirely on the redshift estimate reported by the CfA survey \citep[$cz=600$~km~s$^{-1}$;][]{cfa}. A team of russian scientists \citep[][K08 hereafter]{k08}, triggered by the apparent partial resolution of VV124 into stars in Sloan Digital Sky Survey images \citep[SDSS,][]{sdss}, carefully searched public databases and the literature, eventually making the case for a much lower recession velocity and a much smaller distance for VV124. They followed up this smart intuition with deep V,I photometry and low-resolution spectroscopy and were actually able (a) to resolve the galaxy into individual stars down to $\sim 2$ mag below the Red Giant Branch (RGB) tip, thus obtaining a direct distance estimate of $D=1.1$~Mpc from the tip itself, and (b) to obtain a new estimate of the heliocentric velocity, much lower than the CfA value, $V_{\rm h}=-70\pm 15$~km~s$^{-1}$ (throughout the paper $V_{\rm h}$ stands for heliocentric radial velocity). This meant that K08 had found a new member of the Local Group (LG), since with the newly determined distances and velocity VV124 is found to lie near the turn-around radius of the LG, and, in fact, being its most isolated member. The galaxy has a remarkable elliptical shape and ranks among the brightest LG dwarf spheroidal/transition type galaxies ($M_B=-11.6$). The stellar budget of the galaxy seems dominated by old stars (RGB; age $\ga 2$~Gyr) with colors compatible with low metallicity (Z$\simeq0.001$). However, a sprinkle of bright blue stars, and the identification of an \HII\ region, led K08 to conclude that VV124 is a transition type between dwarf irregulars (dIrr) and dwarf spheroidals (dSph), like Phoenix, Antlia or LGS3 \citep{mateo}. The results by K08 were further discussed in more detail in \citet[][T10 hereafter]{tik}. According to K08, the location and the peculiar velocity of VV124 indicate that it has never been a satellite of a major galaxy of the LG, hence it evolved in full isolation for a Hubble time. Therefore, VV124 may contain a fossil record of precious information on the initial conditions of dwarf galaxies. It may be considered as a possible progenitor of the gas-less amorphous dSphs found in the vicinity of the Milky Way or M31, whose evolution has been likely largely driven by the strong interaction with the large galaxy they are orbiting around \citep[see][and references therein]{mateo,lokas_struc}. In particular, \citet{lucio1,lucionat} have developed a detailed model, within a modern cosmological context, in which dSphs are produced by the morphological transformation of dwarf {\em disk} galaxies by tidal stirring and ram-pressure stripping during their path through the halo of the main galaxy they are gravitationally bound to \citep[see also][and references therein for a more general view on the nature and origin of dSphs]{korme}. As we shall see, VV124 may possibly share some remarkable characteristics with the precursors of modern dSphs envisaged in this model. The general interest in isolated galaxies as objects of undisturbed evolution is witnessed, for example, by the large Hubble Space Telescope (HST) programme LCID \citep{lcid}, aimed at the determination of the star formation history in the center of six isolated LG dwarfs of various morphological types. Here we are more interested in the structure and dynamics of a galaxy that should be untouched by the interactions with other large galaxies since the beginning of time. In particular, the image presented in Fig.~1 of K08 suggests that the galaxy may be more extended than what could be enclosed into the $6\arcmin \times 6\arcmin$ field studied by those authors. For these reasons, we acquired much deeper observations on a much wider field with the $2\times 8.4$~m Large Binocular Telescope (LBT, Mt. Graham - AZ). A beautiful color image derived from these data is presented in Fig.~\ref{imaC}, giving also an idea of the number and variety of background galaxies that can be found in such deep LBT images. In this paper we describe and discuss the results of these observations, as well as those from deep \HI\ data obtained with the Westerbork Synthesis Radio Telescope (WSRT) and from low resolution optical spectroscopy obtained with the Telescopio Nazionale Galileo (TNG). The plan of the paper is the following: in Sect.~\ref{phot} we present the LBT observations, we describe the reduction of these data and the artificial stars experiments. The process of surface photometry of the innermost regions of the main body of the galaxy is described, and the adopted system of local coordinates is also introduced. In Sect.~\ref{cmd} we discuss the derived color-magnitude diagrams (CMD), we provide a revised estimate of the distance to VV124 and we analyze the stellar content of the galaxy. Sect.~\ref{struc} is devoted to the analysis of the surface brightness profile and the surface density distribution, while in Sect.~\ref{HIanalysis} the results of the \HI\ observations are discussed in detail; the derived \HI\ velocity field is compared from the velocities obtained from low resolution optical spectroscopy (Sect.\ref{LRS}). Finally, the overall results are summarized and discussed in a broader context in Sect.~\ref{disc}. A few days before this manuscript was ready for submission, a preprint was posted on the {\em astro-ph} archive \citep[][hereafter J10]{jacobs}, presenting deep HST / Advanced Camera for Surveys (ACS) photometry of VV124. This study turns out to be complementary to ours, as it focuses on the star formation history (SFH) in the innermost $\simeq 40\arcsec$ of VV124, a region essentially out of reach of our photometry because of the extreme crowding (see Sect.~\ref{comple}). We will briefly refer to the results by J10 in the following, when appropriate, but we do not discuss them in detail. In general, for the issues treated in both papers, the results of the two studies are in good agreement.

\label{disc} We have presented the results of deep wide-field photometry, optical low-resolution spectroscopy and \HI\ observations of the dwarf galaxy VV124=UGC4879, recently recognized as lying in the outer fringes of the LG. The main results of our analysis can be summarized as follows: \begin{itemize} \item While a sparse population of young stars (age$\la 500$~Myr) is observed in the inner region of the galaxy, the dwarf is dominated by an old population. There are indications of the presence of age/metallicity gradients, with older and more metal-poor stars being more prevalent at larger distances from the center of the galaxy. \item We used a very clean detection of the RGB tip to obtain a new and more accurate distance estimate with respect to K08 and T10, $D=1.3\pm 0.1$~Mpc, in excellent agreement with the very recent estimate obtained by J10 from HST data. In spite of the larger distance, VV124 can be considered as a (present or future) member of the Local Group of galaxies, according to the criterion adopted by \citet[][see also the discussion and references on the method]{mateo}, and illustrated in Fig.~\ref{apex}, here. \item Independently of the actual membership to the LG, it is confirmed that VV124 is the most isolated nearby galaxy, very likely never disturbed by even a weak interaction with other dwarf or giant galaxies during its whole lifetime (see K08 and J10, for further discussion). For example, the relatively high recession velocity ($V_g$=+100~km~s$^{-=1}$) of the gas-poor isolated dwarf Tucana \citep{tuc} may indicate that it has been ejected to large distances by a three-body encounter, as envisaged by \citet{sales}. Such a scenario cannot be invoked in the case of VV124, which has a much smaller galactocentric velocity \citep[$V_g$=+17~km~s$^{-=1}$; see also][for other scenarios for the transformation of dwarf galaxies and discussion]{donghia,kaza}. %-----------------------LG----------------------------------- \begin{figure} \centering \includegraphics[width=\columnwidth]{apex.pdf} \caption{Heliocentric velocities versus the cosine of the angle between a galaxy and the apex of the Sun motion with respect to the center-of-mass of the LG \citep[see][and references therein]{mateo}, for galaxies within 2~Mpc from the barycenter of the LG ($D_{LG}<2$~Mpc) in the catalogue of \citet{tully}, plus VV124 (data from the present paper). The parameters of the solar motion are also taken from Tully et al., while the position of the barycenter is from \citet{vdb00}. Filled squares are galaxies with $D_{LG}\le 1.0$~Mpc, open squares have $1.0$~Mpc~$<D_{LG}\le 1.5$~Mpc, $\times$ symbols have $D_{LG}> 1.5$~Mpc. VV124 has $D_{LG}\simeq 1.3$~Mpc. The continuous line is the locus of rest with respect to the center-of-mass of the LG, the dotted lines are at $\pm 1\sigma$ and $\pm 2\sigma$ from that locus \citep[$\sigma=60$~km~s$^{-1}$, from][and in good agreement with what found here from galaxies having $D_{LG}\le 1.3$~Mpc]{mateo}.} \label{apex} \end{figure} %%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%% \item The main visible body of VV124 is remarkably elliptical ($\epsilon\simeq 0.44$), regular and smooth, if the asymmetrically distributed minor young component is ignored. We were able to follow the SB profile of the galaxy out to a distance never reached before, $r_{\epsilon}\sim 1.9$~kpc, demonstrating that the galaxy is much more extended than previously believed. \item In addition, we obtained surface density maps of RGB stars revealing the presence of two relatively thin symmetric wings, emanating from the western and eastern edges of the inner elliptical body and aligned along its major axis, extending out to $\simeq 3$~kpc from the center of the galaxy. These low SB features ($\mu_r\ga 30$mag/arcsec$^2$) have no counterpart in other known galaxies of the LG (but see below). \item \HI\ emission was detected for the first time by our high-sensitivity WSRT observations. The \HI\ total mass ($\sim 1\times 10^6~M_{\sun}$) is small with respect to the overall stellar luminosity (mass), making VV124 more gas-poor than typical isolated dwarfs. The density peak of the gas is slightly offsett with respect of the optical center (by $\sim 400$~pc) and the distribution shows a tail in the SE direction, also corresponding with the region with the most negative velocity. This may be indicative of some outflow/inflow process possibly tied to the most recent star formation event (see J10), since the youngest stars are also preferentially found to the SE of the galaxy center. \item It is important to recall that {\em the observed \HI\ contours do not extend beyond the optical body of the galaxy as traced by RGB star counts}. Also, the \HI\ distribution do not show any obvious correlation with the surface density stellar wings described above. \item The velocity field of the \HI\ does not show any sign of overall rotation over the main optical body of the galaxy, while a velocity gradient is observed in the tail-like structure in the SE side. The velocity gradient may indicate either infall or outflow of \HI. The projected size of the tail and the projected difference in velocity with the main body give a timescale of several times $10^7$ yr. This timescale is very uncertain (due to projection effects), but could be connected to some elevated activity in the galaxy some time ago (see J10). It could be an outflow due to the most recent episode of star formation, or the remnant of some \HI\ inflow which triggered the star formation. \item The gas associated with VV124 lies in the velocity range -20~km~s$^{-1} \ge V_{\rm h} \ge$~-40~km~s$^{-1}$, with a mean systemic velocity of $V_{\rm h}=-25\pm 4$~km~s$^{-1}$. The interstellar medium is found in two phases, a narrower component associated with the inner regions (at higher column density), and broader component found over the whole body of the \HI\ distribution, having a typical dispersion of 11~km~s$^{-1}$. \item The correlation between the velocity of various sources associated to VV124, as derived from optical spectra by us, K08 and T10, and the \HI\ velocity field is quite poor. %While the results we obtained from our new optical spectroscopy observations partially %alleviates the problem, the comparison is still unsatisfying. While the observed differences can be accommodated within the uncertainties, it is a bit disturbing that the large majority of optical estimates lie at $V_{\rm h}<-40$~km~s$^{-1}$, beyond the lower side of the range of \HI\ velocities, in particular if one considers that they are now taken from two independent sources, i.e. T10 and this work. This means that one should consider the possibility that there is a real difference between the systemic velocities of the stars and of the gas, due to the hypothesized gas flow suggested by the asymmetric structure of the \HI. Coming Keck-DEIMOS observations of RGB stars in the galaxy will hopefully settle this issue. \item Fig.~\ref{scale} shows that the structural parameters of VV124 are consistent with the scaling laws obeyed by dSphs and dIrrs galaxies \citep[see][and references therein]{korme,tht}. It is interesting to note that it lies at the upper envelope of the $M_V - \mu_V$ relation, i.e. it has the brightest SB for its total luminosity, and in an intermediate position between dIrrs and dSphs in $M_V - r_h$ relation. According to \citet{pena}, the evolution in a relatively strong tidal field would have led to a larger scale radius, a lower luminosity and a lower surface density, thus driving VV124 toward the loci of ''genuine'' dSphs, in these planes. This is in qualitative agreement with the possibility that VV124 may be a precursor of modern dSphs that did not enter in the interaction-driven evolutionary path that produced the latter class of dwarf galaxies. %If VV124 belongs to the population of precursors of dwarf spheroidals, one can guess %what it should have been its present-day structure if it would have suffered the %strong interaction with a larger galaxy that (possibly) transformed its former %''sisters'' into dSphs by using the ``evolutionary vectors'' provided by % \citet{pena}. \end{itemize} %-----------------------scale laws----------------------------------- \begin{figure} \centering \includegraphics[width=\columnwidth]{scale.pdf} \caption{VV124 (black filled circle) compared to other LG dwarf galaxies brighter than $M_V=-7.0$ in the $M_V$ vs. logarithm of the half-light radius (lower panel) and $M_V$ vs. central surface brightness (upper panel) planes. Different symbols are used for different morphological types, according to the classification criteria by \citet{mateo}. Absolute magnitudes and half-light radii are taken from from \citet{walk10} and from \citet{kalirai}; when lacking in these sources they have been drawn from \citet{mateo}. Surface brightness estimates have been taken from \citet{mateo}, \citet{mccon}, \citet{mart09}, \citet{zuck}, \citet{mike}, and \citet{whiting}, with this priority. All the values of surface brightness are corrected for extinction. } \label{scale} \end{figure} %%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%% Probably the most intriguing features of VV124 revealed by the present study are the low SB stellar wings; in particular if they are interpreted as an ancient stellar disk seen edge-on and considering the extreme isolation in which the galaxy evolved. As mentioned in the introduction, the idea that present-day dwarf spheroidals originated from dIrr galaxies that were deprived from their gas at early times, as a result of tides and/or ram pressure stripping within the potential of a much larger galaxy or of a group/cluster, is not new \citep[see][and references therein]{mateo,tht}. The model recently developed by \citet{lucio1,lucionat} explicitly postulates dwarf disk galaxies as the progenitors of dSphs \cite[see][for a recent and thorough discussion]{kaza}. Indeed, observational evidence of the existence of the expected intermediate stages of such transformation process is beginning to emerge, from study of various environments \citep[see, for example, the disky dEs identified in the Virgo cluster by][]{lisker}. In the context of the LG, \citet{n205ivo} have recently reviewed the evidence for the presence of the relics of a disk in NGC~205, suggesting that the process of transformation of the disk into a spheroid is currently ongoing in this satellite of M31. Two independent theoretical studies have shown that some relevant observational properties of the disrupting Sagittarius dSph and of the associated tidal stream can be more easily explained if a disk galaxy is adopted as the progenitor of Sgr \citep{pena_sgr,lokas_sgr}. NGC~205 and Sgr may simply represent two different stages on a similar evolutionary path driving the transformation of a similar low luminosity disky progenitor into a dSph via the interaction with the main galaxy. In this framework, VV124 would be akin to such disky progenitors but {\em never} entered the transformation path. Rather, it evolved passively in isolation, thus preserving its disk intact until the present day. Assuming we are looking at a rotationally supported stellar system, the disk of VV124 would be relatively thick, but this is normally seen in the faintest dIrrs in the Local Group and nearby clusters (Sanchez-Janssen et al. 2010). The thickness of the disk does not reflect environmental effects, but is rather the result of pressure support from either internal feedback from star formation and/or the cosmic ultraviolet ionizing background becoming increasingly more important for the energy balance towards increasingly lower masses. Kaufmann, Wheeler \& Bullock (2007) have shown how, for galaxies having circular velocities $< 30$ km/s hosted in halos with typical spin parameters ($\lambda < 0.05$), an effective temperature (i.e. thermal + turbulent) of the ISM of a few times $10^4$ K, or equivalently an ISM velocity dispersion of $\sim 10-12$ km/s, is sufficient to produce a substantially thick disk (with major-to-minor axis ratio in the range $0.2-0.4$) as an equilibrium configuration. This is because in such systems the gas velocity dispersion is already close to the virial temperature (for circular velocities below $30$ km/s the virial temperature is $< 5 \times 10^4$ K), which forces the gas to acquire a high pressure scale height at equilibrium. Stars in such a system form out of a pressure supported turbulent gas disk and, being collisionless, have no way to dissipate such motions later on. Again, in the Local Group, at the lowest luminosity end of dIrrs, there are a number of examples of dwarfs with \HI\ velocity dispersions around $10$ km/s (SagDig, Leo A and GR8 being some of these), and whose low rotation velocities ($< 15$ km/s) implies a halo circular velocity well below $30$ km/s, consistent with the picture just outlined. Another very important issue is to understand how the galaxy became gas-poor in absence of stripping mechanisms due to the interaction with other galaxies. Photo-evaporation of the gas after re-ionization \citep{ioniz} is a possibility certainly worth further investigation \citep[see, e.g.][]{susa}. Supernovae feedback might also play a role; recent cosmological simulations of the formation of a gas rich dwarf that are finally able to produce a realistic exponential disk with no bulge (Governato, Brook, Mayer et al. 2010) show that outflows at high redshift ($z > 1$) can remove more than $2/3$ of the baryons even in dwarfs with circular velocities exceeding $V_{circ}$ = 30 km/s. Finally, given the fragility of such a low mass galaxy, we cannot exclude that the interaction with intergalactic gas clouds in the Local Group could have caused stripping of at least a fraction of an already diffuse, loosely bound interstellar medium. In this context it is very interesting to recall that also the Dark Matter (DM) halo expected to embed VV124 should be virtually untouched since the epoch of its collapse: the kinematics of the stars in the main body and in the wings of the galaxy should probe the mass profile of this pristine halo out to a remarkably large radius ($\sim 3$~kpc), possibly opening a crucial window on the initial conditions of DM halos of this low-mass scale. %In any case, it is quite clear that this neglected system may hold the key to understand several %long standing issues about the nature and the evolution of dwarf galaxies by highlighting %cleanly the %effect of processes that are blurred by strong environmental effects when we study at objects %well inside %or even at the outskirts of groups and clusters. This is the reason why this "pristine" system %deserves to be studied in the deepest detail. %%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%

1012.3757

We present a detailed study of the dwarf galaxy VV124 (UGC4879), recently recognized as a remarkably isolated member of the Local Group. We have obtained deep (r ≃ 26.5) wide-field (23' × 23') g,r photometry of individual stars with the LBC camera at the Large Binocular Telescope under sub-arcsec seeing conditions. The color-magnitude diagram suggests that the stellar content of the galaxy is dominated by an old, metal-poor population, with a significant metallicity spread. A very clean detection of the RGB tip allows us to derive an accurate distance of D = 1.3 ± 0.1 Mpc. Combining surface photometry with star counts, we are able to trace the surface brightness profile of VV124 out to ~5' ≃ 1.9 kpc radius (where μ<SUB>r</SUB> ≃ 30 mag/arcsec<SUP>2</SUP>), showing that it is much more extended than previously believed. Moreover, the surface density map reveals the presence of two symmetric flattened wings emanating from the central elongated spheroid and aligned with its major axis, resembling a stellar disk seen nearly edge-on. We also present H i observations obtained with the Westerbork Synthesis Radio Telescope (WSRT), the first ever of this object. A total amount of ≃ 10<SUP>6</SUP> M<SUB>⊙</SUB> of H i gas is detected in VV124. Compared to the total luminosity, this gives a value of M<SUB>HI</SUB>/L<SUB>V</SUB> = 0.11, which is particularly low for isolated Local Group dwarfs. The spatial distribution of the gas does not correlate with the observed stellar wings. The systemic velocity of the H i in the region superposed to the stellar main body of the galaxy is V<SUB>h</SUB> = -25 km s<SUP>-1</SUP>. The velocity field shows substructures typical of galaxies of this size but no sign of rotation. The H i spectra indicates the presence of a two-phase interstellar medium, again typical of many dwarf galaxies.

2063731

2010AN....331.1010M

1012

1012.4267_arXiv.txt

Red giants are cool stars with an extended convective envelope, which can, as in main sequence solar-like stars, stochastically excite pressure modes of oscillation. Although stochastic oscillations in a few red giants have already been detected from ground and space observations \citep[e.g.][]{Frandsen02, Joris06, Barban07} it has been only after the photometric space mission COROT \citep{Baglin02} that an unambiguous detection of radial and non-radial modes in a large number of red-giant stars have been achieved \citep{Joris09, Hekker09, Carrier10}. That confirmation has opened the way to the seismic study of the structure and evolution of these objects that play a fundamental role in fields such as stellar age determination and chemical evolution of galaxies. About 2000 of the targets observed by CoRoT in the two first runs of 150 days have been identified as red giants with solar-like oscillations in the frequency domain expected from theoretical scalings by \cite{KB95}. Their spectra show regular patterns that allowed \cite{Mosser10} to derive precise values of the large frequency separation. The analysis of the light curve of the sismo-CoRoT target HR 7349 has revealed a rich solar-like spectrum with 19 identified modes of degrees $\ell=0$, 1 and 2 \citep{Carrier10}. Moreover, the first 34 days of science operations of the KEPLER satellite \citep{kepler09}, have also revealed the presence of solar-like oscillations in a sample of 50 stars whose frequency at maximum power ($\nu_{\rm max}$) indicates that they are low-luminosity red-giants \citep{Beddingetal10}. The mean large ($\Delta \nu$) and small frequency separations ($\delta\nu_{02}$), classically used in the asymptotic interpretation of solar-like oscillations, were also derived for that sample. All these new and high quality data, together with all those we expect in the next years thanks to KEPLER and CoRoT missions, are the motivation for the present study about the physical interpretation of the oscillation spectrum of red-giant stars. In the next sections we present: first the properties of the stellar models we computed for different sets of fundamental parameters and input physics; second, the properties of the corresponding adiabatic spectra; and third, the predicted values for the frequency separations of acoustic modes, as well as the inferences about the physical properties of red-giant stars that we can extract from these separations .

In this paper we presented the properties of the oscillation spectrum of solar-like oscillations during the RGB and core He-burning phases of red-giant evolution, and analyzed the behaviour of large and small frequency separations derived for modes well trapped in the acoustic cavity of these stars. The main results of this global overview are the followings: \begin{itemize} \item Independently of the evolutionary state, $\ell=2$ modes trapped in the acoustic cavity have an inertia of the same order as that of the corresponding radial mode and behave as ``p-modes'' with frequencies regularly spaced by $<\Delta\nu>$. As a consequence, the scatter of $\ell=2$ modes in the folded \'echelle diagrams is rather small. \item The trapping of $\ell=1$ modes in the acoustic cavity depends on the evolutionary state. While a regular pattern of dipole modes is expected in more centrally condensed models, the scatter significantly increases for models in the core He-burning phase. The scatter of $\ell=1$ ``p-modes'' decreases as model concentration increases. Therefore the regularity of $\ell=1$ spectrum could be used to discriminate between different evolutionary phases. \item $\langle\delta\nu_{02}\rangle$ depends almost linearly on the large separation, hence on the mean density of the model, with a slope that slightly depends on the mass. The value of $\langle\delta\nu_{02}\rangle/\langle\Delta\nu_0\rangle$ for a given mass increases with the density contrast and thus with luminosity, and for a given luminosity, it decreases as the mass of the model increases. \item $\langle\delta\nu_{01}\rangle$ seems to reflect the distance between the $\ell=1$ turning point and the bottom of convective envelope. It takes negative (or small) values if tp1 is well inside the convective envelope, what occurs in RGB models. \item The theoretical predictions based on stellar models are in good agreement with the observational results obtained in the first 34 days of KEPLER observations. Comparison of their $\nu_{\rm max}$ values with the population simulations presented in \cite{MiglioPop09} suggests that these red giants are in fact stars with masses lower than 2~\msun in the low-luminosity part of the ascending RGB. On the other hand, the same simulations concluded that the sample of red giants analysed in the CoRoT exo-field are ``Red Clump'' ones. That is, the sample is dominated by low-mass stars in the core He-burning phase. From the theoretical analysis of standard models presented in that paper, we would expect that the corresponding folded \'echelle diagram will show a significant scatter for $\ell=1$ modes. \end {itemize}

1012.4267

The clear detection with CoRoT and Kepler of radial and non-radial solar-like oscillations in many red giants paves the way to seismic inferences on the structure of such stars. We present an overview of the properties of the adiabatic frequencies and frequency separations of radial and non-radial oscillation modes, highlighting how their detection allows a deeper insight into the properties of the internal structure of red giants. In our study we consider models of red giants in different evolutionary stages, as well as of different masses and chemical composition. We describe how the large and small separations computed with radial modes and with non-radial modes mostly trapped in the envelope depend on the stellar global parameters and evolutionary state, and we compare our theoretical predictions and first Kepler data.Finally, we find that the properties of dipole modes constitute a promising seismic diagnostic of the evolutionary state of red-giant stars.

1481902

2012ApJ...748...86O

1012

1012.1218_arXiv.txt

On October 17, 2010 \citet{ATel2959} reported on an outburst from a new source, named \MAXI, revealed by the Gas Slit Camera \citep[GSC;][]{GSC} of the Monitor of All-sky X-ray Image \citep[MAXI;][]{MAXI} experiment aboard the International Space Station. The 4--10~keV flux decreased from its initial value of $\sim$40~mCrab on October 17 to $\sim$30~mCrab on the following day. On October 20 {\em Swift} \citep{Swift} began a target of opportunity observation of the \MAXI\ error circle (0\fdg2 radius). \citet{ATel2962} used the X--Ray Telescope \citep[XRT;][]{XRT} to pinpoint the position of the new source with an estimated 90\% confidence level uncertainty of 1\farcs9 and coordinates RA(J2000): 14$^{\rm h}$~08$^{\rm m}$~02\fs56, Dec(J2000): $-$61\degr~59\arcmin~00\farcs3. Its spectrum is fit by an absorbed power law, with photon index $\Gamma=-0.5^{+0.1}_{-0.6}$. A follow-up XRT observation on November 30 found the source $\sim$7 times brighter than on October 20, with an average 0.3--10 keV flux of $7\times 10^{-10}$ erg~cm$^{-2}$~s$^{-1}$ \citep{ATel3060}. Only in this latter observation a $\sim$500~s pulsation was detected, with a 42\% sinusoidal rms modulation. The source was detected since October 18 also in the 15--50~keV energy band by the {\em Swift} Burst Alert Telescope \citep[BAT;][]{BAT} at the level of $\sim$30~mCrab \citep{Kennea11}. Neither catalogued radio nor X--ray sources were present in the {\em Swift} error circle, while a 2MASS IR star, the likely IR counterpart of the X--ray transient, lays 2\farcs1 from the XRT position \citep{ATel2962}. The source was also observed on October 22 by the Proportional Counter Array \citep[PCA;][]{PCA} aboard {\em RXTE} \citep{XTE}. The X--ray spectrum shows a strong 6.5~keV Iron line (E.W.\ 200~eV), and a continuum modeled by reflection or partially covering absorption \citep{ATel2969}. The 2--10~keV flux was about $10^{-10}$ erg~cm$^{-2}$~s$^{-1}$, and was strongly variable on time-scales of hundreds of seconds. An observation performed on December 4 found \MAXI\ at a flux level 6--7 times higher than the October observation \citep{Atel3070}, and confirmed the 500~s pulse period observed by {\em Swift}. From observations performed on December 2 and 3 by the GLAST Burst Monitor \citep[GBM;][]{GBM} aboard the {\em Fermi} satellite, the presence of a 500~s double peaked pulsating signal was confirmed \citep{ATel3069}. Following the \MAXI\ localization by {\em Swift,} we searched for \SAX\ \citep{SAX} archival observations with the transient position within the Narrow Field Instruments (NFIs) Field of View (FoV). We found that an observation performed in the framework of a Galactic plane survey contained the \MAXI\ position \citep{ATel2965}. In this Paper we report on the spectral analysis of this \SAX\ observation, integrated with high-energy data obtained from the offset fields of the PDS \citep{PDS} instrument. We further present the analysis of a newly analyzed {\em ASCA} observation performed during a Galactic plane survey.

We found five sets of observations containing the position of \MAXI\ in our \SAX/PDS archive, performed in 1997, 2000, and 2001, and one in the {\em ASCA} archive (March 1998). In all our observations the source was in a low state, with 15--100~keV fluxes in the range $\sim$2--8~mCrab, and no spectral variability during the observations. For comparison, an integrated exposure (over 5 years) of 2.4~Ms by {\em INTEGRAL\/}/IBIS provides 2$\sigma$ upper limits on the persistent quiescent emission of 0.2 and 0.4~mCrab in the 20--40 and 40--100~keV energy bands, respectively \citep{ATel2965}. When assuming the source fluxes in outburst as measured by {\em Swift} \citep{ATel2962,Kennea11} and {\em RXTE} \citep{Atel3070}, from our low state measurement we can infer a dynamic range of 400 in 15--50~keV, and of 300 in 2--10~keV. The discovery of CRFs in the low state spectrum of \MAXI, together with its $\sim$500~s pulsations, unambiguously identify the source as an accreting X--ray binary pulsar. Only few sources show multiple CRFs, and only two show resonances above the (1:2). Three CRFs were observed during the 2004--2005 outburst of the X--ray pulsar V0332+53 \citep{Coburn05,Kreykenbohm05,Tsygankov06}, and five (possible six) CRFs were discovered in the spectrum of 4U~0115+63 during its 1999 giant outburst \citep{Santangelo99,Heindl99,Ferrigno09}. In both cases, deviations from a pure harmonic ratio among the CRFs were observed, and were explained in terms of departures from a classical dipolar structure of the magnetic field in the line-forming region \citep{Nishimura05,Nishimura08}. The magnetic field strength at the neutron star surface corresponding to $E_{\rm cyc}\!=\!44$~keV is $3.8\times10^{12} (1+z)$~G \citep{Canuto77}, where $z$, the gravitational redshift, for a typical neutron star of mass 1.4~M$_\odot$ and radius 10~Km, is about 0.3. When left as free parameters in the fit, we found that the best fit CRF line energies do not follow the harmonic relation $E_n = n\,E_1$. A slightly non-harmonicity is expected when relativistic effects are taken into account \citep[see, e.g.,][]{Meszaros} \begin{equation} E_n = m_ec^2\, \frac{\displaystyle\sqrt{1+2n(B/B_{\rm crit})\sin^2\theta} -1}{\sin^2\theta}\, \frac{1}{1+z} \label{eq:en-rel} \end{equation} \noindent where $m_e$ is the electron rest mass, $c$ the speed of light, $\theta$ the angle between the photon and the magnetic field direction, and $B_{\rm crit}=4.414\times 10^{13}$~G is the critical magnetic field strength where the cyclotron energy equals the electron rest mass. The observed (1:2) and (1:3) ratios of the line energies with respect to the fundamental, 1.7$\pm$0.2 and 2.9$\pm$0.3, cannot be explained in terms of Eq.~(\ref{eq:en-rel}), as it is evident from Figure~\ref{harmonics}. A fit to the harmonic relation, shown as the dashed line, gives $E_1=41\pm3$~keV (in agreement with the best fit value $E_{\rm cyc}=41\pm1$~keV found when imposing the harmonic relation in the spectral fit), but with poor significance ($\chi^2_\nu=2.27$ for 2 dof). In the same figure we also show the harmonic relation from Eq.~(\ref{eq:en-rel}) for different values of the magnetic field and the angle $\theta$. Taking into account relativistic effects, we found that the magnetic field responsible for the (1:1) and (1:3) CRFs is about 20\% higher than that responsible for the (1:2) CRF. \begin{figure} \centering \epsscale{1.2}\plotone{fig06.ps} \caption{Harmonicity of the \MAXI\ CRF line energies. The dashed black line corresponds to a linear fit (that is, $E_n=n\,E_1$) to the data, while the colored strips take into account relativistic effects, as detailed by Eq.~(\ref{eq:en-rel}). Each strip corresponds to a fixed value of the magnetic field strength $B_{12}=B\times 10^{12}$~G, and to a range 0.0001--1 for $\sin^2\theta$, where $\theta$ is the angle between the photon direction and that of the magnetic field. While $E_1$ and $E_3$ are consistent with $B_{12}\sim3.8$, $E_2$ is consistent with a magnetic field about 20\% lower.} \label{harmonics} \end{figure} Two points are worth noticing: it is always the (1:2) resonance that shows the larger disagreement with the harmonic relation, and this could be due to the fact that this CRF is due to pure absorption \citep{Nishimura03}, while for the other resonances other effects, like multiple scattering and photon spawning, enter into play \citep[see, e.g.,][and references therein]{Schonherr07}. Unfortunately, because we are not able to reconstruct the CRF profiles, we cannot extract more information, like the electron temperature and the geometry of the emitting region. Second, at variance with the V0332+53 and 4U~0115+63 observations, both performed during giant outbursts, our \SAX\ observations were performed while \MAXI\ was in a low state. This did not allow the study of the dependence of the CRF parameters as a function of luminosity, an important tool for study of the physical conditions in the line-forming region \citep{Mihara04,Nakajima06,Klochkov11}. The likely early-type optical counterpart and the 500~s pulsation makes \MAXI\ a HMXB pulsar. According to the nature of the secondary star we have two possibilities: the source is a supergiant fast X--ray transient \citep[SFXT;][]{Sguera05,Negueruela06} or a Be/HMXB. In favor of the former interpretation is the highly reddened supergiant as possible counterpart, the typical outburst X--ray luminosity of $2\times10^{37}$ erg~s$^{-1}$ (assuming a distance of 14.5~kpc), and the dynamic range of more than two orders of magnitude ($\sim$300) that is typical of the so-called ``intermediate'' SFXT \citep{Sguera07,Clark10}. Against we have that the source active phase, about two months long \citep{ATel3067,Kennea11}, is significantly longer that that typical of SFXT \citep[see, e.g.,][]{Sidoli09}. If this were the case, then our magnetic field measurement would rule out the magnetar nature for SFXTs \citep{Bozzo08}. The observed properties of \MAXI\ are also in agreement with those observed in other Be/HMXBs, like 1A~1118$-$615, a 400~s X--ray pulsar with a hard X--ray spectrum ($\Gamma\!\sim\!1$), a CRF at $\sim$55~keV, and long (tens of years) periods of quiescence interrupted by giant (Type-II) outbursts lasting weeks to months, in which the X--ray luminosity increases by a factor $\sim$200 \citep[see, e.g.,][]{Rutledge07}. This would put the source a factor $\sim$2--3 closer. \begin{figure} \centering \epsscale{1.15}\plotone{fig07.ps} \caption{Error box of AGL~J1410$-$6143 (big red circle) superimposed on the 20--100 keV {\em INTEGRAL}/IBIS deep mosaic image ($\sim$2.3 Ms exposure). The \SAX/MECS position of \MAXI\ is marked by the green square. The two small yellow ellipses represent the Fermi $\gamma$--ray sources 2FGL~J1409.9$-$6129 and 2FGL~J1413.4$-$6204, respectively. White contours (from 50\% to 99\%) refer to the {\em EGRET} source \EGR.} \label{integral-mosaic} \end{figure} Finally, we note that \MAXI\ is located within the $0\fdg5$ error box of the unidentified transient MeV source AGL~J1410$-$6143 (see Fig.~\ref{integral-mosaic}) discovered on February 21, 2008 by the $\gamma$--ray satellite {\em AGILE} \citep{AGILE} during a bright MeV flare lasting only about one day \citep{ATel1394,ATel1419}. The {\em AGILE} large position uncertainty makes very difficult the identification of its lower energy counterpart responsible for the $\gamma$--ray emission. Despite this drawback, it is intriguing to note that the flaring source \MAXI\ is the only catalogued hard X--ray source above 20 keV (20--100 keV) to be located inside the {\em AGILE} error box, according to all the available catalogs in the HEASARC database. This spatial correlation is further supported by a similar transient nature for both \MAXI\ and AGL~J1410$-$6143. We also point out that this is not a unique case. To date, a few HMXBs have been unambiguously detected as flaring MeV sources lasting only a few days \citep{Sabatini10,Tavani09,Abdo09}. In addition, there are several other HMXBs proposed as best candidate counterparts of unidentified transient MeV sources located on the Galactic plane \citep{Sguera09a,Sguera11,Sguera09b}. For the sake of completeness, we note that within the large {\em AGILE} error box there are a number of high energy MeV sources (see Fig.~\ref{integral-mosaic}): i) 2FGL~J1409.9$-$6129 and 2FGL~J1413.4$-$6204 have been reported in the second Fermi source catalog (Abdo et al. 2011) as firmly identified $\gamma$--ray pulsars and this unambiguously excludes their association with the transient AGL~J1410$-$6143, ii) \EGR\ is still unidentified although it has been likely associated with the $\gamma$--ray pulsar 2FGL~J1409.9$-$6129 \citep{OBrien08}. However, \citet{Wallace00} reported a possible MeV flare from \EGR\ lasting a few days on November 1991. This behavior is at variance with the proposed association with the $\gamma$--ray pulsar 2FGL~J1409.9$-$6129 while it is more compatible with the flaring nature of AGL~J1410$-$6143. Further multi-wavelength studies (radio, NIR, X--ray and $\gamma$--ray) of the sky region are strongly needed and encouraged to shed more light on the nature of such high energy emitters.

1012.1218

The transient 500 s X-ray pulsar MAXI J1409-619 was discovered by the slit cameras aboard Monitor of All-sky X-ray Image (MAXI) on 2010 October 17, and soon after accurately localized by Swift. We found that the source position was serendipitously observed in 2000 during BeppoSAX observations of the Galactic plane. Two sources are clearly detected in the Medium-Energy Concentrator Spectrometer (MECS): one is consistent with the position of IGR J14043-6148 and the other one with that of MAXI J1409-619. We report on the analysis of this archival BeppoSAX/MECS observation integrated with newly analyzed observation from ASCA and a set of high-energy observations obtained from the offset fields of the BeppoSAX/PDS instrument. For the ON-source observation, the 1.8-100 keV spectrum is fit by an absorbed power law with a photon index Γ = 0.87<SUP>+0.29</SUP> <SUB>-0.19</SUB>, corresponding to 2-10 and 15-100 keV unabsorbed fluxes of 2.7 × 10<SUP>-12</SUP> and 4 × 10<SUP>-11</SUP> erg cm<SUP>-2</SUP> s<SUP>-1</SUP>, respectively, and a 2-10 keV luminosity of 7 × 10<SUP>34</SUP> erg s<SUP>-1</SUP> for a 15 kpc distance. For a PDS offset field observation, performed about one year later and showing a 15-100 keV flux of 7 × 10<SUP>-11</SUP> erg cm<SUP>-2</SUP> s<SUP>-1</SUP>, we clearly pinpoint three spectral absorption features at 44, 73, and 128 keV, resolved both in the spectral fit and in the Crab ratio. We interpret these not harmonically spaced features as due to cyclotron resonances. The fundamental energy of 44 ± 3 keV corresponds to a magnetic field strength at the neutron star surface of 3.8 × 10<SUP>12</SUP>(1 + z) G, where z is the gravitational redshift. We discuss the nature of the source in the light of its possible counterpart.

3681173

2011MNRAS.412L..30R

1012

1012.0028_arXiv.txt

} Supergiant fast X--ray transients (SFXTs) are a new class of High Mass X--ray Binaries (HMXBs) discovered by \inte{} \citep[e.g.\ ][]{Sguera2005} that are associated with OB supergiant stars via optical spectroscopy. In the X--rays they display outbursts significantly shorter than those of typical Be/X--ray binaries characterized by bright flares with peak luminosities of 10$^{36}$--10$^{37}$~erg~s$^{-1}$ which last a few hours \citep[as observed by \inte; ][]{Sguera2005,Negueruela2006:ESASP604}. As their quiescence is characterized by a luminosity of $\sim 10^{32}$~erg~s$^{-1}$ \citep[e.g.\ ][]{zand2005,Bozzo2010:quiesc1739n08408}, their dynamic range is of 3--5 orders of magnitude. While in outburst, their hard X--ray spectra resemble those of HMXBs hosting accreting neutron stars, with hard power laws below 10\,keV combined with high energy cut-offs at $\sim 15$--30~keV, sometimes strongly absorbed at soft energies \citep{Walter2006,SidoliPM2006}. So, even if pulse periods have only been measured for a few SFXTs, it is tempting to assume that all SFXTs might host a neutron star. The mechanism producing the outbursts is still being debated, and it is probably related to either the properties of the wind from the supergiant companion \citep{zand2005,Walter2007,Negueruela2008,Sidoli2007} or to the presence of a centrifugal or magnetic barrier \citep[][]{Grebenev2007,Bozzo2008}. \src\ was discovered during {\it ASCA} observations of the Scutum arm region performed on 1994 April 12, and 1999 October 3--4 as a flaring source which exhibited flux increases by a factor of 10 (up to $\sim 10^{-10}$ erg cm$^{-2}$ s$^{-1}$) with rising times on the order of 1\,hr \citep{Bamba2001}, a strong absorption $N_{\rm H} =3\times10^{22}$ cm$^{-2}$, and coherent pulsations with a period of $4.7394\pm0.0008$\,s. A {\it Chandra} observation on 2004 May 12, which provided the coordinates refined to arcsecond accuracy [RA(J2000$)=18^{\rm h} 41^{\rm m} 0\fs54$, Dec(J2000$)=-5^{\circ}$ $35^{\prime} 46\farcs8$, \citealt{Halpern2004:18410-0535b}], found the source at a much fainter level ($4\times 10^{-12}$ erg cm$^{-2}$ s$^{-1}$), and with a spectrum that was fit with an absorbed power-law model [$\Gamma=1.35\pm0.30$, $N_{\rm H} =(6.1\pm1.0)\times10^{22}$ cm$^{-2}$]. A newly discovered source, IGR~J18410$-$0535, was observed to flare by \inte\ on 2004 October 8 \citep[][]{Rodriguez2004:18410-0535}, as it reached $\approx 70$\,mCrab in the 20--60\,keV energy range (integrated over 1700\,s) and 20 mCrab in the 60--200\,keV range. The source was also detected in the 20--60\,keV energy range in subsequent observations, at a flux roughly half that of the initial peak. \citet{Halpern2004:18410-0535b} identified IGR~J18410$-$0535 as \src. \citet{Halpern2004:18410-0535a} established that the IR counterpart was 2MASS 18410043$-$0535465, a reddened star with a weak double-peaked H$\alpha$ emission line, initially classified as a Be star, which \citet{Nespoli2008} later reclassified as B1 Ib type star; this corroborated the evidence that \src\ is a member of the SFXT class, as proposed by \citet{Negueruela2006:ESASP604}. \citet{Sguera2009} presented the first broad-band spectrum of this source, obtained with \inte\ (IBIS$+$JEM-X), that they fit with an absorbed power-law with $\Gamma=2.5\pm0.6$, $N_{\rm H}=23^{+19}_{-14} \times 10^{22}$ cm$^{-2}$. In 2007 \sw\ \citep{Gehrels2004mn} observed the outburst of the periodic SFXT IGR~J11215$-$5952 \citep[][]{Romano2007}, which allowed us to discover that the accretion phase during the bright outbursts lasts much longer than a few hours, as seen by lower-sensitivity instruments. This is contrary to what was initially thought at the time of the discovery of this new class of sources. Between 2007 October 26 and 2008 November 15, \src\ was observed by \sw\ as part of a sample of 4 SFXTs which included IGR~J16479$-$4514, XTE~J1739--302, and \igr. The main aims were to characterize their long-term behavior, to determine the properties of their quiescent state, to monitor the onset of the outbursts and to measure the outburst recurrence period and duration \citep[][]{Sidoli2008:sfxts_paperI,Romano2009:sfxts_paperV,Romano2010:sfxts_paperVI}. Approximately two observations per week were collected with the X--ray Telescope \citep[XRT, ][]{Burrows2005:XRTmn} and the UV/Optical Telescope \citep[UVOT, ][]{Roming2005:UVOTmn}. During such an intense and sensitive monitoring, \src\ was the only SFXT that did not go through a bright outburst, although several on-board Burst Alert Telescope \citep[BAT, ][]{Barthelmy2005:BATmn} detections have been recorded \citep[][]{Romano2009:sfxts_paperV}. In this paper we report on the observations of the first outburst of \src\ observed by \sw\ on 2010 June 5 and we compare its properties with those of the prototype of the SFXT class, \igr, which went into a bright outburst on 2010 March 04. \begin{figure} \begin{center} \centerline{\includegraphics[height=8.7cm,angle=-90]{figure1.ps}} \end{center} \caption{{\bf Left:} \sw/XRT (0.2--10\,keV) light curve of the 2007--2008 monitoring campaign \citep[2007 October 26 to 2008 November 15; ][]{Romano2009:sfxts_paperV}. {\bf Right:} light curve of the 2010 June 5 outburst in the same time-scale. The downward-pointing arrows are 3$\sigma$ upper limits. } \label{axj1841fig:lcv_campaign} \end{figure} \begin{figure} \begin{center} \vspace{-0.5truecm} \centerline{\includegraphics*[angle=270,width=8.5cm]{figure2.ps}} \end{center} \vspace{-0.5truecm} \caption{XRT and BAT light curves of the initial orbit of data of the 2010 June 5 outburst of \src\ in units of count s$^{-1}$ and count s$^{-1}$ detector$^{-1}$, respectively. The empty circles correspond to BAT in event mode (S/N$=5$), filled circles to BAT survey mode data. } \label{axj1841fig:lcv_allbands} \end{figure} %

} % In this paper we report our analysis of the 2010 June 5 outburst of \src\ and the 2010 March 04 outburst of the SFXT prototype \igr. While in the first case, the image trigger was a very long one and NFI data could be collected only $\sim 1700$\,s after the trigger, when the source was relatively dim, in the second case, the slew occurred immediately after the trigger, while \igr\ was still very bright. Figure~\ref{axj1841fig:best_sfxts} (panels e and g) shows the full light curves of the outbursts of \src\ and \igr\ as they were observed by \sw\ for 11 and 2 days after the trigger, respectively. The \src\ XRT light curve shows a decreasing trend from the initial bright flare from a maximum of $\sim 8$ counts s$^{-1}$ down to $\sim 0.01$ counts s$^{-1}$ during the first day, with several flares superimposed, hence yielding a dynamic range of approximately 900 during this outburst. Then, after three days, the source count rate rose again and reached $\sim 1$ counts s$^{-1}$. We estimate that the observed dynamical range of this source in the XRT band, considering the historical data we collected during our monitoring campaign \citep[][see Fig.~\ref{axj1841fig:lcv_campaign}]{Sidoli2008:sfxts_paperI,Romano2009:sfxts_paperV} is $\approx 1600$, hence placing it well in the customary range for SFXTs. The outburst of \igr\ has similar characteristics to the one observed on 2008 March 31, as the XRT light curve shows a peak at about 25 counts s$^{-1}$, decreases to about 0.5 counts s$^{-1}$ and then increases again up to about 20 counts s$^{-1}$ at the end of the first orbit (Fig.~\ref{axj1841fig:17544lcv_allbands}). This behaviour was previously observed in \igr\ and, most notably, in IGR~J08408$-$4503 \citep{Romano2009:sfxts_paper08408} and SAX~J1818.6$-$1703 \citep{Sidoli2009:sfxts_sax1818}, so that this multiple-peak structure of the light curve could be considered a defining characteristic of the SFXT class and it is likely due to inhomogeneities within the accretion flow \citep[e.g.][]{zand2005}. Figure~\ref{axj1841fig:best_sfxts} compares the light curves of \src\ and \igr\ with the outbursts of SFXTs as observed during our monitoring campaigns with \sw. The most complete set of X--ray observations of an outburst of a SFXT is the one of the periodic SFXT IGR~J11215$-$5952 \citep[][]{Romano2007,Sidoli2007,Romano2009:11215_2008}, which was surprisingly long. We now know that such a length of the outburst (hence the length of the accetion phase) is a common characteristic of the whole sample of SFXTs followed by \sw, and in this respect \src\ fits right in, as its outburst lasted several days. \begin{figure} \begin{center} \vspace{-0.5truecm} \centerline{\includegraphics[width=8.5cm,angle=0]{figure6.ps}} \end{center} \vspace{-1truecm} \caption[XRT light curves]{Light curves of the most representative outbursts of SFXTs followed by {\it Swift}/XRT referred to their respective BAT triggers (IGR~J11215$-$5952 did not trigger the BAT, so it is referred to MJD 54139.94). Points denote detections, triangles 3$\sigma$ upper limits. Red data points (panels e, g) refer to observations presented here for the first time, while grey points to data presented elsewhere. Where no data are plotted, no \sw\ data were collected. Vertical dashed lines mark time intervals equal to 1 day, up to a week. References: IGR~J08408--4503 \citep[2008-07-05, ][panel a]{Romano2009:sfxts_paper08408}; IGR~J11215$-$5952 \citep[2007-02-09, ][panel b]{Romano2007}; IGR~J16479$-$4514 \citep[2005-08-30, ][panel c]{Sidoli2008:sfxts_paperI}; XTE~J1739$-$302 \citep[2008-08-13, ][panel d]{Sidoli2009:sfxts_paperIV}; SAX~J1818.6$-$1703 \citep[2009-05-06, ][panel f]{Sidoli2009:sfxts_sax1818}. Panels e and g report the 2010-03-04 outburst of IGR~J17544$-$2619 and the 2010-06-05 outburst of AX~J1841.0$-$0536, respectively (this work). } \label{axj1841fig:best_sfxts} \end{figure} We have presented the broad-band (0.3--100\,keV) simultaneous spectroscopy of \src. This allows us to make a comparison with the findings on the other SFXTs that were observed in the same fashion. The soft X--ray spectral properties observed during this flare are generally consistent with those observed with {\it ASCA} during the 1999 flare \citep[][$N_{\rm H} =3\times10^{22}$ cm$^{-2}$, $\Gamma=1$]{Bamba2001}. As \src\ was observed relatively late after the trigger, no meaningful information can be derived on variability of the soft spectral parameters during the outburst, such as the absorbing column density. However, we note that the value of $\Gamma$ in outburst follows the same trend of `harder when brighter' as reported in table~4 of \citet[][]{Romano2009:sfxts_paperV}, which was based on out-of-outburst emission. For the joint BAT$+$XRT spectrum during the 2010 June 5 outburst, an absorbed power-law model is an inadequate description, and more curvy models are required. We considered an absorbed power-law model with an exponential cutoff and an absorbed power-law model with a high energy cut-off, models typically used to describe the X--ray emission from accreting neutron stars in HMXBs. We obtained a good fit of the 0.3--100\,keV spectrum, characterized by high absorption $N_{\rm H} \sim 2\times10^{22}$ cm$^{-2}$, a hard power law below 10\,keV, and a high energy cutoff. These properties of \src\ are reminiscent of those of the prototypes of the SFXT class, \igr\ [whose data we have presented here and in \citet{Sidoli2009:sfxts_paperIII,Sidoli2009:sfxts_paperIV,Romano2010:sfxts_paperVI}], and XTE~J1739$-$302 (\citealt{Sidoli2009:sfxts_paperIII,Sidoli2009:sfxts_paperIV}). Although no statistically significant pulsations were found in the present data, \src\ is one of the 4 SFXTs with known pulse period \citep{Bamba2001}, $P_{\rm spin} =4.7394\pm0.0008$\,s, the others being IGR~J11215$-$5952 \citep[186.78$\pm$0.3\,s, ][]{Swank2007:atel999}, IGR~J16465$-$4507 \citep[228$\pm$6\,s, ][]{Lutovinov2005}, and IGR~J18483$-$0311 \citep[21.0526$\pm$0.0005\,s, ][]{Sguera2007}. While lacking the detection of cyclotron lines, which would yield a direct measurement of the magnetic field $B$ of the neutron star, an indirect estimate can be obtained by considering the {\sc highecut} fit to the broad-band spectrum of \src\ in outburst. Our value of the high energy cutoff $E_{\rm c} < 16$\,keV, although loosely constrained, yields a $B\la 3\times$10$^{12}$~G \citep{Coburn2002}. This value for $B$, which is indeed similar to the one derived for the prototype of the SFXT class \igr, is inconsistent with a magnetar nature of \src. In conclusion, we have shown how AX~J1841.0$-$0536 nicely fits in the SFXT class, based on the observed properties of \src\ during the 2010 June 5 outburst: a large dynamical range in X--ray luminosity, the similarity of the light curve length and shape to those of the prototype of the class \igr, and the X--ray broad-band spectrum, which we show here for the first time down to 0.3\,keV, thus constraining both the absorption and the cutoff energy. \vspace{-0.7cm}

1012.0028

Swift observed an outburst from the supergiant fast X-ray transient (SFXT) AX J1841.0-0536 on 2010 June 5, and followed it with X-ray Telescope (XRT) for 11 d. The X-ray light curve shows an initial flare followed by a decay and subsequent increase, as often seen in other SFXTs, and a dynamical range of ∼1600. Our observations allow us to analyse the simultaneous broad-band (0.3-100 keV) spectrum of this source, for the first time down to 0.3 keV, which can be fitted well with models usually adopted to describe the emission from accreting neutron stars in high-mass X-ray binaries, and is characterized by a high absorption (N<SUB>H</SUB>∼ 2 × 10<SUP>22</SUP> cm<SUP>-2</SUP>), a flat power law (Γ∼ 0.2) and a high-energy cut-off. All of these properties resemble those of the prototype of the class, IGR J17544-2619, which underwent an outburst on 2010 March 4, whose observations we also discuss. We show how well AX J1841.0-0536 fits in the SFXT class, based on its observed properties during the 2010 outburst, its large dynamical range in X-ray luminosity, the similarity of the light curve (length and shape) to those of the other SFXTs observed by Swift and the X-ray broad-band spectral properties.

12224835

2011PhRvD..83d4037G

1012

1012.5111_arXiv.txt

\label{intro} The two-body problem in relativity when one of the bodies is much more massive than the other is of great interest both theoretically and astrophysically. In this limit, the orbit of the smaller body is approximately geodesic on short time scales. Deviations from the geodesic trajectory arise from the back-reaction on the orbit of the spacetime perturbation created by the object, but can also arise from external factors such as gravitational interactions with other bodies, gaseous material in the spacetime and so forth. In all these situations, the orbit can be described as a geodesic acted on by a perturbing force, which is in general small. In this article, we describe techniques for integrating the Kerr geodesic equations in the presence of an arbitrary forcing term, which can be applied to any of these problems. For the back-reaction on the orbit, the perturbing force, called the self-force, is of the order of the mass ratio $\mu/M$ and it can be computed by a perturbation expansion in this small parameter. Computing the linearized metric perturbation sourced by the compact object and hence the self-force is not an easy task and it has taken more than a decade to solve this problem for a nonspinning compact object moving in a Schwarzschild background \cite{Vega:2009qb, Barack:2010tm, Dolan:2010mt, Barack:2009ux}. The conventional approach treats the compact object as a test mass which leads to a divergence of the field at the position of the particle and this must be dealt with using a regularization procedure. The extension to Kerr orbits is underway. The techniques described in this paper will be a useful tool in the future for constructing trajectories evolving under gravitational radiation-reaction. The problem of the motion of two bodies with very different masses is relevant for present and future gravitational wave detectors. Systems with mass ratios of 1:100 (intermediate-mass-ratio inspirals) could be detected by the advanced generation of ground-based detectors that are currently under construction~\cite{LIGOIMRI}. The proposed space-based detector LISA~\cite{SRD} is expected to detect $\sim 10-100$ extreme-mass-ratio inspiral (EMRI) events per year~\cite{gairEMRIrate}. These result from the capture of a compact stellar-mass object (a white dwarf, neutron star or black hole) by a massive black hole (MBH) from a surrounding cusp of stars in a galactic nucleus. The captured object generates a large number of gravitational wave cycles while it is orbiting in the strong field of the MBH, which makes these very good sources to use as probes of strong-field gravity~\cite{EMRIrev}. For both of these classes of source, techniques for evolving the orbit under the influence of both gravitational back-reaction and other perturbing forces are essential for constructing accurate waveform templates and for understanding how external perturbations can leave an imprint on the inspiral trajectory We present two implementations that can be used to integrate geodesic motion in a Kerr background with an external force. We use the method of osculating elements extending previous work \cite{pound08} for Schwarzschild orbits to the Kerr background. The problem of motion under a small perturbation is well studied in celestial mechanics and is regularly applied to model the motion of satellites and small planets. A geodesic in Newtonian mechanics, or relativity, is uniquely characterized either by the three components of the particle position vector, ${\bf r}$, and the three components of the particle velocity, ${\bf \dot{r}}$, at any time or by six orbital constants (three orbital constants of the motion and three initial phases). There is a one-to-one correspondence between the two characterizations. This means that {\it any} trajectory can be instantaneously identified with a geodesic that has the same values of ${\bf r}$ and ${\bf \dot{r}}$. Of course, at two different instances of time, the geodesics will differ, but one can smoothly evolve the geodesic parameters to reproduce any nongeodesic trajectory. There are several approaches to do so and we describe these in the next subsection. \subsection{Osculating Elements or variation of constants} As mentioned above we can describe a bound stable geodesic by six parameters, which we denote by $I$. In the nonrelativistic case these parameters are simply $I = ( {\bf r}, {\dot {\bf r}})$, while for geodesic motion in Kerr we can take $I = \{E, L_z, Q, \psi_0, \phi_0, \chi_0\}$. Here $E$ is the energy, $L_z$ the azimuthal angular momentum, $Q$ is the Carter constant, and the remaining phases are defined in Sec.\ \ref{Kerr} below. At each instant we can therefore identify the true trajectory with a corresponding geodesic such that ${\bf r}$ and ${\bf\dot{r}}$ are the same. This imposes a particular choice of parameters, $I$, at each instance of time, and the whole trajectory is thus described by a sequence in the geodesic phase space, e.g., $I(t) = \{E(t), L_z(t), \iota(t), \psi_0(t), \phi_0(t), \chi_0(t)\}$. These are referred to as the {\it osculating orbital elements} at the {\it osculation epoch} $t$ \cite{bookPert}. Another name for this approach used in the Hamiltonian description is a {\it variation of constants}. We preserve the form of the equations of motion for a geodesic but slowly vary what used to be constants of motion in the unperturbed case. There are well known techniques for tackling such problems which are widely used in Newtonian celestial mechanics and can be extended to the relativistic regime. This was demonstrated by Pound and Poisson~\cite{pound08} for the trajectory of a particle in a Schwarzschild background under the action of (post-Newtonian) radiation reaction. When we have a perturbed system of the form \be {\bf \ddot{r}} = {\bf f}_{\rm geo} + \delta{\bf f}~, \en we can describe the perturbed trajectory using the osculating elements referred to the orbits of the geodesic system ${\bf \ddot{r}} = {\bf f}_{\rm geo}$. From the chain rule, any one of the osculating elements evolves as \be \dot{I} = \nabla_rI\cdot{\bf\dot{r}} + \nabla_v I\cdot{\bf\ddot{r}}~, \en in which the subscripts $r$ and $v$ denote derivatives with respect to the orbital position and velocity respectively. In the absence of the perturbing force, each osculating element is constant, so $\dot{I}=\nabla_rI\cdot{\bf\dot{r}} + \nabla_v I\cdot{\bf f}_{\rm geo}\equiv0$. The perturbation equations thus take the rather simple form \be \dot{I} = \nabla_vI \cdot \delta{\bf f}~. \label{gpe} \en Given an explicit expression for the perturbing force we can integrate these equations. The osculating element method can be formulated in several different ways. There is freedom in the parameterization of the geodesic solution that is used as a basis for deriving the osculating element equations, and in the basis used to prescribe the force. It is also possible to treat the orbital phase constants either as constants of the motion that are evolved explicitly or as part of a total phase variable which satisfies new equations that depend on the perturbation. We will describe two methods for treating the Kerr problem: (i) evolution of $E, L_z, Q$ and the \emph{full} orbital phases with the force prescribed with respect to the Kinnersley orthonormal tetrad; (ii) evolution of the orbital constants of motion $E, L_z, Q$ and the initial phases, with the force prescribed by its Boyer-Lindquist components. In the Hamiltonian approach we start with an unperturbed Hamiltonian, $H_0$ and write the equations of motion in terms of the constant canonical coordinates and momenta, $X^{\alpha}, P_{\alpha}$ (Hamilton-Jacobi approach), which are closely related, if not exactly the same, as the six constants of motion introduced above, $I$ \cite{Goldstein}. If we can describe the perturbation as a small addition $\delta H$ to the unperturbed Hamiltonian, then we can describe the equations of motion in the same generalized coordinates and momenta, which are no longer constants. The derivatives of the perturbation $\delta H$ give the equations for the evolution of $X^{\alpha}, P_{\alpha}$. Quite often those equations are solved iteratively starting with an assumption that the orbit is unperturbed in the right-hand side (in the $\delta H$). This is similar to the adiabatic solution to the osculating element equations which we will describe below. The Hamiltonian approach (if it can be formulated) would give equations equivalent to approach (ii) mentioned above. We note that an obvious method of computing inspirals is to numerically integrate the second-order forced geodesic equations directly, taking the fundamental variables to be the Boyer-Lindquist coordinates and their derivatives with respect to proper time. The key advantage of the methods discussed in this paper over second-order integrations is that they mesh much more naturally with the adiabatic approximation and more generally with two-time-scale approximation techniques \cite{FH}. For extreme-mass-ratio inspirals driven by radiation reaction, the orbital evolution time scale is much longer than the orbital time scale for most of the inspiral, until the orbit becomes close to the innermost stable orbit. The adiabatic approximation to the motion gives the motion as an expansion in the ratio of the time scales, and then there are various postadiabatic corrections to this. Although it is not possible yet to compute numerically the full first-order self-force for generic orbits in Kerr, it is possible to compute the averaged, dissipative piece of this force, which is sufficient to compute leading-order adiabatic inspirals \cite{FH}. The two-time-scale expansion also allows one to go beyond the adiabatic evolution and compute the small, rapidly oscillating perturbations to the evolution of the orbital variables, as well as the slow secular changes to higher order. The two-time-scale method cannot be easily applied to the second-order, forced geodesic equations, but it can be applied to the equations derived in this paper, as we discuss in Secs.~\ref{NLosc} and \ref{Kerr} below. In particular, the osculating elements method allows us to explicitly estimate the orbital average rate of change of the orbital elements. This gives us a physical insight into the effect of a perturbing force on the orbit which is otherwise obscured in the integration of the second-order equations of motion. Estimation of these secular changes also allows us to construct the adiabatic evolution of the orbit in the regime where it is applicable. \subsection{Numerical ``kludge'' waveform} Another application for the results described in this paper is for the construction of numerical kludge waveforms. The numerical kludge waveform for EMRIs is a fast and accurate way to compute the long waveforms \cite{Babak:2006uv} that will be needed for EMRI data analysis. These are built in a not entirely consistent way, but the basic philosophy is to model the underlying trajectory of the inspiralling object as accurately as possible in order to obtain the best possible phase match between the true and approximate waveforms. The approximation is based on geodesic motion in the MBH's spacetime, combined with a flat spacetime waveform generation expression. In the most recent version of the numerical kludge~\cite{Gair:2005ih}, the instantaneous geodesic orbit was updated by evolving the three constants of the motion $E, L_z, Q$ \cite{Glampedakis:2002cb} only. The evolution of the constants was obtained by combining post-Newtonian results with fits to numerical fluxes obtained by solving the Teukolsky equation \cite{Gair:2005ih}. However this method of evolving the geodesics is not complete, as we described above, since we need to evolve the (initial) orbital phases together with the orbital constants $E, L_z, Q$. In particular, the natural (and incorrect) way to evolve the phase constants, which is to fix them at some initial point, leads to significantly different evolutions in a time or frequency domain implementation of the kludge. The desire to resolve this apparent discrepancy between the two implementations was one of the initial motivations for the work described here. This article outlines the correct way to evolve geodesics under the self-force which could be used to further improve the numerical kludge waveforms in both time and frequency domain descriptions of them. \subsection{Main results and the structure of the paper} In this paper we will give a detailed description of the osculating elements approach applied to an arbitrary perturbing force acting on an object undergoing geodesic motion in the Kerr spacetime. As an introduction to the three dimensional relativistic problem of perturbed geodesic motion we will first consider a toy problem in Sec.~\ref{NLosc}. We look at the one-dimensional nonlinear oscillator acted on by an external force. The external force is chosen to have two components: a dissipative part and a conservative part (which just redefines the energy of the system). As we will see later this problem is a very good model for the main problem of perturbed motion in the Kerr spacetime. We show how two implementations of the osculating elements approach work in this simplified model and compare the exact evolution with the adiabatic approximation. The second of these two implementations [in which we evolve the energy and the initial time defined as $x(t_0) = 0$] allows us to treat the problem analytically in terms of Jacobi elliptic functions. This one-dimensional example allows the reader to understand the main approach which we then extend to the problem of forced geodesic motion in the Kerr spacetime in Sec.~\ref{Kerr}. We start that section with an introduction to our notation, before describing the osculating elements approach using the Kinnersley tetrad and ``Hughes'' variables (in terms of the orbital constants and the total phase variables).We then describe the forced geodesic equations in Boyer-Lindquist coordinates, evolving the orbital constants and the initial conditions, which is a direct extension to Kerr of the Schwarzschild results described in~\cite{pound08}. In both cases, we show how we can explicitly avoid the appearance of an apparent divergence in the osculating equations of motion at turning points. In Sec.~\ref{gasdrag}, we illustrate our techniques with a problem in which the perturbing force is a ``gas-drag'' force proportional to the velocity of the inspiralling compact object. This is a toy model for an object inspiralling in a gaseous environment around a MBH. We show that the different approaches give identical results, and once again compare the exact and adiabatic solutions to the problem. The influence of the drag force is to drive the inspiral of the object, but it also tends to increase the eccentricity of the orbit and decrease the orbital inclination. Although we primarily use this problem for illustrative purposes, the increase in eccentricity is an interesting result that could have observational consequences. The increase in eccentricity is a purely relativistic effect, and is to be expected generically, as we discuss in more detail in Appendix~\ref{A:SchDrag}, in the context of a drag force acting on an object in a Schwarzschild background. We summarize and discuss our findings in the concluding section \ref{S:sum}. Some detailed mathematical calculations are included in additional appendixes. \def\alt{ \mathrel{\raise.3ex\hbox{$<$}\mkern-14mu\lower0.6ex\hbox{$\sim$}} }

\label{S:sum} We have described two methods for integrating the equations of motion for bound, accelerated orbits in the Kerr spacetime, which are based on identifying the orbit with a geodesic at each point. The first method parametrizes the position and velocity of the orbit in terms of the conserved quantities (energy, axial angular momentum and Carter constant) in addition to three angular variables which increase monotonically and correspond to relativistic generalizations of the anomalies of Keplerian motion. The second method is the traditional ``osculating element'' technique which parametrizes the position and velocity of the orbit in terms of the geodesic with the same position and velocity. Practically, the second method differs from the first only in the treatment of the three phase variables, which are split up into a geodesic piece and a ``phase offset'' piece that is constant for geodesics. To illustrate the methods, we first analyzed, as a simpler model, a forced anharmonic oscillator. This was written in terms of a set of phase space coordinates. The forced equations of motion contained an apparent divergence at the turning points, but it was possible to reformulate the equations to eliminate the problematic terms and thus obtain equations of motion in a form without divergences. We discussed the adiabatic prescription for computing the leading order motion, which corresponds to a gradual evolution of the oscillator's amplitude and fundamental frequency driven by the phase space averaged forcing function for the amplitude. We presented an alternative analysis of this toy problem analogous to the osculating orbit method in terms of the analytic solution to the un-forced motion. By numerically integrating the equations we verified that both parametrizations gave the same results and compared these to the adiabatic approximation to the solution. Next, we showed that the equation of forced motion in the Kerr spacetime could be reformulated in a similar fashion. For the first method, it was advantageous to parametrize the force in terms of its components on the Kinnersley tetrad instead of using the instantaneous time derivatives of the conserved quantities. We derived a formulation of the equations of motion in terms of phase variables that was manifestly divergence-free at the turning points. We then generalized the second method, of osculating orbits, to generic orbits in the Kerr spacetime and showed how we could write down a divergent-free form of equations of this type without explicit simplification. As an application of our results, we considered the case of a simple force that could represent a gas drag. Numerical integrations of the equations of motion for a choice of parameters verified that the two methods of parametrizing the motion gave the same results. We identified a key observational signature of the presence of a drag force, namely, a decrease in the orbital inclination and an increase in eccentricity, which is opposite to the increase in inclination and decrease in eccentricity characteristic of the gravitational radiation reaction forces during the early stage of an inspiral. The first of our two methods has been applied to the study of transient resonances that occur in the radiation-reaction-driven inspirals of point particles into spinning black holes, using approximate post-Newtonian expressions for the self-force \cite{2010arXiv1009.4923F}. Other applications of this work will include the construction of accurate trajectories for orbits evolving under the action of the self-force, once self-force data for generic orbits are available. This will be essential for the construction of accurate gravitational waveforms for EMRIs, which will be needed for LISA data analysis. The formalism can also be used to estimate the magnitude of any secular changes in the orbital parameters that arise from the action of external perturbing forces. These could arise from gravitational perturbations from distant objects, such as stars or a second massive black hole, or from the presence of other material in the spacetime, such as the gas-drag which we considered in a simple way here. It will be very important to have a quantitative understanding of the importance of all these effects if intermediate-mass-ratio inspirals or EMRIs are to be used to carry out high-precision mapping of the spacetime around Kerr black holes and for tests of general relativity. Finally, the results described here will be useful to augment existing kludge models for inspiral waveforms. In particular, these methods will allow us to extract the secular part of the evolution of both the orbital constants of the motion and the phase constants, from self-force calculations. It is straightforward to include secular changes to the orbital parameters in the kludge framework~\cite{Babak:2006uv}, and by doing this it should be possible to ensure that the kludge waveform stays in phase with the true waveform for long stretches of the inspiral. It will be important to have accurate but cheap-to-calculate waveform models available when the data from gravitational wave detectors are analyzed, as this data analysis will rely heavily on matched filtering using template waveforms.

1012.5111

We present two methods for integrating forced geodesic equations in the Kerr spacetime. The methods can accommodate arbitrary forces. As a test case, we compute inspirals caused by a simple drag force, mimicking motion in the presence of gas. We verify that both methods give the same results for this simple force. We find that drag generally causes eccentricity to increase throughout the inspiral. This is a relativistic effect qualitatively opposite to what is seen in gravitational-radiation-driven inspirals, and similar to what others have observed in hydrodynamic simulations of gaseous binaries. We provide an analytic explanation by deriving the leading order relativistic correction to the Newtonian dynamics. If observed, an increasing eccentricity would thus provide clear evidence that the inspiral was occurring in a nonvacuum environment. Our two methods are especially useful for evolving orbits in the adiabatic regime. Both use the method of osculating orbits, in which each point on the orbit is characterized by the parameters of the geodesic with the same instantaneous position and velocity. Both methods describe the orbit in terms of the geodesic energy, axial angular momentum, Carter constant, azimuthal phase, and two angular variables that increase monotonically and are relativistic generalizations of the eccentric anomaly. The two methods differ in their treatment of the orbital phases and the representation of the force. In the first method, the geodesic phase and phase constant are evolved together as a single orbital phase parameter, and the force is expressed in terms of its components on the Kinnersley orthonormal tetrad. In the second method, the phase constants of the geodesic motion are evolved separately and the force is expressed in terms of its Boyer-Lindquist components. This second approach is a direct generalization of earlier work by Pound and Poisson [A. Pound and E. Poisson, Phys. Rev. DPRVDAQ1550-7998 77, 044013 (2008).10.1103/PhysRevD.77.044013] for planar forces in a Schwarzschild background.

1244482

2010evn..confE..80G

1012

1012.2663_arXiv.txt

Narrow line Seyfert 1 (NLS1) are Active Galactic Nuclei (AGN) characterized by unusual optical spectra, with H$\beta$ line FWHM $< 2000$ km/s, a line intensity ratio $[$OIII$]/$H$\beta < 3$, and a bump in FeII \citep{Pogge2000}. A small fraction of NLS1 is known to be radio loud and in these cases the flat radio spectra and VLBI variability suggest that several of them could host relativistic jets \citep{Komossa2006,Doi2006}. A conclusive evidence for the existence of relativistic jets in these AGN has been obtained thanks to the Large Area Telescope (LAT) on board {\it Fermi}, which has revealed bright gamma-ray emission from the radio loud NLS1 PMN J0948+0022 just a few week after the launch, with flux at $E>200$ MeV is $4.0\pm0.3 \times 10^{-8}$ ph cm$^{-2}$ s$^{-1}$ and photon index $\Gamma = 2.6\pm0.1$ \citep{discovery}, making it one of the brightest gamma-ray AGNs at high latitude \citep{lbas}; using multi wavelength (MWL) data, it has been possible to estimate some basic parameters for this source, such as the size of the emitting region, the magnetic field value, and the particle energy range \citep{discovery,Foschini2010}. A more accurate view could only be attained through a dedicated study based on simultaneous broad band data and therefore a large campaign was organized between March and July 2009, covering gamma-rays, X-rays, UV, optical, and radio; indeed, although radio emission is generally produced somewhat out of the gamma-ray zone, single dish observations can still yield valuable complementary information on correlated flux and spectral variability; in addition, VLBI observations give independent constraints on the jet Doppler factor and viewing angle. The main results from such extended campaign were reported by \citet{mwl}, while in this contribution we present some additional facts about the global e-VLBI observations specifically organized for this project. Further details and a more extensive interpretation will also be given in a dedicated forthcoming publication.

The census of extragalactic gamma-ray sources has been dominated by blazars (flat spectrum radio quasars and BL Lac type objects) through all the EGRET era. A handful of radio galaxies were also detected. The discovery of gamma-ray emission from radio-loud narrow-line Seyfert1 nuclei is therefore of great importance for the study of relativistic jets in AGNs \citep{Abdo2009d}. First, it has permitted to seal the issue whether such relativistic structures do exist in radio loud NLS1, as proposed by \citet{Zhou2003} and \citet{Doi2006}. Moreover, thanks to the sensitivity and the surveying capability of {\it Fermi}, it has triggered the organization of a campaign to measure simultaneously in different bands to construct the spectral energy distribution of the source, as well as at studying its variability across the electromagnetic spectrum. In this framework, high frequency Global VLBI observations are of great value, as they permit us to observe the most compact regions of the radio jet above the self-absorption frequency. Clearly, the maximum return can be obtained with real time VLBI, both because it delivers its results in a suitable time for coordinated multi-wavelength studies and because it permits to observers to monitor the array in real time and directly address problems in the system -- which are likely to arise when a Global array is used. The overall success of the MWL campaign \citep{mwl} and in particular of the Global e-VLBI observations presented in this contribution is highly encouraging for the continuation of the synergy between high energy astrophysics and VLBI. More generally, Global e-VLBI has proven to be an excellent and reliable tool and we expect it to become a widely used resource in the coming years.

1012.2663

The detection of gamma-ray emission by Fermi-LAT from the radio loud Narrow Line Seyfert 1 PMN J0948+0022 (Abdo et al. 2009, ApJ 699, 976) triggered a multi-wavelength campaign between March and July 2009. Given its high compactness (Doi et al. 2006, PASJ 58, 829), inverted spectrum, and 0deg declination, the source was an ideal target to observe at 22 GHz with a Global VLBI array extending from Europe to East Asia and Australia. In order to deliver prompt results to be analysed in combination with the other instruments participating in the campaign, the observations were carried out with real time VLBI, for the first time on a Global scale. Indeed, the main results have been published just a few months after the campaign (Abdo et al. 2009, ApJ 707, 727). Here we present additional details about the e-VLBI observations.

12225704

2011PhRvD..83h3512K

1012

1012.2039_arXiv.txt

Observations of clusters of galaxies, rotations curves of individual galaxies, cosmic microwave background anisotropies, and many other methods suggest the existence of dark matter. A possible realization of dark matter might be in the form of Weakly Interacting Massive Particles (WIMPs). A huge effort is being undertaken by experimentalists to directly detect WIMPs in underground or space experiments, as well as by theorists to incorporate them into viable theories beyond the Standard Model. The situation experimentally is still not clear, as the majority of the experiments have not detected WIMPs so far. Direct search experiments with Earth based detectors like CDMS~\cite{Ahmed:2009zw} and Xenon~\cite{Angle:2008we} have imposed constraints on the WIMP-nuclei cross sections, assuming the local dark matter density around the Earth as inferred from the cosmological and other data (see e.g. Ref.~\cite{Dunkley:2008ie} for the determination of the amount of dark matter from the WMAP data). On the other hand, DAMA experiment~\cite{Bernabei:2010mq} claims dark matter detection with parameters that contradict other experiments if taken at face value. Given the still unclear picture regarding the nature of dark matter, it is of crucial importance to constrain as much as possible the WIMP candidates, including their mass and interactions. Several such candidates exist in the market depending on what theory beyond the Standard Model one chooses, ranging from supersymmetry ~\cite{Jungman:1995df,Bertone:2004pz} and hidden sectors~\cite{Pospelov:2007mp,Hambye:2008bq}, to Technicolor~\cite{Gudnason:2006yj,Kouvaris:2007iq,Belotsky:2008vh,Kouvaris:2008hc,Ryttov:2008xe,Frandsen:2009mi,Kainulainen:2010pk}. The WIMPs can be classified according to their properties, i.e. if they are produced thermally, if they are asymmetric~\cite{Barr:1990ca,Gudnason:2006yj,Belyaev:2010kp}, if they have spin-dependent or spin-independent cross section with the nuclei, if their collisions with the nuclei are elastic or inelastic~\cite{TuckerSmith:2001hy,TuckerSmith:2004jv,Fargion:2005ep,Khlopov:2007ic,Khlopov:2008ty}, and/or whether they are self-interacting~\cite{Spergel:1999mh,Dave:2000ar,Zentner:2009is,Frandsen:2010yj}. Apart from direct searches, constraints on the properties of the WIMPs might arise from astrophysical observations as for example in~\cite{Frandsen:2010mr}. Concentration of the WIMPs within stars can affect, under certain circumstances, the evolution of the latter, and/or products of WIMP annihilation within the stars could be directly or indirectly detected. The capture of WIMPs in the Sun and the Earth~\cite{Press:1985ug,Gould:1987ju,Gould:1987ww} has been used to predict a possible signature for an indirect detection of dark matter based on neutrino production due to WIMP co-annihilation~\cite{Jungman:1994jr,Nussinov:2009ft}. Constraints on the dark matter properties and the dark matter profile can also be imposed due to the effect of dark matter on the evolution of low mass stars~\cite{Casanellas:2009dp,Casanellas:2010sj}, and main sequence stars~\cite{Scott:2008ns}, on possible gravitational collapse of neutron stars~\cite{Goldman:1989nd}, and on the cooling process of compact objects such as neutron stars and white dwarfs ~\cite{Kouvaris:2007ay,Sandin:2008db,Bertone:2007ae,McCullough:2010ai,Kouvaris:2010vv,deLavallaz:2010wp}. In particular, the authors of~\cite{Goldman:1989nd} have investigated under what conditions a neutron star can collapse gravitationally due to accretion of WIMPs, providing an upper bound for the WIMP masses given the local dark matter density and the time of accretion. Although the bound for the mass of a bosonic WIMP was low $\sim 10$~MeV, the upper bound for fermionic WIMPs was quite high $\sim 10^5$~TeV (and therefore not relevant for physics at the TeV scale). In this paper we investigate possible constraints that can arise from stars that accrete WIMPs during their lifetime and then collapse into a more compact object, white dwarf or a neutron star, inheriting the accumulated dark matter. Depending on the location of the star and the WIMP-nuclei cross section, it might be possible to impose constraints on the mass of the WIMP, excluding in some cases candidates that are lighter than TeV. Such constraints improve significantly the existing ones, and may become relevant for LHC physics. More specifically, we consider two different cases. In the first case, we examine the accretion of WIMPs with spin-dependent interactions with protons onto a Sun-like star. We deduce under what circumstances the accumulated WIMPs can trigger a gravitational collapse once the star has turned into a white dwarf. In the second case, a supermassive star accretes WIMPs which have a spin-dependent cross section with nucleons (protons and neutrons). A typical supermassive star of 15 solar masses lives about $10^7$~yr and then explodes forming a neutron star. Under certain assumptions, the WIMPs inherited by the neutron star from its progenitor will thermalize, sink to the center and, for some range of parameters, collapse further into a black hole. Thus, the mere existence of neutron stars might impose constraints on the mass or the cross section of the WIMPs. In all cases we will assume cross sections that are compatible with experimental constraints from the direct dark matter searches.

We derived constraints on the spin-dependent cross section of asymmetric fermionic dark matter WIMPs based on the existence of white dwarfs and neutron stars in globular clusters. Our constraints are competitive to direct dark matter search experiments, excluding a large parameter space of cross sections and masses as low as TeV (or slightly lower than TeV). In the case of white dwarfs, we were able to exclude a range of spin-dependent cross sections and WIMP masses because for these parameters, WIMPs that have been captured during the lifetime of the progenitor have enough time to concentrate within the core of the star that is inherited by the white dwarf, and eventually collapse gravitationally forming a black hole that destroys the star. This constraint is robust in the sense that it depends only on the local dark matter density of the globular cluster and no other hypothesis. We demonstrated that asymmetric WIMP candidates with only spin-dependent interactions with masses even lower than one TeV, trapped during the lifetime of the progenitor can easily thermalize inside the white dwarf due to expected small abundances of isotopes of carbon or oxygen that carry spin. If we now relax the strict condition of thermalization, i.e. to assume that WIMPs are gravitationally trapped by the progenitor but are not necessarily confined within the radius of a white dwarf, our constraints can be extended to higher masses and lower cross sections. However, in this case we have to make an extra assumption of a mechanism that redistributes the WIMP velocities in order for WIMPs which are on orbits around the white dwarf to intersect with it. Such a mechanism can be possibly provided by binaries or WIMP-WIMP interactions. Further investigation of this possibility is needed. In the case of neutron stars, we exclude an area of WIMP masses and cross sections due to direct accretion of WIMPs in the neutron star. The constraints depend only on the local dark matter density of the globular cluster and the age of the neutron star. Upon assuming an extra mechanism of WIMP velocity redistribution, extra parameter space may be excluded down to the TeV scale.

1012.2039

We put constraints on asymmetric dark matter candidates with spin-dependent interactions based on the simple existence of white dwarfs and neutron stars in globular clusters. For a wide range of the parameters (WIMP mass and WIMP-nucleon cross section), weakly interacting massive particles (WIMPs) can be trapped in progenitors in large numbers and once the original star collapses to a white dwarf or a neutron star, these WIMPs might self-gravitate and eventually collapse forming a mini-black hole that eventually destroys the star. We impose constraints competitive to direct dark matter search experiments, for WIMPs with masses down to the TeV scale.

12205305

2011IAUS..274..110R

1012

1012.0720_arXiv.txt

Open questions in stellar physics led to the idea that a dynamo operates in the radiative cores of early-type stars (Spruit 2002). Even the helioseismologic observation of rigid rotation of the solar interior shows in this direction. The angular momentum transport by the large-scale magnetic field pattern (fossil field plus toroidal field induced by differential rotation) does {\em not} lead to a solid-body rotation unless the viscosity of the plasma exceeds the molecular value by a few orders of magnitude (R\"udiger \& Kitchatinov 1996, Eggenberger et al. 2005). Other examples are given by the evolution of the fast rotating early-type stars which can only be understood if i) there is a basic transport of angular momentum outwards and ii) the radial mixing of chemicals remains weak (Yoon et al. 2006, Brott et al. 2008). Hence, if a (magnetic-induced) instability existed in the radiative stellar cores then the corresponding Schmidt number ${\rm Sc}= \nu/D$ must be rather large. We have shown that the kink-type instability (or Tayler instability, TI) of toroidal magnetic fields forms a much-promising candidate for the instability. A Schmidt number larger than ten results as the ratio of the effective viscosity and the diffusion coefficient (R\"udiger et al. 2009). The unstable modes of the TI are basically nonaxisymmetric driven by the energy of the electrical current which produces the toroidal field. Interestingly enough, there exists even an instability of a toroidal magnetic field which in the fluid is current-free. In this case the energy comes from a differential rotation which itself is stable but which is unstable under the influence of the (current-free) toroidal field. We have called this instability as Azimuthal MagnetoRotational Instability (AMRI) as -- like for the standard MRI (with axial fields) -- the field itself is current-free and does not exert forces. In opposition to the standard MRI the AMRI is always nonaxisymmetric and it is, therefore, much more interesting for the dynamo theory. For complicated radial profiles of the toroidal field we shall always have a mixture of TI and AMRI. Generally, the latter is more important for fast rotation ($\Omega > \Omega_{\rm A}$) and v.v. Here the Alfv\'{e}n frequency $\Omega_{\rm A}$ for the toroidal field is used which derives from the Alfv\'{e}n velocity $v_{\rm A}= B_\phi/\sqrt{\mu_0 \rho}$ as the related frequency. Between two cylinders with different radii the toroidal field profile with $B_\phi = A R + B/R$ ($R$ radius) is free of dissipation. The `perfect' AMRI appears for $A=0$ while the `perfect' TI results for $B=0$. One can compute the necessary electrical currents to excite both sorts of instabilities in a columnar Taylor-Couette experiment with gallium as fluid conductor. As we have shown the critical Hartmann numbers for self-excitation of axi- and nonaxisymmetric perturbation modes do not depend on the magnetic Prandtl number of the fluid which is as small as $10^{-3}$ for stellar plasma and $10^{-5}$ for liquid sodium (R\"udiger \& Schultz 2010). It is typical for the nonaxisymmetric TI and AMRI that always the two modes with $m=\pm 1$ are excited for the same critical Hartmann number and also -- if supercritical -- with the same growth rates (Fig.~\ref{fig1}). Despite their simultaneous existence they can be excited as singles with different initial conditions. However, if the initial conditions are as neutral as possible with respect to a preferred helicity, in the majority of the cases one of the modes dominates after our experiences. This behavior may have dramatic consequences with respect to the dynamo theory. Both the modes with $m=\pm 1$ have opposite helicity with the same total amount. The mode with $m=-1$ is identical to the mode with $m=1$ but in a left-hand system. The helicity of $m=1$ in the right-hand system equals the helicity of $m=-1$ in the left-hand system. So it is obvious that in one and the same coordinate system the sum of the helicity of $m=-1$ and $m=1$ is zero. As a consequence the instability of a toroidal field can only develop helicity if by some reasons one of the modes with $m=\pm 1$ dominates the other. There are the two possibilities that i) one mode dominates the other by chance (so as the matter dominates the antimatter) or ii) the existence of a poloidal field prefers one of the modes. We have shown that indeed in stellar radiation zones -- if the background field has a positive current helicity $\bar{\vec{B}} \cdot {\rm curl} \bar{\vec{B}}$ -- the resulting kinetic helicity $\langle \vec{u}' \cdot {\rm curl} \vec{u}'\rangle$ of the fluctuations is always negative (Gellert et al. 2011). A positive current helicity of the background field results if an axial field and an axial electrical current are parallel. A negative current helicity of the background field results if an axial field and an axial electric current are antiparallel. Hence, the resulting kinetic helicities on the basis of current-driven instabilities have, therefore, opposite signs in opposite hemispheres of the model.

Unstable toroidal fields alone are not able to produce helicity and $\alpha$-effect. It has been shown, however, that helicity and $\alpha$-effect are produced by unstable magnetic large-scale field patterns which themselves possess current helicity $\bar{\vec{B}}\cdot {\rm curl} \bar{\vec{B}}$. This is insofar understandable as helicity and $\alpha$-effect are both pseudo-scalars which can only be nonvanishing if in the global setup a nonvanishing large-scale pseudo-scalar like $\bar{\vec{B}}\cdot {\rm curl} \bar{\vec{B}}$ exists. We want to stress, however, that the current helicity of the background field is not the only possible pseudo-scalar existing in magnetized stellar radiation zones on which other forms of helicity and $\alpha$-effect may base.

1012.0720

It is shown that the magnetic current-driven (`kink-type') instability produces flow and field patterns with helicity and even with α-effect but only if the magnetic background field possesses non-vanishing current helicity B. curl B by itself. Fields with positive large-scale current helicity lead to negative small-scale kinetic helicity. The resulting α-effect is positive. These results are very strict for cylindric setups without z-dependence of the background fields. The sign rules also hold for the more complicated cases in spheres where the toroidal fields are the result of the action of differential rotation (induced from fossil poloidal fields) at least for the case that the global rotation is switched off after the onset of the instability.

2524222

2011ApJ...728L..11R

1012

1012.3723_arXiv.txt

\label{sec:intro} Jets and outflows from protostellar objects are fundamental aspects of the current star formation paradigm, and are observed anywhere star formation is ongoing. The mechanism proposed by \citet{bp82}, in which jets are launched from accretion discs by gravitational, magnetic, and centrifugal forces, has been extensively studied numerically (\eg, \citealt{us85, megpl97, op97a, op97b, op99, klb99, vjo02, vR03, ocp03, pf10, snopc10}). By treating the accretion disc as a boundary condition (\eg, \citealt{ukrcl95}), one can study jet dynamics independently of the disc (\eg, \citealt{pofb07_ppv}) though, in order to resolve the launching mechanism, numerical simulations have not followed the jet beyond 100 AU (\eg, \citealt{alkb05}). In stark contrast, protostellar jets are $\gtrsim10^4$ AU long \citep{brd07_ppv}, and only recently have observations reached within 100 AU of the source (\eg, \citealp{hep04,cbp08}). This large scale difference between observations and simulations makes direct comparisons difficult and, in this work, we aim to close this gap. We present axisymmetric (2.5-D) simulations of protostellar jets launched from the inner AU of a Keplerian disc, and follow the jet well into the observational domain (2500 AU). These calculations allows us to address the efficacy of the magnetocentrifugal mechanism, and to relate conditions near the disc with directly observable properties of the jet. The simulations presented herein are performed with an adaptive mesh refinement (AMR) version of \zeus \citep{c96,c10} called \azeus (Adaptive Zone Eulerian Scheme). The \zeus family of codes are among the best tested, documented, and most widely used astrophysical MHD codes available, though this is the first attempt to couple \zeus with AMR\footnote{\textsl{ENZO\/}, a hybrid N-body Eulerian code \citep{osheaetal04}, links AMR with the \emph{hydrodynamical} portion of \textsl{ZEUS-2D}.}. We have implemented the block-based method of AMR detailed in \citet{bc89} and \citet{bbsw94}. Significant effort was spent minimising errors caused by passing waves across grid boundaries, which is of particular importance to this work. A full description of the code and the changes required for AMR on a fully-staggered mesh will appear in Ramsey \& Clarke (in preparation).

\label{sec:discuss} We have presented the first MHD simulations of protostellar jets that start from a well-resolved launching region ($\Delta{z}_{\rm min}=0.00625$~AU) and continue well into the observational domain (2500 AU). On the AU scale, each jet shows the characteristic and near steady-state knotty behaviour first reported by \citet{op97b}, though the origin of our knots is quite different. On the 1000 AU scale, each jet develops into a ballistic, supersonic ($8\lesssim{M}\lesssim11$) outflow led by a magnetically confined ``nose-cone" \citep{cnb86} and a narrow bow shock, consistent with what is normally observed. On comparing Tables \ref{tab:genchar} and \ref{tab:simchar}, our simulations comfortably contain virtually all observed protostellar jets on these four important quantities. We note that these tables would \emph{not} have been in agreement had we stopped the jet at, say, 100 AU and measured these values then. \emph{It is only because our jets have evolved over five orders of magnitude in length scale that we can state with some confidence that the magnetocentrifugal launching mechanism is, by itself, capable of producing jets with the observed proper motion, rotational velocity, radius, and mass outflow rate.} Indeed, our jets are still very young, having evolved to only 100 yr, and allowing them to evolve over an additional one or two orders of magnitude in time may still be useful. For example, it would be interesting to know whether $\langle\beta\rangle$ rises above unity for any of the jets (Fig.\ \ref{fig:pm}a), and thus enter into a hydrodynamically dominated regime. It would also be interesting to see how long it takes for the power laws in jet radius and mass flux as a function of $B_{\rm i}$ to reach their asymptotic limits. Our jet widths tend to be higher than those observed, particularly when one considers that the values for $r_{\rm jet}$ in Table \ref{tab:simchar} are at $t=20$~yr\footnote{Some simulations had not reached $t=100$~yr at the time of this writing.}, and that $r_{\rm jet}$ continues to grow in time (\eg, for the $\beta_{\rm i}=40$ jet, $2\,r_{\rm jet}\sim100$~AU by $t=100$~yr). As our jet radii mark the locations of the contact discontinuity while observed radii mark hot, emitting regions, our widths should be considered upper limits. That our values \emph{contain} all observed jet widths is a success of these simulations. Similarly, our numerical mass fluxes are higher than observed values by at least an order of magnitude. Since observed mass-loss rates account only for emitting material (\eg, in forbidden lines; \citealt{hmr94}), and thus temperatures in excess of $10^4$ K (\citealt[p.\ 104]{dw97_book}), our mass fluxes are necessarily upper limits as well. Indeed, if we measure our mass fluxes near the jet tip (instead of at 200 AU for Table \ref{tab:simchar}) and restrict the integration to fluid above $10^4$ K, our mass fluxes drop by a factor of 10--100, in much better agreement with Table \ref{tab:genchar}.

1012.3723

We present the first 2.5-dimensional magnetohydrodynamic (MHD) simulations of protostellar jets that include both the region in which the jet is launched magnetocentrifugally at scale lengths &lt;0.1 AU and where the propagating jet is observed at scale lengths &gt;10<SUP>3</SUP> AU. These simulations, performed with the new adaptive mesh refinement MHD code AZEuS, reveal interesting relationships between conditions at the disk surface, such as the magnetic field strength, and direct observables such as proper motion, jet rotation, jet radius, and mass flux. By comparing these quantities with observed values, we present direct numerical evidence that the magnetocentrifugal launching mechanism is capable, by itself, of launching realistic protostellar jets.

1536342

2011ApJ...728..106S

1012

1012.3515_arXiv.txt

Over the last two decades, our view of the stellar halo has changed from a diffuse, featureless cloud of stars surrounding the Galaxy, to one crossed by many large-scale features such as the tidal tails of the Sagittarius dwarf galaxy \citep{1994Natur.370..194I,1995MNRAS.277..781I,2003ApJ...599.1082M}, the Virgo overdensity \citep{2008ApJ...673..864J} the Triangulam-Andromeda structure \citep{2004ApJ...615..732R,2004ApJ...615..738M,2007ApJ...668L.123M} and the low latitude Monocerous ring \citep{2002ApJ...569..245N,2003ApJ...588..824Y}. The mapping of these low surface brightness structures can be attributed to the advent of large scale surveys such as the Two Micron All Sky Survey (2MASS) and the Sloan Digital Sky Survey (SDSS) that have large numbers of stars in their catalogues. Future surveys, such as GAIA \citep{perryman02}, LSST \citep{2009AAS...21346003I} SkyMapper \citep{2007PASA...24....1K} and PanSTARRS, will map the stellar halo in ever more detail. The presence of these structures lends support to the $\Lambda$CDM model of galaxy formation in which the stellar halo is built up, at least in part, hierarchically through mergers of smaller satellite systems \citep{2001ApJ...548...33B}. However, simulating the stellar halo in a cosmological context to test this picture is a challenging task for two reasons. First, the stellar halo is intrinsically faint, containing only about $1$ per-cent of the total Milky Way stars. Hence, to simulate the faint structures within the stellar halo with adequate resolution requires enormous computation power. For example, if a satellite with a stellar mass of $10^5 \Msun$ is simulated with at least $10^3$ particles, then to simulate a whole galaxy having stellar mass $10^{11} \Msun$ with the same mass resolution will require a simulation with more than $10^9$ stellar particles. Second, the physics of star formation and its feedback effects are complex phenomena that have not yet been fully understood. In recent years progress has been made in tackling both these issues. Hydrodynamical simulations of galaxies including star formation and feedback recipes are now being done \citep{2003ApJ...591..499A, 2004ApJ...606...32R,2004ApJ...612..894B,2007MNRAS.374.1479G, 2008MNRAS.389.1137S,2009ApJ...702.1058Z}. However, the highest resolution simulations only have a stellar mass resolution of $10^4-10^5 \Msun$, which is not enough to resolve the features in the stellar halo. Alternatively, assuming that the stellar halo is built up entirely by means of accretion events, hybrid techniques have been developed that use collisionless simulations to follow the dynamical evolution of stars in an analytical potential and a semi-analytical prescription to incorporate star formation processes \citep{2005ApJ...635..931B}. Although these hybrid techniques are not fully self-consistent they can produce realistic models of the stellar halos with very high resolution and can resolve even the lowest luminosity structures. Recent improvements in hybrid techniques have implemented the semi-analytical star formation recipes directly into cold dark matter simulations \citep{2009MNRAS.397L..87L, 2008MNRAS.391...14D,2010MNRAS.406..744C}, so that with current dark matter simulations like Via Lactea-II \citep{2008Natur.454..735D} and Aquarius \citep{2008Natur.456...73S}, reaching a resolution of over $10^9$ dark matter particles within the virial radius, the future looks promising. With the tremendous progress in both theory and observations it is now possible to compare the two quantitatively. Since substructures in a system are fluctuations in the density field one way to make this comparison is to analyze the statistics of fluctuations. For example, \cite{2008ApJ...680..295B} looked at the fractional root mean square deviation of the stellar density from a smooth triaxial model in the main sequence turn off stars (hereafter MSTO) of SDSS data. They compared the radial dependence of these deviations with those apparent in the $11$ simulated stellar halos of \cite{2005ApJ...635..931B} and found the observations and simulations to be in rough agreement. One could in principle also do a Fourier analysis of the fluctuations or study the angular correlation function. While such analyses are useful for studying the global statistical properties of structures they cannot say much about individual structures. Also when calculating the rms deviations it is not clear if the deviations are due to structures or is it due to the inability of the analytical model to describe the smooth component of the halo. An alternative approach to comparing fluctuations apparent in observed and simulated stellar halos is to use group finding or clustering algorithms. The strength of the group-finding approach for this particular problem is that, in addition to simple comparisons of group properties, the results have a direct and simple physical interpretation. It was shown in \cite{2008ApJ...689..936J} that if substructures standing out above the background density can be identified (i.e. as groups) then the distribution of the properties of these structures can in principle be related to a galaxy's accretion history in terms of the characteristic epoch of accretion and the mass and orbits of progenitor objects. Hence, group finding can be used to constrain the accretion history of our Galaxy and compare it to general expectations of the $\Lambda$CDM paradigm. For example, in \citet{2010ApJ...722..750S} we applied a hierarchical clustering algorithm \citep{2009ApJ...703.1061S} to a sample of M-giants selected from the 2MASS catalogue and recovered sixteen structures, many of which were previously known and associated with individual accretion events. In this paper, we go on to compare the properties of these recovered structures with those found in synthetic surveys -- designed to mimic the observations in the 2MASS catalogue -- of simulated stellar halos \citep[taken from][]{2005ApJ...635..931B}, in order to assess how similar they are. We then ask what type of accretion events (in terms of satellite accretion time, luminosity and orbit) are recovered in the synthetic surveys, in order to assess how sensitive we expect the 2MASS survey be to our Galaxy's own history. The group finder is also applied to synthetic surveys of the simulated stellar halos chosen to mimic current and near-future photometric catalogues of other stellar tracers. For example, SDSS, which has a magnitude limit of $r<22.5$, was able to map the stellar halo with MSTO stars out to $36 \kpc$ with about 4 million MSTO stars \citep{2008ApJ...680..295B}. A similar MSTO generated from LSST -- which will observe the sky in six photometric bands, $ugrizy$, with a single visit limiting magnitude of $r\sim24.5$ (or $r\sim27.5$ using co-added maps) -- will be capable of probing the halo out to $100 \kpc$ with more than 20 million stars. In addition, both LSST and PanSTARRS will explore the transient optical sky and should be able to detect variable RR Lyrae stars out to $400 \kpc$ and beyond, which is more than the expected edge of the stellar halo. For group finding we use the code EnLink \citep{2009ApJ...703.1061S}, which is a density-based hierarchical group-finder. As noted in \citet{2010ApJ...722..750S}, this code is ideally suited for this application for four reasons. First, structures in the stellar halo have arbitrary shapes and a density-based group-finder is able to identify just such irregular groups. Second, while many density-based group-finders consider only groups formed above a fixed isodensity contour, EnLink's clustering scheme can identify groups at all density levels. This is essential because stellar halo structures span a wide range of densities that cannot be separated by a single isodensity contour. Third, halo structures can have nested substructures and EnLink's organizational scheme allows the detection of this full hierarchy of structures. Finally, the group-finder gives an estimate of the significance of the groups, so spurious clusters can be ignored. The paper is organized as follows. \sec{datasets} describes the data sets used in the paper. \sec{methods} discusses the methods employed for analyzing the data--- i.e., group finding. In \sec{results1} we ask how much any photometric survey can tell us about substructures in the stellar halo by looking at an idealized survey, simulated in the absence of observational errors. In \sec{results2} we look at groups derived from simulated surveys constructed under more realistic circumstances. In \sec{results3} we discuss the implications of these results for groups recovered from current surveys of M-giant stars (from 2MASS) and MSTO stars (from SDSS) in the real Universe. Finally we summarize our findings in \sec{summary}.

\label{sec:summary} In this paper, we have explored the power of a group finding algorithm to recover structures from photometric surveys of the stellar halo and interpret their properties in terms of Galactic accretion history. We first applied our group finder to idealized synthetic stellar surveys, which were generated from our $\Lambda$CDM models without accounting for observational errors. We find a simple dependence for the probability of detecting debris as a group on the parameters of its progenitor accretion event: the probability is highest for recent (small $t_{\rm acc}$) and high luminosity (large $L$) events accreted along circular orbits (large $\epsilon$). The strongest dependence is on $t_{\rm acc}$. The properties of recovered groups --- the number of and fraction of material in groups, along with distribution of the stellar mass and radial distance of the groups --- can in principle place constraints on the {\it recent} accretion history of a halo. Ancient accretion events ($> 10\Gyr$ ago) are not recovered as groups even in the absence of observational limitations because they are too phase-mixed to appear as distinct structures. We then applied our group finder to synthetic surveys that contained more realistic observational errors. Our results emphasize that the capability of a photometric survey to discover structures depends upon its sample size, the distance uncertainty, the depth and the relative sampling probability of different stellar populations. The broadest constraints on accretion history will come by combining the results of current and future surveys. For example, M-giants selected from 2MASS are intermediate age, high metallicity stars and hence are good tracers of relatively recent, high-luminosity accretion events with little contamination from older events. An LSST MSTO survey would contain a range of stellar populations whose properties depend on the severity of the color cut made to select the sample. Limiting the sample to the very bluest MSTO stars increases the dominance of low-metallicity stars in the sample, and hence the sensitivity to low-luminosity accretion events. RR Lyraes selected from LSST are bright enough to probe beyond 100 kpc where the apocenters of the more eccentric orbits lie, and hence will find a more fair sampling of the orbital properties of accretion events. Finally, a quantitative comparison of the results of applying the group-finder to the real 2MASS M-giant sample with those from the mock (synthetic) 2MASS M-giant surveys shows the number and properties of substructures in 2MASS M-giant survey to be roughly in agreement with simulated $\Lambda$CDM stellar halos. These groups most likely correspond to satellites accreted more recently than $10 \Gyr$, with luminosity higher than $5 \times 10^6 \Lsun$ and preferably on orbits of low eccentricity. Overall we conclude that current and near-future photometric surveys are poised to provide a complete census of our Galaxy's recent accretion history. The current results from 2MASS alone map the highest-luminosity recent events, future deep MSTO surveys will fill in the lower end of the luminosity function and RR Lyrae surveys will find debris structures that may be currently missing because the progenitor satellites were on highly radial orbits. Reconstructing more ancient accretion will require additional dimensions of data, such as velocity (proper motions and radial velocities of stars) and chemical abundance information.

1012.3515

The Sloan Digital Sky Survey (SDSS) and the Two Micron All Sky Survey (2MASS) provided the first deep and global photometric catalogs of stars in our halo and not only clearly mapped its structure but also demonstrated the ubiquity of substructure within it. Future surveys promise to push such catalogs to ever increasing depths and larger numbers of stars. This paper examines what can be learned from current and future photometric databases using group-finding techniques. We compare groups recovered from a sample of M-giants from 2MASS with those found in synthetic surveys of simulated ΛCDM stellar halos that were built entirely from satellite accretion events and demonstrate broad consistency between the properties of the two sets. We also find that these recovered groups are likely to represent the majority of high-luminosity (L &gt; 5 × 10<SUP>6</SUP> L <SUB>sun</SUB>) satellites accreted within the last 10 Gyr and on orbits with apocenters within 100 kpc. However, the sensitivity of the M-giant survey to accretion events that were either ancient from low-luminosity objects or those on radial orbits is limited because of the low number of stars, bias toward high-metallicity stars, and the shallow depth (distance explored only out to 100 kpc from the Sun). We examine the extent to which these limitations are addressed by current and future surveys, in particular catalogs of main-sequence turnoff (MSTO) stars from SDSS and the Large Synoptic Survey Telescope (LSST), and of RR Lyrae stars from LSST or PanSTARRS. The MSTO and RR Lyrae surveys are more sensitive to low-luminosity events (L ~ 10<SUP>5</SUP> L <SUB>sun</SUB> or less) than the 2MASS M-giant sample. Additionally, RR Lyrae surveys, with superior depth, are also good at detecting events on highly eccentric orbits whose debris tends to lie beyond 100 kpc. When combined we expect these photometric surveys to provide a comprehensive picture of the last 10 Gyr of Galactic accretion. Events older than this are too phase mixed to be discovered. Pushing sensitivity back to earlier accretion times would require additional dimensions of information, such as velocity and metallicity or abundance measurements.

12231056

2011PThPh.126..435H

1012

1012.5455_arXiv.txt

\label{sec:intro} Cosmological observations have established the existence of dark matter. Nowadays, the energy density of dark matter is precisely determined using the Wilkinson Microwave Anisotropy Probe (WMAP) satellite \cite{Hinshaw:2008kr}. The standard model of particle physics, however, can explain neither its existence nor its nature, which has been a mystery in particle physics and cosmology. The Weakly Interacting Massive Particles (WIMPs) in models beyond the standard model are a good candidate for dark matter. The relic abundance is naturally consistent with the cosmological observation if WIMPs have TeV-scale mass and they are thermally produced in the early universe. This is the so-called thermal relic scenario. In this scenario, the dark matter is originally nonrelativistic and acts as cold dark matter in the era of the structure formation of the universe. The leading candidate of WIMP dark matter is the lightest neutralino in the minimal supersymmetric standard model, which is a Majorana fermion. In addition, other models beyond the standard model at the TeV scale predict the existence of stable vector particles, such as the Kaluza-Klein (KK) photon \cite{Servant:2002aq,Cheng:2002ej} in the universal extra dimensions (UED) \cite{Appelquist:2000nn,Cheng:2002iz} and the $T$-odd heavy photon \cite {Hubisz:2004ft,Birkedal:2006fz,asano} in the Littlest Higgs model with $T$-parity \cite{ArkaniHamed:2001nc,Cheng:2004yc}. Now, the WIMP dark matter scenario has been tested in collider and direct detection experiments. In the Large Hadron Collider (LHC), (pair) production of the WIMP dark matter with TeV-scale mass is expected. On the other hand, XENON100 \cite{Aprile:2011hi}, which is the largest-volume detector ever, is in operation to detect it as scattering signals with nuclei on the earth. If the WIMP-scattering events are discovered and its properties are measured in the direct detection experiments, the measurements would be tested using the data that indicates dark matter production at the LHC. With such recent progress on the experimental side, the theoretical prediction of dark matter signals must have better accuracy. The collider signatures of dark matter (such as in supersymmetric or UED models) have been intensively studied. On the other hand, the scattering cross section of dark matter with nuclei in the direct detection experiments has also been calculated more accurately. In recent works \cite{Hisano:2010fy,Hisano:2010ct}, it was pointed out that a gluon-WIMP effective interaction is one of the leading contributors to the elastic scattering cross section, and it was also shown how to evaluate it precisely. (In those works, the results are shown for supersymmetric dark matter.) For the other models, in contrast, the gluon contribution and other leading terms are not taken into account correctly in the calculation of the scattering cross section (\cite{Cheng:2002ej, Servant:2002hb, Arrenberg:2008wy} for the UED model and \cite{Birkedal:2006fz} for the Littlest Higgs model with $T$-parity). In this paper, we assume that WIMP dark matter is composed of vector particles, and evaluate its elastic scattering cross section with nucleons, which is relevant to the direct detection experiments. Here, we consider the case in which vector dark matter interacts with standard-model quarks and exotic heavy fermions, which is a fundamental representation of $SU(3)_c$. As an application of our formulae, we discuss the direct detection of the first KK photon in the minimal universal extra dimension (MUED) model. We found that the spin-independent scattering cross section with a proton is increased by up to a factor of ten, compared with those in the previous works. We also show that the spin-independent scattering cross section of the KK photon with a proton ranges from about $3 \times 10^{-46} \text{cm}^2$ to $5 \times 10^{-48} \text{cm}^2$ on the parameter region with the thermal relic abundance consistent with the cosmological observations. Future direct detection experiments with ton-scale detectors might cover the range. This paper is organized as follows. In the next section, we briefly review the formulation for the scattering cross section of general vector dark matter with nuclei. Next, in \S\ref{sec:EffLag}, we derive the effective Lagrangian describing the interaction of general vector dark matter with quarks and gluons in target nuclei. In \S\ref{KK}, we review the MUED model and summarize the mass spectra of KK particles, and apply the formulae that are generally derived to the case of the first KK photon dark matter. Then, we numerically calculate the scattering cross section of the first KK photon dark matter with/without higher KK mode contributions, and discuss the significance of each effective operator. Section \ref{sec:conc} is devoted to the conclusion.

\label{sec:conc} In this paper, we assume that WIMP dark matter is composed of vector particles and have evaluated an elastic cross section with nucleons in the direct detection experiments. The vector dark matter is predicted in models beyond the standard model, such as the KK photon in the UED model and the $T$-odd heavy photon in the Littlest Higgs model with $T$-parity. On the other hand, however, the cross section had not yet been consistently evaluated, even at the leading order of $\alpha_s$. We have derived the general formulae for the cross section of general vector dark matter with nucleons. As an application of our formulae, we discussed the direct detection of the first KK photon in the MUED model. We found that the cross section is larger than those in the previous works by up to a factor of 10. We showed that the spin-independent cross section of the KK photon with a proton ranges from nearly $ 3 \times 10^{-46}~{\rm cm}^2$ to $5 \times 10^{-48}~{\rm cm}^2$ on the parameter region with the thermal relic abundance consistent with the cosmological observations. The future direct detection experiments with ton-scale detectors might cover the range. ~\\~\\ {\it Note Added:} While preparing the manuscript, we became aware of a paper by G.~Belanger, M.~Kakizaki and A.~Pukhov~\cite{Belanger:2010yx}. In that paper, the elastic cross section of the KK photon with nucleons in the MUED model was calculated using microOMEGA \cite{Belanger:2006is}. In the calculation, the tree level contribution was included in the scattering process whereas the one-loop contributions were incomplete.

1012.5455

In this paper, we complete formulae for the elastic scattering cross section of general vector dark matter with nucleons in the direct detection at the leading order of the strong coupling constant α_s, assuming that the dark matter is composed of vector particles and interacts with heavy fermions with color charge as well as standard-model quark. As an application of our formulae, the direct detection of the first Kaluza-Klein photon in the minimal universal extra dimension model is discussed. It is found that the scattering cross section is larger than those in the previous works by up to a factor of ten.

12202937

2011CQGra..28m4001M

1012

1012.2872_arXiv.txt

In the last five years, the field of gravitational wave (GW) astronomy has witnessed a remarkable convergence of experimental and theoretical achievement. The current generation of GW interferometers, particularly the Laser Interferometer Gravitational Wave Observatory (LIGO) and VIRGO detectors, has positioned this community on the cusp of discovery, by achieving their design sensitivities and making the first direct detection of GWs a real possibility. Meanwhile, the GW signature from the merger of a black-hole binary (BHB), expected to be the most common source for these and future observatories, went from being a complete unknown to a reasonably well-understood, surprisingly smooth transition between the earlier inspiral and the final ringdown, thanks to tremendous advancements in the field of numerical relativity (NR). In this article, we will limit our scope to the advancements that have occurred in just the past year, with only a minimal amount of background information, in order to provide a snapshot of the most current research in the field. For a review of progress in the preceding year, we refer the reader to \cite{Hinder:2010vn}, and for a broader review of NR achievements, please see \cite{Hannam:2009rd,Centrella:2010mx}. We consider it an indicator of the healthy progress of the field that, in addition to the general review articles typical of any field, in recent years the field of NR has progressed rapidly enough to warrant an annual review of the past year's progress. In Sec.~\ref{sec:background}, we give a brief review of the equations being solved by numerical relativists, and the methods that are most typically applied to solve them. In Sec.~\ref{sec:res}, we discuss the most recent achievements in NR methodology, and the simulations of numerically challenging systems of astrophysical interest that have been facilitated. In Sec.~\ref{sec:theory}, we briefly mention novel theoretical developments resulting from the interpretation of numerical results. We summarize and conclude in Sec.~\ref{sec:conc}.

\label{sec:conc} We have presented a brief overview of the current state of black-hole binary simulations in numerical relativity. While progress has slowed somewhat since the ``gold rush'' period following the breakthroughs in 2005--2006, there have still been a number of novel results of significant astrophysical and theoretical interest. In the past year, progress has been made in simulating longer waveforms overall, as well as systems with more extreme mass ratios and larger spins. Advances in methodology have significantly enhanced the accuracy and efficiency achievable by state-of-the-art codes. Progress has also been made in using the body of available numerical simulations to inform a greater theoretical understanding of the strong-field behavior of black-hole binaries. The inclusion of matter and electromagnetic fields in black-hole binary simulations has led to unexpected discoveries, which will have significant astrophysical implications in the years to come. The field of numerical relativity has now entered a period of more modest, but hopefully sustainable growth, with a substantial amount of discovery space that remains to be explored.

1012.2872

The advent of long-term stability in numerical relativity has yielded a windfall of answers to long-standing questions regarding the dynamics of space-time, matter, and electromagnetic fields in the strong-field regime of black-hole binary mergers. In this review, we will briefly summarize the methodology currently applied to these problems, emphasizing the most recent advancements. We will discuss recent results of astrophysical relevance, and present some novel interpretation. Although we primarily present a review, we also present a simple analytical model for the time-dependent Poynting flux from two orbiting black holes immersed in a magnetic field, which compares favorably with recent numerical results. Finally, we will discuss recent advancements in our theoretical understanding of merger dynamics and gravitational waveforms that have resulted from interpreting the ever-growing body of numerical relativity results.

12133663

2010arXiv1012.5380D

1012

1012.5380_arXiv.txt

Many important features of the observable universe can be understood as the result of out of equilibrium processes during the early stages of its history, when it was filled with a hot primordial plasma. This includes the decoupling of the cosmic microwave background, the creation of light elements, dark matter production, baryogenesis, and, in inflationary cosmology, the production of particles altogether during reheating. In many cases the relevant energy scales by far exceed those that can be realised in any human made experiments and provide an excellent laboratory to test particle physics models beyond the Standard Model (SM). Nonequilibrium dynamics also play a crucial role in the understanding of signals from heavy ion colliders. Thus, a quantitative understanding of nonequilibrium processes is crucial for cosmology as well as particle physics. Boltzmann equations and their matrix valued generalisations can accurately describe many nonequilibrium phenomena. They are known to suffer from uncertainties when the coherence lengths are large, but are usually believed to be accurate in absence of such effects. On the other hand, it is known that in gauge theories at high temperature resummations are necessary because processes involving many quanta, which naively are of higher order in the coupling, contribute to the relaxation rate at leading order. How does this accord with the use of single particle distribution functions in the kinetic equations? Medium effects are often included in effective kinetic equations by the use of thermal masses for quasiparticles. What are the limits of the validity of this procedure? To which degree can the quasiparticles be treated as particles? Do the thermal masses act as kinetic masses? All these problems are related to the assumption that the system should be described in terms of the properties of individual particles as asymptotic states. The definition of these is, however, ambiguous in a dense plasma. We show that the above questions can be answered in a consistent and intuitive way when swapping the single particle phase space distribution functions as dynamical variables for correlation functions of quantum fields. This formalism in addition allows for a full quantum treatment of coherent oscillations and quantum memory effects. In this work, we study the relaxation of scalar and fermionic quantum fields in a large thermal bath. Depending on the initial state of the out-of-equilibrium fields, they either gain energy from or dissipate it into the bath. This resembles a large number of interesting phenomena including thermal production of particles from a plasma, propagation in dense matter and cosmological freezeout processes. It is, apart from reporting a number of new results summarised below, one of the objectives of this article to explain and promote the formalism employed here to a wider audience. It provides powerful tools to treat nonequilibrium quantum systems in terms of quantities that have a clear physical interpretation, without semiclassical (on-shell) approximations or a gradient expansion. We aim to make the connection between nonequilibrium quantum field theory and effective kinetic equations transparent also to readers without background knowledge in the former. Therefore we try to avoid technicalities and restrain to language commonly used in particle physics. \newpage \begin{center} \textbf{Outline of this Article} \end{center} In section \ref{QuantumFieldsAndParticles} we introduce our notation and review the formalism we use, including the physical interpretation of the various quantities. Section \ref{scalarsectioN} is devoted to the relaxation of scalar fields in a thermal bath. In section \ref{trilinearsection} we study in detail the kinematics of the relaxation via a trilinear interaction with a bath of other scalars. In the first part we recall the interpretation of the known result in the quasiparticle approximation. In the second part we derive a formula that includes corrections beyond this approximation, using resummed perturbation theory. These appear as off-shell contributions and give rise to apparent violation of energy conservation in the quasiparticle picture. In section \ref{ScalarwithYukawa} we consider the case that the scalar is coupled to a bath of fermions with gauge interactions via a Yukawa coupling. In section \ref{quarticsection} we discuss the kinematic differences between relaxation via 3-vertices and 4-vertices, using the simple example of a quartic interaction. In section \ref{comparisonsection} we compare the contributions to the relaxation rate from decays and scatterings of real quasiparticles to off-shell contributions. Section \ref{sectionaboutfermion} is devoted to the relaxation of a fermion with Yukawa coupling. In \ref{fermioncorrelators} we give exact expressions for the nonequilibrium two-point functions of a fermion in a thermal bath. In \ref{fermionviayukawa} we compute an analytic expression for the rate of relaxation via a Yukawa coupling in the quasiparticle regime. Details of the calculation are given in appendix \ref{YukawaSelfEnergy}. In section \ref{comparisontoBE} we compare the time evolution of the energy density in the Boltzmann and field theoretical approaches. In section \ref{InflatonSection} we apply our results to reheating of the universe via perturbative inflaton decay. We discuss the possibility of an upper bound on the temperature due to closure of the phase space for decays by large thermal masses. Further implications and possible extensions are discussed in section \ref{discussion}. In appendix \ref{KBEessentials} we review some basic ingredients of nonequilibrium field theory used in the analysis, and in appendix \ref{YukawaSelfEnergy} we derive an analytic expression for the relaxation rate of a fermion with Yukawa coupling in a thermal plasma that has, to the best of our knowledge, not been reported in the literature.\\ The main new results reported in this article are the discussion of the nonequilibrium propagators in sections \ref{relaxratesection} and \ref{comparisontoBE}, the explicit computation of higher order contributions in section \ref{OffShellSection} and the related discussion in section \ref{analsection}, the numerical comparison between leading order and resummed results in section \ref{comparisonsection}, the expression for the nonequilibrium propagator for Dirac fermions in section \ref{sectionaboutfermion}, the analytic formula for the fermion relaxation rate in section \ref{fermionviayukawa}, the comparison of the energy density in the Boltzmann and quantum field theoretical approaches in section \ref{comparisontoBE} and the application to cosmic reheating in section \ref{InflatonSection}. We embed these into detailed discussion that aims to draw an intuitive physical picture that is coherent without previous knowledge in nonequilibrium quantum field theory.

\label{discussion} We studied the relaxation of scalar and fermionic fields in a large thermal bath from first principles of nonequilibrium quantum field theory. This situation is realised in various interesting physical systems, including thermal production of particles from a plasma, propagation in dense matter or cosmological freezeout processes. Depending on the initial state of the out-of-equilibrium fields, relaxation can mean dissipation of energy into the bath or thermal production of quanta. The formalism we used allows a quantum description of both processes. We performed all computations in terms of correlation functions for quantum fields, avoiding the ambiguities in Boltzmann type kinetic equations related to quantum coherence, the definition of asymptotic states or particle numbers in a dense plasma and the assumption of molecular chaos. Basis of our analysis are exact expressions for nonequilibrium correlation functions for arbitrarily large deviations from thermal equilibrium that are valid for all times. They were obtained without semiclassical approximation or a gradient expansion. The main goals of our study were to explore which effect the modified dispersion relations in the plasma have on the relaxation rate and under which circumstances a description in terms of quasiparticles is suitable for a relativistic quantum system. \subsection{Summary of Results}\label{SummmaryOfResults} At high temperature and density, relaxation in a plasma is not only driven by decays and scatterings of individual particles. Processes involving many quanta, that naively are of higher order in the coupling, have to be taken account. Their contribution is negligible in a dilute plasma, but grows with the occupation numbers. Therefore consistent results can only be obtained using resummed perturbation theory. In our setup the contribution from processes involving many quanta can be parametrised by the widths of the resonances in the plasma. \\ If the widths of all resonances are small, these can be interpreted as quasiparticles with dispersion relations given by the poles of the resummed propagators. Then the exchange of energy between the modes of different fields can effectively be described in terms of decays and scatterings amongst individual quasiparticles. This is an essential condition for the validity of effective kinetic equations. However, even in the quasiparticle regime it is generally not valid to obtain such equations by simply replacing the vacuum masses by thermal masses because the dispersion relations can be more complicated. In addition, collective excitations may appear as propagating states. Though effective kinetic equations for quasiparticles describe the energy transfer between different degrees of freedom with good accuracy, they do not capture all properties of the system. We explicitly showed that the energy density of the out-of-equilibrium field is not that of a quasiparticle gas and does not follow a Boltzmann equation for quasiparticles. The modified dispersion relations in a plasma lead to a number of interesting effects. If the out-of-equilibrium field couples via 3-vertices, the leading order contribution to the energy exchange with the bath can effectively be switched off in some temperature regime. This happens when, due to the temperature dependent dispersion relations, none of the quasiparticles connected by the vertex can kinematically decay into the other two. This can considerably delay the relaxation and have interesting cosmological implications, for instance during reheating or for the thermal production of particles. Interactions by 4-vertices always allow for relaxation at leading order as these mediate $2\rightarrow 2$ scatterings. However, since they require real quanta from the plasma in the initial state they give only small contributions at low temperature.\\ The widths of the quasiparticles parametrise the effect of processes that involve more quanta and vertices. These appear to be energy violating in the quasiparticle picture, but indeed only reflect the fact that relaxation can also happen via off-shell processes and multiple scatterings with many quanta from the bath. At high temperature there are various effects including induced transitions for bosons, the large density of scattering partners or the Landau-Pomeranchuk-Migdal effect which can make these contributions comparable to or even larger than the rate at leading order. Mathematically this manifests in the poor convergence of the perturbative series and the need for resummation. In the quasiparticle picture, this appears as an increase of the width or \textit{melting} of the resonance peaks at high temperature. Physically it means that in a dense plasma the picture of individual particles travelling freely between isolated interactions does not hold. When the widths of the resonances are not small, the quasiparticle interpretation and description in terms of kinetic equations break down. For instance, even in a temperature region where all leading order processes involving real quasiparticles are kinematically forbidden, the relaxation rate may be as big as in the region where decays and inverse decays are allowed. Thus, naive estimates based on kinematics involving thermal masses fail to even qualitatively reproduce the behaviour. In the quasiparticle picture this can be interpreted in terms of off-shell processes, mediated by virtual quasiparticles. At low temperatures, these effects can be included into the Boltzmann approach by adding matrix elements for higher order processes. At large temperatures, this would require the inclusion of infinitely many diagrams with large numbers of particles in the initial and final state.\\ We computed analytic expressions for the relaxation rates of scalars and fermions via a trilinear and a Yukawa interaction, respectively, for arbitrary energies and momenta in the quasiparticle regime. We furthermore derived a formula for the rate of relaxation via the trilinear scalar interaction, expressed as a one-dimensional integral, which includes the effects of higher order processes. Comparing to the quasiparticle result, we found that the corrections can be large even for a moderate width $\Gamma_{i}/M_{i}\sim 10^{-2}$. We applied our results to a particular cosmological problem, namely the question whether large thermal masses of the decay products can close the phase space for perturbative inflaton decay and therefore impose an upper bound on the reheating temperature.\\ Finally, we would like to emphasise that in reality the details of the relaxation process can be far more complicated than in our model. We assumed instantaneous equilibration of the bath degrees of freedom. There may be different time scales related to this process. The word ``quasiparticles'' has been used for propagating states with definite momentum, the modes of a homogeneous field, with properties defined by averaging over a statistical ensemble. They are not localised particles (wave packets). When the relaxation is pictured as a statistical process, driven by the competition of decays and inverse decays (and scatterings) of quasiparticles, each individual of these processes is localised in time and space. Even though the relaxation time of $\phi$ is much larger than that of the other fields, the details of the thermalisation process and local deviations from equilibrium do affect the decay of individual quasiparticles. However, we are not interested in the fate of individual quasiparticles here, but in the overall relaxation of the system. Therefore we believe that the thermal ensemble average is justified and gives correct results as long as the relaxation time in the plasma is much shorter than $1/\Gamma$ \subsection{Conclusions and Outlook} Boltzmann equations and their matrix valued generalisations can accurately describe many nonequilibrium phenomena in a medium. They are known to suffer from uncertainties when the coherence lengths are large, but are usually assumed to be accurate in absence of such effects. On the other hand, it has been known for long that resummation is necessary in gauge theories at high temperature because processes involving many quanta that naively are of higher order in the coupling in fact contribute at leading order. It is important to understand in which cases and how these many-particle amplitudes can be incorporated into the framework of single particle distribution functions. The inconsistencies are related to the assumption that the system should be described in terms of individual particles as asymptotic states. We showed that they can be removed in an intuitive way when swapping the phase space distribution functions as dynamical variables for correlation functions of quantum fields. The Schwinger-Keldysh formalism used in this work allows for a full quantum treatment of nonequilibrium phenomena in the framework of quantum field theory. A consistent perturbation theory can be formulated in the framework of the $n$PI effective action. It has previously been applied to problems in which all degrees of freedom are far from equilibrium, including preheating after inflation and equilibration after heavy ion collisions. However, the equations of motion, coupled second order integro-differential equations, are considerably more complicated than the semiclassical Boltzmann equations. In most cases they can only be solved numerically in toy models. The use of exact equations comes at a price, the loss of the simplicity, transparency and intuitive physical interpretation that make the Boltzmann equations appealing. In this work we studied systems in which only few degrees of freedom are out of equilibrium and weakly coupled to a large thermal bath. In this setup most of the computations can be performed analytically and the quantities in the field theoretical approach have an intuitive physical interpretation. The analytic solutions allowed us to study in detail the effect of quantum and higher order corrections, parametrise them in quasiparticle properties and understand the limits of this approximation. The semiclassical description in terms of effective kinetic equations can be recovered in the limit of a dilute gas. The main simplification that allowed a widely analytical treatment lies in the negligible backreaction. We have neglected backreaction in two ways: by ignoring the effect of the out-of-equilibrium fields on the bath temperature and by dropping contributions to the self energies that have nonequilibrium propagators in the loops. This allows to use resummation techniques known from the theory in equilibrium. There are interesting systems in which the second assumption is not justified while the first one still holds. In leptogenesis, for instance, the baryon asymmetry of the universe originates from $CP$-violating interactions of heavy right handed neutrinos that are out of equilibrium with a thermal bath of SM fields. Dropping the small Cabibbo-Kobayashi-Maskawa-phase in the SM, their Yukawa couplings are the only source of $CP$-violation. When computing the $CP$-asymmetry, the suppression of contributions from diagrams that have the out-of-equilibrium fields in the loop by the large number of degrees of freedom in the bath does not apply as none of the other interactions is $CP$-violating. Indeed, their inclusion is crucial as the generation of an asymmetry comes from a quantum interference between the two leading orders in the Yukawa couplings. Therefore it is important to develop techniques that allow a consistent evaluation of higher order diagrams in which the nonequilibrium fields appear in the loop while making maximal use of the simplifications due to the weak coupling to a thermal bath and negligible backreaction on the temperature. The nonequilibrium propagators used in this work can be used for a perturbative treatment and the formulation of Feynman rules is straightforward. However, practical evaluation of higher order diagrams is technically difficult due to the dependence on centre of mass time, and it remains to be seen how to perform the resummation of all stronger (gauge) interactions in the bath.

1012.5380

Boltzmann equations and their matrix valued generalisations are commonly used to describe nonequilibrium phenomena in cosmology. On the other hand, it is known that in gauge theories at high temperature processes involving many quanta, which naively are of higher order in the coupling, contribute to the relaxation rate at leading order. How does this accord with the use of single particle distribution functions in the kinetic equations? When can these effects be parametrised in an effective description in terms of quasiparticles? And what is the kinematic role of their thermal masses? We address these questions in the framework of nonequilibrium quantum field theory and develop an intuitive picture in which contributions from higher order processes are parametrised by the widths of resonances in the plasma. In the narrow width limit we recover the quasiparticle picture, with the additional processes giving rise to off-shell parts of quasiparticle propagators that appear to violate energy conservation. In this regime we give analytic expressions for the scalar and fermion nonequilibrium propagators in a medium. We compare the efficiency of decays and scatterings involving real quasiparticles, computed from analytic expressions for the relaxation rates via trilinear scalar and Yukawa interactions for all modes, to off-shell contributions and find that the latter can be significant even for moderate widths. Our results apply to various processes including thermal production of particles from a plasma, dissipation of fields in a medium and particle propagation in dense matter. We discuss cosmological implications, in particular for the maximal temperature achieved during reheating by perturbative inflaton decay.

12134250

2010ASPC..438..229R

1012

1012.2934_arXiv.txt

Despite significant advances in the understanding of the Galaxy's magnetic field, the large scale topology is still not well understood. Part of the difficulty in determining the structure of the Galactic magnetic field is due to the vast distances involved. With current propulsion methods it would take about 80,000 years for a small probe to reach the nearest star \citep{spacecraft}. Therefore, direct measurements of the Galaxy's magnetic field with magnetometers would take an unacceptably long time. Furthermore, magnetic fields themselves do not radiate. Consequently, indirect observational methods must be used. Most of the information we have about the Galactic magnetic field has been derived from radio observations of polarised sources as follows. Since a significant fraction of the interstellar medium consists of a plasma with a frozen-in magnetic field \citep{katia}, it acts as a birefringent material to radio waves that pass through it. As a result, the polarisation angle of a wave will rotate as it propagates through the medium. This effect, known as Faraday rotation, is quantified using the aptly named Rotation Measure (RM), defined as \begin{equation} \label{rmequation} {\rm{RM}} = 0.812 \int n_e \; {\bf{B}} \cdot {\rm{d{\bf{l}}}}, \end{equation} where $n_e$ is the electron density, ${\bf{B}}$ is the magnetic field, and d{\bf{l}} is the pathlength element, which is always defined to be from the source to the receiver. The amount of rotation a radio wave typically experiences is linearly dependent on the square of the wavelength ($\lambda$), with the slope of the dependence being the RM such that \begin{equation} \label{pangle} \tau = \tau_\circ + \lambda^2 {\rm{RM}} \end{equation} where $\tau$ is the observed polarisation angle, and $\tau_\circ$ is the emitted polarisation angle. Assuming all emissions from a given source are at the same $\tau_\circ$, measuring a source of polarised radiation at multiple wavelengths will allow us to easily determine the RM for a given source. If we know something about the distribution of free electrons along a given line of sight, we can then work backwards to determine what the magnetic field might be to create the RMs we observe. Using RMs from sources both inside (puslars) and outside (extragalactic sources or EGS) the Galaxy allows us to further constrain the possible topology of the field. The more sources we have, the more accurate the inferred magnetic field is likely to be.

We have determined reliable rotation measures for 1316 extragalactic sources, calculated from the four-band polarisation observations carried out by the Synthesis telescope at DRAO as part of the Canadian Galactic Plane Survey (CGPS). These data represent more than a factor of 20 increase in the number of sources in the same region published prior to the introduction of the CGPS. Using these data, we have shown that the magnetic field in the outer disk of the Galaxy has a very small (almost zero) pitch angle, strongly suggesting that the field is not aligned with the spiral arms in this region. In addition, we have estimated the rotation measure scale height in the longitude region of the CGPS latitude extension to be on the order of 1.2 kpc. This value, however, requires more investigation, as the dominant contribution to the RMs (namely the Perseus arm) has not been considered separately in this calculation. It is expected that new observations south of the Galactic disk, which are currently underway, will contribute considerably to the understanding of the disk-halo transition for the magnetic field in this region.

1012.2934

The Galactic magnetic field is important in the dynamics of our Galaxy. It is believed to play a role in star formation and influence the structure of the Galaxy. In order to understand how the Galactic magnetic field originally formed or how it is evolving, we must first determine its present topology. To this end, we have used observations from the Canadian Galactic Plane Survey (CGPS) to calculate the highest source density of rotation measures (RM) to date in the disk of the Galaxy. Using these data, we estimate the Galactic longitude of the RM null point in the outer Galaxy (where the RMs of extragalactic sources are observed to pass through zero, on average, with increasing Galactic longitude). We have also examined the RM scale height using the CGPS latitude extension. The values of these parameters offer critical constraints for modeling the large-scale magnetic field in the Galactic disk.

12144586

2010LRSP....7....6P

1012

1012.5513_arXiv.txt

\label{section:introduction} Solar cycle prediction is an extremely extensive topic, covering a very wide variety of proposed prediction methods and prediction attempts on many different timescales, ranging from short term (month--year) forecasts of the runoff of the ongoing solar cycle to predictions of long term changes in solar activity on centennial or even millennial scales. As early as 1963, Vitinsky published a whole monograph on the subject, later updated and extended \citep{Vitinsky:book1, Vitinsky:book2}. More recent overviews of the field or aspects of it include \cite{Hathaway:prediction.review}, \cite{Kane:pred23rev}, and \cite{Pesnell}. In order to narrow down the scope of the present review, we constrain our field of interest in two important respects. Firstly, instead of attempting to give a general review of all prediction methods suggested or citing all the papers with forecasts, here we will focus on those aspects of the solar cycle prediction problem that have a bearing on dynamo theory. We will thus discuss in more detail empirical methods that, independently of their success rate, have the potential of shedding some light on the physical mechanism underlying the solar cycle, as well as the prediction attempts based on solar dynamo models. Secondly, we will here only be concerned with the issue of \textit{predicting the amplitude} (and optionally the epoch) {of an upcoming solar maximum no later than right after the start of the given cycle.} This emphasis is also motivated by the present surge of interest in precisely this topic, prompted by the unusually long and deep recent solar minimum and by sharply conflicting forecasts for the maximum of the incipient solar cycle~24. As we will see, significant doubts arise both from the theoretical and observational side as to what extent such a prediction is possible at all (especially before the time of the minimum has become known). Nevertheless, no matter how shaky their theoretical and empirical backgrounds may be, forecasts \textit{must} be attempted. Making verifiable or falsifiable predictions is obviously the core of the scientific method in general; but there is also a more imperative urge in the case of solar cycle prediction. Being the prime determinant of space weather, solar activity clearly has enormous technical, scientific, and financial impact on activities ranging from space exploration to civil aviation and everyday communication. Political and economic decision makers expect the solar community to provide them with forecasts on which feasibility and profitability calculations can be based. Acknowledging this need, the Space Weather Prediction Center of the US National Weather Service does present annually or semiannually updated ``official'' predictions of the upcoming sunspot maximum, emitted by a Solar Cycle Prediction Panel of experts, starting shortly before the (expected) minimum \citep{SWPC2009}. The unusual lack of consensus during the early meetings of this panel during the recent minimum, as well as the concurrent more frequently updated but wildly varying predictions of a NASA MSFC team \citep{MSFC2009} have put on display the deficiencies of currently applied prediction techniques; on the other hand, they also imply that cycle 24 may provide us with crucial new insight into the physical mechanisms underlying cyclic solar activity. While a number of indicators of solar activity exist, by far the most commonly employed is still the smoothed relative sunspot number $R$; the ``Holy Grail'' of sunspot cycle prediction attempts is to get $R$ right for the next maximum. We therefore start by briefly introducing the sunspot number and inspecting its known record. Then, in Sections~\ref{sect:precursor}, \ref{sec:Extrapolation-Methods}, and \ref{sec:Model-Based-Predictions} we discuss the most widely employed methods of cycle predictions. Section~\ref{sec:Summary-Evaluation} presents a summary evaluation of the past performance of different forecasting methods and collects some forecasts for cycle~24 derived by various approaches. Finally, Section~\ref{sec:Epilogue} concludes the paper. \subsection{The sunspot number} \label{sect:wolfno} Despite its somewhat arbitrary construction, the series of relative sunspot numbers constitutes the longest homogeneous global indicator of solar activity determined by direct solar observations and carefully controlled methods. For this reason their use is still predominant in studies of solar activity variation. As defined originally by \cite{Wolf:relno.def}, the relative sunspot number is \begin{equation} R_W=k(10g+f) \end{equation} where $g$ is the number of sunspot groups (including solitary spots), $f$ is the total number of all spots visible on the solar disc, while $k$ is a correction factor depending on a variety of circumstances, such as instrument parameters, observatory location and details of the counting method. Wolf, who decided to count each spot only once and not to count the smallest spots, the visibility of which depended on seeing, used $k=1$. The counting system employed was changed by Wolf's successors to count even the smallest spots, attributing a higher weight (i.e.~$f>1$) to spots with a penumbra, depending on their size and umbral structure. As the new counting naturally resulted in higher values, the correction factor was set to $k=0.6$ for subsequent determinations of $R_W$ to ensure continuity with Wolf's work, even though there was no change in either the instrument or the observing site. This was followed by several further changes in the details of the counting method (\citealp{Waldmeier:book}; see \citealp{Kopecky+:relativeno}, \citealp{Hoyt+Schatten:GSN} and \citealp{Hathaway:LRSP} for further discussions on the determination of $R_W$). In addition to introducing the relative sunspot number, \cite{Wolf:cycle.length.activity} also used earlier observational records available to him to reconstruct its monthly mean values since 1749. In this way, he reconstructed 11-year sunspot cycles back to that date, introducing their still universally used numbering. (In a later work he also determined annual mean values for each calendar year going back to 1700.) In 1981, the observatory responsible for the official determination of the sunspot number changed from Z\"urich to the Royal Observatory of Belgium in Brussels. The website of the SIDC (originally Sunspot Index Data Center, recently renamed Solar Influences Data Analysis Center), \url{http://sidc.oma.be}, is now the most authoritative source of archive sunspot number data. But it has to be kept in mind that the sunspot number is also regularly determined by other institutions: these variants are informally known as the American sunspot number (collected by AAVSO and available from the National Geophysical Data Center, \url{http://www.ngdc.noaa.gov/ngdc.html}) and the Kislovodsk Sunspot Number (available from the web page of the Pulkovo Observatory, \url{http://www.gao.spb.ru}). Cycle amplitudes determined by these other centers may differ by up to 6--7\% from the SIDC values, NOAA numbers being consistently lower, while Kislovodsk numbers show no such systematic trend. These significant disagreements between determinations of $R_W$ by various observatories and observers are even more pronounced in the case of historical data, especially prior to the mid-19th century. In particular, the controversial suggestion that a whole solar cycle may have been missed in the official sunspot number series at the end of the 18th century is taken by some as glaring evidence for the unreliability of early observations. Note, however, that independently of whether the claim for a missing cycle is well founded or not, there is clear evidence that this controversy is mostly due to the very atypical behaviour of the Sun itself in the given period of time, rather than to the low quality and coverage of contemporary observations. These issues will be discussed further in section \ref{sect:evenodd}. Given that $R_W$ is subject to large fluctuations on a time scale of days to months, it has become customary to use annual mean values for the study of longer term activity changes. To get rid of the arbitrariness of calendar years, the standard practice is to use 13-month boxcar averages of the monthly averaged sunspot numbers, wherein the first and last months are given half the weight of other months: \begin{equation} R=\frac1{24}\left(R_{\mathrm{m},-6}+2\sum_{i=-5}^{i=5} R_{\mathrm{m},i} +R_{\mathrm{m},6}\right) \label{eq:Rdef} \end{equation} $R_{\mathrm{m},i}$ being the mean monthly value of $R_W$ for $i$th calendar month counted from the present month. It is this running mean $R$ that we will simply call ``the sunspot number'' throughout this review and what forms the basis of most discussions of solar cycle variations and their predictions. In what follows, $\Rmax^{(n)}$ and $\Rmin^{(n)}$ will refer to the maximum and minimum value of $R$ in cycle $n$ (the minimum being the one that starts the cycle). Similarly, $\tmax^{(n)}$ and $\tmin^{(n)}$ will denote the epochs when $R$ takes these extrema. \subsubsection{Alternating series and nonlinear transforms} Instead of the ``raw'' sunspot number series $R(t)$ many researchers prefer to base their studies on some transformed index $R'$. The motivation behind this is twofold. (a) The strongly peaked and asymmetrical sunspot cycle profiles strongly deviate from a sinusoidal profile; also the statistical distribution of sunspot numbers is strongly at odds with a Gaussian distribution. This can constitute a problem as many common methods of data analysis rely on the assumption of an approximately normal distribution of errors or nearly sinusoidal profiles of spectral components. So transformations of $R$ (and, optionally, $t$) that reduce these deviations can obviously be helpful during the analysis. In this vein, e.g., Max Waldmeier often based his studies of the solar cycle on the use of logarithmic sunspot numbers $R'=\log R$; many other researchers use $R'=R^\alpha$ with $0.5\leq\alpha<1$, the most common value being $\alpha=0.5$. (b) As the sunspot number is a rather arbitrary construct, there may be an underlying more physical parameter related to it in some nonlinear fashion, such as the toroidal magnetic field strength $B$, or the magnetic energy, proportional to $B^2$. It should be emphasized that, contrary to some claims, our current understanding of the solar dynamo does \textit{not} make it possible to guess what the underlying parameter is, with any reasonable degree of certainty. In particular, the often used assumption that it is the magnetic energy, lacks any sound foundation. If anything, on the basis of our current best understanding of flux emergence we might expect that the amount of toroidal flux emerging from the tachocline should be $\int |B-B_0|\,dA$ where $B_0$ is some minimal threshold field strength for Parker instability and the surface integral goes across a latitudinal cross section of the tachocline \citep[cf.][]{Ruzmaikin:biasing}. As, however, the lifetime of any given sunspot group is finite and proportional to its size \citep{Petrovay+vDG:decay1, Henwood+}, instantaneous values of $R$ or the total sunspot area should also depend on details of the probability distribution function of $B$ in the tachocline. This just serves to illustrate the difficulty of identifying a single physical governing parameter behind $R$. One transformation that may still be well motivated from the physical point of view is to attribute an alternating sign to even and odd Schwabe cycles: this results in the the \textit{alternating sunspot number series} $R_\pm$. The idea is based on Hale's well known polarity rules, implying that the period of the solar cycle is actually 22 years rather than 11 years, the polarity of magnetic fields changing sign from one 11-year Schwabe cycle to the next. In this representation, first suggested by \cite{Bracewell}, usually odd cycles are attributed a negative sign. This leads to slight jumps at the minima of the Schwabe cycle, as a consequence of the fact that for a 1--2 year period around the minimum, spots belonging to both cycles are present, so the value of $R$ never reaches zero; in certain applications, further twists are introduced into the transformation to avoid this phenomenon. After first introducing the alternating series, in a later work \cite{Bracewell:trafo} demonstrated that introducing an underlying ``physical'' variable $R_B$ such that \begin{equation} R_\pm = 100\left( R_B/83\right)^{3/2} \label{eq:Bracewell} \end{equation} (i.e.~$\alpha=2/3$ in the power law mentioned in in item (a) above) significantly simplifies the cycle profile. Indeed, upon introducing a ``rectified'' phase variable\epubtkFootnote{The more precise condition defining $\phi$ is that $\phi=\pm\pi/2$ at each maximum and $\phi$ is quadratically related to the time since the last minimum.} $\phi$ in each cycle to compensate for the asymmetry of the cycle profile, $R_B$ is a nearly sinusoidal function of $\phi$. The empirically found 3/2 law is interpreted as the relation between the time-integrated area of a typical sunspot group vs. its peak area (or peak $R_W$ value), i.e.~ the steeper than linear growth of $R$ with the underlying physical parameter $R_B$ would be due to to the larger sunspot groups being observed longer, and therefore giving a disproportionately larger contribution to the annual mean sunspot numbers. If this interpretation is correct, as suggested by Bracewell's analysis, then $R_B$ should be considered proportional to the total toroidal magnetic flux emerging into the photosphere in a given interval. (But the possibility must be kept in mind that that the same toroidal flux bundle may emerge repeatedly or at different heliographic longitudes, giving rise to several active regions.) \subsection{Other indicators of solar activity} Reconstructions of $R$ prior to the early 19th century are increasingly uncertain. In order to tackle problems related to sporadic and often unreliable observations, \cite{Hoyt+Schatten:GSN} introduced the \textit{Group Sunspot Number} (GSN) as an alternative indicator of solar activity. In contrast to $R_W$, the GSN only relies on counts of sunspot groups as a more robust indicator, disregarding the number of spots in each group. Furthermore, while $R_W$ is determined for any given day from a single observer's measurements (a hierarchy of secondary observers is defined for the case if data from the primary observer were unavailable), the GSN uses a weighted average of all observations available for a given day. The GSN series has been reproduced for the whole period 1611--1998 (Figure~\ref{fig:GSNrecord}) and it is generally agreed that for the period 1611--1818 it is a more reliable reconstruction of solar activity than the relative sunspot number. Yet there have been relatively few attempts to date to use this data series for solar cycle prediction. One factor in this could be the lack of regular updates of the GSN series, i.e. the unavailability of precise GSN values for the past decade. \epubtkImage{}{% \begin{figure}[htbp] \centerline{\includegraphics[width=\textwidth]{GSNrecord}} \caption{13-month sliding averages of the monthly average relative sunspot numbers $R$ (green) and group sunspot numbers $R_G$ (black) for the period 1611--1998.} \label{fig:GSNrecord} \end{figure}} Instead of the sunspot number, the total area of all spots observed on the solar disk might seem to be a less arbitrary measure of solar activity. However, these data have been available since 1874 only, covering a much shorter period of time than the sunspot number data. In addition, the determination of sunspot areas, especially farther from disk center, is not as trivial as it may seem, resulting in significant random and systematic errors in the area determinations. Area measurements performed in two different observatories often show discrepancies reaching $\sim 30\%$ for smaller spots \citep[cf.\ the figure and discussion in Appendix~A of][]{Petrovay+:decay2}. A number of other direct indicators of solar activity have become available from the 20th century. These include e.g., various plage indices or the 10.7~cm solar radio flux -- the latter is considered a particularly good and simple to measure indicator of global activity (cf.\ Figure~\ref{fig:radioflux}). As, however, these data sets only cover a few solar cycles, their impact on solar cycle prediction has been minimal. \epubtkImage{}{% \begin{figure}[htbp] \centerline{\includegraphics[width=\textwidth]{F10_6}} \caption{Monthly values of the 10.7~cm radio flux in solar flux units for the period 1947--2009. The solar flux unit is defined as $10^{-22}\,$W$/$m$^2\,$Hz. The green curve shows $R_{\mathrm{m}}+60$ where $R_\mathrm{m}$ is the monthly mean relative sunspot number. (The vertical shift is for better comparison.) Data are from the NRC Canada (Ottawa/Penticton).} \label{fig:radioflux} \end{figure}} Of more importance are \textit{proxy indicators} such as geomagnetic indices (the most widely used of which is the $aa$ index), the occurrence frequency of aurorae or the abundances of cosmogenic radionuclides such as $^{14}$C and $^{10}$Be. For solar cycle prediction uses such data sets need to have a sufficiently high temporal resolution to reflect individual 11-year cycles. For the geomagnetic indices such data have been available since 1868, while an annual $^{10}$Be series covering 600 years has been published very recently by \cite{Berggren+:10Be_600yrs}. Attempts have been made to reconstruct the epochs and even amplitudes of solar maxima during the past two millennia from oriental naked eye sunspot records and from auroral observations \citep{secular.book, Nagovitsyn:reconstr}, but these reconstructions are currently subject to too many uncertainties to serve as a basis for predictions. Isotopic data with lower temporal resolution are now available for up to 50,000 years in the past; while such data do not show individual Schwabe cycles, they are still useful for the study of long term variations in cycle amplitude. Inferring solar activity parameters from such proxy data is generally not straightforward. \epubtkImage{}{% \begin{figure}[htbp] \centerline{\includegraphics[width=0.9\textwidth]{spotcycle}} \caption{The variation of the monthly smoothed relative sunspot number $R$ during the period 1749--2009, with the conventional numbering of solar cycles.} \label{fig:SSNrecord} \end{figure}} \subsection{The solar cycle and its variation} \label{sect:cycle} The series of $R$ values determined as described in Section~\ref{sect:wolfno} is plotted in figure~\ref{fig:SSNrecord}. It is evident that the sunspot cycle is rather irregular. The mean length of a cycle (defined as lasting from minimum to minimum) is 11.02~years (median 10.7~years), with a standard deviation of 1.2~years. The mean amplitude is 113 (median 115), with a standard deviation of 40. It is this wide variation that makes the prediction of the next cycle maximum such an interesting and vexing issue. It should be noted that the period covered by the relative sunspot number record includes an extended interval of atypically strong activity, the so called Modern Maximum (see below), cycles 17--23. Excluding these cycles from the averaging, the mean and median values of the cycle amplitude are very close to 100, with a standard deviation of 35. The mean and median cycle length then become 11.1 and 11.2~years, respectively, with a standard deviation of 1.3~years. \subsubsection{Secular activity variations} \label{sect:secular} Inspecting Figure~\ref{fig:SSNrecord} one can discern an obvious long term variation. For the study of such long term variations, the series of cycle parameters is often smoothed on time scales significantly longer than a solar cycle: this procedure is known as \textit{secular smoothing}. One popular method is the so-called \textit{Gleissberg filter} or \textit{12221 filter} \citep{Gleissberg:12221}. For instance, the Gleissberg filtered amplitude of cycle $n$ is given by \begin{equation} \langle \Rmax\rangle_{\mathrm{G}}^{(n)} = \frac 18\left( \Rmax^{(n-2)} +2\Rmax^{(n-1)} +2\Rmax^{(n)} +2\Rmax^{(n+1)} +\Rmax^{(n+2)}\right) . \end{equation} \epubtkImage{}{% \begin{figure}[htbp] \centerline{\includegraphics[width=\textwidth]{Gleissbgfilt}} \caption{Amplitudes of the sunspot cycles (dotted) and their Gleissberg filtered values (blue solid), plotted against cycle number. } \label{fig:Gleissbgfilt} \end{figure}} The Gleissberg filtered sunspot number series is plotted in Fig.~4. One long-term trend is an overall secular increase of solar activity, the last six or seven cycles being unusually strong. (Four of them are markedly stronger than average and none is weaker than average.) This period of elevated sunspot activity level from the mid-20th century is known as the ``Modern Maximum''. On the other hand, cycles 5, 6 and 7 are unusually weak, forming the so-called ``Dalton Minimum''. Finally, the rather long series of moderately weak cycles 12--16 is occasionally referred to as the ``Gleissberg Minimum'' -- but note that most of these cycles are less than $1\sigma$ below the long-term average. While the Dalton and Gleissberg minima are but local minima in the ever changing Gleissberg filtered SSN series, the conspicuous lack of sunspots in the period 1640--1705, known as the Maunder Minimum (Figure~\ref{fig:GSNrecord}) quite obviously represents a qualitatively different state of solar activity. Such extended periods of high and low activity are usually referred to as \textit{grand maxima} and \textit{grand minima.} Clearly, in comparison with the Maunder Minimum, the Dalton Minimum could only be called a ``semi-grand minimum'', while for the Gleissberg Minimum even that adjective is undeserved. A number of possibilities have been proposed to explain the phenomenon of grand minima and maxima, including chaotic behaviour of the nonlinear solar dynamo \citep{Weiss+:chaotic.dynamo}, stochastic fluctuations in dynamo parameters \citep{Moss+:stochastic.dynamo1, Moss+:stochastic.dynamo2} or a bimodal dynamo with stochastically induced alternation between two stationary states \citep{Petrovay:bimodal}. The analysis of long-term proxy data, extending over several millennia further showed that there exist systematic long-term statistical trends and periods such as the so called secular and supersecular cycles (see Section~\ref{sect:spectral}). \subsubsection{Does the Sun have a long term memory? \label{sect:memory}} Following customary usage, by ``memory'' we will refer to some physical (or, in the case of a model, mathematical) mechanism by which the state of a system at a given time will depend on its previous states. In any system there may be several different such mechanisms at work simultaneously ---if this is so, again following common usage we will speak of different ``types'' of memory. A very mundane example are the RAM and the hard disk in a computer: devices that store information over very different time scales and the effect of which manifests itself differently in the functioning of the system. There is no question that the solar dynamo (i.e. the mechanism that gives rise to the sunspot number series) does possess a memory that extends at least over the course of a single solar cycle. Obviously, during the rise phase solar activity ``remembers'' that it should keep growing, while in the decay phase it keeps decaying, even though exactly the same range of $R$ values are observed in both phases. Furthermore, profiles of individual sunspot cycles may, in a first approximation, be considered a one-parameter ensemble \citep{Hathaway:periodampl}. This obvious effect will be referred to here as {\it intracycle memory.} As we will see, correlations between activity parameters in different cycles are generally much weaker than those within one cycle, which strongly suggests that the intracycle memory mechanism is different from longer term memory effects, if such are present at all. Referring back to our analogy, the intracycle memory may work like computer RAM, periodically erased at every reboot (i.e. at the start of a new cycle). The interesting question is whether, in addition to the intracycle memory effect, any other type of memory is present in the solar dynamo or not. To what extent is the amplitude of a sunspot cycle determined by previous cycles? Are subsequent cycles essentially independent, randomly drawn from some stochastic distribution of cycle amplitudes around the long term average? Or, in the alternative case, for how many previous cycles do we need to consider solar activity for successful forecasts? The existence of long lasting grand minima and maxima suggests that the sunspot number record must have a {\it long-term memory} extending over several consecutive cycles. Indeed, elementary combinatorical calculations show that the occurrence of phenomena like the Dalton minimum (3 of the 4 lowest maxima occurring in a row) or the Modern maximum (4 of the 5 highest maxima occurring within a series of 5 cycles) in a random series of 24 recorded solar maxima has a rather low probability (5\,\% and 3\,\%, respectively). This conclusion is corroborated by the analysis of long-term proxy data, extending over several millennia, which showed that the occurrence of grand minima and grand maxima is more common than what would follow from Gaussian statistics \citep{Usoskin+Solanki}. It could be objected that for sustained grand minima or maxima a memory extending only from one cycle to the next would suffice. In contrast to long-term (multidecadal or longer) memory, this would constitute another kind of short-term ($\la 10$ years) memory: a cycle-to-cycle or {\it intercycle} memory effect. In our computer analogy, think of system files or memory cache written on the hard disk, often with the explicit goal of recalling the system status (e.g. desktop arrangement) after the next reboot. While these files survive the reboot, they are subject to erasing and rewriting in every session, so they have a much more temporary character than the generic data files stored on the disk. The intercycle memory explanation of persistent secular activity minima and maxima, however, would imply a good correlation between the amplitudes of subsequent cycles, which is not the case (cf.\ Sect.\ \ref{sect:minimax} below). With the known poor cycle-to-cycle correlation, strong deviations from the long-term mean would be expected to be damped on time scales short compared to e.g.\ the length of the Maunder minimum. This suggests that the persistent states of low or high activity are due to truly long term memory effects extending over several cycles. Further evidence for a long-term memory in solar activity comes from the persistence analysis of activity indicators. The parameter determined in such studies is the Hurst exponent $0<H<1$. Essentially, $H$ is the steepness of the growth of the total range $\cal R$ of measured values plotted against the number $n$ of data in a time series, on a logarithmic plot: ${\cal R}\propto n^H$. For a Markovian random process with no memory $H=0.5$. Processes with $H>0.5$ are persistent (they tend to stay in a stronger-than-average or weaker-than-average state longer), while those with $H<0.5$ are anti-persistent (their fluctuations will change sign more often). Hurst exponents for solar activity indices have been derived using rescaled range analysis by many authors \citep{Mandelbrot+:Hurst, Ruzmaikin:Hurst, Komm:DopplerHurst, Oliver+Ballester:Hurst, Kilcik+}. All studies coherently yield a value $H=0.85$--$0.88$ for time scales exceeding a year or so, and somewhat lower values ($H\sim 0.75$) on shorter time scales. Some doubts regarding the significance of this result for a finite series have been raised by \cite{Oliver+Ballester:noHurst}; however, \cite{Qian+Rasheed} have shown using Monte-Carlo experiments that for time series of a length comparable to the sunspot record, $H$ values exceeding 0.7 are statistically significant. A complementary method, essentially equivalent to rescaled range analysis is detrended fluctuation analysis. Its application to solar data \citep{Ogurtsov:Hurst} has yielded results in accordance with the $H$ values quoted above. The overwhelming evidence for the persistent character of solar activity and for the intermittent appearance of secular cyclicities, however, is not much help when it comes to cycle-to-cycle prediction. It is certainly reassuring to know that forecasting is not a completely idle enterprise (which would be the case for a purely Markovian process), and the long-term persistence and trends may make our predictions statistically somewhat different from just the long-term average. There are, however, large decadal scale fluctuations superposed on the long term trends, so the associated errors will still be so large as to make the forecast of little use for individual cycles. \epubtkImage{}{% \begin{figure}[htbp] \centerline{\includegraphics[width=\textwidth]{waldmeier}} \caption{Monthly smoothed sunspot number $R$ at cycle maximum plotted against the rise time to maximum (left) and against cycle length (right). Cycles are labeled with their numbers. In the plots the red dashed lines are linear regressions to all the data, while the blue solid lines are fits to all data except outliers. Cycle~19 is considered an outlier on both plots, cycle~4 on the right hand plot only. The corresponding correlation coefficients are shown.} \label{fig:waldmeier} \end{figure}} \subsubsection{Waldmeier effect and amplitude--frequency correlation} \label{sect:Waldmeier} \begin{nquote} {\sl Greater activity on the Sun goes with shorter periods, and less with longer periods. I believe this law to be one of the most important relations among the Solar actions yet discovered.}\\ \strut\hfill{\citep{Wolf:cycle.length.activity}} \end{nquote} It is apparent from Figure~\ref{fig:SSNrecord} that the profile of sunspot cycles is asymmetrical, the rise being steeper than the decay. Solar activity maxima occur 3 to 4 years after the minimum, while it takes another 7--8 years to reach the next minimum. It can also be noticed that the degree of this asymmetry correlates with the amplitude of the cycle: to be more specific, the length of the rise phase anticorrelates with the maximal value of $R$ (Figure~\ref{fig:waldmeier}), while the length of the decay phase shows weak or no such correlation. Historically, the relation was first formulated by \cite{Waldmeier:effect} as an inverse correlation between the rise \textit{time} and the cycle amplitude; however, as shown by \cite{Tritakis:Waldmeier}, the total rise time is a weak (inverse logarithmic) function of the rise rate, so this representation makes the correlation appear less robust. (Indeed, when formulated with the rise time it is not even present in some activity indicators, such as sunspot areas -- cf.\ \cite{Dikpati+:Waldmaier}.) As pointed out by \cite{Cameron+:Waldmeier}, the weak link between rise time and slope is due to the fact that in steeper rising cycles the minimum will occur earlier, thus partially compensating for the shortening due to a higher rise rate. The effect is indeed more clearly seen when the rate of the rise is used instead of the rise time \citep{Lantos, Cameron+:Waldmeier}. The observed correlation between rise rate and maximum cycle amplitude is approximately linear, good (correlation coefficient $r\sim 0.85$), and quite robust, being present in various activity indices. Nevertheless, when coupled with the nearly nonexistent correlation between the decay time and the cycle amplitude, even the weaker link between the rise time and the maximum amplitude is sufficient to forge a weak inverse correlation between the total cycle length and the cycle amplitude (Figure~\ref{fig:waldmeier}). This inverse relationship was first noticed by \cite{Wolf:cycle.length.activity}. A stronger inverse correlation was found between the cycle amplitude and the length of the {\it previous} cycle by \cite{Hathaway:periodampl}. This correlation is also readily explained as a consequence of the Waldmeier effect, as demonstrated in a simple model by \cite{Cameron+:prediction}. Note that in a more detailed study \cite{Solanki+:cycle.length} find that the correlation coefficient of this relationship has steadily decreased during the course of the historical sunspot number record, while the correlation between cycle amplitude and the length of the {\it third} preceding cycle has steadily increased. The physical significance (if any) of this latter result is unclear. In what follows, the relationships found by \cite{Wolf:cycle.length.activity}, \cite{Hathaway:periodampl}, and \cite{Solanki+:cycle.length}, discussed above, will be referred to as ``$\Rmax$--$t_{\mathrm{cycle},n}$ correlations'' with $n=0$, $-1$ or $-3$, respectively. Modern time series analysis methods offer several ways to define an instantaneous frequency $f$ in a quasiperiodic series. One simple approach was discussed in the context of Bracewell's transform, Equation~(\ref{eq:Bracewell}), above. \cite{Mininni+:vanderpol} discuss several more sophisticated methods to do this, concluding that G\'abor's analytic signal approach yields the best performance. This technique was first applied to the sunspot record by \cite{Palus+Novotna}, who found a significant long term correlation between the smoothed instantaneous frequency and amplitude of the signal. On time scales shorter than the cycle length, however, the frequency--amplitude correlation has not been convincingly proven, and the fact that the correlation coefficient is close to the one reported in the right hand panel of Figure~\ref{fig:waldmeier} indicates that all that the fashionable gadgetry of nonlinear dynamics could achieve was to recover the effect already known to Wolf. It is clear from this that the ``frequency--amplitude correlation'' is but a secondary consequence of the Waldmeier effect. On the left hand panel of Figure~\ref{fig:waldmeier}, within the band of correlation the points seem to be sitting neatly on two parallel strings. Any number of faint hearted researchers would dismiss this as a coincidence or as another manifestation of the ``Martian canal effect''. But \cite{Kuklin:Waldmeier} boldly speculated that the phenomenon may be real. Fair enough, cycles 22 and 23 dutifully took their place on the lower string even after the publication of Kuklin's work. This speculation was supported with further evidence by \cite{Nagovitsyn:reconstr} who offered a physical explanation in terms of the amplitude--frequency diagram of a forced nonlinear oscillator (cf.\ Section~\ref{sect:oscillator}). Indeed, an anticorrelation between cycle length and amplitude is characteristic of a class of stochastically forced nonlinear oscillators and it may also be reproduced by introducing a stochastic forcing in dynamo models \citep{Stix:Waldmeier, Ossendrijver+:stoch.dynamo, Charbonneau+Dikpati}. In some such models the characteristic asymmetric profile of the cycle is also well reproduced \citep{Mininni+:vanderpol, Mininni+:vanderpol2}. The predicted amplitude--frequency relation has the form \begin{equation} \log\Rmax^{(n)} = C_1+ C_2f \,. \label{eq:stochRf} \end{equation} Nonlinear dynamo models including some form of $\alpha$-quenching also have the potential to reproduce the effects described by Wolf and Waldmeier without recourse to stochastic driving. In a dynamo with a Kleeorin--Ruzmaikin type feedback on $\alpha$, \cite{Kitiashvili+:nonlin.dynamo} are able to qualitatively reproduce the Waldmeier effect. Assuming that the sunspot number is related to the toroidal field strength according to the Bracewell transform, Equation~(\ref{eq:Bracewell}), they find a strong link between rise time and amplitude, while the correlations with fall time and cycle length are much weaker, just as the observations suggest. They also find that the form of the growth time--amplitude relationship differs in the regular (multiperiodic) and chaotic regimes. In the regular regime the plotted relationship suggests \begin{equation} \Rmax^{(n)} = C_1-C_2\left(\tmax^{(n)}-\tmin^{(n)}\right) , \end{equation} while in the chaotic case \begin{equation} \Rmax^{(n)} \propto {\left[1/\left(\tmax^{(n)}-\tmin^{(n)}\right)\right]} . \end{equation} Note that based on the actual sunspot number series Waldmeier originally proposed \begin{equation} \log\Rmax^{(n)} =C_1-C_2\left(\tmax^{(n)}-\tmin^{(n)}\right) , \end{equation} while according to \cite{Dmitrieva:Waldmeierdynamo} the relation takes the form \begin{equation} \log\Rmax^{(n)} \propto {\left[1/\left(\tmax^{(n)}-\tmin^{(n)}\right)\right]} . \end{equation} At first glance, these logarithmic empirical relationships seem to be more compatible with the relation (\ref{eq:stochRf}) predicted by the stochastic models. These, on the other hand, do not actually reproduce the Waldmeier effect, just a general asymmetric profile and an amplitude--frequency correlation. At the same time, inspection of the the left hand panel in Figure~\ref{fig:waldmeier} shows that the data is actually not incompatible with a linear or inverse rise time--amplitude relation, especially if the anomalous cycle~19 is ignored as an outlier. (Indeed, a logarithmic representation is found not to improve the correlation coefficient -- its only advantage is that cycle~19 ceases to be an outlier.) All this indicates that nonlinear dynamo models may have the potential to provide a satisfactory quantitative explanation of the Waldmeier effect, but more extensive comparisons will need to be done, using various models and various representations of the relation. \newpage

\label{sec:Summary-Evaluation} The performance of various forecast methods in cycles 21--23 was discussed by \cite{Li+:pred.rev} and \cite{Kane:pred23rev}. \textit{Precursor methods} stand out with their internally consistent forecasts for these cycles which for cycles 21 and 22 proved to be correct. For cycle~23 these methods were still internally consistent in their prediction, mostly scattering in a narrow range between 150 and 170; however, the cycle amplitude proved to be considerably lower ($\Rmax=121$). It should be noted, however, that one precursor based prediction, that of \cite{Schatten+:pred23} was significantly lower than the rest ($138\pm 30$) and within $0.6\sigma$ of the actual value. Indeed the method of Schatten et al.\ (\citeyear{Schatten+:pred22, Schatten+:pred23}), has consistently proven its skill in all cycles. As discussed in Sect.~\ref{sect:polar}, this method is essentially based on the polar magnetic field strength as precursor. \textit{Extrapolation methods} as a whole have shown a much less impressive performance. Overall, the statistical distribution of maximum amplitude values predicted by ``real'' forecasts made using these methods (i.e.\ forecasts made at or before the minimum epoch) for any given cycle does not seem to significantly differ from the long term climatological average of the solar cycle quoted in Section~\ref{sect:cycle} above ($100\pm 35$). It would of course be a hasty judgement to dismiss each of the widely differing individual approaches comprised in this class simply due to the poor overall performance of the group. In particular, some novel methods suggested in the last 20 years, such as SSA or neural networks have hardly had a chance to debut, so their further performance will be worth monitoring in upcoming cycles. One group of extrapolation methods that stands apart from the rest are those based on the even--odd rule. These methods enjoyed a relatively high prestige until cycle~23, when they coherently predicted a peak amplitude around 200, i.e., $\sim70\%$ higher than the actual peak. This can only be qualified as a miserable failure, independently of the debate as to whether cycle~23 is truly at odds with the even--odd rule or not. In this context it may be worth noting that the double peaked character and long duration of cycle~23 implies that its \textit{integrated} amplitude (sum of annual sunspot numbers during the cycle) is much less below that of cycle~22 than the peak amplitude alone would indicate. This suggests that forecasts of the integrated amplitude (rarely attempted) could be more robust than forecasts of the peak. Nevertheless, one has to live with the fact that for most practical applications (space weather) it is the peak amplitude that matters most, so this is where the interest of forecasters is naturally focused. Finally, \textit{model based methods} are a new development that have had no occasion yet to prove their skill. As discussed above, current dynamo models do not seem to be at a stage of development where such forecasts could be attempted with any confidence, especially before the time of the minimum. (The method of \citealp{Choudhuri+:prediction}, using polar fields as input near the minimum, would seem to be akin to a version of the polar field based precursor method with some extra machinery built into it.) The claimed good prediction skills of models based on data assimilation will need to be tested in future cycles and the roots of their apparent success need to be understood. Table~\ref{tab:1} presents a collection of forecasts for the amplitude of cycle~24, without claiming completeness. \citep[See, e.g.,][for a more exhaustive list.]{Pesnell} The objective was to include one or two representative forecasts from each category. \begin{table} \caption{A selection of forecasts for cycle 24.} \label{tab:1} \begin{tabular}{lrll} \toprule Category & Peak amplitude & Link & Reference \\ \midrule Precursor methods\\ \bek Minimax & $80 \pm 25$ & Eq.~\ref{eq:minimax} & \cite{Brown:minimax}; \cite{Brajsa+:gleissbg}$^*$\\ \bek Minimax3 & $69 \pm 15$ & Eq.~\ref{eq:minimax3} & \cite{Cameron+:prediction}$^*$\\ \bek Polar field & $75 \pm 8$ & Sect.~\ref{sect:polar} & \cite{Svalgaard+:prediction24} \\ \bek Polar field & $80 \pm 30$ & Sect.~\ref{sect:polar} & \cite{Schatten+:pred24}\\ \bek Geomagnetic (Feynman) & $150$ & Sect.~\ref{sect:geomg} & \cite{Hathaway+Wilson}\\ \bek Geomagnetic (Ohl) & $93\pm 20$ & Sect.~\ref{sect:geomg} & \cite{Bhatt+}\\ \bek Geomagnetic (Ohl) & $101\pm 5$ & Sect.~\ref{sect:evenodd} & \cite{Ahluwalia:pred24}\\ \bek Geomagnetic (interpl.)& $97\pm 25$ & Sect.~\ref{sect:geomg} & \cite{Wang+Sheeley:geomg.precursor}\\ \bek Field reversal & $94\pm 14$& Eq.~\ref{eq:tlatov} & \cite{Tlatov:polar.precursors}$^*$\\ Extrapolation methods\\ \bek Linear regression & $90\pm 27$ & Sect.~\ref{sect:regression} & \cite{Brajsa+:gleissbg}\\ \bek Linear regression & $110\pm 10$ & Sect.~\ref{sect:regression} & \cite{Hiremath}\\ \bek Spectral (MEM) & $90\pm 11$ & Sect.~\ref{sect:spectral} & \cite{Kane:MEM24}\\ \bek Spectral (SSA) & $117$ & Sect.~\ref{sect:spectral} & \cite{Loskutov:SSA}\\ \bek Spectral (SSA) & $106$ & Sect.~\ref{sect:spectral} & \cite{Kuzanyan:predict.poster}\\ \bek Attractor analysis & $87$ & Sect.~\ref{sect:nldpro} & \cite{Kilcik+}\\ \bek Attractor analysis & $65\pm 16$ & Sect.~\ref{sect:nldpro} & \cite{Aguirre+}\\ \bek Attractor analysis & $145\pm 7$ & Sect.~\ref{sect:nldpro} & \cite{Crosson+Binder}\\ \bek Neural network & $145$ & Sect.~\ref{sect:neural} & \cite{Maris+Oncica}\\ \bek Neural network & $117.5\pm 8.5$ & Sect.~\ref{sect:neural} & \cite{Uwamahoro}\\ Model based methods\\ \bek Explicit models & $167\pm 12$ & Sect.~\ref{sect:explicit} & \cite{Dikpati+:prediction}\\ \bek Explicit models & $\sim 80$ & Sect.~\ref{sect:explicit} & \cite{Choudhuri+:prediction}\\ \bek Explicit models & $\sim 85$ & Sect.~\ref{sect:explicit} & \cite{Jiang+:prediction}\\ \bek Truncated models & $\sim 80$ & Sect.~\ref{sect:truncated} & \cite{Kitiashvili+:prediction}\\ \bottomrule \end{tabular} {\footnotesize References marked with $^*$ are to the basic principle used in the given prediction method while the actual numerical evaluation for cycle~24 was done by the author. The application for forecast purposes does not necessarily reflect the original intention of the basic principle, as laid out in the cited publications.} \end{table} The incipient cycle~24 may be a milestone for solar cycle forecasting. Current evidence indicates that we are at the end of the Modern Maximum when the Sun is about to switch to a state of less intense long term activity. The appearance of a number of novel prediction methods, in particular the model based approach, as well as the unusually large discrepancy between forecasts based on the precursor approach imply that, whichever course solar activity will take in the coming years, we have a lot to learn from the experience. \newpage

1012.5513

A review of solar cycle prediction methods and their performance is given, including forecasts for cycle 24. The review focuses on those aspects of the solar cycle prediction problem that have a bearing on dynamo theory. The scope of the review is further restricted to the issue of predicting the amplitude (and optionally the epoch) of an upcoming solar maximum no later than right after the start of the given cycle.

12163338

2011A&A...532A..14C

1012

1012.1106_arXiv.txt

\begin{table*} \begin{minipage}[htbp]{\textwidth} \caption{Parameters of the electron number density profiles.}\label{table_neCC} \centering \renewcommand{\footnoterule}{} % \begin{tabular}{lcccccc} \hline Cluster & $n_{e0}$ & $\beta$ & $\theta_{c1}$ & $\theta_{c2}$ & $f$ \\ & (10$^{-2}$cm$^{-3}$) & & (arcsec) & (arcsec) & \\ \hline \hline A1413 & $3.89\pm0.54$ & $0.535\pm0.016$ & $6.7\pm1.4$ & $40.1\pm4.1$ & $0.760\pm0.020$ \\ A1689 & $4.15\pm0.31$ & $0.871\pm0.040$ & $21.6\pm1.0$ & $104.5\pm5.3$ & $0.870\pm0.010$ \\ A1835 & $11.3\pm0.4$ & $0.802\pm0.015$ & $9.3\pm0.2$ & $63.8\pm1.6$ & $0.940\pm0.001$ \\ A2204 & $20.4\pm1.1$ & $0.716\pm0.028$ & $7.5\pm0.3$ & $67.6\pm1.9$ & $0.959\pm0.004$ \\ A2261 & $4.07\pm0.59$ & $0.631\pm0.024$ & $10.2\pm1.8$ & $39.1\pm5.9$ & $0.760\pm0.050$ \\ MS1358.4+6245 & $9.63\pm0.79$ & $0.676\pm0.017$ & $3.3\pm0.2$ & $37.0\pm1.8$ & $0.934\pm0.003$ \\ RXJ1347.5-1145 & $28.5\pm1.4$ & $0.632\pm0.009$ & $4.0\pm0.2$ & $23.3\pm1.6$ & $0.942\pm0.004$ \\ ZW3146 & $16.9\pm0.3$ & $0.669\pm0.005$ & $4.4\pm0.1$ & $25.8\pm0.6$ & $0.882\pm0.004$ \\ \hline \end{tabular} \end{minipage} \end{table*} Clusters of galaxies are the largest gravitationally-bound objects arising thus far from the process of hierarchical structure formation (\cite{Voit05}). As the most recent and most massive objects of the Universe, clusters are excellent probes for studying its formation and evolution. The observed state of gas within a cluster is determined by a combination of shock heating during accretion, radiative cooling, and thermal feedback produced by the cooling itself, so the density and temperature of the ICM represent the full thermal history of clusters' formation. To better understand the physics of ICM, it is necessary to have sufficient knowledge of the gas density and temperature distributions. Though clusters are the ideal target objects for X-ray observations of the hot ICM, millimeter and sub-millimeter measurements provide independent and complementary tools for studying the same ICM by exploiting the Sunyaev Zel'dovich (SZ) effect (\cite{Sunyaev72}). The SZ effect is the Comptonization of the cosmic microwave background (CMB) photons, coming from the last scattering surface, by the hot electrons population of the ICM. The photon energy variation, which is caused by the scattering process, can be expressed as CMB temperature variations \begin{equation}\label{eq_SZ} \Delta T_{SZ}=yT_{CMB}f(x)(1+\delta_n(x,\theta_e))+\Delta T_{kin} \end{equation} where \begin{equation}\label{eq_y} y=\int\theta_ed\tau_e=\int\left(\frac{k_BT_e}{m_ec^2}\right)\sigma_Tn_edl\propto\int P_edl \end{equation} represents the comptonization parameter, $x=(h\nu)/(k_BT_{CMB})$ the dimensionless frequency, $h$ and $k_B$ are respectively the Planck and Boltzmann constants, $T_{CMB}$, $m_e$ and $\sigma_T$, the CMB temperature at $z=0$, the electron mass at rest and the Thomson cross section, $\theta_e$ represents the dimensionless thermal energy of the ICM, $\tau_e$ is the electron optical depth. The parameters $n_e$, $T_e$, and $P_e$ are the electron number density, temperature, and pressure of the ICM, $\delta_n(x,\theta_e)=f_n(x)\theta_e^n/f(x)$ is the relativistic correction term that accounts for the thermal energy of the electrons involved in the scattering processes, where $f(x)=x[(e^x+1)/(e^x-1)-4]$ is a dimensionless quantity that describes the spectral signature of the effect, and the subscript $n$ indicates the maximum order of the relativistic correction ($n=4$ in this work, \cite{Nozawa95}). The last term of Eq. \ref{eq_SZ} is the kinematic component of the SZ effect, which contains the contribution from the bulk motion of the electron population with respect to the last scattering surface reference frame. For the purpose of this paper, this term is omitted, assuming that it is disentangled from the thermal component by multi-frequency observations, together with the signal from the primary CMB emission. The SZ effect is redshift independent and, for this reason, it is possible to detect distant clusters without any existing X-ray or optical observations. This is the case of the ongoing ground-based experiments such as SPT (\cite{Ruhl04}), ACT (\cite{Kosowsky03}), and the all sky survey like Planck (\cite{Planck11a}) or the upgraded MITO (\cite{DePetris07}) and OLIMPO (\cite{Masi08}) with new spectroscopic capabilities and the proposed 30-m diameter C-CAT (\cite{Sebring06}). However, some assumptions on cluster physics still have to be made in order to directly extract cluster observables. Estimates of cluster's total mass can be derived by SZ observations when X-ray or lensing measurements are available or by empirically calibrated scaling relations linking the SZ flux to the total mass (e.g. \cite{Vikhlinin09, Arnaud10, Planck11b, Comis11}). Total mass can also be determined by SZ observations alone when applying thermal energy constraints (\cite{Mroczkowski11}). To accurately reproduce the gas inside the cluster, an ICM universal model is mandatory (e.g. \cite{Nagai07, Arnaud10}). In this paper we confirm that the simple isothermal \textit{beta}-model is clearly an inappropriate cluster representation for total mass recovery by SZ observations, particularly in the presence of relaxed cool core (CC) clusters. These objects show a well studied peaked density profile with a temperature decrement in the core region (\cite{Jones84}). In the local universe this class of clusters is observationally a significant percentage of the total cluster population (\cite{Eckert11}). Even if the X-ray estimated CC fraction is biased by selection effects in flux-limited samples, recently a 35\% of clusters have picked up in the SZ high signal-to-noise ratio Planck early cluster data-set (\cite{Planck11c}) are CC clusters. Large scatter in mass estimates of CC clusters has been highlighted previously using numerical simulations by Hallman et al. (2006) and Hallman et al. (2007). We investigate the bias on the mass in a limited sample of eight nearby ($0.1<z<0.5$) and high-mass ($M>10^{14}$ $M_{\odot}$) CC clusters observed by Chandra. The SZ maps of these clusters, which are expressed in thermodynamic temperature units and convolved with several instrumental beams, are dealt with by applying the isothermal \textit{beta}-model. The total mass is derived in three different ways: by assuming hydrostatic equilibrium and a fixed gas fraction and by applying a self similar scaling relation. To focus only on the consequences of the employed ICM model, in our analysis we neglect all the contaminants present in the sky by assuming in this way the best situation to recover cluster total mass. In Sect. \ref{sec_ne_Te}, we discuss the electron number density radial profile and follow self-similar studies to characterize a universal electron temperature radial profile of a limited sample of eight CC clusters observed by Chandra. In Sect. \ref{sec_MCMC} we generate maps of the SZ effect in thermodynamic temperature. In Sect. \ref{sec_total_mass} we evaluate cluster total mass under different sets of assumptions. The bias on the recovered mass is described in Sect. \ref{sec_MB}, which discusses the main contributions. Conclusions are summarized in Sect. \ref{sec_conclusions}.

We studied the bias that affects the estimate of cluster total mass by SZ observations when an isothermal \textit{beta}-model is assumed to describe the ICM physical properties, specifically in the case of CC clusters when X-ray and/or lensing information is missing. It is well known that, rather than the central Comptonization parameter $y_0$, affected by the choice of cluster profile modeling, an integrated quantity, like the parameter $Y$, appears to be a more robust mass proxy. Nevertheless, we have shown that CC clusters can generate observed $y$ maps in which the ICM morphology could still be substantially hidden, even for the current most sensitive experiments operating from the largest available mm/submm telescopes. While a general assumption of cluster morphology is efficient at detecting them in blind SZ maps, the possible mismatch with the actual cluster profile results in a mass bias. In fact simple ICM models, like the isothermal \textit{beta}-model, applied to SZ observations can wrongly estimate cluster total mass in the presence of peculiar ICM dynamics as in CC clusters, which are studied in the current analysis, and mergers. We analyzed the mass bias as derived in a limited sample of eight CC clusters observed by Chandra, both nearby ($0.1<z<0.5$) and with high mass ($M>10^{14}$ $M_{\odot}$). Under the assumption of an isothermal \textit{beta}-model, the cluster total mass was derived applying three different approaches: the hydrostatic equilibrium equation, a fixed gas fraction, and a self-similar $M_{tot}-Y$ relation. Assuming we had no information from X-ray observations, we reported the bias on the derived total mass as dependent on electron gas temperature. Only in the case of hydrostatic equilibrium does this bias appear almost constant for the considered clusters in the range of 50-80 \%. Incidentally, we notice that an electron temperature value exists for which the FGF and SL mass biases vanish. This could be the only case in which a simple isothermal \textit{beta}-model accurately reproduces the mass of CC clusters. The large biases on total cluster mass recovery in CC clusters represent another reason to definitely discard the isothermal \textit{beta}-model for this purpose and to firmly support more sophisticated models, with universal pressure profiles (e.g. \cite{Arnaud10}). This is already employed for modeling cluster atmospheres in almost all the present blind-survey data reduction (SPT and Planck), and it is planned in the next future for ACT observations.

1012.1106

The Sunyaev Zel'dovich (SZ) effect from galaxy clusters is one of the most powerful cosmological tools for investigating the large-scale Universe. The big advantage of the SZ effect is its redshift independence, which is not the case for visible and X-ray observations. It allows us to directly estimate the cluster's total mass from the integrated comptonization parameter Y, even for distant clusters. However, not having a full knowing intra-cluster medium (ICM) physics can affect the results. By taking self-similar temperature and density profiles of the ICM into account, we studied how different ICM morphologies can affect the cluster total mass estimation. With the help of the high percentage of cool core (CC) clusters, as observed so far, the present analysis focuses on studying this class of objects. A sample of eight nearby (0.1 &lt; z &lt; 0.5) and high-mass (M &gt; 10<SUP>14</SUP> M<SUB>⊙</SUB>) clusters observed by Chandra was considered. We simulated SZ observations of these clusters through X-ray derived information and analyzed the mock SZ data again with the simplistic assumption of an isothermal beta-model profile for the ICM. The bias on the recovered cluster total mass using different sets of assumptions is estimated to be 50% higher in the case of hydrostatic equilibrium. Possible contributions to the total bias due to the line-of-sight integration and the considered ICM template are taken into account. The large biases on total mass recovery firmly support, if still necessary, cluster modeling based on more sophisticated universal profiles as derived by X-ray observations of local objects and hydrodynamical simulations.

12207189

2011JCAP...03..051C

1012

1012.4515_arXiv.txt

\label{introduction} Cosmology and astrophysics provide several convincing evidences of the existence of Dark Matter~\cite{JungmanReview, BertoneReview, EinastoReview}. The observation that some mass is missing to explain the internal dynamics of galaxy clusters and the rotations of galaxies dates back respectively to the '30s and the '70s~\cite{rotation}. The observations from weak lensing~\cite{lensing}, for instance in the spectacular case of the so-called `bullet cluster'~\cite{bullet}, provide evidence that there is mass where nothing is optically seen. More generally, global fits to a number of cosmological datasets (Cosmic Microwave Background, Large Scale Structure and also Type Ia Supernovae) allow to determine very precisely the amount of DM in the global energy-matter content of the Universe at $\Omega_{\rm DM} h^2 =0.1123 \pm 0.0035$~\cite{cosmoDM}\footnote{Here $\Omega_{\rm DM} = \rho_{\rm DM}/\rho_c$ is defined as usual as the energy density in Dark Matter with respect to the critical energy density of the Universe $\rho_c = 3 H_0^2/8\pi G_N$, where $H_0$ is the present Hubble parameter. $h$ is its reduced value $h = H_0 / 100\ {\rm km}\, {\rm s}^{-1} {\rm Mpc}^{-1}$.}. All these signals pertain to the gravitational effects of Dark Matter at the cosmological and extragalactical scale. Searches for explicit manifestation of the DM particles that are supposed to constitute the halo of our own galaxy (and the large scale structures beyond it) have instead so far been giving negative results, but this might be on the point of changing. \medskip Indirect searches for Dark Matter aim at detecting the signatures of the annihilations or decays of DM particles in the fluxes of cosmic rays, intended in a broad sense: charged particles (electrons and positrons, protons and antiprotons, deuterium and antideuterium), photons (gamma rays, X-rays, synchrotron radiation), neutrinos. Pioneering works have explored this as a promising avenue of discovery since the late-70's: gamma rays from annihilations were first considered in~\cite{Gunn,Stecker,Zeldovich} and then revisited in~\cite{Ellis}, antiprotons in~\cite{SilkSrednicki,SteckerRudazWalsh} and then more systematically in~\cite{Ellis,SteckerRudaz,SteckerTylka}, positrons in~\cite{SilkSrednicki,Ellis,SteckerRudaz,TurnerWilczek}, antideuterons have been first discussed in~\cite{pioneerDbar,Dbar2,followupDbar}, radio-waves from synchrotron radiation from DM in~\cite{BerezinskyGurevich,BerezinskyBottino,Gondolo,BertoneSiglSilk} and later in~\cite{AloisioBlasiOlinto} (which questions the approach in~\cite{Gondolo,BertoneSiglSilk}), extragalactic gamma rays have been first discussed in~\cite{BergstomEdsjoUllio}. Inverse Compton gamma rays from DM have been only relatively recently considered as a possible signal (see e.g.~\cite{BaltzWai,Cholis, Zhang}). In general, a key point of all these searches is to look for channels and ranges of energy where it is possible to beat the background from ordinary astrophysical processes. This is for instance the basic reason why searches for charged particles focus on fluxes of antiparticles (positrons, antiprotons, antideuterons), much less abundant in the Universe than the corresponding particles. A well spread theoretical prejudice wants the DM particles to be thermal relics from the Early Universe. They were as abundant as photons in the beginning, being freely created and destroyed in pairs when the temperature of the hot plasma was larger then their mass. Their relative number density started then being suppressed as annihilations proceeded but the temperature dropped below their mass, due to the cooling of the Universe. Finally the annihilation processes also froze out as the Universe expanded further. The remaining, diluted abundance of stable particles constitutes the DM today. As it turns out, particles with weak scale mass ($\sim 100\, {\rm GeV} - 1\, {\rm TeV}$) and weak interactions \cite{JungmanReview, BertoneReview} could play the above story remarkably well, and their final abundance would automatically (miracolously?) be the observed $\Omega_{\rm DM}$. While this is not of course the only possibility, the mechanism is appealing enough that a several-GeV-to-some-TeV scale DM particle with weak interactions (WIMP) is often considered as the most likely DM candidate. In any case, this mass range (TeV-ish DM) has the best chances of being thoroughly explored in the near future by charged particle and photon observatories, also in combination with direct DM searches (aiming at detecting the nuclear recoil produced by a passing DM particle in ultra-low background underground detectors) and, possibily, production at LHC collider. It is therefore the focus of our attention. Supposing (and hoping) therefore that anomalous features are detected in the fluxes of cosmic rays, it will be crucial to be able to `reverse engineer' them to determine which Dark Matter is at their origin. Moreover, it will be useful to be able to quickly compute which other associated signals are implied by a possible positive detection and have to be looked for in other channels. Only via a cross-correlation of multi-messenger signals a putative detection of DM will be confirmed or disproved. More generally, in order to compute the predicted signatures of a given model of Dark Matter, a number of particle physics and astrophysics ingredients are needed. These ingredients are what we aim to provide. \bigskip This work does not contain any new theoretical proposal nor any new study of (existing or foreseen) data. It contains instead all the phenomenological ingredients that allow to perform the analyses sketched above in the most general possible way. More precisely, the rest of this compilation is organized as follows. In Section~\ref{DMdistribution} we start by recalling the most commonly used DM distribution profiles in the Milky Way, that we will adopt for the computation of all signals. In Section~\ref{primary} we discuss the production of the fluxes of Standard Model particles from DM annihilations (and decays): we compare the {\sc Pythia} and {\sc Herwig} Monte Carlos and quantify the uncertainties. Section~\ref{charged} deals with the propagation in the Galaxy and the resulting fluxes of charged cosmic rays from DM: electrons, positrons, antiprotons, antideuterons. Section~\ref{promptgamma} deals with the basics of prompt gamma rays from Dark Matter annihilations (or decays). Section~\ref{ICSgamma} discusses the `secondary' radiation from $e^\pm$ produced by DM annihilations or decays: Inverse Compton Scattering (ICS) $\gamma$-rays and synchrotron radiation. Section~\ref{extragalactic} presents the results on extragalactic gamma rays. One signature that we do not discuss here is that of neutrino fluxes from the annihilation/decays of DM accumulated in the center of the Sun: we refer the reader to~\cite{DMnu,wimpsim}, where they have been discussed in a spirit very similar to the one of the present work, and we intend to upgrade those former results in the light of the current developements in upcoming work. Several of these parts contain innovations with respect to the existing literature. For instance, the comparison among Monte Carlo generators; the propagation halo functions for $e^\pm$, which allow to take into account, in a semi-analytic way, point-dependent energy losses and therefore to compute much more precisely the fluxes of charged cosmic rays and (above all) ICS photons; the introduction of a formalism in terms of (other) halo functions to compute the flux of IC $\gamma$ rays; the study of the impact of different choices for the model of extragalactic background light on the predicted fluxes of extragalactic gamma rays... \medskip All our numerical results are available at the \myurl{www.marcocirelli.net/PPPC4DMID.html}{website} referenced in~\cite{website}. So finally in Section~\ref{summary} we give a summary of these provided numerical ingredients and we list the entry points in the text for the main recipes. Of course, many refined numerical tools which allow to (directly or indirectly) compute Dark Matter indirect detection signatures have been developed in the latest decades. Among them, GALPROP~\cite{galprop}, DarkSUSY~\cite{darksusy}, MicrOMEGAs~\cite{micromegas}, IsaTools~\cite{isatools}, WimpSim~\cite{wimpsim}... Rather than focusing on a particular DM model, we try to be model-independent and parameterize the observables in terms of the DM mass, of the DM decay or annihilation rates and channels, as well as in terms of a few uncertain astrophysical assumptions. We prefer, whenever possible, a semi-analytical treatment that allows us to keep track of the approximations and choices that we make along the way. Also, we aim at providing the reader with ready-to-use final products, as opposed to the generating code. We make an effort to extend our results to large, multi-TeV DM masses (recently of interest because of possible multi-TeV charged cosmic ray anomalies) and small, few-GeV DM masses (recently discussed because of hints from DM direct detection experiments), at the edge of the typical WIMP window. Above all, our aim is to provide a self-consistent, independently computed, comprehensive set of results for DM indirect detection. Whenever possible, we have compared with existing codes, finding good agreement or improvements.

\label{summary} \subsection{Ingredients} On the~\myurl{www.marcocirelli.net/PPPC4DMID.html}{website}~\cite{website}, we provide the following: \begin{enumerate}[label=\fbox{\arabic*}] \item \label{dlNdlxEW} {\tt dlNdlxIEW[primary->final][DMmass,log$_{10}$x]}: spectrum $d\ln N/d\ln x$ in $x = E/M_{\rm DM}$ of {\tt final} particles generated by DM annihilations into a pair of {\tt primary} particles.\\ Also in terms of numerical tables. \item \label{dlNdlxPythia} {\tt dlNdlxI[primary->final][DMmass,log$_{10}$x]}: same as \ref{dlNdlxEW}, but without EW corrections.\\ Also in terms of numerical tables. \item \label{b[E,r,z]} {\tt b[E,r,z]}: energy loss coefficient function $b(E,\vec x)$ for $e^\pm$ of energy {\tt E} at the position {\tt (r,z)} in the Galaxy. \item \label{ElectronHaloFunctGalaxyAnnI} {\tt ElectronHaloFunctGalaxyAnnI[halo,propag][log$_{10}$x,log$_{10}$E$_s$,r,z]}: generalized halo functions $I(E,E_{\rm s},r,z)$ for $e^\pm$ for annihilations, in any given point ({\tt r}, {\tt z}) of the Galaxy, expressed as a function of $x = E/E_{\rm s}$.\\ Analogous for decay. \item \label{ElectronHaloFunctEarthAnnI} {\tt ElectronHaloFunctEarthAnnI[halo,propag][log$_{10}$x,log$_{10}$E$_s$]}: generalized halo functions $I(E,E_{\rm s},\vec r_\odot)$ for $e^\pm$ for annihilations, at the location of the Earth, expressed as a function of $x = E/E_{\rm s}$.\\ Analogous for decay. \item \label{fitcoefficientspos} Tables of {\tt fit coefficients} for the reduced halo functions for $e^\pm$ for annihilations $\mathcal{I}(\lambda,\vec r_\odot)$ at the location of the Earth.\\ Analogous for decay. \item \label{fitcoefficientspbar} Tables of {\tt fit coefficients} for the propagation functions for $\bar p$ for annihilations $R(K)$ at the location of the Earth.\\ Analogous for decay. \item \label{fitcoefficientsdbar} Tables of {\tt fit coefficients} for the propagation functions for $\bar d$ for annihilations $R(K_d/n)$ at the location of the Earth.\\ Analogous for decay. \item \label{ElectronFluxAnn} {\tt ElectronFluxAnn[primary,halo,propag][DMmass,$\sigma$v,log$_{10}$E$_e$]}: differential flux $\displaystyle\frac{d\Phi_{e^\pm}}{dE}$ at Earth.\\ Analogous for decay. \item \label{ProtonFluxAnn} {\tt ProtonFluxAnn[primary,halo,propag][mass,$\sigma$v,log$_{10}$K]}: differential flux $\displaystyle\frac{d\Phi_{\bar p}}{dK}$ at Earth.\\ Analogous for decay. \item \label{DeuteronFluxAnn} {\tt DeuteronFluxAnn[primary,halo,propag][mass,$\sigma$v,log$_{10}$K$_d$]}: differential flux $\displaystyle\frac{d\Phi_{\bar d}}{dK_d}$ at Earth.\\ Analogous for decay. \item \label{Jave} {\tt Jave[halo][Log$_{10}\theta$]}: factor $J(\theta)$ for prompt gamma rays.\\ Analogous for decay. \item \label{IIC} {\tt IICAnnI[halo,propag][Log$_{10}E_s$,Log$_{10}$E$_\gamma$,Log$_{10}\ell$,Log$_{10}b$]}: halo functions $I_{\rm IC}(E_\gamma,E_{\rm s},b,\ell)$ for Inverse Compton, for annihilations.\\ Analogous for decay. \item \label{ICcode} {\tt Code bite to compute} $\displaystyle \frac{d\Phi_{\rm IC \gamma}}{dE_\gamma}$. \item \label{BoostF} {\tt BoostF[z,minhalomass,cmodel]}: cosmological boost factor $B(z)$ as a function of redshift {\tt z} for a choice of $M_{\rm min}$ and $c(M)$. \item \label{ETau} {\tt ETau[E,z',UVmodel]}: optical depth of the Universe $e^{\tau(E,z')}$ for a choice of UV background. \item \label{EGgammaFluxAnn} {\tt EGgammaFluxAnn[primary,minhalomass,cmodel,UVmodel][mass,$\sigma$v,lE]}: differential flux $\displaystyle \frac{d\Phi_{{\rm EG}\gamma}}{dE_\gamma}$.\\ Analogous for decay. \end{enumerate} The {\sc Mathematica}$^{\tiny{\textregistered}}$ {\tt InterpolationFunction}s and the code bite provided in \ref{ICcode} have been produced with {\sc Mathematica}$^{\tiny{\textregistered}}$ version 7.0.0. The {\tt InterpolationFunction}s are expected to be compatible with version 2 and any later version. The code bite employs solution methods included in {\sc Mathematica}$^{\tiny{\textregistered}}$ version 4 and later. \bigskip Table~\ref{tab:choices} lists the discrete choices for the variables employed in the functions above, together with a reference to the corresponding discussion in the text (when available). \begin{table} \begin{center} \begin{tabular}{c|c|c} Variable & Values or range & Refer to\\ \hline {\tt primary} & $ \begin{array}{c} {\tt eL,\ eR,\ \mu L,\ \mu R,\ \tau L,\ \tau R,}\\[1mm] {\tt q, \ c, \ b, \ t, \ \gamma,\ g, WL,\ WT,\ ZL,\ ZT, }\\[1mm] {\tt h_{115}, \ h_{135}, \ h_{170}, \ h_{200}, \ h_{300}, \ h_{400}, \ h_{500},} \\[1mm] {\tt \nu e, \ \nu \mu, \ \nu \tau, V \to e, \ V \to \mu, \ V \to \tau } \end{array}$ & Eq. (\ref{primarychannels})\\ \hline {\tt final} & e, p, $\gamma$, d, $\nu$e, $\nu \mu$, $\nu \tau$ & Sec.~\ref{primary} \\ \hline {\tt DMmass} & $ \begin{array}{c} 5 \ {\rm GeV} \to 100 \ {\rm TeV\ (annihilation)} \\ 10 \ {\rm GeV} \to 200 \ {\rm TeV\ (decay)} \end{array}$ & Sec.~\ref{fluxesresults} \\ \hline {\tt halo} & {\tt NFW, Ein, EiB, Iso, Bur, Moo} & Fig.~\ref{fig:DMprofiles} \\ \hline {\tt propag} & {\tt MIN, MED, MAX} & Table \ref{tab:proparam} \\ \hline {\tt $\sigma$v} or {\tt $\Gamma$} & any & \\ \hline {\tt minhalomass} & {\tt 10}$^{{\tt -6}}$, {\tt 10}$^{{\tt -9}}$ & Sec.~\ref{subsec:EGclusering} \\ \hline {\tt cmodel} & {\tt maccio, powerlaw} & Sec.~\ref{subsec:EGclusering} \\ \hline {\tt UVmodel} & {\tt noUV, minUV, maxUV} & Sec.~\ref{subsec:absorption} \\ \end{tabular} \end{center} \caption{\em \small {\bfseries Variables of the numerical functions and their admitted values}. \label{tab:choices}} \end{table} \subsection{Recipes} The main recipes for computing DM indirect detection signals are: \begin{enumerate}[label=\Roman*.] \item {\bf Computing $\displaystyle \frac{d\Phi_{e^\pm}}{dE}$: the differential flux of $e^\pm$ at Earth}:\\[2mm] eq.~(\ref{eq:positronsflux}) with \ref{b[E,r,z]}, \ref{dlNdlxEW} and \ref{ElectronHaloFunctEarthAnnI}. Or use \ref{ElectronFluxAnn}. \item {\bf Computing $\displaystyle \frac{d\Phi_{e^\pm}}{dE}$ with approximated energy losses}:\\[2mm] eq.~(\ref{eq:positronsfluxnox}) with \ref{dlNdlxEW} and \ref{fitcoefficientspos}. \item {\bf Computing $\displaystyle \frac{d\Phi_{\bar p}}{dK}$: the differential flux of antiprotons at Earth}:\\[2mm] eq.~(\ref{eq:fluxpbar}) with \ref{dlNdlxEW} and \ref{fitcoefficientspbar} in eq.~(\ref{eq:fitantiprotons}). Or use \ref{ProtonFluxAnn}. \item {\bf Computing $\displaystyle \frac{d\Phi_{\bar d}}{dK_d}$: the differential flux of antideuterons at Earth }:\\[2mm] eq.~(\ref{eq:fluxdbar}) with \ref{dlNdlxEW} and \ref{fitcoefficientsdbar} in eq.~(\ref{eq:fitantideuterons}). Or use \ref{DeuteronFluxAnn}. \item {\bf Computing $\displaystyle \frac{d\Phi_{\gamma}}{dE_\gamma}$: the differential flux of prompt $\gamma$ rays}:\\[2mm] eq.~(\ref{gammafluxI}) with \ref{dlNdlxEW} and $\bar J$ from table~\ref{tab:Jfactors} or from eq.~(\ref{eq:formulaeJfactors}) with \ref{Jave}. \item {\bf Computing $\displaystyle \frac{d\Phi_{\rm IC \gamma}}{dE_\gamma}$: the differential flux of galactic ICS $\gamma$ rays}:\\[2mm] eq.~(\ref{eq:summaryICI}) with eq.~(\ref{eq:summaryIC}), with \ref{dlNdlxEW} and \ref{IIC}. Or use directly \ref{ICcode}. \item {\bf Computing $\displaystyle \nu \frac{dW_{\rm syn}}{d\nu \, d\Omega}$: the differential flux of galactic synchrotron radiation}:\\[2mm] eq.~(\ref{synspectrumfinal}), with \ref{dlNdlxEW} and \ref{ElectronHaloFunctGalaxyAnnI}. \item {\bf Computing $\displaystyle \frac{d\Phi_{\rm EG \gamma}}{dE_\gamma}$: the differential flux of extragalactic $\gamma$ rays}:\\[2mm] eq.~(\ref{eq:EGfluxtoday}) with eq.~(\ref{jEGprompt}) and eq.~(\ref{jEGICS}), and with \ref{BoostF}, \ref{ETau} and \ref{dlNdlxEW}. Or use directly \ref{EGgammaFluxAnn}. \end{enumerate} \paragraph{Acknowledgements} We thank Gianfranco Bertone, C\'eline Combet, David Maurin, Fabrizio Nesti, Michele Papucci, Stefano Pozzorini, Pierre Salati, Peter Skands, Pasquale D. Serpico, Torbj\"orn Sj\"ostrand and Marco Taoso for useful discussions. We are particularly grateful to Caner \"Unal (Middle East Technical University and CERN Summer Student Program 2010) for his important feedback on the numerical results. We also thank the anonimous JCAP referee for his/her useful and thorough comments on the first version of this paper. This work was supported by the ESF Grant 8090, ESF Grant 8499, Estonian Ministry of Education and Research project SF0690030s09 and European Social Fund (Mobility Grant MJD52). We also thank the EU Marie Curie Research \& Training network ``UniverseNet" (MRTN-CT-2006-035863) for support. \bigskip \appendix \footnotesize \begin{multicols}{2}

1012.4515

We provide ingredients and recipes for computing signals of TeV-scale Dark Matter annihilations and decays in the Galaxy and beyond. For each DM channel, we present the energy spectra of at production, computed by high-statistics simulations. We estimate the Monte Carlo uncertainty by comparing the results yielded by the Pythia and Herwig event generators. We then provide the propagation functions for charged particles in the Galaxy, for several DM distribution profiles and sets of propagation parameters. Propagation of e<SUP>±</SUP> is performed with an improved semi-analytic method that takes into account position-dependent energy losses in the Milky Way. Using such propagation functions, we compute the energy spectra of e<SUP>±</SUP>,bar p and bar d at the location of the Earth. We then present the gamma ray fluxes, both from prompt emission and from Inverse Compton scattering in the galactic halo. Finally, we provide the spectra of extragalactic gamma rays. All results are <A href="http://www.marcocirelli.net/PPPC4DMID.html">available in numerical form</A> and ready to be consumed.

3830987

2011MNRAS.414.3014C

1012

1012.3103_arXiv.txt

Over the last few years, a number of observational discoveries have brought magnetars (ultra-magnetized isolated neutron stars) to the forefront of researchers attention. These extreme objects comprise the Anomalous X-ray Pulsars (AXPs; 10 objects) and the Soft Gamma-ray Repeaters (SGRs; 5 objects), which are observationally very similar classes in many respects. They are all slowly rotating X-ray pulsars with spin periods clustered in a narrow range ($P\sim$ 2--12\,s), relatively large period derivatives ($\dot P \sim 10^{-13}-10^{-10}$s\,s$^{-1}$), spin-down ages of $10^3-10^4$\,yr, and magnetic fields, as inferred from the classical magnetic dipole spin-down formula, of $10^{14}-10^{15}$\,G (for a recent review see \cite{2008A&ARv..15..225M}). SGRs undergo periods of activity during which recurrent bursts with sub-second duration and peak luminosities of $\sim 10^{38}-10^{41}$ erg/s are observed. SGRs also show, on rare occasions, extreme events known as giant flares. These are characterized by an initial spike of duration comparable to that of recurrent bursts, but many orders of magnitude higher luminosity. Only three giant flares have so far been observed in over 30 yr of monitoring. According to the magnetar model \citep{1993ApJ...408..194T,2001ApJ...561..980T} energy is fed impulsively to the neutron star magnetosphere when local ``crustquakes'' let magnetic helicity propagate outwards, giving rise to recurrent bursts with a large range of amplitudes. Giant flares are believed to originate from large-scale rearrangements of the inner field or catastrophic instabilities in the magnetosphere \citep{2001ApJ...561..980T,2003MNRAS.346..540L}. Most of this energy breaks out of the magnetosphere in a fireball of plasma expanding at relativistic speeds which results in the initial spike of giant flares. The decaying, oscillating tail that follows the spike displays many tens of cycles at the neutron star spin rate. This is interpreted as being due to a ``trapped fireball'' which remains anchored inside the magnetosphere and cools down in a few minutes. The total energy released in this tail is $\sim 10^{44}$ erg in all three events detected so far. A power spectrum analysis of the high time resolution data from the 2004 Dec 27 event of SGR1806-20, observed with the X-Ray Timing Explorer (RXTE), led to the discovery of fast Quasi Periodic Oscillations (QPOs) in the X-ray flux of the decaying tail of SGR (\cite{I2005}). QPOs with different frequencies were detected, some of which were active simultaneously and displayed highly significant QPO signals at about 18, 26, 30, 93, 150, 625 and 1840 Hz \citep{WS2006a}. A re-analysis of the decaying tail data from the 1998 giant flare of another magnetar, SGR 1900+14, revealed QPOs around frequencies of 28, 54, 84 and 155 Hz \citep{SW2006}. Hints for a signal at $\sim$ 43 Hz in the March 1979 event from SGR 0526-66 were reported as early as 1983 \citep{1983A&A...126..400B}. All QPO signals show large amplitude variations with time and especially with the phase of the stars rotational modulation. QPOs have also been argued to provide independent evidence for superstrong magnetic fields in SGRs \cite{2007ApJ...661.1089V}. Numerous explanations have been proposed for the origin of the QPOs including the torsional oscillations of the crust alone or even as global seismic vibration modes of magnetars \citep{2005ApJ...634L.153P,2007MNRAS.375..261S,SA2007,2009PhRvL.103r1101S}. Moreover, \cite{2006MNRAS.368L..35L} argued that the QPOs may be driven by the global mode of the magneto-hydrodynamic (MHD) fluid core of the neutron star and its crust, rather than the mechanical mode of the crust. Following this idea, \cite{2008MNRAS.385L...5S,2009MNRAS.396.1441C,2009MNRAS.397.1607C} recently made two-dimensional numerical simulations (both linear and non-linear) and found that the Alfv\'en oscillations form continua which may explain the observed QPOs. The weak point of these simulations is the absence of the crust. Moreover, some of the expected eigenfrequencies \citep{1998ApJ...498L..45D,I2005,2005ApJ...634L.153P,2007MNRAS.375..261S,2008MNRAS.385L...5S,2009MNRAS.396.1441C,2009MNRAS.397.1607C} match the observed QPO frequencies found by \cite{SW2006}. These findings have opened the field of magnetar seismology, which provides an unprecedented probe of the star's crust and interior. In retrospective, predictions made in \cite{2009MNRAS.396.1441C} led to re-analysis of the 2004 Dec 27 event and revealed some new frequencies \citep{2011A&A...528A..45H}. In this work we move to a more general model for neutron star i.e. we assume a fluid interior and a solid crust which are threatened by a dipole magnetic field. In the interior the perturbations will be dominated by Alfv\'en waves propagating along the field lines forming a continuum spectrum as it has been seen in \citep{2008MNRAS.385L...5S,2009MNRAS.396.1441C,2009MNRAS.397.1607C} while in the crust torsional oscillations dominate \citep{2007MNRAS.375..261S}. Our results differ partially from the ones by \cite{2011MNRAS.410L..37G}, where the authors study a magnetised neutron stars with the presence of an elastic crust: in fact, we discover the presence of a type of ``discrete Alfv\'en modes'' which are not present in the absence of a crust. These modes have been seen by \cite{2011MNRAS.410.1036H} in the gaps between two contiguous continua (for this reason they call them also ``gaps modes"). However, we are able to resolve these modes also inside the continua.

We studied the torsional oscillations of a magnetar in a general relativistic framework, consisting of a relativistic neutron star with a solid crust and a fluid core. The core and the crust are coupled by the strong magnetic field and this coupling is defined by the appropriate boundary conditions. We find that the presence of a crust and its coupling with the core partially alters the earlier picture presented in \cite{2009MNRAS.396.1441C}. In particular, the presence of a crust makes the spectrum denser and thus offers an explanation of the nature of all the observed frequencies. We can distinguish, using the eigenfunction, between global modes (i.e. modes for which both the crust and the core oscillate), crust modes and discrete Alfv\'en modes (i.e. modes that have a discrete nature but that scale with the magnetic field). Unlike the paper by \cite{2011MNRAS.410.1036H}, we find that also in the case of intermediate magnetic field ($B<10^{15}$G), the oscillations is not confined to the core but also the crust is excited. We find a family of modes, the discrete Alfv\'en modes, in the gaps between two contiguous continua (as found by \cite{2011MNRAS.410.1036H}) and, in addition, we can resolve them also inside the continua. We can identify all the frequencies observed in SGR 1806-20 and SGR 1900+14 and this allows us to put some constrains on radius, mass, crust thickness and magnetic field strength. In particular, the frequencies observed in SGR 1806-20 can be fitted uniquely with the APR stellar model with a mass $M=1.4 \; M_{\odot}$, a radius $R=11.57$km, a compactness $M/R=0.178$ and a crust thickness $\Delta r/R=0.099$. The magnetic field strength at the pole is $B=4\times 10^{15}$G. Contrary, the SGR 1900+14 cannot be strongly constrained because of the limited information from the observations, i.e just three QPOs were identified. For this reason, the SGR 1900+14 cannot be reproduced uniquely by a single model, as it has been done for the SGR 1806-20, since at least two stellar models could reproduce its observed frequencies with a good accuracy: the APR stellar model with a mass $M=1.4 \; M_{\odot}$, a radius $R=11.57$km, a compactness $M/R=0.178$ and a crust thickness $\Delta r/R=0.099$ but with a magnetic field strength $B=4.25\times 10^{15}$G and the WFF model with a mass $M=1.4M_{\odot}$, a radius $R=10.91$km, a compactness $M/R=0.189$, crust thickness $\Delta r/R=0.085$ and a magnetic field strength $B=4\times 10^{15}$G. Note that, in all the models, the crust is described by the NV equation of state.

1012.3103

We study the axisymmetric perturbations of neutron stars endowed with a strong magnetic field (magnetars), considering the coupled oscillations of the fluid core with the solid crust. We recover discrete oscillations based mainly on the crust and a continuum in the core. We also confirm the presence of 'discrete Alfvén modes' in the gap between two contiguous continua and, in addition, we can resolve some of them also inside the continua. Our results can explain both the lower and the higher observed quasi-periodical oscillations (QPOs) in SGR 1806-20 and SGR 1900+14 and put constraints on the mass, radius and crust thickness of the two magnetars.

2674579

2011AdAst2011E...1D

1012

1012.3045_arXiv.txt

Both dark matter problem and the cosmic acceleration problem challenge physicists' understanding of the universe. In the standard $\Lambda$CDM model, two mixed fluids, dark matter and dark energy fluid, are assumed. These two fluids influence the cosmic evolution separately. However, present gravitational probe does not have the ability to differentiate these two fluids. This is the dark degeneracy problem \cite{dg1} \cite{dg2}. It is reasonable to model dark matter and dark energy with single fluid or single field assumption. Some unified models have been proposed to detect the possibility of this unified assumption, like unified dark fluid model \cite{sf1} \cite{sf2} \cite{sf3} \cite{sf4} \cite{sf5}, which assumes the single fluid equation of state; Chaplygin gas model and generalized Chaplygin gas \cite{cp} \cite{cp2} \cite{cp3} \cite{cp4} \cite{cp5}, which discuss the cosmology consequences of an exotic equation of state; scalar field method \cite{field1} \cite{field2} \cite{field3}. The introduction of viscosity into cosmology has been investigated from different view points \cite{G} \cite{PC}. There are some recent developments like dark energy model \cite{vd1} \cite{vd2} \cite{vd3} \cite{vd4}, the cosmic singularity \cite{vs1}. In this review, we give a brief introduction to unify dark matter and dark energy with viscosity medium. In such models, the universe is assumed to be filled with viscous single fluid \cite{sv1} \cite{sv2} \cite{sv3} \cite{sv4} \cite{sv5} \cite{sv6} \cite{sv7}. The cosmic density is not separated as dark energy part and dark matter part. The bulk viscosity contributes to the cosmic pressure, and plays the role as accelerating the universe. After considering the bulk viscosity, the cosmic pressure can be written as \begin{equation} p=(\gamma-1)\rho-3\zeta H \end{equation} Where $\gamma$ parameterizes the equation of state. Generally the form of bulk viscosity is chosen as a time-dependent function. In \cite{v1} \cite{v2} \cite{v3} \cite{v4}, a density-dependent viscosity $\zeta=\alpha\rho^{m}$ coefficient is investigated extensively. For modeling the unified dark matter and dark energy, it is often assumed that the parameter $\gamma=1$, that the pressure of the viscosity fluid is zero and the viscosity term contributes an effective pressure. There raises some problems here. From the observational results \cite{w}, the cosmic density nearly equals to the cosmic pressure. In the viscosity model, the viscosity term dominates the cosmic pressure, and surpasses the pressure contributions from other cosmic matter constitutions, which contradicts the traditional fluid theory. \cite{ni1} \cite{ni2} propose non-standard interaction mechanism to solve this problem. Obviously, it is important to build solid foundation for the research of the viscous cosmology. Equation of state $w<-1$ lies in the phantom region. It is shown that cosmology models with such equation of state possess the so-called the future singularity called the Big Rip \cite{bigrip}. The larger viscosity model parameter space can help to solve the cosmic singularity problem and produces different kinds of evolution mode of the future universe, for more details \cite{sv6}. The rest of this review is organized as follows: In the next chapter, general framework of the viscosity model will be reviewed. In Sec. \textbf{III}, we discuss the modeling of the unified model with viscosity. In this section, two concrete models are analyzed. In Sec. \textbf{IV}, data fitting method is introduced briefly.

In this review, we discuss three aspects of the viscosity model, \begin{itemize} \item General framework for viscosity modeling. General form of Hubble parameter is presented. This general form is convenient for comparing different scale factor(or redshift) dependent viscosity models. \item Two kinds of viscosity models are used to model unified models. \item Observation constraint is necessary for model building. We can see the fitting results are consistent with data. It is prospected that more accurate direct measurements of Hubble constant will provide a new constraint on cosmological parameters \cite{hp}. \end{itemize} Especially we focus on its application on modeling the unified dark energy and dark matter. In the cosmic background level, dynamical analysis can be performed. The statefinder method is useful for discriminating different models \cite{stf1} \cite{stf2} \cite{stf3} \cite{stf4} \cite{stf5}. Compared with $\Lambda$CDM model, evolution of the statefinder of the viscosity model is different and can be discriminated easily, more details can be found in \cite{stfv1} \cite{sv6}. More plentiful and accurate data will improve the power of the statefinder method, which will give enough constraint on the late universe model. We review the viscosity model which is on the level of zero order. The perturbation analysis and the large scale structure are especially useful for the model building. The model predictions need to be consistent with CMB and LSS data. Some works has investigated the perturbation aspects of the viscosity model \cite{p1} \cite{v4}. After corresponding the model parameters, the viscosity model has the connection with the Chaplygin gas model. Though the Chaplygin gas model can fit the SNe Ia data well, in the perturbation level it is found the Chaplygin gas model does not behave in a satisfactory way. Whether the viscosity models could behave well needs further investigation.

1012.3045

The bulk viscosity is introduced to model unified dark matter. The viscous unified model assumes the universe is filled with a single fluid with the bulk viscosity. We review the general framework of the viscous cosmology. The Hubble parameter has a direct connection with the bulk viscosity coefficient. For concrete form of the bulk viscosity, the Hubble parameter which has the scaling relation with the redshift can be obtained. We discuss two viscosity models and the cosmological evolution to which they lead. Using SNe Ia data, the viscosity model can be fitted. We briefly review the fitting method here.

1964311

2011ApJ...730....7T

1012

1012.1871_arXiv.txt

\label{sec:intro} The local Universe provides ample evidence for the existence of Super-Massive Black Holes (SMBHs) in the centers of most galaxies. Typical masses are in the range $\mbh\sim10^6-10^9\, \Msun$, with few exceptionally massive objects reaching $\sim10^{10}\, \Msun$. As first argued by Soltan (1982), the total local BH mass is consistent with the total radiation emitted by accreting SMBHs in active galactic nuclei (AGNs). The accumulation of mass onto SMBHs can be traced back through cosmic history by analyzing the redshift-dependent quasar luminosity function. Such studies suggest that the peak epoch of SMBH growth was at $z\sim2-3$ (e.g., Miyaji \et 2001; Hasinger \et 2005; Silverman \et 2008; Croom \et 2009). However, this statistical approach does not provide sufficient information about the mass of individual SMBHs at various redshifts. A more detailed evolutionary study requires such measurements, in combination with a reliable estimate of the bolometric luminosity (\lbol) and hence the light-to-mass ratio or, equivalently, the normalized accretion rate $\lledd\equiv\lbol/L_{\rm Edd}$.% Several recent studies suggest that the more massive BHs experience most of their growth at very early epochs ($z\gtsim3$), while those observed as AGNs in the the local Universe tend to have lower \mbh\ (``downsizing''; e.g., Marconi \et 2004; Shankar \et 2009). A similar effect is suggested for the distributions of \lledd, such that at $z\gtsim2$ most AGNs accrete close to their Eddington limit (Merloni 2004; Shankar \et 2009). It is also predicted that there should be an anticorrelation between \mbh\ and \lledd\ at all redshifts. These trends are in general agreement with observations (e.g., McLure \& Dunlop 2004; Netzer \& Trakhtenbrot 2007; Shen \et 2008). Detailed simulations of galaxy mergers predict that the instantaneous AGN luminosity and SMBH accretion rate vary greatly on very short timescales (Di Matteo \et 2005; Hopkins \et 2006; Sijacki \et 2007). The overall BH accretion period following major mergers lasts $\sim1\, {\rm Gyr}$, out of which the central source would appear as a luminous, unobscured AGN for at most a $few \times\,100\, {\rm Myr}$. \mbh\ may grow by as much as a factor $\sim1,000$ during such mergers. The merger history of SMBH hosts can also be traced in cosmological simulations of structure formation, by following halo merger trees (e.g., Volonteri et 2003). Almost all these models require (or assume) that at $z\gtsim3$ most AGN would accrete close to, or indeed at their Eddington limit. The fast growth has to last almost continuously from very early epochs ($z\sim20$) and involve massive seed BHs ($\mseed\gtsim10^3\,\Msun$), to account for the very massive BHs detected at $z\sim6$ (e.g., Fan \et 2006; Volonteri 2010, and references therein). The combination of structure formation models predictions at $z\sim3-6$ with the recently observed high clustering of high-redshift luminous AGNs (Shen \et 2007) suggests that the typical duty cycles $-$ the fraction of the total time involving fast accretion $-$ should remain above 0.5 and probably reach unity (e.g. White \et 2008; Wyithe \& Loeb 2009; Shen \et 2010; Shankar \et 2010a; Bonoli \et 2010). These studies use various prescriptions to link AGNs with their dark matter halos. At lower redshifts (i.e. $z\ltsim2$), the accretion can be more episodic with a typical duty cycle between $\sim10^{-3}$ and $\sim0.1$ (e.g., Marconi \et 2004; Merloni 2004; Shankar \et 2009) A comprehensive, up-to-date review of many of these issues is given in Shankar (2009). In order to test these scenarios, \mbh\ and \lledd\ ought to be measured in large, representative samples. For unobscured, type-I AGNs, this is usually achieved by using ``single epoch'' (or ``virial'') \mbh\ determination methods, which are based on the results of long-term reverberation mapping campaigns. These methods are based on estimating the size of the broad line region (BLR) and involve an empirical relation of the form $R_{\rm BLR} \propto \left(\lambda L_{\lambda}\right)^{\alpha}$, where $\lambda L_{\lambda}$ is the monochromatic luminosity in a certain waveband. Combining $R_{\rm BLR}$ with the assumption of virialized motion of the BLR gas, we get $\mbh=f G^{-1} L^{\alpha} V_{\rm BLR}^2$, where $f$ is a geometrical factor of order unity (e.g., Kaspi \et 2000; Vestergaard \& Peterson 2006; Bentz \et 2009). The effect of radiation pressure force on such estimates is still a matter of some discussion (Marconi \et 2008), but recent work suggests it is not very important (Netzer 2009a; Netzer \& Marziani 2010). Single epoch mass estimate methods based on the \hb\ and \MgII\ lines (e.g., Kaspi \et 2005 and McLure \& Dunlop 2004, respectively) were used to estimate \mbh\ up to $z\simeq2$ in large optical surveys (e.g., Corbett \et 2003; McLure \& Dunlop 2004; Netzer \& Trakhtenbrot 2007; Fine \et 2008; Shen \et 2008). Much smaller samples of $z>2$ sources were studied by observing \hb\ or \mgii\ in one of the NIR bands (Shemmer \et 2004, hereafter S04; Kurk \et 2007, hereafter K07; Netzer \et 2007, hereafter N07; Marziani \et 2009; Willott \et 2010, hereafter W10). \mbh\ can also be estimated from the broad \CIV\ line, using specifically calibrated relations (e.g., Vestergaard \& Peterson 2006). This would potentially enable the study of large samples of AGN at high redshifts. However, there is clear evidence that \civ-based estimates of \mbh\ are unreliable. Baskin \& Laor (2005) found that the \civ\ line is often blue-shifted with respect to the AGN rest-frame, which suggests that the dynamics of the \civ-emitting gas may be dominated by non-virial motion. Several studies of large samples clearly demonstrate that the relation between the widths of the \civ\ line and of the lower ionization lines (\hb\ and \mgii) is weak and shows considerable scatter (e.g., Shen \et 2008; Fine \et 2010) that is inconsistent with the virial assumption used in such mass estimators. Finally, our own study (N07) of luminous \znetprev\ AGNs shows a complete lack of correlation between \civ-based and \hb-based estimates of \mbh. These discrepancies become crucial at high redshifts and large \mbh\ and lead to the conclusion that only \hb-based and \mgii-based mass estimates are reliable enough to infer the properties of such sources. Our previous project (S04 \& N07) presented the largest sample of \znetprev\ type-I AGNs for which \mbh\ and \lledd\ were reliably measured using NIR \hb\ spectroscopy. The distribution of \lledd\ at \znetprev\ was found to be broad, and about half of the sources had $\lledd<0.2$, inconsistent with several of the models mentioned above. In particular, the typically low accretion rates and the very high masses (up to $\lmbh\simeq10.5$) also mean that $\sim60\%$ of the sources did not have enough time to grow to the observed \mbh\ by continuous accretion at the observed rates. These findings suggest that an epoch of faster SMBH growth must have occurred, for most objects, at $z>3.5$. In order to probe such redshifts, the \mgii\ line must be observed in either the $H$ or the $K$-bands. Practically, this corresponds to focusing on \zfpe\ or \zsix\ sources. In this paper we present a systematic study of \mbh\ and \lledd\ in a large, well-defined sample of \zfpe\ type-I AGNs. This is based on new \hband\ spectroscopic observations, which enable the measurement of the \mgii\ line. We describe the sample selection and the observations in \S\ref{sec:sample_obs} and the way we deduced \mbh\ and \lledd\ in \S\ref{sec:fit_mbh_lledd}. The main results of these measurements are presented in \S\ref{sec:results} and discussed in \S\ref{sec:discussion} where we compare these results to those of other high-redshift samples. The main findings are summarized in \S\ref{sec:summary}. Throughout this work we assume a standard $\Lambda${\sc CDM} cosmology with $\Omega_{\Lambda}=0.7$, $\Omega_{M}=0.3$ and $H_{0}=70$\,\kms\,Mpc$^{-1}$.

\label{sec:discussion} \subsection{Comparison with other studies} \label{sec_sub:dis_compare} Our main goal is to identify the epoch at which most SMBHs experienced the first episode of fast growth. We focus on a comparison of our \zfpe\ sample to several $z>2$ samples with reliably measured \mbh\ and \lledd. In S04 and N07 we studied a sample of 44 \znetprev\ type-I AGNs using NIR spectroscopy to measure \Lop\ and FWHM(\hb). The study showed that a large fraction of the high \mbh\ sources accrete at a rate that is well below the Eddington limit. Assuming these sources accrete at constant \lledd, such accretion rates cannot explain their measured masses and are in contradiction with several theoretical predictions. Having obtained new data on the \zfpe\ sample, we can now compare several groups of AGNs, at several epochs, in an attempt to follow their growth all the way from \zsix\ to \ztpf. The samples we consider here are the new one at \zfpe\ and our earlier (S04 \& N07) samples at \znetprev. We also consider a small number of sources at \zsix\ from the samples studied by K07 and W10, which have 5 and 9 reliable \mbh\ \& \lledd\ measurements, respectively. These small samples do \textit{not} have a common flux limit and simply represent the up-to-date collection of the sources discovered at those redshifts that were also observed in follow-up NIR spectroscopy. As such, they are not representative of the AGN population at \zsix. This, along with the small size of the samples, limits their statistical usefulness. To use the data from the \zsix\ samples, we re-calculated \mbh\ and \lbol\ using the methods described above. First, we corrected the published FWHM(\mgii) according to the assumption of only one of the doublet components. This reduces \fwmg\ by a mean factor of $\sim1.2$ and the reported \mbh\ by a mean factor of $\sim1.44$. Second, the smaller \fboluv\ adopted here reduces the reported \lbol\ by a factor of $\sim1.49$. Regarding \lledd, the two corrections almost completely cancel out. \begin{figure} \includegraphics[width=8.5cm]{props_vs_z_z_2_6_20101107.eps} \caption{\lbol, \mbh\ and \lledd\ vs. redshift, for samples of different redshifts discussed in the text: the new \zfpe\ sample presented here (black squares), the \znetprev\ samples of S04 and N07 (magenta and blue triangles) and the combined \zsix\ sample from K07 and W10 (red circles).} \label{fig:all_z_props} \end{figure} \begin{figure} \includegraphics[width=8.5cm]{Mbh_cdf_z_2_6_20101107.eps} \includegraphics[width=8.5cm]{LLedd_cdf_z_2_6_20101107.eps} \caption{The cumulative distribution functions (CDFs) of \mbh\ (top) and \lledd\ (bottom) for the samples discussed in the text.} \label{fig:all_z_mbh_lledd_cdf} \end{figure} In Figure~\ref{fig:all_z_props} we present \lbol, \mbh\ and \lledd\ for the new \zfpe\ sample, the \znetprev\ samples of S04 and N07, and the combined sample of \zsix\ AGNs. The four samples suggest an increasing \mbh\ and decreasing \lledd\ with cosmic time. This is better illustrated in Figure \ref{fig:all_z_mbh_lledd_cdf}, where we present the cumulative distribution functions (CDFs) of \mbh\ and \lledd. The increase in the median \mbh\ and the decrease in the median \lledd\ as a function of redshift are evident. In particular, there is a clear shift of $\sim0.5$ dex between the median \mbh\ values of the \zfpe\ and \ztpt\ samples, in the sense of a lower \mbh\ at \zfpe. There is also an opposite shift between the median \lledd\ values. The differences between the \ztpt\ and the \ztpf\ samples are much smaller, although the \ztpf\ sample includes several sources with extremely massive BHs ($\lmbh>10$) which are not observed at earlier epochs. Regarding \mbh, only $\sim14\%$ of the \ztpt\ sample (2 sources) lie below the median value of the \zfpe\ sample ($\lmbh=8.92$). Similarly, only $\sim13\%$ of the \zfpe\ sample (5 sources) lie above the median value of the \ztpt\ sample ($\lmbh=9.37$) and even less ($\sim6\%$; 2 sources) above the median value of the \ztpf\ sample ($\lmbh=9.59$). We have further tested the significance of these differences by performing a series of two-sample Kolmogorov-Smirnov tests. The null hypothesis that the observed distributions of \lledd\ at \zfpe\ and at \ztpt\ (or at \ztpf) are drawn from the same parent distribution is rejected with significance levels $>99\%$. A similar test was applied to the \zfpe\ and \zsix\ samples, and could not reject the null hypothesis, suggesting that the distributions of \lledd\ at \zfpe\ and \zsix\ are statistically similar. We obtain similar results when comparing the distribution of \mbh\ at \zfpe\ to that of the \znetprev\ samples. Due to the large fraction of \zsix\ sources with $\lmbh\ltsim8.5$, we can reject the hypothesis that the \mbh\ values at \zfpe\ and \zsix\ are drawn from the same distribution. Some of the above results may be biased by the fact that the four samples cover a different range of \lbol, which originate from the different target selection criteria. We thus repeated the statistical tests aforementioned, focusing on subsamples which share a common range of $46.4 < \llbol < 47.4$ (i.e., matching the \lbol\ range of the \zfpe\ sample). All but one comparison result in the same conclusions, with similar confidence levels. The one exception are the distributions of \mbh\ at \zfpe\ and \zsix, for which the null hypothesis now \textit{cannot} be rejected, suggesting that these distributions represent the same parent population. This is not surprising given the lower luminosities of most of the \zsix\ sources, the dependence of \mbh\ on source luminosity, and the incompleteness of the \zsix\ sample. We conclude that there is strong evidence for a rise in \mbh\ and a drop in \lledd\ of about a factor of 2.8 between \zfpe\ and \ztpt. This strong trend is \textit{not} observed with respect to neither lower nor higher-redshift samples, which span similar periods of time \footnote{For the adopted cosmology, the physical time between $z=4.8$ and $z=3.3$ is $\sim680\,{\rm Myr}$ and between $z=3.3$ and $z=2.4$ is $\sim790\,{\rm Myr}$.}. It thus seems that the most massive BHs, associated with the most luminous AGNs, started the episode of fast BH growth at redshifts above about $z\sim5$. By $z\sim2-3$, these SMBHs reach their peak (final) mass ($\>10^{10}\,\Msun$) and their mass accretion is less efficient. \subsection{The growth of active SMBHs from \zfpe\ to \ztpt\ and \ztpf} \label{sec_sub:dis_z65_z2} Given the trends and differences in \mbh\ and \lledd, as well as the short $e$-folding times of the \zfpe\ AGNs (\S\S\ref{sec_sub:lifetimes}), it is possible to think of the \zfpe, \ztpt\ and \ztpf\ samples as representing different evolutionary stages of the same parent population of SMBHs. In what follows, we focus on the \zfpe\ and \ztpt\ samples (separated by $\sim 680\, {\rm Myr}$) to test this evolutionary interpretation, assuming various scenarios. We consider two growth scenarios: constant accretion rate (i.e. constant \lbol) and constant \lledd \footnote{There are, of course, many other possible scenarios involving, for example, host related evolution. These are beyond the scope of the present paper.}. As explained, the assumption of constant \lledd\ results in an exponential growth of \mbh\ and given our chosen value of $\eta$, the mass growth $e$-folding time is $\tau\simeq45\left(\lledd\right)^{-1}\,{\rm Myr}$, which translates to $\tau\ltsim 240\,{\rm Myr}$ ($\lledd>0.18$) for the \zfpe\ sources. Thus, even the lowest \lledd\ SMBHs at \zfpe\ could have increased their \mbh\ by a factor of $\sim20$ by \ztpt, which is much larger than the typical difference in \mbh\ between the two samples. The assumption of a constant \lbol\ results in much slower growth. The luminosities of the \zfpe\ sources translate to $4\ltsim \dot{M}_{\rm BH} \ltsim 37 \, \mpyr$. Even this much slower growth scenario results in very high \mbh\ at \ztpt. This scenario also produces very low accretion rates ($\lledd<0.1$) at \ztpt. In Figure \ref{fig:m_growth_for} we present evolutionary tracks for our \zfpe\ AGNs to \ztpt\ and eventually to $z=2$. Clearly, continuous constant \lbol\ growth results in too large \mbh\ and cannot reproduce the lower \mbh\ sources at \znetprev. Specifically, the calculated distribution of \mbh\ for the \zfpe\ sources, when evolved to \ztpt, has a median value of $\lmbh\simeq10$, larger by $\sim0.6$ dex. than the \textit{observed} median of the \ztpt\ sample. This scenario also fails to reproduce the observed range of \lledd, since the lowest \lledd\ sources at \ztpt\ have $\lledd\gtsim0.1$. As the dotted evolutionary tracks in Fig~\ref{fig:m_growth_for} demonstrate, the constant \lledd\ scenario produces even larger masses, and is only feasible if all the \zfpe\ sources cease their accretion shortly after their observed active phase. This scenario assumes a constant \lledd, hence the significant difference between the \lledd\ distributions at \zfpe\ and at \ztpt\ ($\sim0.5$ dex ; see Fig~\ref{fig:all_z_mbh_lledd_cdf}) is not resolved. We also note that in both scenarios the \zfpe\ sources would have been easily observed at \ztpt, since they would have \lbol\ which is either similar (linear growth) or much higher than (exponential growth) the sources observed by S04 and N07 at \znetprev. From these two simplistic growth scenarios we conclude that the fast growth of the \zfpe\ sources must either experience a shut down before \ztpt, or accrete in several short episodes, with duty cycles which are much smaller than unity. \begin{figure} \includegraphics[width=8.5cm]{Mbh_t_growth_for_20101107.eps} \caption{Evolution scenarios for the \zfpe\ SMBHs. Symbols are identical to those in Fig.~\ref{fig:all_z_props}. Solid lines describe the growth of \mbh\ under the assumption of a constant \lbol, while dotted lines represent the constant \lledd\ scenario.} \label{fig:m_growth_for} \end{figure} To further constrain the evolution of the observed SMBHs, we ran a series of calculations with different duty cycles for each of the two evolutionary scenarios. In each calculation we assembled the distribution of calculated \mbh\ at exactly $z=3.3$ and $z=2.4$. Several of these distributions are shown in Figure \ref{fig:mbh_cdf_evolve_for}. For the fast, constant \lledd\ scenario, the only calculated distributions which resemble the observed distribution of \mbh\ at \ztpt\ are those with a duty cycles in the range $7.5-12.5\%$. The observed distribution at \ztpf, on the other hand, can only be achieved by duty cycles of $5-7.5\%$. For the constant \lbol\ scenario, the \ztpt\ distribution can be matched by assuming a duty cycle in the range of $15-25\%$, while the \ztpf\ distributions can be partially explained by assuming duty cycles of $\sim5-25\%$. The ranges of duty cycles which account for the observed distribution of \lledd\ are somewhat different: $10-15\%$ for the \lledd\ distribution at \ztpt\ and $5-10\%$ for the one at \ztpf. All the above calculations assumed $\eta=0.1$. Naturally, lower (higher) radiative efficiencies will require shorter (longer) duty cycles, to reproduce the same calculated CDFs at \znetprev. This means that for any assumed $\eta$ in the range $0.05\leq\eta\leq0.3$ the aforementioned ``acceptable'' duty cycles can change by up to a factor of $\sim2$, where the exact factor scales as $\eta /(1- \eta)$. We note that Fig.~\ref{fig:mbh_cdf_evolve_for} suggests that, in both evolution scenarios, the more massive BHs at \znetprev\ seem to have grown at higher duty cycles than the less massive ones. Alternatively, this can be interpreted as an increase in radiative efficiency with increasing resultant \mbh\ at \znetprev. In particular, the observed distributions of \mbh\ and \lledd\ at \ztpf\ are much more complex than the calculated ones and no single, fixed duty cycle can account for the shape of the observed distributions. The reason for this apparent discrepancy might also be the way the \ztpf\ sources were selected in the S04 and N07 studies. \begin{figure} \includegraphics[width=8.5cm]{M_cdf_growth_dc_4panels_20101208.eps} \caption{Observed and calculated \mbh\ CDFs for different evolutionary scenarios. In all panels solid black lines show the observed \mbh\ CDF at \zfpe, while magenta lines (left panels) and blue lines (right panels) show the CDFs at \ztpt\ and at \ztpf, respectively. Dashed black lines are the calculated CDF of the \zfpe\ sample assuming different growth scenarios and duty cycles. Top panels assume a constant \lbol\ scenario, and duty cycles of 5, 20, and 50\% (from left to right). Bottom panels assume a constant \lledd\ scenario, and duty cycles of 5, 10 and 15\% (from left to right). The dotted lines in the bottom panels illustrate how a duty cycle of 50\% produces extremely over-massive BHs at \ztpt\ and \ztpf. All the calculations assume $\eta=0.1$. Assuming the extreme cases of $\eta=0.05$ (or $0.3$) would mean that the plotted calculated CDFs correspond to duty cycles which are a factor of $\sim2$ smaller (or larger; see text for details). } \label{fig:mbh_cdf_evolve_for} \end{figure} To illustrate how the above duty cycles facilitate an evolutionary connection between the three samples, we present in Figure \ref{fig:mbh_t_duty_cycle} the evolutionary tracks of the \zfpe\ sample, similar to Fig~\ref{fig:m_growth_for}, but this time assuming duty cycles of 10\% and 20\%, for the constant \lledd\ and constant \lbol\ scenarios, respectively and assuming again $\eta=0.1$. We also verified that evolving the \mbh\ of the \znetprev\ sources \textit{backwards}, under the assumption of constant \lledd\ and a duty cycle of $10\%$ results in a distribution of \mbh\ which is not critically different than the one directly observed at \zfpe. However, we again find that the few highest \mbh\ sources at \ztpf\ probably require higher duty cycles. \begin{figure} \includegraphics[width=8.5cm]{Mbh_t_growth_for_dc_20101107.eps} \caption{Evolution scenarios for the \zfpe\ SMBHs with various duty cycles. Symbols are identical to those in Fig.~\ref{fig:all_z_props}. Solid lines describe \mbh\ growth under the assumption of constant \lbol\ and a duty cycle of 20\%. Dotted lines represent the constant \lledd\ scenario and a duty cycle of 10\%.} \label{fig:mbh_t_duty_cycle} \end{figure} We conclude that all the observed measurements of \mbh, \lledd\ and \lbol\ are consistent with duty cycles of about 10-20\%. Considerably longer duty cycles can be consistent with observations only if we assume an extremely high radiative efficiency ($\eta\simeq0.3$) for all the \zfpe\ sources. Such low duty cycles are in good agreement with the models presented in Shankar \et (2009), which are able to reproduce the growth of a similar population of SMBHs, in terms of redshift and range of \mbh\ (c.f. their Fig.~7).. Our constrains on duty cycles are, however, in contrast with several other studies that suggest that the duty cycle at high redshift should reach unity, based on the clustering of high-redshift AGNs (e.g. White \et 2008; Wyithe \& Loeb 2009; Bonoli et al. 2010, and references therein). All this leads us to suggest that for the most massive $z>2$ type-I AGNs, an epoch of fast SMBH growth took place before $z\sim4$, which we partially observe at \zfpe. This epoch is efficient enough to produce the very massive ($\lmbh\gtsim10$) BHs observed at \ztpf. The fast growth slows down before \ztpt, and the sources observed at \ztpt\ show much lower \lledd. In addition, there is no significant rise in \mbh\ between \ztpt\ and \ztpf. The growth of \mbh\ during this epoch can only be explained by assuming that mass accretion proceeded in short episodes, lasting an order $10-20\%$ of the total period. This translates to accretion episodes which (cumulatively) last $\sim70-140\, {\rm Myr}$ ( $\sim150-300\, {\rm Myr}$) over the period between \zfpe\ and \ztpt\ (\zfpe\ and \ztpf), respectively. Another possibility is that the \zfpe\ SMBHs have experienced an even faster shut down, and so at \ztpt\ their inactive relics have masses which do not differ significantly from those we observe at \zfpe. Major mergers between massive, gas-rich galaxies are capable of supplying large amounts of cold gas directly to the innermost regions, to be accreted by the central SMBHs. Detailed simulations of such mergers suggest that these events can fuel significant BH growth over periods of $\sim1\, {\rm Gyr}$. The mass accretion rate and source luminosity during such mergers may vary on time-scales of $few \times 10\, {\rm Myr}$ (Di Matteo \et 2005; Hopkins \et 2006; Hopkins \& Hernquist 2009). According to Hopkins \et (2006), AGNs would appear to have $\lbol\gtsim2.7\times10^{46}\,\ergs$ (i.e., matching the range we observe at \zfpe) for typically $\ltsim100\,{\rm Myr}$. Numerical studies suggest that massive dark matter halos may undergo more than one major mergers per ${\rm Gyr}$ at \zfpe\ (e.g. Genel \et 2009 and references therein). Thus, during the $\sim680\, {\rm Myr}$ between \zfpe\ and \ztpt, the hosts of the \zfpe\ SMBHs may have undergone a single merger, during which the period of significant accretion by the SMBHs would last for $\ltsim100\, {\rm Myr}$. This is in good agreement with the duty cycles of $10-20\%$ we find here, which correspond to $70-140\,{\rm Myr}$. The decline in activity observed for the most massive BHs at \ztpt\ and \ztpf\ may then be associated with the decline in the major merger rate, which drops by a factor of $\sim4$ between \zfpe\ and \ztpf\ (Genel \et 2009). However, it is likely that the hosts of the \ztpf\ SMBHs have experienced an additional major merger during the $\sim790\,{\rm Myr}$ between \ztpt\ and \ztpf, especially if these SMBHs reside in the more massive dark matter halos. In such a scenario, the largest BHs at \ztpf\ may have gathered their high mass during more efficient (i.e., higher duty-cycle) accretion episodes. This requirement of higher duty-cycles for higher \mbh\ sources is also reflected in our analysis (see upper-right panel of Fig.~\ref{fig:mbh_cdf_evolve_for}). If major mergers are indeed the main drivers of SMBH accretion history at $z\sim2-5$, the results presented here predict that the host galaxies of the \zfpe\ sources would be found in the early stages of major mergers, while the hosts of \znetprev\ sources would appear to be in either later stages, or in merging systems which have a lower mass ratio. There are several other mechanisms to make large amounts of cold gas available for BH accretion. Most of these might also trigger intense star formation (SF) activity in the hosts of SMBHs. Indeed, many high-luminosity AGNs show evidence for intense SF (e.g. Netzer 2009b; Lutz \et 2010, and references therein). New observations by \textit{Herschel} suggest that SF in the hosts of the most luminous AGN peaks at $z\sim3$ and quickly decreases at later epochs (Serjeant \et 2010). If correct, it will indicate that the amount of gas available for both BH accretion and SF has depleted by $z\sim2-3$, consistent with the decrease in SMBH accretion activity we find here. Future \textit{Herschel} and ALMA observations of the hosts of the \zfpe\ AGNs may be able to reveal the presence of such SF activity, and perhaps the amount and dynamical state of the cold gas. This will enable a better understanding of the galaxy-scale processes which drive the accretion history of the fast-growing \zfpe\ SMBHs. We present new \hband\ spectroscopy for a flux-limited sample of \Ntot type-I SDSS AGNs at \zfpe. The sample covers $\sim1/4$ of all the (spectroscopically observed) SDSS sources at that redshift band, and thus about $\sim1/20$ of the total population of $\lbol > 2.75\times10^{46} \ergs$ sources (over the entire sky). The main results of our study are: \begin{enumerate} \item The \zfpe\ AGNs have, on average, higher accretion rates and lower massees than those observed at lower redshifts. The accretion rates and masses are comparable to those of the small, incomplete samples of \zsix\ AGNs. \item We have observed an epoch of fast SMBH growth, probably the very first such phase for most SMBHs. Assuming continuous growth from about $z=20$, these observations provide the very first look at the distribution of seed BHs in the early universe. About 65\% of the SMBHs at \zfpe\ have had enough time to grow to their observed \mbh, assuming continuous accretion at the observed \lledd. About $40\%$ of the sources could have started their growth from BH seeds which are stellar remnants ($\mseed < 100\,\Msun$). For the minority the sources, those with small \lledd, there might have been an ever earlier epoch of faster accretion or, perhaps, they started their growth from a much larger seeds ($\mseed\gtsim10^{6}\, \Msun$). \item The \zfpe\ sources can be regarded as the progenitor population of the most massive ($\mbh\gtsim10^{10}\,\Msun$) BHs observed at \znetprev. The growth rate of those massive BHs seems to be much slower between \ztpt\ and \ztpf. \item A comparison of the observed distributions of \mbh\ at \zfpe, \ztpt\ and \ztpf\ indicates that the \zfpe\ sources either completely stop their accretion shortly after \zfpe, or that their accretion proceeds in relatively short episodes. We find that for mass growth rates that follow either constant \lledd\ or constant \lbol\ scenarios, duty cycles of either $\sim10\%$ or $\sim20\%$, respectively, give reasonable agreement to the observed distributions of \mbh. \end{enumerate}

1012.1871

We present new H-band spectroscopy for a flux-limited sample of 40 z ~= 4.8 active galactic nuclei, selected from the Sloan Digital Sky Survey. The sample probably contains the most massive active black holes (BHs) at this redshift and spans a broad range in bolometric luminosity, 2.7 × 10<SUP>46</SUP> erg s<SUP>-1</SUP> &lt; L <SUB>bol</SUB> &lt; 2.4 × 10<SUP>47</SUP> erg s<SUP>-1</SUP>. The high-quality observations and the accurate fitting of the Mg II λ2798 line enable us to study, systematically, the distribution of BH mass (M <SUB>BH</SUB>) and normalized accretion rate (L/L <SUB>Edd</SUB>) at z ~= 4.8. We find that 10<SUP>8</SUP> M <SUB>sun</SUB> &lt;~ M <SUB>BH</SUB> &lt;~ 6.6 × 10<SUP>9</SUP> M <SUB>sun</SUB> with a median of ~8.4 × 10<SUP>8</SUP> M <SUB>sun</SUB>. We also find that 0.2 &lt;~ L/L <SUB>Edd</SUB> &lt;~ 3.9 with a median of ~0.6. Most of these sources had enough time to grow to their observed mass at z ~= 4.8 from z = 20, assuming a range of seed BH masses, with ~40% that are small enough to be stellar remnants. Compared to previously studied samples at z ~= 2.4 and sime3.3, the masses of the z ~= 4.8 BHs are typically lower by ~0.5 dex and their L/L <SUB>Edd</SUB> is higher by a similar factor. The new z ~= 4.8 sample can be considered as the progenitor population of the most massive BHs at z ~= 2.4 and sime3.3. Such an evolutionary interpretation requires that the growth of the BHs from z ~= 4.8 to z ~= 3.3 and z ~= 2.4 proceeds with short duty cycles, of about 10%-20%, depending on the particular growth scenario. <P />Based on observations collected at the European Organisation for Astronomical Research in the Southern Hemisphere, Chile, as part of programs 081.B-0549, 082.B-0520, and 085.B-0863, and at the Gemini Observatory, as part of programs GN-2007B-Q-56 and GN-2008B-Q-75.

1758907

2011MNRAS.412L..63B

1012

1012.5254_arXiv.txt

Among the radio pulsars that are found in binary systems, four belong to a subgroup where the pulsar has the characteristics of a young, isolated pulsar, but is orbiting a massive companion ($\ga3$\,M$_\odot$) in a rather wide (50--2000\,day orbital period) and eccentric ($e>0.5$) orbit \citep{sta04,kbjj05}. The companion to the first of these systems, PSR\,B1259$-$63 \citep{jml+92} was identified as a $V=10$ early-type B2e star. PSR\,J0045$-$7319, located in the SMC, is the second system and has a $V=16$ B1\,V companion \citep{kjb+94}. No information on a possible counterpart has been published for the fourth system, PSR\,J1638$-$4725 \citep{lfl+06}. The third system, PSR J1740$-$3052, was discovered by \citet{sml+01} in the Parkes Multibeam Survey \citep{mlc+01} as a young (350\,kyr) pulsar. Subsequent timing observations showed that the pulsar is part of a binary system with an orbital period of 231\,d and an eccentricity of 0.57. The mass function of the pulsar yielded a lower limit to the mass of the companion of 11\,M$_\odot$. A candidate counterpart to PSR\,J1740$-$3052 was identified in near-infrared observations. From $K$-band spectra, the counterpart was classified as late-type, having a spectral type between K5 and M3 \citep{sml+01}. This identification is at odds with the binary parameters which set the Roche lobe radius significantly smaller than that expected for a late-type star at the estimated distance. Subsequent phase-resolved spectroscopic observations by \citet{tsw+10} confirmed the unlikeliness of the star being the binary companion, as no significant radial velocity variations consistent with the expected binary orbit were found. The nature of the massive binary companion remains unclear, and could be either an early-type main-sequence star or a compact remnant. Though \citet{sml+01} reported small variations in the dispersion measure of PSR\,J1740$-$3052 near periastron, which would be consistent with a stellar companion, they could not conclusively rule out a stellar mass black hole as the binary companion. In this paper we report on interferometric radio observations to obtain an accurate position of the PSR\,J1740$-$3052 and adaptive optics corrected near-infrared observations to search for a stellar counterpart to PSR\,J1740$-$3052. Section\,2 describes the interferometric observations, while Section\,3 details the astrometry and photometry of the near-infrared imaging. We discuss the results in Section\,4 and conclude in Section\,5.

Since its discovery, it was unclear if the binary companion of PSR\,J1740$-$3052 was an early-type main-sequence star, a late-type giant, or a stellar mass black hole, though the data favored a main-sequence type companion. Using accurate astrometric radio interferometry and adaptive optics corrected near-infrared imaging observations, we show that the late-type star located near the pulsar position is located at the edge of the 95\% confidence error ellipse of the interferometric radio position of the pulsar. This further strengthens the case made by \citet{sml+01} and \citet{tsw+10} that it is not the binary companion to PSR\,J1740$-$3052. In the wings of the late-type star, and near the center of the error ellipse, we find a counterpart whose observed $K_\mathrm{s}$-band magnitude is consistent with the expected mass and age of the binary companion to PSR\,J1740$-$3052 at the estimated distance and reddening to the system. We argue that this counterpart is the binary companion to PSR\,J1740$-$3052. Our observations show that the companion is on the main-sequence, which is consistent with the observations of both the small variations in the dispersion measure near periastron of the binary orbit, and the periastron advance being due to tidal and spin quadrupoles \citep{sml+01}. These observations rule out the need for a black hole or a giant as the binary companion. At $K_\mathrm{s}=15.87$ the companion is within reach of near-infrared spectrographs. Adaptive optics would be a necessity to resolve the companion in the wings of the late-type star which is almost 6\,mag brighter in $K_\mathrm{s}$. To extract the spectrum of the companion an algorithm such as that of \citet{hyn02} would be needed to resolve the blend. Spectra of the companion could be used to further confirm the identification through radial velocity variations and the spectral classification. With a massive main-sequence companion, the PSR\,J1740$-$3052 system likely experienced a previous phase of mass transfer in which the progenitor of the pulsar lost its hydrogen envelope and the companion accreted a significant amount of mass (for a review, see \citealt{th06}). The subsequent (type Ib) supernova caused the orbit to become eccentric. In the next few million years the companion will evolve off the main-sequence and tidal forces will circularize the orbit when the star approaches Roche lobe overflow. The ensuing common envelope phase will likely lead to a merger and the formation of a Thorne-Zytkow object \citep{tz75} as there is not enough energy in the orbit to unbind the envelope of the subgiant (e.g.\,\citealt{dt10}).

1012.5254

We report on the identification of a near-infrared counterpart to the massive (&gt;11 M<SUB>⊙</SUB>) binary companion of pulsar J1740-3052. An accurate celestial position of PSR J1740-3052 is determined from interferometric radio observations. Adaptive optics corrected near-infrared imaging observations show a counterpart at the interferometric position of the pulsar. The counterpart has K<SUB>s</SUB>= 15.87 ± 0.10 and J-K<SUB>s</SUB> &gt; 0.83. Based on distance and absorption estimates from models of the Galactic electron and dust distributions, these observed magnitudes are consistent with those of a main-sequence star as the binary companion. We argue that this counterpart is the binary companion to PSR J1740-3052 and thus rule out a stellar mass black hole as the pulsar companion.

12201943

2011ASPC..448...61M

1012

1012.2642_arXiv.txt

The open cluster $\eta$ Chamaeleontis is one of the closest ($d\sim94$~pc) and youngest ($t\sim8$~Myr) stellar aggregates in the Solar neighbourhood. A census of its stellar population currently stands at 18 systems, covering spectral types B8--M5.5 \citep{Mamajek99,Mamajek00,Lawson02,Song04b,Luhman04a,Lyo04,Murphy10}. At high and intermediate masses the cluster Initial Mass Function (IMF) follows that of other star-forming regions and young stellar groups, but there is a clear deficit of members at masses $<$0.15~$M_{\odot}$. Comparing the observed mass function of the cluster to other young groups, \citet{Lyo04} predict that an additional 20 stars and brown dwarfs in the mass range $0.025<M<0.15~M_{\odot}$ remain to be discovered. Efforts to observe this hitherto unseen population have so far failed to find any additional members at either larger radii from the cluster core \citep[][to 1.5~deg, 4 times the radius of known membership]{Luhman04} or to low masses in the cluster core \citep[][to $\sim$13~$M_{\rm Jup}$]{Lyo06}. Failure to find these low--mass members raises a fundamental question: has the cluster's evolution been driven by dynamical interactions which dispersed the stars into a diffuse halo at even larger radii, or does $\eta$~Cha possess an abnormally top-heavy IMF deficient in low--mass objects? The latter result would seemingly be at odds with the growing body of evidence that suggests the IMF is universal and independent of initial star-forming conditions \citep*[for an excellent review see][]{Bastian10}. \citet*{Moraux07} have attempted to model the observed properties of $\eta$~Cha using $N$-body simulations of the cluster's dynamical evolution starting with standard initial conditions. They are able to replicate the current configuration of the cluster assuming a log-normal IMF and 30--70 initial members. New calculations incorporating binaries (Becker \& Moraux, 2010, in prep.) show almost identical results can be obtained starting with $\sim$20 binary systems. This suggests the deficit of low--mass objects seen in the present day cluster may not be due to a peculiar IMF but to dynamical evolution instead. These simulations predict there should exist a diffuse halo of cluster ejectees beyond the radius currently surveyed. To test the dynamical evolution hypothesis we have undertaken a detailed search for this putative halo of low-mass objects surrounding $\eta$~Cha. Our survey methods and results are described in \citet{Murphy10}, with an extensive follow-up spectroscopy campaign for two candidates thought to be harboring accretion disks presented in \cite{Murphy10b}. Definitive information can be found in these two papers -- in the following contribution we give a precis of our work to date. \begin{figure}[t] \centering \includegraphics[width=0.7\textwidth]{Murphy_S_fig1} \caption{Selection criteria for the 81 photometric candidates from \citet{Murphy10}. Contours show the cumulative total of stars enclosed. Known KM-type cluster members are shown as filled stars. We select candidates (filled circles) within 1.5~mag of the empirical cluster isochrone, having $i-J>1.5$. Proper motion candidates are denoted by open circles. The 6 intermediate-gravity stars are labeled.} \label{fig:cmd} \end{figure}

1012.2642

We have identified several lithium-rich low-mass (0.08 &lt; M &lt; 0.3 M<SUB>⊙</SUB>). stars within 5.5 deg of the young open cluster η Chamaeleontis, nearly four times the radius of previous search efforts. We propose 4 new probable cluster members and 3 possible members requiring further investigation. Candidates were selected on the basis of DENIS and 2MASS photometry, NOMAD astrometry and extensive follow-up spectroscopy. Several of these stars show substantial variation in their Hα emission line strengths on timescales of days to months, with at least one event attributable to accretion from a circumstellar disk. These findings are consistent with a dynamical origin for the current configuration of the cluster, without the need to invoke an abnormally top-heavy Initial Mass Function, as proposed by some authors.

5042552

2011ASPC..448..683M

1012

1012.5915_arXiv.txt

\label{intro} The origin of the stellar initial mass function (IMF) is one of the major issues in astrophysics. The low-mass end of the IMF, in particular, has been subject of numerous observational and theoretical studies over the past decade (see \citealt{bonnell07}). Star forming regions and young clusters harbor a large population of free-floating objects below the stellar-mass boundary, found to be at $\sim$0.075$\,M_{\odot}$. These objects include brown dwarfs (BDs), but also a population of objects with masses comparable to those of massive planets (we refer to these objects as ``planetary-mass objects'' or PMOs). Despite the large advances in our understanding of substellar objects over the past decade, some of the most important questions still remain unanswered. The origin of BDs and PMOs is not clear \citep{whitworth05}; competing scenarios include turbulent fragmentation \citep{padoan&norlund04}, ejection from multiple systems \citep{bate09}, and ejection from fragmenting protoplanetary disks \citep{stamatellos08}. The shape of the IMF at very low masses is the subject of an ongoing debate in the literature \citep{bonnell07,chabrier03}. In nearby star forming regions the total number BDs relative to the number of low-mass stars varies between 3 and 8 with large uncertainties \citep{andersen08}. The number of PMOs and its dependence on environment is even more uncertain. The IMF could be still rising below 0.015$M_{\odot}$ \citep{caballero07}, or declining in this regime \citep{lucas05}. A cutoff in the mass function has not been observed yet. SONYC -- Substellar Objects in Nearby Young Clusters -- is an ongoing project to provide a complete census of the brown dwarf and planetary mass object population in nearby young clusters, and to establish the frequency of substellar mass objects as a function of cluster environment. The resulting catalog of substellar mass candidates will provide the basis for detailed characterization of their physical properties (disks, binarity, atmospheres, accretion, activity). The primary means of identifying candidates is broad-band imaging in the optical and the infrared, thus aiming to detect the photosphere. The survey is also combined with the 2MASS and Spitzer photometry catalogs. Photometric selection results in large samples of candidates and requires extensive spectroscopic follow-up to asses the real nature of the objects. Our observations are designed to reach limiting masses of $\sim\,$0.005$\,M_{\odot}$, well below the deuterium-burning limit at 0.015$\,M_{\odot}$, and thus require us of 4- to 8-m-class telescopes. By probing several star forming regions we want to probe for environmental differences in the frequency and properties of substellar objects. In this contribution, we summarize the results delivered in the framework of SONYC over the past two years. We have surveyed three star forming regions: NGC~1333 \citep{scholz09}, $\rho$~Ophiuchus \citep{geers10}, and Chamaeleon-I (Mu\v{z}i\'c et al., submitted to ApJ).

It is clear that the census of BDs and PMOs in most star forming regions is still incomplete. Based on the existing data we can conclude that: (a) there are hints of regional differences in the mass function at the very-low-mass end, and (b) only a combination of different search techniques can provide a robust picture of the substellar population.

1012.5915

The origin of the lowest mass free-floating objects - brown dwarfs and planetary-mass objects - is one of the major unsolved problems in star formation. Establishing a census of young substellar objects is a fundamental prerequisite for distinguishing between competing theoretical scenarios. Such a census allows us to probe the initial mass function (IMF), binary statistics, and properties of accretion disks. Our SONYC (Substellar Objects in Nearby Young Clusters) survey relies on extremely deep wide-field optical and near-infrared imaging, with follow-up spectroscopy, in combination with Spitzer photometry to probe the bottom end of the IMF to unprecedented levels. Here we present SONYC results for three different regions: NGC 1333, ρ Ophiuchus and Chamaeleon-I. In NGC 1333, we find evidence for a possible cutoff in the mass function at 10-20 Jupiter masses. In ρ Oph we report a new brown dwarf with a mass close to the deuterium-burning limit.

942277

2011A&A...526A..87Z

1012

1012.1707_arXiv.txt

\label{intr} \subsection{Review of observational approaches} \label{rooa} One of the most enduring unknowns in stellar physics has been the inner distribution of the angular momentum in a star. In the past few decades, significant progress has been made in describing theoretically the evolution of rotating stars. This has required an understanding of numerous hydrodynamic and magnetic instabilities triggered by the rotation, as well as the mixing processes of chemical elements unleashed by these instabilities \citep{tass78,tass00,zahn83, zahn92,iau215,mae09}. However, apart from the Sun, reliable observational information about the internal rotation of stars remains scarce or non-existent.\par Nevertheless, many attempts have been made to obtain information on the internal rotation from detailed studies of: a) the position of stars in the HR diagram; b) the evolution of the $V\!\sin i$ parameter during the main sequence (MS) phase; c) the shape of absorption lines, whose characteristics can depend upon the rotational law in layers close to the stellar surface; d) the global stellar geometry described with interferometric data.\par We briefly review these efforts:\par \medskip {\bf a}) The most numerous efforts among those just mentioned are the statistical analysis of photometric data on the rotational spread of the MS, which can be described by \citep{rox66,mae68,coll77,coll91} \begin{equation} \label{mvrot} \Delta M_{\rm V} = k(n)V^n, \end{equation} \noindent where $\Delta M_{\rm V}$ is the deviation in absolute magnitude from the zero-rotation MS, $V$ is the true equatorial rotational velocity, and $k(n)$ is a constant whose value depends on the power $n$. When $n\!=\!2$, $k(n\!=\!2)$ is on the order of $k_o10^{-5}$ mag/(km~s$^{-1})^2$, so that for $k_o\!\lesssim\!1$ the deviations may indicate that the internal rotation is uniform, while for $k_o\!\gtrsim\!1$ the internal rotation can be differential \citep{cott83}. This type of analysis found that stars do not seem to rotate uniformly. However, owing to the measurement uncertainties and difficulties in defining the MS of zero rotation, the available data could not provide any firm evidence of a particular law of non-uniform rotation \citep{strsar66,golay68,mae68,maepe70,smi71,smiwo74,mosmi81}. Furthermore, using detailed model atmospheres for differentially rotating stars, \citet{coll85} concluded that photometry alone can place albeit rather weak constraints on the degree of differential rotation within the stars.\par \medskip {\bf b}) Depending upon the internal angular momentum redistribution and evolutionary rearrangements of the inertial momentum, the surface equatorial rotational velocity of stars changes accordingly. Thus, the study of the variation in the true rotational velocity, $V$, as a function of time was studied by several authors using the ratio \begin{equation} \label{vsinlc} R_{\rm LC} = \frac{\langle V\!\sin i\rangle_{\rm LC}}{\langle V\!\sin i\rangle_{\rm ZAMS}} = \frac{\langle V\rangle_{\rm LC}}{\langle V\rangle_{\rm ZAMS}}, \end{equation} \noindent where $\langle V\!\sin i\rangle_{\rm LC}$ is the average of the $V\!\sin i$ parameters of stars with in principle the same mass and luminosity class (LC), $\langle V\!\sin i\rangle_{\rm ZAMS}$ is the average of $V\!\sin i$ for stars with the same mass, but located near the zero-age-main sequence (ZAMS). These ratios were compared with similar ones predicted theoretically for stars evolving as rotators in two different and extreme ways. On the one hand, the stars were assumed to evolve all their way as uniform rotators, which implies that the angular momentum is entirely redistributed at each evolutionary step. On the other hand, it was assumed that each stellar layer conserved its initial specific angular momentum, i.e., the stars did not undergo any redistribution of its internal angular momentum. Since in many cases the observed ratios $R_{\rm LC}$ were found to be situated in-between the two extreme theoretical predictions, it was suggested that stars should be differential rotators. However, these studies could not provide any information about the characteristics of the internal rotational law \citep{san55,danfa72,zo87col92,zo04iau215}. Somewhat related to this category of inquiries is the study of the evolution of the total angular momentum of B and Be stars carried by \citet{zo90nato316}, who concluded that these objects should undergo some internal angular momentum redistribution to explain the observed evolution of the $V\!\sin i$ parameters.\par \medskip {\bf c$_1$}) The study of the absorption line profiles of MS B-type stars found some evidence of possible surface differential rotation. The angular velocity in the surface of stars was assumed to depend on the colatitude angle, $\theta$, as \begin{equation} \label{omsto} \Omega(\theta) = \Omega_o[1-S\times(1-R(\theta)\sin^2\theta)] \ , \end{equation} \noindent where $\Omega_o$ is the equatorial angular velocity, $R(\theta)$ is the equation of the stellar surface, $S$ is the parameter that testifies to the differential rotation. Using stellar models that are more or less gravity darkened, \citet{sto68} and \citet{sto87} found that in most cases $S\!<\!0$, which suggested that the angular velocity tends to increase from the equator to the pole. Nevertheless, a dependence of the surface angular velocity on the latitude $\Omega(\theta)$ could be due either to an actual differential rotation present under the stellar surface, or simply to zonal atmospheric currents, which could appear in rapidly rotating early-type stars, as speculated by \citet{cra93}. In Sect.~\ref{egosetfr}, we recall that an acceleration of the angular velocity towards the equator, i.e. $S\!>\!0$, can be promoted by a temperature gradient induced by the gravity darkening effect. \par \medskip {\bf c$_2$}) The possibility of detecting surface differential rotation by means of the Fourier analysis of spectral line profiles was discussed by \citet{hua61}, \citet{gray77}, \citet{brun81}, \citet{gar82}, and \citet{rei02}. Evidence of surface differential rotation in late-type stars with $V\!\sin i\!<\!50$ km~s$^{-1}$ were given by \citet{rei03b}, \citet{rei04}, and \citet{reiro04}, but no differential rotation for late-type stars with $V\!\sin i\!>\!50$ km~s$^{-1}$ and A-type stars with $V\!\sin i\!>\!150$ km~s$^{-1}$ were reported by \citet{rei03a} and \citet{gray77}, respectively. It is possible that modest differential rotation is difficult to detect with the Fourier transform technique in slowly rotating A-type stars, because the rotational broadening is not large compared with the broadening caused by other mechanisms such as thermal turbulence and pressure effects \citep{gray77}. However, in those cases where there is some evidence of differential rotation, the parameter $S$ cannot be differentiated from the unknown inclination angle factor $\sin i$. Nevertheless, its sign indicates acceleration of the angular velocity towards the equator. To our knowledge, the Fourier technique for differential rotation has not yet been applied to early-type stars.\par \medskip {\bf c$_3$}) \citet{and80} proposed a method to probe the inner angular velocity of stars based on the use of the rotational splitting of non-radial oscillations. However, owing to uncertainties in the identification of pulsation modes and rough determinations of stellar fundamental parameters, particularly their evolutionary stage, this method has not yet been able to be applied with reliable success. Nevertheless, from the analysis of pulsation modes derived from photometric variations \citep{deup04} and COROT data \citep{degr09}, constraints on the internal rotation of $\beta$~Cep stars have been inferred.\par \medskip {\bf d}) In the past few years, interferometric methods have helped provide remarkable insights into not only the rotational distortion of stars \citep{arm03,vanbel04}, but also the induced gravity darkening effect by means of imaging techniques \citep{mcalis05, auf06,vanb06,zhao09}. New instruments with higher spectral resolutions of up to the 10000 attained by VLTI/AMBER in the J and K bands and an angular resolution of about 1 mas in the K band \citep{petrov07}, or spectral resolution reaching 30000 and angular resolutions as high as 0.3 mas in the visible using the VEGA/CHARA interferometric array \citep{mourard09}, will not only probably enable us to determine with greater detail than in previous studies the global geometry of stars deformed and gravity darkened by the rotation, but also carry out differential interferometry.\par A method based on differential interferometry that requires high spectral and spatial resolution was presented by \citet{arm04a,arm04b,arm04c} to distinguish observationally the parameter controlled by the degree of the surface differential rotation from the inclination angle factor $\sin i$.\par \subsection{Aims of the present attempt} \label{aotpa} Most theoretical predictions about the evolution of rotating early-type stars and the mixing of chemical species triggered by the instabilities set up by the rotation, come from calculations performed in the framework of two significant assumptions: a) the global rotational energy stored by the stars in the ZAMS is lower than the limit allowed by the rigid rotation in the critical regime; b) the internal angular velocity undergoes an instantaneous ``shellular" redistribution at each evolutionary step. However, \citet{clem79} using a cylindrical (conservative) distribution of the angular velocity, and \citet{mae08} basing their calculation on a ``shellular" distribution, showed in a more detailed way that in rapidly rotating early-type stars the envelope layers beneath the surface, may have wider convective zones in radius than in non-rotating stars. In the Sun, only the layers unstable to convection rotate differentially with a non-shellular pattern. This motivates the inquiry of whether in massive and intermediate mass stars some coupling may also exist between convection and rotation beneath their surface. In this case, the characteristics of the rotational law in the external stellar layers should differ from those currently assumed in the above evoked stellar models.\par As demonstrated by many authors, the global geometry of a star depends not only on the total amount of angular momentum stored by the star, but also on its internal distribution \citep{boden71,Zo86,smi92,uryu94,uryu95}. This geometry mostly relies on the stellar surface rotation, which acts as an imprint of its properties in the layers beneath the surface. In this case, we should not exclude the resulting mixing of chemical elements in the stellar atmosphere being more or less dependent on the characteristics of the external rotational law, upon which the description of the stellar structure, based on the abundance determination of chemical elements, should also rely. Therefore, to provide new information and/or constraints to test the global assumptions currently made to calculate models of stellar structure with rotation and thus help deepen our understanding of the properties of early-type fast rotators, we might ask: 1) what can be deduced, using first principles, about the properties the rotation laws can have beneath the surface as a consequence of the coupling between rotation and convection; 2) what are the parameters needed to characterize these stars that may be accessible to observations; 3) whether the combined interpretations of spectroscopic and interferometric data of rapidly rotating early-type stars enable us to determine these parameters. \par In this attempt, the most interesting information might be the indication of some differential rotation in the stellar surfaces and the sign of its latitudinal gradient. Both pieces of data can be obtained, as much as possible, in a consistent way by taking into account the stellar geometrical deformation produced by this rotation and the concomitant gravitational darkening effect that responds to possible non-conservative rotation laws.\par \medskip The present paper is organized as follows. In Sect.~\ref{rlie}, we use first principles to infer possible rotation laws in the convective layers beneath the stellar surface of early-type rapid rotators. Sect.~\ref{eotss} presents the equation of the stellar surface of stars with non-conservative rotation laws. A discussion of the gravity darkening effect for non-conservative rotation laws is presented in Sect.~\ref{tdss}. The discussion about the validity of the Roche approximation in representing the gravitational potential is presented in Sect.~\ref{mrots}. This discussion is based on 2D models of rotating stars where the evolutionary stages are taken into account in a simplified way. We briefly comment on the determination of the rotational profile in the stellar envelope in Sect.~\ref{trpe}. In Sect.~\ref{attain}, we summarize the attainable information on rapidly rotating early-type stars with external differential rotation from the combined analysis of spectroscopic and interferometric data. Our conclusions are presented in Sect.~\ref{conclus}.\par

\label{conclus} Recent theoretical works suggest that rapidly rotating early-type stars should have rather deep convective layers in the envelope. This situation may favor some coupling between the convective region and the differential rotation, as happens in the Sun. We have considered the possible properties the iso-rotation curves can have in the envelope immediately under the stellar surface by assuming several conditions for the specific entropy $S$ as a function of the: 1) specific angular momentum $j$, as implied by marginal Solberg-H\o iland stability condition; 2) squared angular velocity $\Omega^2$, which accounts for the Solar differential rotation in the convection zone; 3) specific kinetic energy $\varpi^2\Omega^2$, which reproduces the calculated rotational profiles in radiative envelopes. In all cases, the function that describes the gradient $\partial\Omega/\partial\theta$ in the stellar surface must be specified in advance. We thus assumed a simple Maunder formula with a unique differential rotation-parameter $\alpha$ to represent the imprint on the stellar surface produced by the differential rotation beneath the surface. In spite of the simplicity of this relation, for many theoretical reasons it would be highly interesting to infer from observations reliable orders of magnitude of $\alpha$ and its sign. While a value $\alpha\neq0$ can imply that the rotation beneath the surface is neither rigid nor shellular, neither spectroscopy nor interferometry can say anything about their actual functional nature beneath the surface. Other piece of information is thus needed, perhaps non-radial pulsation data of pulsation modes excited in the intermediate and upper stellar layers.\par In early-type fast rotators, we have shown that the surface temperature gradient might induce a positive gradient in the surface angular velocity, i.e., $\partial\Omega/\partial\theta>0$.\par In this paper, we have summarized the kind of information that it is possible to extract from the spectroscopic and interferometric data about the differential rotation of rapidly rotating early-type stars. The equation of the surface of a star with a non-conservative rotation law was discussed. It has been shown that a differential rotation in the surface induces measurable stellar deformations. These deformations carry a surface gravitational darkening effect that needs to be studied consistently with the induced geometrical deformations. We have shown that the effects induced by the surface differential rotation can be studied consistently and in a reliable way by a combined analysis of spectroscopic and interferometric data. From spectroscopy, estimates of the differential rotation parameter $\alpha$ with uncertainties $\delta\alpha\lesssim0.15$ and of the inclination angle $i$ with $\delta i\lesssim5^o$ can be obtained. The interferometry can provide information about the true ratio $R_{\rm e}/R_{\rm p}$ of the equatorial to the polar radii to within less than 3\%, provided that a reliable estimate of the inclination angle is given from spectroscopy. The VEGA/CHARA interferometric instrument is capable of providing useful information about rotationally induced stellar deformations for nearly 60 OB rapidly rotating stars ($V\!\sin i>200$ km~s$^{-1}$) whose spectra are not marred by circumstellar emission/absorption. Differential interferometry can help distinguish the $\alpha$ and $i$ parameters and provide reliable independent estimates of each.\par By calculating simplified models of stars at different evolutionary stages in the MS with internal rigid and differential conservative rotation laws, we have tested the use of the Roche approximation to represent the gravitational potential of rotating stars. We have concluded that we can safely use the Roche approximation for most, if not all stellar objects. Only in highly centrifugally deformed objects may this approximation not be applicable, but doubts exist as to whether these cases actually exist in Nature.\par Uncertainties and constraints carried by the spectroscopic and interferometric modeling as well as the application of these methods to real stars will be presented in separate papers.\par

1012.1707

Context. Since the external regions of the envelopes of rapidly rotating early-type stars are unstable to convection, a coupling may exist between the convection and the internal rotation. <BR /> Aims: We explore what can be learned from spectroscopic and interferometric observations about the properties of the rotation law in the external layers of these objects. <BR /> Methods: Using simple relations between the entropy and specific rotational quantities, some of which are found to be efficient at accounting for the solar differential rotation in the convective region, we derived analytical solutions that represent possible differential rotations in the envelope of early-type stars. A surface latitudinal differential rotation may not only be an external imprint of the inner rotation, but induces changes in the stellar geometry, the gravitational darkening, the aspect of spectral line profiles, and the emitted spectral energy distribution. <BR /> Results: By studying the equation of the surface of stars with non-conservative rotation laws, we conclude that objects undergo geometrical deformations that are a function of the latitudinal differential rotation able to be scrutinized both spectroscopically and by interferometry. The combination of Fourier analysis of spectral lines with model atmospheres provides independent estimates of the surface latitudinal differential rotation and the inclination angle. Models of stars at different evolutionary stages rotating with internal conservative rotation laws were calculated to show that the Roche approximation can be safely used to account for the gravitational potential. The surface temperature gradient in rapid rotators induce an acceleration to the surface angular velocity. Although a non-zero differential rotation parameter may indicate that the rotation is neither rigid nor shellular underneath the stellar surface, still further information, perhaps non-radial pulsations, is needed to determine its characteristics as a function of depth. <P />Table 5 is only available in electronic form at <A href="http://aanda.org">http://aanda.org</A>

12167558

2011ApJ...728...99S

1012

1012.3702_arXiv.txt

Magnetic fields are known to play an important role in star formation \citep[e.g.,][]{mc99,mck07}. The line-of-sight component of the magnetic field is often measured using the Zeeman effect or Faraday rotation \citep[e.g.,][]{cht03}. The angle of the magnetic field with respect to the plane of the sky can be deduced through continuum polarization measurements. Far-infrared continuum polarization is due to emission from elongated dust grains that align perpendicularly to the magnetic field \citep[e.g.,][]{hil88}. For the same reason, optical light that has undergone partial absorption by dust grains exhibits polarization parallel to the magnetic field \citep[e.g.,][]{dav51}. In general the magnetic field lines of the Milky Way follow the direction of the spiral arms \citep[often measured via Faraday rotation of pulsar signals; e.g.,][]{han06}. Since dust polarization vectors are orthogonal to magnetic field lines, the Galactic magnetic field should induce a tendency for mapped polarization angles to be perpendicular to the Galactic disk. External galaxies also show magnetic fields following the spiral arms \citep[e.g.,][]{sof86}, but on small scales (about 20\,pc in the case of NGC 6946) localized processes (e.g., star formation) dominate, tangling magnetic field lines \citep{bec07}. The 350\,$\mu$m Hertz polarimeter module was located at the Caltech Submillimeter Observatory and operated from 1994 to 2003. The Hertz polarimeter was well-suited for polarization measurements that probe the dense environments around forming stars, specifically clumps and cores. Since the decommissioning of Hertz, \citet{dot10} have published an archive of 56 different objects, 52 of which are Galactic star forming regions. In this paper, we use this relatively large dataset of star-forming regions to calculate a single large-scale (up to an angular diameter of about 10 arcminutes) average degree of polarization ($P$, percentage of wave that is polarized), position angle ($\theta$), and flux density ($I$) for each of the 52 star-forming Hertz datasets. The data are then explored to test whether the magnetic fields of these regions have any statistical relationship to the Galactic magnetic field. \citet{eri96} applied a similar analysis on visible polarization of 3000 stars (from the \citet{mat70} dataset) that are in the Galactic plane. They found that the measured optical polarization angles from most stars, particularly those at further distances, are parallel to the Galactic magnetic field. \citet{fos02} found that optical polarization measurements had a sinusoidal dependence with respect to Galactic longitude. However, we note that the observations probe very different physical regions: submillimeter observations probe dense regions ($A_V \gtrsim 30$) while optical observations probe diffuse regions ($A_V \lesssim 5$). Although dust polarization of star formation regions has been compared to the Galactic plane magnetic field before, the studies have been limited to small sample sizes. Continuum polarimetry at $\lambda$~=~0.8 and 1.3 mm of approximately 10 star formation sources (at scales of $\sim$1\arcmin) revealed no correlation for the magnetic field direction with respect to the Galactic plane \citep{gle99}. On the other hand, \citet{li06} analyzed the dust polarization of 4 giant molecular clouds (capturing most of the clouds on scales of $\sim$10\arcmin) at 450\,$\mu$m and found that 3 of these have a significant field aligned within $15^\circ$ of the Galactic plane. Our dataset probes similar regions, but we have a much larger sample size that will constrain, with much higher confidence, the relationship of the magnetic field to the Galactic plane. Our results do not show any similar correlation for degree of polarization or angle with Galactic location. An in-depth analysis was applied for $\theta$ by binning data based on Galactic longitude and spiral arm locations, again suggesting no correlation. This paper is organized as follows: in \S \ref{oda} we discuss our data and the analysis methods used. It also shows that a single polarization angle for each dataset is typically meaningful. Our results are presented in \S \ref{results} and discussed in \S \ref{discussion}.

\label{discussion} We have used Hertz \citep{dot10} polarization measurements in dense, star-forming clouds, to investigate whether the ordered component of the magnetic field in these clouds is correlated with location in the Galaxy. In order to obtain a single polarization percentage, angle, and intensity for each of the 52 Hertz objects (as well as the 22 complexes), the polarization information for each beam was combined into a large-scale average. With each object represented by these values, correlations with respect to location in the Galaxy (both by Galactic coordinates and spiral arm locations) were investigated. There are three primary results for this paper: \begin{itemize} \item A meaningful polarization angle can be determined for most objects and complexes. \item No evidence was found in our data for a correlation between the polarization angle and location within the galaxy. \item The polarization angle for an object or a complex on the sky is consistent with a random distribution. \end{itemize} The fact that a meaningful mean direction can be identified for the magnetic field in most objects implies the existence of an ordered, large scale component of the field {\em within} the dense, star-forming clouds we have studied. This is consistent with the continuity of magnetic field direction on different scales within these clouds, discussed by \cite{li09}. However, since these are star-forming clouds, feedback processes from newly formed stars can generate appreciable scatter in the magnetic field directions within each object. This may be, in part, the source of the observed angle dispersion within each dataset. We have found that there is no significant evidence for the existence of any correlation between mean polarization angle and location, which is consistent with the results by \citet{gle99} (see \S1). The fact that objects do not seem to avoid polarization angles aligned with the Galaxy's magnetic field implies that the polarization angles detected are almost entirely created by the analyzed object rather than a large-scale, external field. This suggests that complexes as a whole may become their own dynamical system that is separate from the Galaxy. The results in this paper imply that cloud cores usually have a meaningful net field (which may correlate with other cores in the same complex, \citep{li09}) that has no preferred direction within the Galaxy, yet are embedded in a diffuse medium in an ordered, Galactic large-scale field. \citet{fis03} did a similar analysis as this paper, but with a different methodology. They analyzed the line-of-sight magnetic fields through Zeeman splitting of OH masers in massive star-forming regions; on sub-kiloparsec scales (about 0.5\,kpc), two sources often had opposite line-of-sight field directions, suggesting multiple cores in a complex tangle the magnetic fields. In some areas of the Galaxy, \citeauthor{fis03} found some line-of-sight field alignment in parts of the Sagittarius Arm and Norma Arm on scales of about 2\,kpc. Still, they also found no evidence for correlations of magnetic field directions in star-forming regions with the Galactic field or with the spiral arms on larger scales. The cloud formation process involves instabilities on Galactic scales \citep{shettyo08,mkc09,tt09}, which are responsible for the accumulation of enough mass to form the clouds. At the same time, these instabilities generate turbulence in the ISM of the Galaxy. The cloud magnetic field is thus expected to decouple from the Galactic field during the cloud formation process. Stellar feedback is an additional mechanism driving the cloud magnetic field away from alignment with the Galactic field direction. These effects are likely responsible for the dichotomy between the arm/interarm regions in terms of the ratio of strengths between the ordered and tangled components of the magnetic field observed in external galaxies \citep{bec05}.

1012.3702

We investigate the effect of the Milky Way's magnetic field in star-forming regions using archived 350 μm polarization data on 52 Galactic star formation regions from the Hertz polarimeter module. The polarization angles and percentages for individual telescope beams were combined in order to produce a large-scale average for each source and for complexes of sources. In more than 80% of the sources, we find a meaningful mean magnetic field direction, implying the existence of an ordered magnetic field component at the scale of these sources. The average polarization angles were analyzed with respect to the Galactic coordinates in order to test for correlations between polarization percentage, polarization angle, intensity, and Galactic location. No correlation was found, which suggests that the magnetic field in dense molecular clouds is decoupled from the large-scale Galactic magnetic field. Finally, we show that the magnetic field directions in the complexes are consistent with a random distribution on the sky.

2586275

2011ApJ...741..111Q

1012

1012.3191_arXiv.txt

1012.3191

The Q/U Imaging ExperimenT (QUIET) employs coherent receivers at 43 GHz and 94 GHz, operating on the Chajnantor plateau in the Atacama Desert in Chile, to measure the anisotropy in the polarization of the cosmic microwave background (CMB). QUIET primarily targets the B modes from primordial gravitational waves. The combination of these frequencies gives sensitivity to foreground contributions from diffuse Galactic synchrotron radiation. Between 2008 October and 2010 December, over 10,000 hr of data were collected, first with the 19 element 43 GHz array (3458 hr) and then with the 90 element 94 GHz array. Each array observes the same four fields, selected for low foregrounds, together covering ≈1000 deg<SUP>2</SUP>. This paper reports initial results from the 43 GHz receiver, which has an array sensitivity to CMB fluctuations of 69 μK\sqrt{s}. The data were extensively studied with a large suite of null tests before the power spectra, determined with two independent pipelines, were examined. Analysis choices, including data selection, were modified until the null tests passed. Cross-correlating maps with different telescope pointings is used to eliminate a bias. This paper reports the EE, BB, and EB power spectra in the multipole range \ell = 25-475. With the exception of the lowest multipole bin for one of the fields, where a polarized foreground, consistent with Galactic synchrotron radiation, is detected with 3σ significance, the E-mode spectrum is consistent with the ΛCDM model, confirming the only previous detection of the first acoustic peak. The B-mode spectrum is consistent with zero, leading to a measurement of the tensor-to-scalar ratio of r = 0.35<SUP>+1.06</SUP> <SUB>-0.87</SUB>. The combination of a new time-stream "double-demodulation" technique, side-fed Dragonian optics, natural sky rotation, and frequent boresight rotation leads to the lowest level of systematic contamination in the B-mode power so far reported, below the level of r = 0.1.

12137753

2010idm..confE..98H

1012

1012.4587_arXiv.txt

The dark matter (DM) particle is fantastically stable. Its decay lifetime has to be obviously larger than the age of the Universe, $\sim 10^{18}$~seconds. In most cases it has to be even larger than about $10^{26}$ seconds, in order that the decay doesn't cause fluxes of cosmic positrons, antiprotons and $\gamma$ larger than the observed ones. It turns out nevertheless that in most DM models this astonishing fact is not explained but rather assumed by hand, typically by assuming a discrete $Z_2$ symmetry. That, of course, does not mean at all that these models couldn't be correct, but certainly that their completeness is questionable. The origin of this stability is most probably preponderant in determining the structure of the DM interactions. Besides providing, for example, an explanation for the DM relic density, it would be highly desirable that DM models also provide a fundamental explanation for this stability. Here, by fundamental, we mean from a dynamical reason, resulting from a gauge symmetry and the particle field content of the model. A strong motivation to look for such stabilization mechanism is the situation of the Standard Model (SM). In the SM there is a series of stable particles whose stability results in all cases from the gauge symmetries and particle content, rather than from an ad-hoc symmetry. The photon is stable because it is the massless gauge boson of the exact electromagnetic $U(1)_{QED}$ gauge symmetry. The $e^-$ is stable because it is the lightest particle charged under this gauge group. The lightest neutrino is stable because of Lorentz invariance, since it is the lightest fermion. Finally the proton is stable because of conservation of baryon number, which is not imposed by hand, but results accidentally from the SM gauge symmetries and the gauge charges assigned to the SM particles. In the following we review briefly the various simple stabilization mechanisms which can be invoked. We stress that the stabilization mechanism has indeed a preponderant influence on the structure of the model and therefore on the associated phenomenology. We will distinguish 2 kinds of mechanisms, the low energy ones, which similarly to the SM allow to understand the particle stability directly at low energy, and the high energy ones which rely on the existence of some explicit UV physics, typically on the existence of a high energy gauge group. The latter possibility is less interesting phenomenologically because it relies on an assumption difficult to probe experimentally, but is certainly theoretically attractive too, especially if it is connected to grand-unification.

The fact that the DM particle is stable on cosmological time scales is a peculiar property. It definitely needs an explanation. On the basis of a gauge principle there are quite a few possible origins for this stability. Along these mechanisms DM can be a scalar, a fermion, or even a gauge boson (or higher spin object such as the gravitino). The gauge symmetry invoked can be abelian or non abelian (both in their confined or unconfined phase). Each mechanism leads to a characteristic phenomenology. In particular, 1) the existence of an exact gauge group results in extra radiation which might have many astrophysical effects, 2) accidental symmetry mechanisms are expected to lead to a rich indirect detection phenomenology from DM decay, such as intense $\gamma$ ray lines in the hidden vector DM setup, 3) high energy stabilization mechanism, based in particular on $SO(10)$, can have DM as the missing piece for unification of gauge couplings, 4) if no extra gauge group are assumed the DM must be part of a high $SU(2)_L$ multiplet with definite properties. Apart for the models where DM lies in the multi-TeV range, and for the models based on extra-radiation, all models can lead to specific signatures at the LHC.

1012.4587

From the particle physics point of view, the most peculiar property of the dark matter particle is its stability on cosmological time scales. We briefly review the possible origins of this characteristic feature for candidates whose relic density results from the thermal freeze-out of their annihilation. We emphasize that each stabilization mechanism implies an all specific phenomenology. The models reviewed include supersymmetric and non-supersymmetric models where the stability is a consequence of grand-unification, models where stability is due to an unbroken gauge group and models where the DM stability is accidental. The latter possibility includes minimal dark matter, hidden vector dark matter and composite DM models.

12169735

2011ApJS..195....8C

1012

1012.4064_arXiv.txt

\setcounter{footnote}{0} \begin{figure*}[ht] \resizebox{\textwidth}{!}{\includegraphics{f1.pdf}} \caption{Fourth (left) and third (right) quadrant field-differenced \flow\ Stokes \I\ map smoothed to the beam scale ($5'$), with color scale in MJy/sr. The solid black lines indicate the survey coverage. } \label{fig:fdmap} \end{figure*} Millimeter (mm), sub-millimeter (sub-mm) and far-infrared (FIR) observations are ideal for studying the properties of star-forming regions in the galaxy, in particular the cool envelopes of dust and gas which host sites of potential and active star formation. By spanning the peak in the spectra of these objects, measurements between the mm and FIR can tightly constrain the parameters of the thermal radiation produced by the dust. In particular, the mass of a star-forming core and its surrounding envelope is well-traced by its measured flux in these bands, since this radiation is optically thin at sub-mm and longer wavelengths. Surveys covering large sections of the galaxy have the potential to collect statistical samples of cores in a range of evolutionary states, comparatively free of bias introduced by targetting particular regions. These surveys are ideal to study processes related to star-forming regions, such as measuring the core mass function (from which the stellar initial mass function may be derived), particularly at the high-mass end, which, on account of the short-lived high mass cores, is understudied relative to lower masses~\citep[e.g.][]{enoch2006,young2006,enoch2008}. Combination with infrared (IR) data yields insight into the ages of cores, permitting differentiation between prestellar sub-mm cores, which lack an IR counterpart, and protostellar cores, in which the ultraviolet radiation produced by protostars is re-radiated into the mm, sub-mm and IR by the surrounding envelope. Phenomena associated with later evolutionary phases, such as mass ejection, dissipation of the envelope, and dynamical interactions are not significant in the prestellar or protostellar stage --- the mass and spatial distribution of such cores therefore capture information regarding the fragmentation process~\citep{enoch2006}. Observations of polarized radiation permit a window to study the role of magnetic fields~\citep[e.g.][]{greaves1995,novak1997}, and their role in providing support against collapse. In the mm and sub-mm, polarization is due to emission along the long axis of dust grains partially aligned by the magnetic field, and thus measurements of the dust polarization directly probe local magnetic fields~\cite[e.g.][]{hildebrand1988}. These fields are thought to strongly influence the evolution of molecular clouds, since they provide support preventing the collapse of the gas and subsequent triggering of star formation. Several large-scale surveys are already underway or completed to help address these questions. Herschel~\citep{Herschel} and Planck~\citep{Planck} will provide extensive spectral coverage from the radio to the far infrared, fully characterizing the spectral energy distribution (SED) of star-forming cores over the fully sky; selected existing results in targetted regions include~\citep[e.g.][]{hennemann2010,juvela2010,andre2010}, but are limited to total intensity observations. Ground and balloon instruments also contribute substantially to the literature: ~\cite{schuller2009} present an APEX LABOCA 95 deg$^2$ survey in total intensity with resolution of $19.2''$ at 353 GHz, with the final survey coverage expected to reach 350 deg$^{2}$; Bolocam has mapped 150 deg$^{2}$ of the first galactic quadrant at 1.1mm (268~GHz) with resolution $33''$, with a source catalog presented in~\cite{rosolowsky2009}; BLAST~\citep{olmi2009,netterfield2009} provide a 50 deg$^{2}$ survey of the Vela molecular cloud at 250, 350 and 500 microns (36, 42 and 60 arcsec resolution respectively). Observations at comparable resolution are currently scarce at $\sim$\flow, and yet provide additional constraining power to the Rayleigh-Jeans tail of the thermal dust spectrum, and probe for contributions due to other emission mechanisms which contribute increasingly at lower frequencies (e.g. free-free). Furthermore, there is little high angular resolution polarization data at these frequencies, despite their utility in understanding star-forming regions. In this paper we present a catalog of compact sources found in the QUaD galactic plane survey~\citep{culverhouse2010}, which covers over $\sim800$ square degrees of the low-latitude galactic plane at 100 and \fhigh~with beam FWHM of $5'$ and $3.5'$ respectively, in Stokes \I, \Q~and \U~parameters \footnote[1]{The QUaD maps and source catalogs analyzed in this paper are available for public download at \tt{http://find.spa.umn.edu:/quad/quad\_galactic/}}. A survey of this size, frequency and angular resolution can be used to investigate the polarized and unpolarized properties of both diffuse emission and discrete sources. The QUaD survey was conducted \it{blind}\rm, in that no region was specifically targetted. In principle, this allows the construction of statistical samples of cores, representative of the distribution of core masses and ages in the galaxy as a whole. However, we note that at the $\sim$few arcminute resolution of the survey, the maps do not generally resolve individual cores: Dense cores typically have size $\sim 0.1\mathrm{pc}$~\citep[e.g.][]{Williams2000}, hence for nominal distances of a few hundred pc, the sources presented here should be considered as `clumps' hosting cores rather than individual cores themselves. In addition to the lack of resolution and accurate clump distances, the contribution of free-free emission at \flow\ biases measurements of the flux from the dust component; these caveats prevent reliable mass calculation. Our goals here are therefore to analyze the observed quantities of the sources in the survey, rather than infer their physical properties. Basic information on the instrument, observations and maps is presented in Section~\ref{sec:instobs}. In Section~\ref{sec:srcextraction}, we describe our algorithm for extracting sources in the presence of a diffuse background. The global properties of the catalog are discussed in Section~\ref{sec:data}, with the full catalogs presented in Table~\ref{tab:srccat} (total intensity) and Table~\ref{tab:polsrccat} (polarized intensity). Discussion and Conclusions are found in Section~\ref{sec:conclusions}. Extensive simulations, presented in Appendix~\ref{app:sims}, are used to quantify the effects of mapmaking and source extraction algorithm on the recovered source properties. \begin{figure*}[ht] \resizebox{\textwidth}{!}{\includegraphics{f2.pdf}} \caption{Source segmentation and fitting using a section of QUaD \flow\ \I\ data; colorscale is the same in all panels and runs from -0.0035 to 0.035 MJy/sr. {\it Top Left:} Raw map $m_{r}$. {\it Bottom Left:} Initial estimate of background map $m_{bck}$. {\it Top row, second column:} Initial background subtracted map $m_{s}$. {\it Bottom row, second column:} Background map corrected for discrete source flux ($m_{bck,corr}$). {\it Top row, third column:} Background subtracted map $m_{s}$ using source-corrected background map $m_{bck,corr}$. {\it Bottom row, third column:} Map used for source fitting $m_{r}-m_{bck,corr}$. {\it Top right:} Model of discrete source population using fits to $m_{f}$. {\it Bottom right:} Residual between model sky and input image $m_{r}$. } \label{fig:sourceseg} \end{figure*}

\label{sec:conclusions} We present a catalog of discrete sources extracted from the QUaD galactic plane survey, which spans approximately 245-295$^\circ$ and 315-5$^\circ$ in galactic longitude $l$ and -4 to +4$^\circ$ in galactic latitude $b$ --- a total of $\sim800$ square degrees coverage in Stokes \I, \Q, and \U\ at 100 and \fhigh, with resolution 5 and 3.5 arcminutes respectively. Simulations of a toy model galaxy including spatially clustered point sources and diffuse emission indicate a 90\% completeness flux of 5.9 (2.9) Jy at 100 (150) GHz in \I, and 20.3 (1.1) Jy in polarization at 100 (150) GHz. The high \flow\ completeness in polarization is due to source confusion in the larger beam, and is dependent on the parameters used for the toy galactic model; randomly distributed sources yield a completeness of 1.3 Jy --- see Appendix~\ref{subsec:completeness} for a full discussion. At a signal-to-noise threshold of 5 (3) in total (polarized) intensity, the catalog is 98\% pure in \I\ at both frequencies, and 97\% (92\%) pure in polarization at 100 (150) GHz. Simulations without a diffuse background are used in the total intensity computation because substructure in diffuse emission, detected as discrete sources, biases the purity low. While this could also affect the real catalog, the high percentage of IRAS-PSC counterparts to QUaD sources (97\% and 87\% at 100 and \fhigh\ respectively) indicates that this effect is likely very small. The polarized diffuse background, with fractional polarization $\sim2\%$ (see Map Paper), is faint enough that it does not bias the purity of the catalog at the signal-to-noise threshold of the survey. Instrumental effects prevent detection of polarized sources with polarization fraction $\sim1\%$ or less. In total intensity the catalog contains 505 unique sources, of which 239 are spatially matched between frequency bands, with 50 (216) detected at 100 (150) GHz alone. The \I\ flux distributions are well-approximated by a power law over more than two orders of magnitude above $\sim10$ Jy at both frequencies. We find power-law slopes of $\gamma_{S,100}=-1.8\pm0.4$ at \flow, and $\gamma_{S,150}=-2.2\pm0.4$ at \fhigh; the latter is consistent with~\cite{rosolowsky2009}, who find $-2.4\pm0.1$ at 268~GHz with Bolocam at higher resolution. The flatter slope at \flow\ may be the result of resolution effects due to the larger beam at this frequency. Simulations indicate that if the diffuse background contributes spurious sources, as expected the recovered source flux distribution does not accurately follow the underlying distribution; however, as discussed above, the high percentage of QUaD sources spatially matched to IRAS indicates this effect is insignificant. The spectral index probability distribution of sources in total intensity is found to peak at $\alpha\sim0.25$, flatter than expected for sources whose emission is dominated by thermal dust. Simulations indicate the diffuse background does not strongly influence source spectral indices; the flatness is therefore likely due to free-free emission, which becomes significant at $\sim$\flow\ and below. At this frequency, free-free results in higher fluxes than expected from dust alone, shifting the spectral index distribution to lower values. We explore the clustering of galactic sources by fitting the two-point correlation function to a power-law using the \I\ source locations. Simulations indicate that the underlying correlation function slope can be accurately reconstructed in the range $0.4^{\circ}<\theta<2^{\circ}$, with $\theta$ the angular separation between a pair of sources. The correlation function breaks down at larger angular scales because so few ($<1\%$) of sources are located beyond $|b|=3^{\circ}$. At galactic latitudes smaller than $0.4^{\circ}$, $w(\theta)$ is not well reconstructed because for a power-law $dN/dS\propto S^\gamma_{S}$ with $\gamma_{S}<0$, the survey does not detect most neighbors of a source bright enough to be included --- one must extend the search to large angular separations before enough neighbors are detected for accurate reconstruction. Fitting to the QUaD \I\ catalog data in the range $0.4^{\circ}<\theta<2^{\circ}$, we find power-law slopes of $\gamma_{\theta,100}=-1.21\pm0.04$ and $\gamma_{\theta,150}=-1.25\pm0.04$ at 100 and \fhigh\ respectively. These are consistent with the value found by~\cite{enoch2006} Bolocam observations of the Perseus molecular cloud, $w(\theta)\propto \theta^{-1.25}$, though the results are not directly comparable on account of the different sources probed by QUaD (clumps) and Bolocam (cores) due to their differing angular resolution. 97\% (87\%) of the sources detected at 100 (150) GHz have IRAS-PSC counterparts. These large fractions indicate that most of the clumps detected in the survey are past the prestellar phase and have envelopes heated by protostars. This observation might be expected, given that the QUaD frequency bands lie far from the spectral peak: Only these sources are bright enough in the Rayleigh-Jeans portion of the spectrum to be detected in the QUaD survey, unlike prestellar or starless sources. Since the QUaD survey is sensitive to free-free emission as well as dust, particulary in the \flow\ band, sources might also be detected if their gas is sufficiently ionized to produce free-free but their envelopes are yet to thermalize. However, the small fraction of \flow\ QUaD sources unmatched to IRAS-PSC imply that almost all the detected sources at this frequency do have a thermal component. At \fhigh, the larger unmatched fraction is likely due to single IRAS sources being resolved into two sub-clumps by the the higher QUaD resolution at this frequency --- only one of these sub-clumps can be spatially matched to the IRAS source. Of the sources with an IRAS counterpart, 182 satisfy the Wood-Churchwell criteria for ultracompact \HII\ regions~\citep{wood1989}, providing new spectral constraints on this class of object. Four compact polarized sources were detected by the automated source-finding algorithm: 284.33-0.36 (IRAS 10227-5730 or RCW 49), 359.93-0.06 (Sagittarius A*), 0.18-0.06 (IRAS 17431-2846 or Galactic Center Arc), and 0.63-0.05 (IRAS 17440-2823 or Sagittarius B2). One additional extended source was located `by eye' from the raw \Q\ and \U\ maps; this object appears host to several discrete total intensity sources, including 344.99+1.79, 345.01+1.53, 345.37+1.42, and 345.22+1.03. The brightest polarized source is 0.18-0.06, which does not have an obvious discrete counterpart in total intensity, but has a polarization fraction of $\sim10\%$ if the diffuse background is used as a measure of \I. It has a polarized flux of $7.91\pm0.33$ ($4.90\pm0.32$) Jy at 100 (150)~GHz, and a polarized spectral index of $\alpha_{P}=-1.04\pm0.17$, indicating a synchrotron emission source. Its detection against a polarized background implies that there is a strong local deviation from the galactic magnetic field. Less than 1\% of the sources detected in \I\ have a polarized counterpart. If discrete sources do not harbor strong local magnetic fields or shielding, dust grains in their envelopes will align with the large-scale galactic field. The only way to separate diffuse from discrete polarized emission would then be via morphology (similar to \I) or spectrally, since the orientation of the polarization would be similar for diffuse and discrete sources. Alternatively, the discrete total intensity sources may have fractional polarization $<1\%$, as might be expected from a star-forming clump, in which case instrumental effects prevented detection of their polarized emission here. Discrete sources may therefore not be a significant contributor to the low-latitude galactic polarized emission. More sensitive observations (such as from the Planck satellite) will be needed to better study the polarization of these sources, and the role of magnetic fields in star-forming regions. The QUaD catalog may prove useful for a variety of additional purposes. Total intensity source fluxes could better measure the continuum spectra of clumps in conjunction with independent data sets, improving the separation of different emission components and tightening constraints on dust emissivity and gas temperatures. The maps provide upper limits to source polarization, allowing a statistical study of polarized contribution to anomalous emission similar to~\cite{dickinson2007}. Finally, the catalog provides a cross-check of astrometry and absolute calibration for instruments with access to the southern hemisphere.

1012.4064

We present a catalog of compact sources derived from the QUaD Galactic Plane Survey. The survey covers ~800 deg<SUP>2</SUP> of the inner galaxy (|b| &lt; 4°) in Stokes I, Q, and U parameters at 100 and 150 GHz, with angular resolutions of 5 and 3.5 arcmin, respectively. Five hundred and twenty-six unique sources are identified in I, of which 239 are spatially matched between frequency bands, with 53 (234) detected at 100 (150) GHz alone; 170 sources are identified as ultracompact H II regions. Approximating the distribution of total intensity source fluxes as a power law, we find a slope of γ<SUB> S, 100</SUB> = -1.8 ± 0.4 at 100 GHz and γ<SUB> S, 150</SUB> = -2.2 ± 0.4 at 150 GHz. Similarly, the power-law index of the source two-point angular correlation function is γ<SUB>θ, 100</SUB> = -1.21 ± 0.04 and γ<SUB>θ, 150</SUB> = -1.25 ± 0.04. The total intensity spectral index distribution peaks at α<SUB> I </SUB> ~ 0.25, indicating that dust emission is not the only source of radiation produced by these objects between 100 and 150 GHz free-free radiation is likely significant in the 100 GHz band. Four sources are detected in polarized intensity P, of which three have matching counterparts in I. Three of the polarized sources lie close to the Galactic center, Sagittarius A*, Sagittarius B2, and the Galactic Radio Arc, while the fourth is RCW 49, a bright H II region. An extended polarized source, undetected by the source extraction algorithm on account of its ~0fdg5 size, is identified visually, and is an isolated example of large-scale polarized emission oriented distinctly from the bulk Galactic dust polarization.

2358220

2011MNRAS.413..813C

1012

1012.1551.txt

\subsection{Scientific background} Observations of high-redshift galaxies and the cosmic microwave background \citep{Spergel2007} have revealed the Universe to be dominated by dark matter and dark energy \citep{Riess1998,Perlmutter1999}, providing a working paradigm for the formation of structure \citep[e.g.][]{Springel2005nat}. However, the mechanisms that form the luminous content of the dark-matter potential (i.e.\ the stars and galaxies that we observe) remain the key unknowns of modern extra-galactic astronomy. These processes are driven by the hydrodynamics and chemistry of the gas, combined with complex radiative feedback processes. High-redshift observations alone are not sufficient to constrain these processes, lacking spectral information and spatial resolution \citep{Faber2007}. It is therefore necessary to complement these studies with detailed analysis of nearby objects, tracing the {\it fossil record} of the formation process. Early-type (elliptical E and lenticular S0) galaxies (ETGs) are especially useful as they are old, have smaller levels of star formation and limited amount of dust, which simplifies the interpretation of the observations. Significant progress has been made in this direction in the past few decades, building on the classic observational works that still capture much of our understanding of the structure of local ETGs \citep[e.g.][]{Hubble1936,faber76,Davies1983,Dressler1987,Djorgovski1987, bender92,kormendy95}. A major step forward was brought by the era of large galaxy surveys. Thanks to the unprecedented sample size, one of the most important contributions of the Sloan Digital Sky Survey \citep[SDSS,][]{York2000} was to firmly establish a statistically significant bimodality in the colour distribution of local galaxies, such that they can be clearly separated in a so-called `blue cloud', generally consisting of star-forming spiral galaxies, and a `red sequence', mostly of non-star-forming ETGs \citep{Strateva2001,Baldry2004}. Accurately quantifying this bimodality, and the realization that it can be traced back in time to higher redshift \citep{Bell2004,Faber2007}, allowed a dramatic improvement in the detailed testing of galaxy formation scenarios. The bimodality can only be explained with the existence of a mechanism, which suppresses episodes of intense star formation by evacuating gas from the system, resulting in a rapid transition of galaxies from the blue cloud to the red sequence \citep{Springel2005,Faber2007}. Many simulation groups have reproduced the bimodality qualitatively, though with rather different assumptions for the star formation and feedback processes \citep{Granato2004,DiMatteo2005,Bower2006,Cattaneo2006,Croton2006}. A generic feature of these models is that red-sequence galaxies form by dissipational `wet mergers' of gas-rich blue-cloud galaxies, followed by quenching of the resulting intense star-formation by rapid ejection of the gas, caused by the feedback from a central supermassive black hole, supernovae winds, by shock heating of the gas in the most massive halos \citep{Keres2005,Dekel2006} or gravitational gas heating \citep{Naab2007,Khochfar2008,Johansson2009}. The merging of the most massive blue galaxies, however, is not sufficient to explain the population of most-massive red-sequence galaxies. Dissipationless `dry mergers' of gas-poor, red-sequence galaxies is therefore also required, evolving galaxies {\em along} the red-sequence as they increase in mass \citep{Khochfar2003,Naab2006kb,Hopkins2009,Khochfar2009,Oser2010}. Both wet and dry major mergers generally produce red, bulge-dominated galaxies when feedback is included in the models. The kinematic structure of the remnants is however very different. In a major (1:1) merger between blue gas-rich galaxies, the gas tends to form a disk, so that the end result of the merger, after the gas has been removed from the system by ejection, heating or conversion to stars, will be a red stellar system dominated by rotation \citep{Cox2006,Naab2006,Robertson2006,Jesseit2009}. In major mergers between red gas-poor galaxies, dissipationless processes dominate, resulting in a red galaxy with little or no net rotation \citep{Barnes1992,Hernquist1992,Naab1999,Naab2003,Cox2006}. Unlike major mergers, minor mergers (1:3 or less) retain more closely the structure of the progenitor, to an extent that depends on the amount of mass and gas accreted, so that the remnant of a spiral galaxy will always display significant rotation \citep{Naab2006,Robertson2006,Bournaud2007,Jesseit2009}. These simulations demonstrate that if galaxies assemble by mergers, the existence of the red/blue galaxies dichotomy therefore also suggests the existence of a {\em kinematical} differentiation {\em within} the red sequence between fast and slow rotating galaxies. Various classic observational indicators of an ETGs dichotomy have been proposed in the past two decades. ETGs have been found to exhibit trends as a function of luminosity in terms of (i) their distribution on the \vse\ diagram, which relates the ratio of ordered $V$ and random $\sigma$ stellar motion to the galaxy ellipticity $\varepsilon$ \citep[e.g.][]{Illingworth1977,Binney1978,Davies1983}, (ii) their isophote shape (disky or boxy) \citep{Bender1989,Kormendy1996}, (iii) the inner slope of their photometric profiles: cored/cuspy \citep{Ferrarese1994,Lauer1995,faber97} or excess/deficit of core light \citep{Graham2004,Ferrarese2006,Kormendy2009}. However, none of these signatures have been able to give clear evidence for a distinction between the two classes of red-sequence galaxies, primarily because they are all essentially secondary indicators of the galaxies' internal kinematic structure. By the application of integral-field spectroscopy to a representative sample of nearby ETGs, the \sauron\ survey \citep{deZeeuw2002} has revealed the full richness of the kinematics of these objects \citep{Emsellem2004,mcdermid06,Krajnovic08}. From the two-dimensional nature of this unique data set, two distinct morphologies of stellar rotation fields are clearly evident, corresponding to the above described fast- and slow-rotators. In two companion papers of that survey a global quantitative measure of this morphology was defined, termed $\lambda_R$, that can be used to kinematically classify these galaxies in a way that is more robust than the \vse\ diagram, is nearly insensitive to projection effects \citep{Emsellem2007,Cappellari2007}. $\lambda_R$ relates directly to their formation, and is precisely reproducible in current cosmological simulations \citep{Jesseit2009,Bois2010}. This is the basic new finding we plan to exploit in the present project to improve our understanding of the structure and formation of ETGs. Additional results of the \sauron\ survey on ETGs include the robustness and empirical `calibration' of the simple virial mass estimator to measure mass in the central parts of ETGs and a determination of their dark matter fraction \citep{Cappellari2006}. The survey found a high incidence of ionized gas in ETGs \citep{Sarzi2006} and explained their ionization mechanism as mainly due to the evolved stellar population \citep{Sarzi2010}. It was shown that the stellar population gradients correlate well with the escape velocity, both locally within galaxies and globally among different ETGs \citep{Scott2009}. Star formation in ETGs only happens in fast rotators and follows two distinct modes: in disks or widespread \citep{Shapiro2010}, where the latter cases are in low-mass systems \citep{Jeong2009,Kuntschner2010}. Disks in fast rotators have enhanced metallicity, while kinematically distinct cores in slow rotators show no stellar population signatures \citep{Kuntschner2006,Kuntschner2010}. \subsection{Goals of the Project} Due to the exploratory character of the \sauron\ survey \citep{deZeeuw2002}, the ETGs were selected to sample, with a relatively small number of objects, a wide range of masses, shapes and morphologies. This was done by selecting galaxies brighter than a total magnitude $M_B<-18$ mag equally subdivided into 24 E and 24 S0. Within each E/S0 subclass the selected objects sample uniformly a grid in the $(M_B,\varepsilon)$ plane. Although that approach was crucial in bringing the fast/slow rotator dichotomy to light and in most of the findings mentioned in the previous section, the selection criteria impose complex biases and do not allow for a quantitative statistical comparisons of galaxy properties with simulations, which is a main goal of the \atl\ project. Moreover, with only 48 galaxies, the statistical uncertainties are large. The power of the kinematic classification based on $\lambda_R$ is to be able to study differences in the formation process along the red-sequence galaxy population. The $\lambda_R$ parameter describes in a compact way the present status of the galaxies, however it is essential to obtain information on the formation history and the detailed dynamical structure as well. The stellar population contains a record of the more distant history (a few Gyr). Recent gas accretion is recorded in the cold atomic gas components, generally detected on galaxy scales with radio observations of \hi, while the ongoing accretion and star formation activity is traced by cold molecular gas (e.g.\ CO), often detected in regular disks in the central regions. For comparison with theoretical predictions one needs to observe all these quantities for a statistically significant, volume-limited sample of galaxies complete to some useful lower limit in mass. With these ideas in mind we carefully selected the \atl\ sample of ETGs and we systematically observed all the above quantities. The \atl\ dataset now provides a complete inventory of the baryon budget and a detailed two-dimensional description of stellar and gaseous kinematics, together with resolved stellar population within the main body of a complete and statistically significant sample of ETGs. Our goal is to use this dataset to perform archeological cosmology by specifically answering the following questions: \begin{enumerate} \item How do slow rotators form? What are the physical processes that determine their kinematic and photometric features? What is the role of major and minor mergers in their formation history? This will be reflected in the kinematics, gas content and stellar population. \item Why are most ETGs fast rotators? There seems to be a dominant formation mechanism that delivers galaxies with quite homogenous rotation properties. Can this be merging? Can significant major merging be excluded? \item How is star formation in ETGs quenched? Is it different for fast and slow rotators ETGs? How does it depend on environment? Can we infer the quenching mechanism from the amount and distribution of the left-over gas, the presence of AGNs or metallicity gradients? The distribution of stellar population and gas properties constitute a stringent test for future galaxy formation models. \item Most past studies have focused on single stellar population models of ETGs, but cosmological models predict more complex histories. Can we infer the star formation history in ETGs for detailed comparison with simulations? \item How do counter rotating cores in massive and old ETGs form and survive to the present time? Are these relics of the very early Universe? \item Can we link the present day properties of ETGs to results form existing and upcoming surveys at higher redshift with respect to e.g. masses, sizes, stellar populations, gas fractions, star formation? Our study will constitute a $z=0$ redshift benchmark to trace the time evolution of galaxies. \end{enumerate} The \atl\ sample includes all nearby ETGs observable from the northern Earth hemisphere, and for this reason we hope its homogeneous dataset will ultimately constitute a legacy for future studies. We trust that our and other groups will exploit our data and sample well beyond what we had originally envisioned. Our first steps in the directions outlined above are presented in the following papers, while the other aspects will be presented in subsequent papers of this series: \begin{enumerate} \item \citet[hereafter Paper II]{Krajnovic2010}, which describes the morphology of the kinematics and the kinematical misalignment in ETGs; \item \citet[hereafter Paper III]{Emsellem2010}, which presents a census of the stellar angular momentum in the central region of ETGs; \item \citet[hereafter Paper IV]{Young2010}, which quantifies the distribution of molecular gas content in ETGs; \item \citet[hereafter Paper V]{Davis2010}, which studies the \citet{Tully1977} relation from the width of the molecular lines in ETGs; \item \citet[hereafter Paper VI]{Bois2011}, which studies the formation of the fast and slow-rotator galaxies via numerical simulations of binary mergers; \item \citet[hereafter Paper VII]{Cappellari2011}, which revisits the morphology of nearby galaxies and presents the {\em kinematic} morphology-density relation. \end{enumerate} Here in Section~2 we discuss the selection criteria for the {\em parent} sample of galaxies, from which the \atl\ sample of ETGs was extracted (Section~3). In Section~4 we present the \sauron\ observing strategy for the survey, the integral-field data, and the kinematic extraction, while other additional datasets and simulations from our project are listed in Section~5. We give a summary in Section~6. In the paper we assume $H_0=72$ \kms\ Mpc$^{-1}$.

We described the motivation and goals of the \atl\ project, which aims at constraining models of galaxy formation by obtaining two-dimensional observations of the distribution and kinematics of the atomic (\hi), molecular (CO) and ionized gas, together with the stellar population and kinematics, for a volume-limited nearly mass-selected ($K_s$-band) sample of ETGs. We defined the selection criteria for the volume-limited ($1.16\times10^5$ Mpc$^3$) {\em parent} sample of 871 galaxies with $D<42$ Mpc and $M_K<-21.5$ mag, satisfying our observability criteria, and investigated possible selection biases, especially due to redshift incompleteness. We found that incompleteness cannot amount to more than a couple of percent, making the sample essentially complete and representative of the nearby population. We additionally tested the representativeness of the sample by comparing its $K_s$-band luminosity function with the one derived from a much larger sample \citep{Bell2003} and found a very good agreement. We described the morphological selection used to extract the 260 ETGs of the \atl\ sample from the parent sample and showed that the ETGs define a narrow red-sequence, on a colour-magnitude diagram, with few objects in transition from the blue cloud. We presented the size-luminosity relation for the \atl\ sample and the full parent sample to illustrate the general characteristic of our galaxies. We described the strategy for the \sauron\ integral-field observations, the data reduction, the extraction of the stellar kinematics and their typical errors. We gave an overview of the additional dataset already available for our sample, which include interferometric observations of the atomic gas as traced by \hi, single-dish and interferometric observations of molecular gas as traced by the CO lines, and multi-band optical photometry. We summarized the ongoing semi-analytic modeling and the cosmological and binary-merger N-body simulations we are performing to interpret our observations. This is the first paper of a series devoted to our understanding of the formation of ETGs. Key additional elements are provided by the kinematics, ages and chemical composition of the stars in the galaxies, the presence of cold atomic or molecular gas, the photometric profiles and the dynamical masses, as a function of environment. The observations for the \atl\ sample will be compared against the model predictions, to test formation scenarios and to tune the model parameter. This will be the topic of future papers of this series. The \atl\ project aims to constitute a zero redshift baseline, against which one can investigate the evolution of galaxy global parameters with redshift, to trace galaxy evolution back time. Future studies should extend this effort to more dense environment than can be explored in the nearby universe, and to increasingly higher redshifts to explore the time evolution of the structure of ETGs.

1012.1551

The ATLAS<SUP>3D</SUP> project is a multiwavelength survey combined with a theoretical modelling effort. The observations span from the radio to the millimetre and optical, and provide multicolour imaging, two-dimensional kinematics of the atomic (H I), molecular (CO) and ionized gas (Hβ, [O III] and [N I]), together with the kinematics and population of the stars (Hβ, Fe5015 and Mg b), for a carefully selected, volume-limited (1.16 × 10<SUP>5</SUP> Mpc<SUP>3</SUP>) sample of 260 early-type (elliptical E and lenticular S0) galaxies (ETGs). The models include semi-analytic, N-body binary mergers and cosmological simulations of galaxy formation. Here we present the science goals for the project and introduce the galaxy sample and the selection criteria. The sample consists of nearby (D &lt; 42 Mpc, |δ- 29°| &lt; 35°, |b| &gt; 15°) morphologically selected ETGs extracted from a parent sample of 871 galaxies (8 per cent E, 22 per cent S0 and 70 per cent spirals) brighter than M<SUB>K</SUB> &lt; -21.5 mag (stellar mass M<SUB>★</SUB>≳ 6 ×10<SUP>9</SUP> M<SUB>⊙</SUB>). We analyse possible selection biases and we conclude that the parent sample is essentially complete and statistically representative of the nearby galaxy population. We present the size-luminosity relation for the spirals and ETGs and show that the ETGs in the ATLAS<SUP>3D</SUP> sample define a tight red sequence in a colour-magnitude diagram, with few objects in the transition from the blue cloud. We describe the strategy of the SAURON integral field observations and the extraction of the stellar kinematics with the pPXF method. We find typical 1σ errors of ΔV≈ 6 km s<SUP>-1</SUP>, Δσ≈ 7 km s<SUP>-1</SUP>, Δh<SUB>3</SUB>≈Δh<SUB>4</SUB>≈ 0.03 in the mean velocity, the velocity dispersion and Gauss-Hermite (GH) moments for galaxies with effective dispersion σ<SUB>e</SUB>≳ 120 km s<SUP>-1</SUP>. For galaxies with lower σ<SUB>e</SUB> (≈40 per cent of the sample) the GH moments are gradually penalized by pPXF towards zero to suppress the noise produced by the spectral undersampling and only V and σ can be measured. We give an overview of the characteristics of the other main data sets already available for our sample and of the ongoing modelling projects.

2537211

2011IAUS..276..225C

1012

1012.0584_arXiv.txt

The $15$ years since the discovery of the first exoplanet around a Solar-like star \citep{1995Natur.378..355M} have seen a revolution in our understanding of planet formation and evolution. Observations and theoretical modeling have worked hand-in-hand to discover and explain new architectures of planetary orbits. It is now well known that many exoplanets have large $e$ compared to our Solar system planetary orbits, indicating an active dynamical history \citep[e.g.,][]{Chatterjee08,Juric08,Nagasawa08}. Recent RM measurements of many transiting planets are putting further constraints on theoretical models of various planet formation and evolution scenarios \citep[e.g.,][]{2010arXiv1008.2353T,2010ApJ...718L.145W,2010arXiv1010.4025M}. Indeed, recent measurements find a large population of highly inclined planetary orbits, and even some retrograde orbits (see contributions from Winn et\ al. and Triaud et\ al. in this volume). Disk--planet migration models generally predict alignment between the planetary orbital angular momentum and the stellar spin axis from an aligned protoplanetary disk. Thus high-inclination orbits suggest dynamical evolution to be important in shaping the exoplanet architectures. Alternatively, inclined orbits might also result if the inner portion of the protoplanetary disk itself had been misaligned \citep[][see also Lai et\ al. in this volume]{2010arXiv1008.3148L}. Recent high-contrast imaging has revealed another class of planets---giant planets at very large $a$ \citep[$\gtrsim50\,\rm{AU}$; e.g.,][]{Fomb_Kalas,HR8799_Marois, 2008A&A...477L...1D,2010arXiv1011.2201I}. Timescale considerations for the core-accretion model of planet formation indicates that forming these planets in situ may be hard \citep[e.g.,][]{2001Icar..153..224L,2009ApJ...707...79D}. Instead, we consider formation at closer orbital separations followed by planet--planet scattering to launch planets in large-$a$ orbits from dynamically active systems \citep[e.g.][]{Chatterjee08,Juric08,Nagasawa08}. However, these simulations predict that these orbits typically also have high $e\gtrsim0.6$. Interestingly, some of the observed large-$a$ systems also have disks \citep[e.g.][]{Fomb_Kalas} and dynamical modeling of these disks indicates moderate values of $e\approx0.3$. We have started to explore the possibility that a planet launched into a large-$a$ and high-$e$ orbit can later be circularized near its apocenter if the planet enters a debris disk during its apocenter excursion. In Section\ \ref{inclination} we summarize our numerical setup and present results for expected inclinations from planet--planet scattering. In Section\ \ref{long} we discuss how planet--planet scattering followed by circularization due to a residual disk may create moderate-$e$ giant planets at large $a$. In Section\ \ref{conclude} we conclude.

\label{conclude} Using numerical simulations we have studied how planet--planet scattering can naturally create high-inclination orbits. We have studied whether giant planets can be launched into large-$a$, but modest-$e$ orbits, similar in architecture to Fomalhaut-b, via planet--planet scattering. We find that planet--planet scattering can naturally create many large-$a$ orbits, however, these orbits are also expected to have high $e$ \citep[e.g.,][]{Chatterjee08}. Nevertheless, if a massive outer disk exists and the scattered giant planet near its apocenter enters the disk, dynamical friction from the disk can raise the planet's $a$ until the planet reduces the disk density significantly via scattering. Then the planet can migrate inwards via scattering some of the remaining disk particles outwards. During the whole process the $e$ of the planetary orbit reduces. We plan to further study this process in detail exploring different disk masses, densities, as well as varying planet masses. We thank the SOC and LOC for arranging this excellent conference and the opportunity to present these results.

1012.0584

Recent observations have revealed two new classes of planetary orbits. Rossiter-Mclaughlin (RM) measurements have revealed hot Jupiters in high-obliquity orbits. In addition, direct-imaging has discovered giant planets at large (~ 100 AU) separations via direct-imaging technique. Simple-minded disk-migration scenarios are inconsistent with the high-inclination (and even retrograde) orbits as seen in recent RM measurements. Furthermore, forming giant planets at large semi-major axis (a) may be challenging in the core-accretion paradigm. We perform many N-body simulations to explore the two above-mentioned orbital architectures. Planet-planet scattering in a multi-planet system can naturally excite orbital inclinations. Planets can also get scattered to large distances. Large-a planetary orbits created from planet-planet scattering are expected to have high eccentricities (e). Theoretical models predict that the observed long-period planets, such as Fomalhaut-b have moderate e ~ 0.3. Interestingly, these are also in systems with disks. We find that if a massive-enough outer disk is present, a scattered planet may be circularized at large a via dynamical friction from the disk and repeated scattering of the disk particles.

12224390

2011PhRvD..83f3504B

1012

1012.0067_arXiv.txt

There is an abundance of evidence to indicate the existence of dark matter (DM), including its necessary contribution to both galactic stability and structure formation in the early Universe~\cite{Bertone:2004pz,Jungman:1995df,Liddle:1993fq}. The standard $\Lambda$CDM cosmological model, in which cold dark matter (CDM) makes up 22\% of the universal energy budget, provides an excellent description of our Universe. However, little is known about the particle properties of dark matter. In addition, some problems with CDM are encountered at small scales. A popular class of CDM candidates is weakly interacting massive particles (WIMPs). In WIMP models the DM couples weakly to standard model (SM) particles, which allows for scattering/annihilation processes. These serve to keep the dark sector in thermal equilibrium with the visible sector in the early Universe and can, with an appropriate choice of coupling, cause the DM to freeze out with the correct relic density. Interaction with the SM is similarly appealing from a detection standpoint, potentially providing both direct~\cite{Ahmed:2008eu,Angle:2007uj,Bernabei:2008yi} and indirect~\cite{Adriani:2008zr,Ackermann:2010rg,Knodlseder:2005yq} signatures of a given model. Nonobservation of these signatures allows for constraints to be placed on parameters such as particle masses or coupling constants. Though the DM mass is unknown, some information can be inferred from observations of large scale structure. For cold dark matter, structure forms hierarchically, with the earliest structures formed on short length scales, which can then merge to form larger structures. This is to be contrasted with hot dark matter in which the largest superclusters form first. Numerical simulation has shown the CDM scenario to fit observations well~\cite{Diemand:2006ik}, while hot dark matter is strongly disfavored. The CDM model is not-problem free, however, as it tends to overproduce small scale power \cite{Diemand:2006ik,Moore:1999gc,Navarro:2003ew,Gentile:2004tb,Salucci:2007tm,Diemand:2006ik,Gilmore:2006iy,Gilmore:2007fy,Gilmore:2008yp,Wyse:2007zw,Governato:2002cv,Klypin:1999uc,Metcalfe:2000nv,Peebles:2001nv,SommerLarsen:1999jx,Tikhonov:2009jq}. Simulations predict cusps in the DM density at the centers of galactic halos in conflict with observation. CDM also over-predicts the number of dwarf galaxies orbiting a Milky Way-sized galaxy by about a factor of $10$. Although simulations do not include visible matter, the gravitational potential wells they predict would promote a level of star formation not observed. Though these issues may be partially alleviated by tidal disruption and other effects, the small scale power problems of $\Lambda$CDM are still poorly understood (see e.g.~\cite{D'Onghia:2009pz,Schneider:2010jr} for recent work). Such issues have led many to consider a ``warm'' DM candidate, with a mass of keV scale, intermediate between hot and cold dark matter. In this work, as in \cite{Finkbeiner:2007kk,ArkaniHamed:2008qn,TuckerSmith:2001hy,Chang:2008gd,Zurek:2008qg,Feldman:2010wy,Profumo:2009tb,Cline:2010kv}, we will consider an alternative hypothesis in which the usual assumption of a single DM candidate is challenged. We consider a scenario with two WIMP candidates, in which one species is unstable to decay into the other. If the mass splitting between the two WIMPS is sufficiently small, the decay process will leave the overall halo mass unaffected, while giving its constituent DM particles a small velocity kick. Such velocity kicks heat the dark matter halos and cause them to expand, softening the central cusps and disrupting small halos~\cite{Melia,Peter:2010jy,Cembranos:2005us,Kaplinghat:2005sy,SanchezSalcedo:2003pb}. Such models are appealing, as they can alleviate the small scale structure problems, while retaining the attractive features of cold dark matter. We shall assume the DM decays predominantly via the channel \bea \x^*\rightarrow\x+l\,,\label{dec} \eea where $\x^*$ and $\x$ denote the heavier and lighter candidate, respectively, and $l$ is some relativistic final state. The mass splitting between $\x^*$ and $\x$ is given by \bea \D m=m_{\x^*}\e\,,\label{delm} \eea where $\e\ll1$. Abdelqader and Melia~\cite{Melia} have shown the dwarf halo problem can be solved for $\e\simeq(5-7)\times10^{-5}$ and a decay lifetime of $(1-30)$ Gyr. The work of Peter, Moody, and Kamionkowski~\cite{Peter:2010jy} has demonstrated that galaxy cusps can be alleviated for a wider range of $\e$ and $\t_{\x^*}$, with the most favored lifetimes in the range $(0.1-100)$ Gyr. Subsequent work by Peter and Benson~\cite{Peter:2010sz} has used properties of galactic subhalos to further constrain the allowed values of $\e$, preferring lower values to those favored in \cite{Melia}. Dark matter decays may be further constrained from analysis of their effect on weak lensing of distant galaxies as in~\cite{Wang:2010ma}. However, at present such analyses have only placed limits on models with much larger values of $\e$ than those considered in this work. An interesting possibility, from an observational standpoint, is a decay mode in which the relativistic final state, $l$, consists of SM particles. This allows the possibility of verifying the model, via the detection of particles produced by decay in our own Galaxy, or of a diffuse flux from decays in halos throughout the Universe. Current astrophysical observations place constraints on the allowed parameters, via comparison of the decay fluxes with relevant astrophysical backgrounds. Reference~\cite{Yuksel:2007dr} placed stringent constraints on the decay parameters for the case in which $l$ is a photon, while Ref.~\cite{Bell:2010fk} derived somewhat weaker constraints for the cases in which $l\,=\overline{\nu}\nu$ or $e^\pm$. For $l\,=\gamma$ or $e^\pm$ the lifetime is restricted to be below about 1 Gyr, while a much larger range of lifetimes is permitted for $l\,=\overline{\nu}\nu$ . The aim of this work is to construct a particle physics model which can realize the decaying dark matter scenario. We shall use the criterion specified by Abdelqader and Melia \cite{Melia} [namely, $\e\simeq(5-7)\times10^{-5}$ and $\t\sim(1-30)$ Gyr] as a reference point for these models, but given the constraints of Ref.~\cite{Peter:2010sz}, we will choose the more restrictive value of $\e\sim10^{-5}$ and $\t\simeq(1$-$10$) Gyr. In Sec.~\ref{sec:models} we introduce and discuss two possible models for decaying dark matter, and outline the DM production mechanism. Section~\ref{sec:constraints} focuses on constraints on the models and the available regions of parameter space. We conclude in Sec.~\ref{sec:conclude}.

\label{sec:conclude} Models for decaying dark matter are interesting in that they maintain the attractive features of the $\Lambda$CDM model, while alleviating the issues pertaining to the over prediction of small scale power. In this work we investigated two examples of the class of DM models in which decay occurs via the process $\x^*\rightarrow\x+l$, where $\x^*$ and $\x$ are nearly degenerate in mass (in this work we chose $\D m/m_{\x^*} \equiv \epsilon \simeq 10^{-5}$) and $l$ is relativistic. In the first scenario, we considered the possibility of decays into SM final states. We demonstrated that through the breaking of a discrete $\mathbb{Z}_4$ symmetry with the real scalar field $\f$, we could both produce two Majorana DM candidates $\x^*$ and $\x$ with nondegenerate mass, and allow for the decay channel $\x^*\rightarrow\x+\textrm{SM}$. The required long lifetime [(0.1-100) Gyr] was naturally achieved, as the the decay rate was suppressed by a high power of $\epsilon$, by small Yukawa couplings, and by the small mixing between SM-sector and dark-sector particles. The only two viable decay modes involving SM final states were $\x^*\rightarrow\x+e^+e^-$ and $\x^*\rightarrow\x+\overline{\n}\n$, where the latter is possible only in the case of Majorana neutrinos. We found that for DM masses in the range (50-$800$) GeV [$\D m\simeq(0.4-8)$ MeV] and for a $\f$-Higgs mixing of $\a\simeq10^{-5}$, all required criteria, including thermal production, could be met if the $\mathbb{Z}_4$ symmetry was broken at the MeV scale, with $\ang{\f}\simeq(3-20)$ MeV ($m_\f\simeq(2-20)$ MeV). Interestingly, in applying the constraints on decays to $e^+e^-$ derived in \cite{Bell:2010fk}, we showed that this final state is almost excluded for DM lifetimes in the (1-30) Gyr range preferred in \cite{Melia}. Dirac neutrinos were unable to fulfill the requirements for decaying DM, as their Yukawa coupling is too small. Thus decays to Majorana neutrinos are preferred by such DM decay models, as they are not constrained to either the short lifetimes applicable for $e^+e^-$ decays nor the small Yukawa couplings of Dirac neutrinos. In the second scenario, we considered the possibility of non-SM final states. This was achieved by replacing the discrete $\mathbb{Z}_4$ symmetry with a continuous $U(1)$ symmetry. Breaking of the $U(1)$ symmetry led to a pseudo-Nambu-Goldstone boson, which became the light final state produced in decays. As the DM decay process was no longer strongly suppressed, we were forced to finely tune model parameters to obtain a DM lifetime in the correct range. A consequence of this fine-tuning was to make DM production via mixing with the SM no longer possible. In this scenario, the dark and visible sectors are almost decoupled from each other. Though aesthetically less appealing, this model demonstrates the feasibility of decaying dark matter, independent of the strength of coupling to the SM. \bigskip

1012.0067

The standard cold dark matter cosmological model, while successful in explaining the observed large scale structure of the Universe, tends to overpredict structure on small scales. It has been proposed this problem may be alleviated in a class of late-decaying dark matter models, in which the parent dark matter particle decays to an almost degenerate daughter, plus a relativistic final state. We construct explicit particle physics models that realize this goal while obeying observational constraints. To achieve this, we introduce a pair of fermionic dark matter candidates and a new scalar field, which obey either a Z<SUB>4</SUB> or a U(1) symmetry. Through the spontaneous breaking of these symmetries, and coupling of the new fields to standard model particles, we demonstrate that the desired decay process may be obtained. We also discuss the dark matter production processes in these models.

12217130

2011PASP..123..130P

1012

1012.3408_arXiv.txt

IRXS J070407+262501 (hereafter RX0704) is a weak hard X-ray source coinciding with a 16th magnitude star (USNO 1125.04825852). \citet[hereafter G05]{gaensicke05} associated this star with the X-ray source, and described its properties: a high-excitation cataclysmic variable (CV) with an orbital period near 4 hours, and a strong optical pulse with a fundamental period of 480 s (and most of the power at 240 s, the first harmonic). This suggests membership among the DQ Hercuilis stars, or ``intermediate polars" as they are also called \citep[for reviews and rosters, see][]{patterson94,hellier96, mukai09}. This has now been confirmed by detection of X-rays pulsed with the same period \citep{anzolin08}, which is interpreted as the spin period of an accreting, magnetic white dwarf. The very high pulse amplitude caught our eye; we had been looking for a star with a pulsed signal sufficiently strong that small-telescope observers could track nearly every pulse -- and thereby accumulate a long and detailed observational record of the periodicity. This was partially successful, although the star is a little faint, and the pulse a little fast, to be tracked in fine detail by small telescopes. Still, we managed to learn the periods precisely, and to track long-term changes in pulse period. We report here on the results of that 2006-2010 campaign.

1. From radial velocities and photometry, we establish a precise binary period of 0.18208(5) d. This period is also manifest as a weak photometric signal, and as the lower orbital sideband ($\omega_{\rm spin} - \omega_{\rm orb}$) of the main pulse frequency, probably indicating a component reprocessed in structures fixed in the orbital frame. 2. The white dwarf spins with P = 480.6700 s, with most power at the first harmonic, very likely signifying a two-pole accretor. Oddly, the signal at the lower orbital sideband frequency (``reprocessed component") appears to be dominated by the fundamental. 3. The X-ray dip occurs 0.12 cycles before the broader optical minimum. 4. The spin period decreases on a timescale of $P/\dot P = 2.5 \times 10^6$ yr. 5. There may be a low-frequency photometric signal at P = 6.3 hours. Future observation should seek to clarify whether this is a common feature of the star. If confirmed, this periodic wave might be a long-period manifestation of a ``superhump", with the additional novelty that it occurs in a magnetic cataclysmic variable. Further theoretical exploration of this possibility is very desirable.

1012.3408

We present a study of the recently discovered intermediate polar 1RXS J070407 + 262501, distinctive for its large-amplitude pulsed signal at P = 480 s. Radial velocities indicate an orbital period of 0.1821(2) days, and the light curves suggest 0.18208(6) days. Time-series photometry shows a precise spin period of 480.6700(4) s, decreasing at a rate of 0.096(9) ms yr<SUP>-1</SUP> : i.e., on a time scale . The light curves also appear to show a mysterious signal at P = 0.263 days, which could possibly signify the presence of a superhump in this magnetic cataclysmic variable.

12167927

2011ApJ...729...13C

1012

1012.4466_arXiv.txt

Stars that wander too close to the MBHs that reside at the center of galaxies are shredded by the tidal gravitational field of the hole. After a tidal disruption event, about half of the debris are spewed into eccentric bound orbits and fall back onto the hole, giving rise to a bright UV/X-ray outburst that may last for a few years \citep[e.g.][]{rees88}. ``Tidal flares" from MBHs may have been observed in several nearby inactive galaxies \citep{komossa02,esquej07}. The inferred stellar disruption frequency is $\sim10^{-5}~{\rm yr^{-1}}$ per galaxy (with an order of magnitude uncertainty, \citealt{donley02}), comparable to the theoretical expectations for single MBHs fed by two-body relaxation \citep{wang04}. Yet, MBHs are not expected to grow in isolation. According to the standard paradigm of structure formation in the universe, galaxies merge frequently during the assembly of their dark matter halos. As MBHs become incorporated into larger and larger halos, they sink to the center of the more massive progenitor owing to dynamical friction from distant stars, and form bound binaries (MBHBs). In a purely stellar background, as the binary separation decays, the effectiveness of dynamical friction slowly declines, and the pair then ``hardens" via three-body interactions, i.e. by capturing stars that pass close to the holes and ejecting them at much higher velocities \citep[e.g.][]{begelman80,quinlan96,volonteri03,sesana06}. If the hardening continues sufficiently far, possibly driven by efficient stellar relaxation processes in a triaxial potential \citep[e.g.][]{merritt04} or in the presence of massive perturbers \citep[e.g.][]{perets07,perets08}, or by dissipative gaseous processes \citep[e.g.][]{colpi09}, gravitational radiation losses finally take over, and the two MBHs will coalesce in less than a Hubble time \citep[e.g.][]{mm05,sesana05,sesana07}. In \citet{chen09}, we used scattering experiments to show that gravitational slingshot interactions between a non-evolving, unequal-mass hard binary and a bound stellar cusp will inevitably be accompanied by a burst of stellar tidal disruptions. Our work differed from those by \citet{ivanov05}, who developed an analytical theory of the secular evolution of stellar orbits in the gravitational field of a MBHB, and by \citet{chen08}, who argued that stellar disruption rates by MBHBs fed by two-body relaxation would be smaller than those expected for single MBHs. Our numerical experiments revealed that a significant fraction of stars initially bound to the primary hole are scattered into its tidal disruption loss cone by resonant interactions with the secondary hole, close encounters that change the stellar orbital parameters in a chaotic way. In this paper we continue our investigations of stellar disruptions by MBHBs embedded in bound stellar cusps. We develop a hybrid model that self-consistently follows over time the shrinking of an MBH binary, the evolution of the stellar cusp, and the stellar disruption rate. The plan is as follows. In \S~\ref{background}, we introduce the basic theory of stellar disruption processes by MBHB systems. We describe our numerical scattering experiments in \S~\ref{experiments}, and discuss our results for different binary parameters as well as the effect of general relativistic corrections in \S~\ref{tests}. A detailed study of the properties of disrupted stars is carried out in \S~\ref{properties}. As a first step towards understanding the dependence of stellar consumptions on the parameters of the system, in \S~\ref{lcfr} we fix the binary semimajor axis and its eccentricity, and calculate the stellar disruption rate in the stationary case. In \S~\ref{sdr}, we present our hybrid model and calculate the disruption rates for an evolving, shrinking MBHB. Finally, we summarize and discuss our results in \S~\ref{dc}.

\label{dc} In this paper, we have studied the tidal disruption rate in a system composed by a MBHB and a bound stellar cusp. We have carried out numerical scattering experiments for a detailed investigation of the mechanisms responsible for the repopulation of the tidal loss cone, and developed a hybrid model to self-consistently solve for the evolutions of the binary, the depletion of the stellar cusp, and the stellar consumption rate. Our main results can be summarized as follows: \begin{enumerate} \item For unequal binaries ($q < 0.1$), the tidal disruption cross section for bound stars, which quantifies the probability of stellar disruption, is three orders of magnitude larger than the cross section for a single MBH fed by two-body relation. Two mechanisms contribute to such enhancement, the Kozai secular effect and chaotic resonant interactions. When the eccentricity of the MBHB is small, stars inside the Kozai wedge repopulate the tidal loss cone on the Kozai timescale, while stars outside the Kozai wedge but inside the interaction loss cone are scattered into the tidal loss cone at random times due to close interactions with the secondary hole. When the eccentricity is large, chaotic loss-cone repopulation becomes dominant over the entire range of stellar semimajor axis $a_*\ga (1-e)a$. \item GR and cusp-induced precession quench the Kozai secular evolution of interacting stars, causing a significant suppression (by a factor of $\sim 10$) of the disruption rate for $q < 0.01$. Therefore, the optimal enhancement of the tidal disruption rate by a MBHB occurs for mass ratios $0.01 < q < 0.1$. Note that even if suppressed by a factor of $\sim 10$, the tidal disruption rate for binaries with $q < 0.01$ is still two order of magnitude larger than that given by standard relaxation processes around a single MBH. \item If a MBHB with mass ratio $q\ll1$ does not evolve significantly during $1/q$ revolutions, tidal disruptions of bound stars could initially persist at a constant rate (``plateau phase") that is four dex higher than the typical rates predicted for single MBHs. After one Kozai timescale (evaluated at $a_*=a$), the tidal loss cone is repopulated mainly by chaotic interaction, and the stellar disruption rate decreases with time. The majority of stars are disrupted during a post-plateau later phase. \item If a MBHB evolves significantly on a timescale of $1/q$ revolution, the plateau phase of stellar disruptions may last longer than a Kozai timescale. Tidal disruptions of bound stars slow down the shrinking of the binary and speed up the growth of binary eccentricity. \end{enumerate} Our results indicate that, after the formation of an unequal-mass MBHB at the center of a dense stellar cusp, the tidal disruption rate may go through three distinct evolutionary phases. The first phase begins shortly after the MBHs become bound, and is characterized by a disruption rate as high as $0.1-1$ stars per year, resulting from the three-body interactions between the binary and the bound stars \citep{chen09}. When the decay timescale of the MBHB becomes longer than the tidal disruption timescales of stars with $a_*\sim a$, a second phase starts, where cusp depletion from slingshot ejections and tidal disruptions causes a sharp drop in the disruption rate. \citet{chen08} showed that, unless stellar relaxation is far more efficient than two-body ``collisions", the tidal disruption rate in this phase is orders of magnitudes lower than typical for a single MBHs. A third phase begins if the MBHB shrinks to the gravitational wave regime and eventually coalesces. The tidal disruption rate then gradually increases to the value typical for single MBHs, $10^{-5}-10^{-4}\,\ndot$, within one stellar relaxation timescale \citep{merritt05}. The number of stars disrupted during phase I is about $10^4-10^5$ for $M_7=1$ and $q=1/81$. The number of stars disrupted in phases II and III depends on the efficiency of stellar relaxation, but would not significantly exceed $\sim10^5-10^6$. If a galaxy formed, on the average, one unequal-mass MBHB following a minor merger in its lifetime, then the above numbers imply that in a sample of tidal flares from MBHs of $\sim10^7\msun$, about $10\%$ of events would be associated to binaries. If a galaxy forms unequal-mass MBHBs multiple times during its lifetime, then the detection rate of tidal events from binaries in transient surveys may be higher. Given the very high rates, there is also the possibility to identify MBHBs in galaxies hosting multiple tidal flares within a years-to-decades time span. Over the next decade, synoptic surveys are expected to detect hundreds of tidal disruption candidates. A tidal flare associated to a MBHB is likely interrupted within one orbital period of the binary \citep{liu09}, therefore is distinguishable from the flares produced by single MBHs, as long as the orbital period of the binary is shorter than the duration of a transient survey. If MBHB-driven disruptions account for $10\%$ of the total rate, then the prospects of identifying MBHBs through tidal flares are promising. Because the predicted rates in the three phases are significantly different from one another, the average stellar disruption rate over the lifetime of a galaxy is sensitive to the infalling rate of secondary MBHs and the relative duration of each phase. A comparison between the observational detection rate of tidal events \citep{donley02,gezari08} and those predicted during the three phases may then shed light on the abundance and dynamical evolution of MBHBs.

1012.4466

Tidal stellar disruptions have traditionally been discussed as a probe of the single, massive black holes (MBHs) that are dormant in the nuclei of galaxies. We have previously used numerical scattering experiments to show that three-body interactions between bound stars in a stellar cusp and a non-evolving "hard" MBH binary will also produce a burst of tidal disruptions, caused by a combination of the secular "Kozai effect" and by close resonant encounters with the secondary hole. Here, we derive basic analytical scalings of the stellar disruption rates with the system parameters, assess the relative importance of the Kozai and resonant encounter mechanisms as a function of time, discuss the impact of general relativistic (GR) and extended stellar cusp effects, and develop a hybrid model to self-consistently follow the shrinking of an MBH binary in a stellar background, including slingshot ejections and tidal disruptions. In the case of a fiducial binary with primary hole mass M <SUB>1</SUB> = 10<SUP>7</SUP> M <SUB>sun</SUB> and mass ratio q = M <SUB>2</SUB>/M <SUB>1</SUB> = 1/81, embedded in an isothermal cusp, we derive a stellar disruption rate \dot{N}_* ∼ 0.2 yr<SUP>-1</SUP> lasting ~3 × 10<SUP>5</SUP> yr. This rate is three orders of magnitude larger than the corresponding value for a single MBH fed by two-body relaxation, confirming our previous findings. For q Lt 0.01, the Kozai/chaotic effect could be quenched due to GR/cusp effects by an order of magnitude, but even in this case the stellar-disruption rate is still two orders of magnitude larger than that given by standard relaxation processes around a single MBH. Our results suggest that gsim10% of the tidal-disruption events may originate in MBH binaries.

12205269

2011IAUS..270..141K

1012

1012.1596_arXiv.txt

\label{sec:introd} \noindent The fundamental dynamical properties of stellar populations are the masses of the stars and their correlation in multiple stellar systems. The distribution of stellar masses at birth, the IMF, is rather well constrained and has (surprisingly) been found to be invariant despite theoretical models predicting systematic variation for example of the mean stellar mass or even of the minimum mass with the physical conditions of star formation. This problematical issue has been discussed at some length by \cite{K08}, where the {\sc IMF Universality Hypothesis} is stated. Equivalently, the question may be raised whether the other distribution functions characterising a stellar population, namely the distribution functions of binary systems, are just as invariant. If this were the case then it would have important bearings on the theory of star formation as the fragmentation length-scale may then not depend much on the physical conditions of the molecular cloud core. A change in the properties of the binary-star distribution functions with mass scale, if found, would yield important clues to the fragmentation and angular momentum re-distribution processes during star formation. The three important distribution functions describing the initial binary population (IBP) are the distribution of periods ($P$, here always in days), or equivalently of semi-major axes ($a$, in AU), the distribution of mass-ratios ($q=m_2/m_1 \le 1$) and the distribution of orbital eccentricities ($e$). These are related by Kepler's third law: $a^3/P_{\rm yr}^2 = m_1+m_2$, where $P_{\rm yr}$ is the orbital period in years ($P=365.25\,P_{\rm yr}$) and $m_1, m_2$ are the primary- and secondary-star masses in $M_\odot$. Because the periods of binary stars range over many orders of magnitude the shorthand $lP\equiv {\rm log}_{10}P$ is used throughout this text.

\label{sec:concs} \noindent It appears that the star-formation outcome in terms of stellar masses and multiple systems can be formulated by the {\sc Star Formation Universality Hypothesis} (\S~\ref{sec:sfuniv}). For stars with $m\simless 2\,M_\odot$ the {\sc Birth Binary Population} (\S~\ref{sec:IBP}) can be defined. This is the outcome of star formation in low to intermediate density ($\rho\simless 10^4\,M_\odot$/pc$^3$) cloud regions (e.g. of an embedded cluster). For $m\simgreat 2\,M_\odot$ stars the pairing rules change (\S~\ref{sec:IBP2}) perhaps reflecting the outcome of star formation in dense regions such as in the cores of embedded clusters ($\rho\simgreat 10^5\,M_\odot$pc$^3$). Brown dwarfs follow entirely separate rules (\S~\ref{sec:IBP2}) being an accompanying but distinct population to stars. It remains to be understood why these changing IBP properties do not correspond to the structure evident in the IMF (\S~\ref{sec:IBP-IMF}). \vspace{3mm} \begin{footnotesize} \noindent{\bf Acknowledgments:} I thank the organisers for a stimulating and memorable conference in Barcelona. My warmest gratitude I wish to express to Sverre Aarseth for his brilliant work on Nbody codes which are freely available and without which this work would not have been possible, and for his unwavering support. This text was written while being a Visitor at ESO/Garching, and I am thankful for the kind hospitality of my colleagues there. \end{footnotesize}

1012.1596

It is possible to extract, from the observations, distribution functions of the birth dynamical properties of a stellar population, and to also infer that these are quite invariant to the physical conditions of star formation. The most famous example is the stellar IMF, and the initial binary population (IBP) seems to follow suit. A compact mathematical formulation of the IBP can be derived from the data. It has three broad parts: the IBP of the dominant stellar population (0.08-2M<SUB>solar</SUB>), the IBP of the more-massive stars and the IBP of brown dwarfs. These three mass regimes correspond to different physical regimes of star formation but not to structure in the IMF. With this formulation of the IBP it becomes possible to synthesize the stellar-population of whole galaxies.

3191364

2011A&A...527A...1N

1012

1012.3841_arXiv.txt

The solution of the non-LTE multilevel-atom radiative transfer problem is a classical one in astrophysics. Indeed, the assumption of non-LTE implies, consistently with the departure of the source functions from Planck functions, that the population density of the atomic or molecular levels considered depart from what can be derived at LTE, in a straightforward manner, using Saha and Boltzmann relations (see e.g., Mihalas 1978). In the non-LTE case, one has on the contrary to solve simultaneously and self-consistently for a set of $N_{\rm T}$ equations of radiative transfer together with $N_{\rm L}$ equations of statistical equilibrium (hereafter ESE) describing the detailed balanced between excitation and de-excitation processes between every atomic or molecular levels. Since absorption and stimulated emission radiative rates depends explicitely on the radiation field, which itself depends on the level populations, this problem is intrinsically a search for the solution of coupled {\em nonlinear} equations. Since the beginning of numerical radiative transfer in the late 60's, the two most popular methods used for tackling this problem have been the complete linearization method of Auer \& Mihalas (1969) and the Accelerated $\Lambda$-Iteration based scheme called MALI (Rybicki \& Hummer 1991). Despite their apparent differences, they have however in common the fact that basically, one is conducted to deal with {\em linearized} equations. An interesting comparative study of these two approaches have been made by Socas-Navarro \& Trujillo Bueno (1997). In this study, we investigate on the use of a quasi-Newton numerical method for the solution of the nonlinear ESE. Our choice went to Broyden's method (1965) whose elements will be presented in \S2. To the best of our knowledge, Koesterke et al. (1992) were the first to bring this numerical scheme into the field of radiation transfer. Their study was presented in the context of the modelling of spherically expanding atmospheres of hot and massive Wolf-Rayet stars. Broyden's method was more recently invoked in the context of the coupled-escape probability method (Elitzur \& Asensio Ramos 2006). Besides from the required algebra and mention to caveats related to the implementation of the method, it remains however difficult to figure out from Koesterke et al. (1992) the actual performances of such an approach. A comparison with another method have also been barely evoked by the authors, who mentioned however a significant speed-up provided by Broyden algorithm for large $N_{\rm L}$ atomic models. In particular, being contemporary with the publication of Rybicki \& Hummer (1991), it is a pity that no comparison with the MALI method could be made yet. Such an evaluation is the scope of the present work.

We propose an alternative method for the solution of the non--LTE multilevel radiation transfer problem. It is based on Broyden's method for the solution of nonlinear systems of equations. The method is easy to implement and it is about of factor of 4.5 times faster than the well-known MALI method. Another advantage is that it does not require any modification of usual formal solvers, as it is the case for GS-SOR methods developed after MALI. It is also potentially very well-suited for parallel computing. Further tests will include the self-consistent treatment of the ionization balance, usually treated together with MALI with a Newton-Raphson scheme. In a next step, we shall consider more demanding models such as H$_{2}$O, for instance.

1012.3841

This study concerns the fast and accurate solution of multilevel non-LTE radiation transfer problems. We propose and evaluate an alternative iterative scheme to the classical MALI method. Our study is instead based on the application of Broyden's method for the solution of nonlinear systems of equations. Comparative tests, in 1D plane-parallel geometry, of the popular MALI method and our alternative method are discussed. The Broyden method is typically 4.5 times faster than MALI. This makes it also fairly competitive with the Gauss-Seidel and Successive Over-Relaxation methods developed after MALI.

2142315

2011JCAP...04..033M

1012

1012.1305_arXiv.txt

1012.1305

In a recent paper, Gurzadyan &amp; Penrose claim to have found directions on the sky centred on which are circles of anomalously low variance in the cosmic microwave background (CMB). These features are presented as evidence for a particular picture of the very early Universe. We attempted to repeat the analysis of these authors, and we can indeed confirm that such variations do exist in the temperature variance for annuli around points in the data. However, we find that this variation is entirely expected in a sky which contains the usual CMB anisotropies. In other words, properly simulated Gaussian CMB data contain just the sorts of variations claimed. Gurzadyan &amp; Penrose have not found evidence for pre-Big Bang phenomena, but have simply re-discovered that the CMB contains structure.

12131788

2010arXiv1012.2901F

1012

1012.2901_arXiv.txt

The Australian Square Kilometre Array Pathfinder (ASKAP; \cite{johnston:2007},\cite{westmeier:2010}), represents a significant advance in radio telescope design. This facility will combine high resolution imaging (through the use of a 36-element aperture synthesis array with a maximum baseline of $6$~km) with a wide field of view (achieved with innovative focal plane array technology) at frequencies between 700 MHz and 1.8 GHz. Installation of the first ASKAP antenna at the Murchison Radio Observatory site, Western Australia, occured in early 2010, and the 6-antenna BETA test array will operate from September 2011-March 2013. It is anticipated that full science operations will be underway by 2014. Processing and data transport requirements for ASKAP are described in \cite{cornwell:2008}, and \cite{quinn:2010} provides an overview of the data infrastructure requirements. WALLABY \cite{wallaby:url} is one of ten ASKAP Survey Science Projects currently in the design study phase. WALLABY aims to significantly enhance our understanding of the extragalactic neutral hydrogen (H{\sc i}) universe. The survey will cover 75\% of the sky, detecting $\sim\!0.5$ million galaxies to redshift $z = 0.26$ (lookback time $\sim\!3$ Gyr). Key science outcomes are studies of galaxy formation and the missing satellite problem in the Local Group, evolution and star formation in galaxies, mergers and interactions, determination of the H{\sc i} mass function and its variation with galaxy density, and the nature of the cosmic web. Unlike previous H{\sc i} surveys, it will not be feasible to keep all of the raw data (i.e. Fourier visibilities) from ASKAP observations for subsequent reprocessing. Instead, pipeline-preprocessed spectral data cubes will be provided for analysis. Each WALLABY spectral cube is anticipated to comprise 6144 $\times$ 6144 spatial pixels and 16,384 spectral channels (i.e. $\sim\,600$ gigavoxels or volume elements in total), requiring 2.5 terabytes (TB) of storage. A total of 1200 cubes will be required to cover the sky south of declination $\delta = +30^\circ$. Likely additional outputs are integrated (moment) maps, continuum images, sub-cubes (individual objects or scaled versions of larger datasets), and full parameterisation of all galaxies, resulting in several petabytes of data products. Such data volumes pose considerable challenges for the existing work practices of astronomers. Indeed, visualisation and analysis (hereafter, ``V+A'') of WALLABY data products will require both evolutionary and revolutionary changes to existing software and hardware, with a likely move away from desktop-only solutions, and a greater reliance on remote services. A brief overview of the WALLABY workflow from data collection to catalogue is as follows: \begin{enumerate} \item Observe field. \item Generate spectral data cube from visibilities. \item Visualise cube as quality control prior to deletion of raw data. \item Transfer preprocessed data cube to archive. \item Perform source finding on data cube. \item Fit models to candidates and perform related quantitative analysis. \item Add parameterised candidates to catalogue. \end{enumerate} Apart from personnel, the main resource for completion of all of these stages is access to appropriate computing infrastructure (hardware and software). As a framework within which to assess the practicalities of achieving each step in the WALLABY workflow, we begin (Section II) by considering desktop and high performance computing (HPC) cluster resources available and used by astronomers today, and project these forward to configurations available by 2014. In Section III, we present five challenges that V+A of WALLABY data products will face in the likely computing environment. We consider tasks that can be done essentially the same way they are now, and those requiring an investment in new technology or the development of new software, in order to deal with data sets orders of magnitude larger than previous extragalactic H{\sc i} survey projects. We make our concluding remarks in Section IV. Throughout, we make comparisons with the existing H{\sc i} Parkes All Sky Survey (HIPASS; \cite{barnes:2001}), conducted with the Parkes Multibeam receiver \cite{staveley:1996}. The southern catalogue, HICAT ($\delta < +2^{\circ}$; \cite{meyer:2004}), comprised 4315 galaxies, and the northern extension, NHICAT ($+2^{\circ} < \delta < +25^{\circ}30'$; \cite{wong:2006}), a further 1002 sources. Russell Jurek (Australia Telescope National Facility; ATNF) has combined the 388 individual southern sky data cubes into a single all-sky cube with $1721 \times 1721 \times 1025 = 3$ gigavoxels, and a file size of 12 GB.

Perhaps the biggest challenge to planning strategies for visualisation and analysis is that no ASKAP data exists yet. We do not know what the exact imaging properties of ASKAP will be. Although simulated data cubes are now being generated, until the full ASKAP system undergoes commissioning, we will not fully understand all of the calibration, noise, interference, etc. issues that will arise with the relatively new technology of focal plane arrays. Testing source finders often includes injecting fake sources, with a given signal level, and then seeing how often they are recovered. With real WALLABY data cubes unavailable until 2014, progress in testing source finders will necessarily be limited. While we can do our best to plan source finders based on existing datasets, and early science data from the BETA configuration (September 2011-March 2013), we may find that our techniques do not work adequately for the full dataset. By considering the various V+A tasks now, and identifying approaches based on new hardware and software that were not available or feasible for earlier surveys, we can hope to minimise the impact of the ``unknown unknowns'' of ASKAP. Graphics processing units offer an intriguing solution to a number of the current desktop-bound problems. Table IV summarises our thoughts regarding the visualisation and analysis tasks that will require either an evolution of existing software and hardware, or a revolution in how they are approached. By planning today, we aim to maximise the scientific return from WALLABY tomorrow.

1012.2901

Visualisation and analysis of terabyte-scale datacubes, as will be produced with the Australian Square Kilometre Array Pathfinder (ASKAP), will pose challenges for existing astronomy software and the work practices of astronomers. Focusing on the proposed outcomes of WALLABY (Widefield ASKAP L-Band Legacy All-Sky Blind Survey), and using lessons learnt from HIPASS (HI Parkes All Sky Survey), we identify issues that astronomers will face with WALLABY data cubes. We comment on potential research directions and possible solutions to these challenges.

12203124

2011crpa.conf...58T

1012

1012.0135_arXiv.txt

Recently several experiments have pointed out interesting features in the measured energy spectrum of the $e^+$ and $e^-$ components of cosmic rays. In particular, we recall here the positron excess detected by Pamela \cite{pam}\!, the excess of electron counts ($e^+ + e^-$) between 100 GeV and 1 TeV detected by Fermi\cite{fermi}\!, and the bump in the cosmic ray electron (CRE) spectrum, centered around $\sim 500$ GeV, detected by ATIC\cite{atic}\! (not confirmed by H.E.S.S. low energy analysis\cite{hess})\!. The observational picture will be probably improved after the observations of AMS-02\cite{ams}\! (which will be operating in February 2011), due to the large acceptance and accuracy of the apparatus and to the long operation time (at least ten years). In this work we do not look for possible interpretations of these data, rather we study the possibility of using radio observations to confirm the above CREs spectral features. In particular, we concentrate on the evaluation of CRE synchrotron emission in a galactic environment, with typical values of the magnetic field strength of $1-5$ $\mu$G. Due to the explorative aims of this study, we limit our attention to the intensity and degree of polarization of radiation produced by an ensemble of electrons, in the hypothesis of homogeneity of the galactic magnetic field on scales greater than the gyro-radius, and discuss the properties of radiation at the source, not taking into account radiative transfer through the interstellar medium (ISM), which affects both intensity and degree of polarization. This issue can be the object of a subsequent paper. Here we point out the relevant frequency bands for possible observations and signatures that may allow one to disentangle their signal from contaminants like thermal dust emission, peaking in the sub-mm range. We show also that expected signals are not a limiting foreground for Cosmic Microwave Background experiments, being expected at $\sim$ 1 THz and above. Finally we discuss a few observational aspects of some astrophysical sources which could be used as targets for this investigation.

We have shown that the features in synchrotron intensity and polarization produced by $>$ 100 GeV electrons fall into the sub-mm wave/far IR regime (therefore, they are not an issue when removing foregrounds from CMB maps). Moreover a bump in the electron spectrum (ATIC) modulates significantly (more than 10$\%$) the degree of synchrotron linear polarization at frequencies around 1 THz. Unfortunately, the Galactic dust, through its grey-body emission, $ I_d(\nu) = k_d (\nu / \nu_0)^{\beta_d} BB(T_d)$, completely dominates the sky brightness in these bands. In Fig.\ref{aba:fig4} we show the expected brightness of a clean region of the sky normalized (through $k_d$ emission coefficient) at the DIRBE \cite{dirbe}\! 100 $\mu$m channel. We assumed a dust temperature $T_d = 20$ K, and grey-body emissivity scaling as $\nu^{\beta_d}$ ($\beta_d = 1.7$). We see that dust signal overcomes the synchrotron one by several orders of magnitude. \begin{figure}[h!] \begin{center} \psfig{file=fig4.eps, width=4.5in} \caption{Synchrotron brightness corresponding to different electron spectra (same line style as in fig.\ref{aba:fig2} and \ref{aba:fig3}.). The filled circles correspond to the emission of cold (20 K) dust in a low contamination region of the sky, normalized to the measurements of the 100 $\mu$m DIRBE channel.} \label{aba:fig4} \end{center} \end{figure} \noindent This is just an example, since in general dust contamination depends on the region observed (different temperature, density and composition of dust). Let's now consider different sources of the observed CRE excess. \emph{A) Dark Matter decay or annihilation in our galaxy.} The distribution of synchrotron radiation is probably similar to the expected Dark Matter distribution in the galactic halo. We expect a smooth distribution with a maximum of intensity in direction of the galactic center decreasing away from it. The best directions of observations are far from the galactic disk, where the synchrotron emission of the background electrons and the thermal emission of the interstellar dust are faint and more uniform. This is the situation we have considered to estimate the signals shown in Fig.\ref{aba:fig4}: above 100 GHz thermal dust emission is the dominating signal also far from the Galactic disk. Therefore it is difficult to imagine a detection of a diffuse, smoothly varying synchrotron signal coming from our own Galaxy, such as what could be associated to a Dark Matter halo. In this case, also considering the typical polarization signatures of synchrotron radiation, the removal of the thermal background would be extremely difficult, despite the small degree of linear polarization of thermal dust emission. \emph{B) Electrons from Pulsar or SNR.} A galactic source like a pulsar or a SNR can be found (much more easily) on the galactic disk. Here both the synchrotron emission of the background electrons and the thermal emission of the interstellar dust are stronger respect to the signals coming from the halo and shown in Fig.\ref{aba:fig4}. Dust emission is expected to increase more than synchrotron. In addition the angular distribution is also more anisotropic. This situation is compensated by a possible \emph{boost factor} enhancing the synchrotron signal coming from these sources, if CRE come out along the direction of observation. In fact electrons radiate only towards the speed vector, and we detect their radiation only if this direction is aligned with our line of sight. Besides electrons originated by these sources come to the Earth position through a diffusion process, because the gyro-radius, in the interstellar magnetic field even at energy up to 1 TeV, is much smaller than the distance of the closest candidate sources. In addition around these sources the magnetic field is far from uniform. This \emph{boost factor} could be large up to the Lorentz factor $\gamma$, but we must take into account also that the detected CRE could come from sources not radiating in the direction of our line of sight. In conclusion one could map a region surrounding a galactic source, like a pulsar or a SNR, looking for an anisotropic signal at small angular scales. In this case the background subtraction could be easier, in particular if we look at the spectral feature in the polarized signal. \emph{C) Electrons in extragalactic sources.} An alternative approach is to observe extragalactic radio sources: spiral or elliptical galaxies. In fact we expect the same phenomenology, regarding cosmic rays and synchrotron radiation, in external galaxies as in our own. Also regions surrounding AGNs could be used as target to observe features in the synchrotron radiation (which here is much more intense), but our understanding of these sources is still uncomplete and CRE producing synchrotron radiation are not observable. \emph{(1)} Spiral galaxies should have synchrotron and dust emissions similar to the Milky Way, and should have a similar population of pulsars and SNR in the disk. We can look for an anisotropic signal against the background at small angular scale, but the presence of the thermal dust emission does not facilitate the background removal. In this case we can not decouple the effect generated by local sources from a Dark Matter signature. \emph{(2)} Elliptical galaxies show a low dust and gas content. This means that both thermal emission and SNR generated by supernova explosions in a low density ISM are no longer important contaminants. Therefore these galaxies can be used as favorite targets for looking at a signature of Dark Matter annihilation or decay from the galactic halo. In order to investigate the features we pointed out in this paper, spectral and polarimetric information, together with good angular resolution, would help. A spectro-polarimeter, or a polarimeter operating in different photometric bands, installed in the focal plane of a large far-IR telescope would be necessary. Because of the frequencies under investigation a space or a balloon-borne experiment could be preferred. To have a more realistic estimate of the signals, in all the situations considered above, detailed calculations have to be done, including radiative transfer effects, in particular concerning polarized signals.

1012.0135

Recent measurements of cosmic ray electron energy spectra suggest that above 10 GeV there may be deviations from a single power law spectrum. There are hints (ATIC) for a bump occurring between 100 GeV and 1TeV, meaning that there might be more high energy electrons than expected. Whether these electrons are produced within pulsar magnetospheres, or due to Dark Matter annihilation or decay, this is still matter of debate. Understanding the nature of these ultra high energy particles is a difficult task that can be fulfilled using all the available astrophysical observables. We investigate how different energy spectra produce different observable manifestations in the radio/microwave/mm-wave domain, where corresponding deviations from a synchrotron power law could appear. We raise the question around the detectability of these possible radio spectral features, which may be interesting for a wide scientific community including astrophysicists and scientists working on foregrounds removal for CMB experiments.

1778550

2011MNRAS.415.3807S

1012

1012.0842_arXiv.txt

Galactic haloes are an excellent testbed for cosmology and galactic dynamics. Their exploration can constrain the early assembly of galaxies as well as the dynamics of accretion of smaller galaxies. Our Milky Way offers an ideal case for those investigations, as we can directly obtain the detailed parameters like kinematics, elemental abundances and physical properties of single stars surrounding us. New material is still being accreted into the Galactic halo, as the numerous streams and newly discovered dwarf galaxies confirm \citep[see e.g.][]{Ibata95, Klement06, Belokurov07}. As metallicity gradients go in lockstep with star formation, young accreted objects may be more metal-poor than old stars from the inner Galaxy and thus metallicity can not be simply used as a cosmic clock. This would also make it plausible that a later accreted halo component could indeed be on average younger and more metal poor than the older parts. The formation of at least parts of the halo by accretion (combined with later adiabatic contraction) could give rise to differences between early more turbulent accretion/collapse processes and later accretion, which might leave an imprint in differences between the inner and outer halo \citep[see e.g.][]{Cooper10}. Another possible source of discrepancies between inner and outer halo is dynamical friction, which could be more efficient for prograde than for retrograde infall \citep{Quinn86, Byrd86}. This could give rise to a different rotational signature for accreted material in the outer Galactic halo compared to the inner regions \citep{Murante10}. Historically (and as well today), halo stars have been extremely difficult to identify, particularly in local samples, e.g. demonstrated by the historic argument between Oort and Str\"omberg \citep[][]{Oort26, Strom27}. Like it is practically impossible to get a clean selection into thin and thick Galactic disc based on kinematics \citep[][]{SB09b}, we face the analogous problem between thick disc and the prograde stars of the halo. Wrong assumptions about the kinematics of the Galactic disc will thus affect results on the halo component. Soon after the existence of the Galactic halo was established \citep[][]{Schwarzschild52, Eggen62}, the central question was raised if the stars of the inner and the outer halo had the same properties or if gradients or even breaks in metallicities or kinematics existed with galactocentric radius. Two main strategies to identify and examine halo stars have been used in the past: either stars in the solar neighbourhood are studied, classified according to their kinematics and then conclusions about the structure further away are drawn by extrapolation \citep[e.g.][]{jesper}, or the surveys concentrate on bright objects in the outer halo regions, such as RR Lyrae variables or globular clusters \citep[e.g.][]{Sandage70}. The second alternative allows to directly map the spatial structure by those standard candles with good distance information \citep[e.g.][]{Saha85}. This strategy implies selection biases: for example the position and presence of RR Lyrae stars on the horizontal branch are correlated with metallicity and age, while it is not known if the formation of globular clusters is representative also for all halo field stars. Claims of differences between the inner and outer Galactic halo are almost as old as the discovery of the halo itself. After \cite{Bergh67} discussed differences in metallicity and the second parameter between the haloes of the Milky Way and those of its neighbouring galaxies (M31, M33), \cite{Searle78} found that Galactic clusters in the outer regions showed a larger scatter in the ratio of blue to red horizontal branch stars than inner halo globular clusters, which they interpreted as a signature of an age spread. \cite{Preston90} found a similar difference in field BHB stars.\footnote{While it is clear that age is one parameter which will cause an older globular cluster to be bluer than a younger one at the same metallicity, whether this is the dominant cause of differences in horizontal branch morphology is still debated (see, for example, \citet{Dotter10} and \cite{VandenBerg00} for opposing views.)} Differences in kinematics have also been suggested between inner and outer halo globular clusters \citep{zinn93}, although precision and reliability of estimates in this respect are limited by the small number of available globular clusters. Various claims of an asymmetry in the halo azimuthal velocity distribution with an extended tail to retrograde orbits have been made \citep[e.g.][]{Norris89}. \cite{Majewski92} even found the entire halo to be on average counterrotating, could, however, not find any significant velocity gradient. \cite{Ryan92} pointed out that measurements of kinematics based on proper motions were particularly vulnerable to distance errors and showed that overestimated distances for halo stars can lead to false identifications of counter-rotating stars. In this paper we will revisit the recent claim by \cite{Carollo07} (hereafter C07) and \cite{Carollo10} (hereafter C10) that the Galactic halo consists of two components: a more metal-poor counterrotating component with larger scaleheights, distinct from a slightly prograde component and starts dominating the halo at high altitudes in their analysis.\footnote{This dominance of a retrograde component in the outer halo has recently been contested by \cite{Deason10} although they find a retrograde motion for metal-poor ($\feh < -2$) and prograde ($\feh > -2$) for metal-rich stars using a BHB sample in the outer Galactic halo. In our view, this issue needs to be further investigated, as they assume constant $g$ band magnitudes for the horizontal branch in a region affected by the blue tail, which spans of order $2 \mag$ in $g$ band luminosity. If a considerable fraction of the halo giants is in the blue tail, their colour and temperature cuts remove a large part of this tail, but still leave $BHB$ members spanning $\sim 0.7$ magnitudes, as we tested it on SDSS photometry of metal-poor globular clusters known to have such a strong blue tail ($M3$, $M13$, $M15$ and $M92$) and the BASTI isochrones.} In particular we will carefully re-examine their distance estimation procedure; we will focus on C10 as this paper deviates from its precursor mostly by the larger sample size. To avoid relying on any of the uncertain available distance calibrations, we apply in parallel both the C10 distances and two native SDSS main sequence distance calibrations, checking results additionally with an isochrone method. In Section \ref{sec:segue} we outline those methods, discuss the SDSS/SEGUE data used for this purpose and describe how the sample cleaning was performed. Thereafter (Section \ref{sec:grav}) we discuss the underlying assumptions and the reliability of gravity estimates used to sort stars into different sequences as well as the actual assumptions for absolute magnitudes by C07 and C10. We will show that their claim to have distances precise to $10-20 \%$ is unsupported and that the C10 sample contains a class of stars with significant distance overestimates by being sorted into unphysical positions in the HR diagram. In Section \ref{sec:sign} we present statistical proofs of distance biases in the sample and in Section \ref{sec:vel} we discuss the implications of different distance schemes on kinematics and the inner-outer halo dissection.

We have described how errors in distance estimates result in an apparent systematic retrograde motion of the Galactic halo, an effect to which the SEGUE/SDSS sample is especially prone by its strong poleward orientation. The general problem of distance biases similarly applies to any study that makes use of proper motion-based estimates. We find that the distance derivation of \cite{Carollo07} and \cite{Carollo10} is flawed by sorting stars into unphysical positions in the HR diagram: objects are placed between the subgiant and dwarf sequences in positions that would require stellar ages in excess of the age of the universe. Despite the elegance of the general idea to sort stars into known sequences according to their estimated gravities, the method itself and the used gravity cuts are not well supported by measurements. Moreover there is no "turn-off"-sequence, but turn-off stars are populating a region that spans of order $1$ magnitude in luminosity. In this light the statement by C10 to have distances precise to about $10-20 \%$ is an unsupported claim. From the distances kindly provided by Carollo et al. we calculated back to their assumed absolute magnitudes and found systematic differences of $\sim 0.2$ to $0.3 \mag$ and a large scatter for metal-poor main sequence stars towards the adopted main sequence calibration as well as towards the age-dependent calibration by \cite{Ivz08}, also far to the red side of the suspected turn-off region. The adopted main sequence calibration is only slightly fainter than the theoretical BASTI isochrones. We have shown in Section \ref{sec:HR} that the claim by C07 and C10 to have found a counter-rotating extended tail of the halo is largely caused by unphysical assumptions about locations of stars in the HR diagram, by magnitude uncertainties in the turn-off stars and by the use of a too bright main sequence calibration. The correctness of this tail can be ruled out by statistical tests, as described in Section \ref{sec:sign}. The tail diminishes when we limit the C10 sample to dwarf stars and disappears when we make use of the better founded \cite{Ivz08} calibrations, which are consistent with fiducials and isochrones. In Section \ref{sec:comp} we demonstrated that the halo distribution for the dwarf samples regardless of the applied distance determination can be fit by a simple Gaussian component. We have also shown that in the DR7 pipeline stars with lower metallicities are shifted towards lower gravities, considerably increasing their fraction among the thought-to-be turn-off stars. Further the magnitude difference for their main sequence stars against the \cite{Ivz08} main sequence calibration grows towards lower metallicities. The stronger prevalence of distance errors at the metal-poor end of the metallicity distribution will thus give any spurious counter-rotating tail members a biased metallicity distribution. Finally we have shown that the claim of C10 that the counter-rotating component members reach to higher altitudes can as well be traced back to distance determinations: the dispersion of the vertical velocity component is significantly increased by their distance errors, though due to the polewards sample orientation this effect would at first order be smaller than for the other velocities. As shown in Section (\ref{sec:sign}) simultaneously the $W$ velocities of the halo stars with distance overestimates are artificially increased by of order $50 \kms$ via spurious velocity cross-over terms from the heliocentric azimuthal velocities in the derivation. The effect is strongest for the most strongly retrograde objects (they have the largest heliocentric $V_h$) and is aggravated by their selection for stars in galactocentric radius $7 \kpc < R < 10 \kpc$. This colludes with their metallicity dependent distance bias (see above) to produce their findings of decreasing metallicities at high altitudes. There is a slight excess of more metal-poor stars in a single velocity bin at high vertical velocities, which is not mirrored by the behaviour at high total kinetic energies. We argue that this is most likely a reflection of a well-known local stream that has been identified by \cite{Helmi99} and \cite{Kepley07}. Another source of error is the modelling of especially the Galactic disc azimuthal velocity distribution by Gaussian components. It was shown by \cite{Strom27} that Gaussian modelling of the Galactic disc lead to unphysical results and the identification of spurious components on the low rotation side because of the extented tail. As the skewed $V$ velocity distribution enforces in most cases the introduction of a second Gaussian component, that - being a mere artifact by wrong assumptions - can then be misinterpreted as physical reality, Gaussian modelling of the Galactic disc in a combined disc and halo sample can wrongly force the halo component into the prograde regime to compensate for the two steeply falling disc terms. Consequently this then creates the need for inference of a retrograde component to compensate for the bias. We also argue that magnitude based distance assessment schemes introduce a velocity bias that resembles the Lutz-Kelker bias: If the error in absolute magnitudes follows a Gaussian distribution, the distance error distribution will thus form an extended tail that grows stronger with increasing dispersions. Via the proper motion part in the determination of space velocities, which is proportional to the estimated distance, measured velocities develop extended tails away from the solar motion. For the $V$ velocity distribution of especially the halo this gives the halo an asymmetric velocity distribution with a longer tail in the retrograde regime, a process that can explain the moderate asymmetries found e.g. by \cite{Norris89}. Finally we note that it is by no means imperative that the halo have a Gaussian velocity distribution. In this light it is rather surprising that our simple Gaussian halo component can fit the data so well. Even if there were deviations from Gaussianity this would alone be no convincing sign for a separate component. A more convincing indication would be a proven difference between the prograde and the retrograde tail of the halo azimuthal velocity distribution, but disentangling this from disc contamination on the prograde side will be difficult. Recently Carollo and collaborators submitted a rebuttal paper \citep[][]{B11} claiming that our analysis presented here be wrong. Instead of discussing all our arguments their revised analysis relies on two central claims: They state that the distance scale adopted by us is wrong and that there is an asymmetric halo azimuthal velocity distribution for their metal-poor stars, neither of which we concur with. Concerning the distance issue we stress that our conclusions are valid for the Carollo et al. (2007, 2010) and both \cite{Ivz08} distance calibrations; our work does not rely on a single distance scale in contrast to claims made by \cite{B11}. Beers et al. criticize us for adopting the incorrect main sequence calibration of \cite{Ivz08} but failed to note that we actually stretched this calibration in the same direction of their preferred one by increasing the luminosities by $0.1$ magnitudes and accounting for alpha-enhancement by increasing the measured metallicities by $0.2$ dex (Sect. 2). Importantly, we have also made use of their preferred \cite{Ivz08} calibration (here denoted IvzA7) and find no significant differences (e.g. Fig. 11). Finally, we have made use of direct isochrone distances, which fully corroborates our findings. As for the azimuthal velocity distribution of the most metal-poor stars, we re-emphasize that any magnitude-based distance scheme invokes a bias in the inferred distances and thus an asymmetric azimuthal velocity distribution by definition; this effect is akin to the well-known Lutz-Kelker (1973) bias and is illustrated in \figref{fig:velos2}. In view of a large magnitude scatter (which their metal-poor stars clearly have, see their Fig. 5) a Gaussian fit is inappropriate due to the missing error handling. In this light it is not surprising that their new revised parameters are quite different from their original results (e.g. for $z_{max} > 5 \kpc$ their „outer halo“ mean velocity rose from $-128 \kms$ (cf. Table 1 in C10) to $-59 \pm 20$ km/s). Finally, we argue that moving a considerable fraction of the wrongly identified turnoff stars up to the subgiant/giant branch as done by Beers et al. will make the distance overestimate for misidentified dwarfs among them even more severe. In summary our criticism of the C07, C10 works remains in full: our in-depth re-analysis of their data with different distance calibrations and a proper error handling reveal no convincing evidence for a dual halo. The systematic distance uncertainties make it dangerous to draw a definitive conclusion for the strength or existence of a possible counter-rotating halo component. All current distance calibrations have problems and need improvement before a method along the lines used by C07 and C10 can be attempted. We would like to stress that we do not and would not want to rely on either of them. Two central conclusions can, however, be drawn without having to trust any of the different distance calibrations: Even on the C10 or C07 sample using their distances, no reliable detection of any non-Gaussianity in the halo, be it a counter-rotating halo or not, is possible on the examined data set in any of the dwarf samples. If a separate component gets detected on a larger sample in the future, it should be significantly weaker than what was claimed by C10.

1012.0842

We examine the kinematics of the Galactic halo based on SDSS/SEGUE data by Carollo et al. We find that their claims of a counter-rotating halo are the result of substantial biases in distance estimates (of the order of 50 per cent): the claimed retrograde component, which makes up only a tiny fraction of the entire sample, prone to contaminations, is identified as the tail of distance overestimates. The strong overestimates also result in a lift in the vertical velocity component, which explains the large altitudes those objects were claimed to reach. Errors are worst for the lowest metallicity stars, which explains the metal-poor nature of the artificial component. We also argue that measurement errors were not properly accounted for and that the use of Gaussian fitting on intrinsically non-Gaussian Galactic components invokes the identification of components that are distorted or even artificial. Our evaluation of the data leads to a revision of the estimated velocity ellipsoids and does not yield any reliable evidence for a counter-rotating halo component. If a distinct counter-rotating halo component exists, then it must be far weaker than claimed by Carollo et al. Finally, we note that their revised analysis presented in Beers et al. does not alleviate our main concerns.

12168041

2011ApJ...731..123K

1012

1012.1180_arXiv.txt

\label{sec-intro} Exoplanets detected using both transits and the radial velocity (RV) technique offer the unique opportunity to measure many of their physical properties. Transits can place limits on the radius of an exoplanet given the radius of the star, and transits can also constrain orbital parameters such as inclination and orbital period. When coupled with RV measurements, we can constrain the mass ($M_{p}$), and the average density of a planet ($\rho_{p}$). We may therefore probe the interiors of transiting exoplanets and constrain bulk composition and formation models \citep{fortney08,torres08,baraffe08}. Here we present and interpret three lightcurves of the transiting planet GJ 1214b \citep{charbonneau09}, a planet unlike any in our Solar System. Prior to 2009, all transiting exoplanets were found with radii consistent with the giant planets in our Solar System, \ie\ they had gaseous envelopes. As surveys improved, they began to reach sensitivities which could detect smaller rocky planets. In our Solar System, Neptune has a mass of 17 \mearth while the Earth is the largest terrestrial planet, suggesting that the transition mass between rocky and gaseous planets lies somewhere between these values. With no such ``transition'' object in the Solar System, we must rely on theory, which predicts that $\sim 10$ \mearth\ is the critical mass \citep{pollack96}, although it could be as low as 2 \mearth \citep{ikoma01} or as high as 16 \mearth \citep{lissauer09}. The 10 \mearth\ limit should therefore be seen as a sort of median of theoretical results, and not a true boundary between rocky and gaseous worlds. The discovery of transiting planets between 1 and 17 \mearth therefore provides critical insight into the planet formation process. Two transiting planets are now known with 1\mearth $\le M_{p} \le$ 10 \mearth; CoRoT-7b \citep{leger09,queloz09}, and GJ 1214b \citet{charbonneau09}. The former appears to be rocky \citep{leger09}, while the latter may contain significant amounts of water or a gaseous envelope \citep{rogersseager10}. In order to address the ambiguities in the planet formation process, the radii of these planets must be measured accurately and precisely. The CoRoT satellite already surpasses all other projects in the ability to make follow-up observations of CoRoT-7b (153 transits reported by \citet{leger09}). The relatively recent detection and hence the dearth of follow-up measurements on GJ 1214b led us to focus on this object; a Super-Earth planet orbiting M dwarf star 13 pc from Earth \citep{charbonneau09}. The planetary nature of this transit has been confirmed by \citet{sada10}. The discovery of such a small-sized planet bodes well for transit searches for habitable planets around M-Dwarfs. The low luminosity of M-Dwarfs mean their habitable zones are very close to the star \citep{kasting93,selsis07}. This increases the transit probability \citep{boruckisummers84}, and hence some of the first planets to be characterized as rocky and in the habitable zone may well be transiting planets around M-Dwarfs. Although GJ 1214b most likely possesses a Hydrogen-rich envelope, the detection of small planets such as this heralds the discovery of rocky habitable worlds if they exist. M-Dwarfs make up a very large fraction of the stellar component of the Milky Way \citep{MS79,reid02}, so the prospect of habitable planets around M-Dwarfs raises the very interesting possibility that life-bearing planets may be fairly common in the galaxy. In this paper we report observations of three transits of GJ 1214b in 2010. In $\S$ \ref{sec-data} we outline our observations and data reduction techniques; in $\S$ \ref{sec-syspars} we describe our lightcurve model and the use of Markov Chain Monte Carlo (MCMC) techniques to constrain system parameters from single and multi-wavelength data. In $\S$ \ref{sec-ttv} we report the absence of strong transit timing variations (TTV). In $\S$ \ref{sec-mr} we describe the derivation of various stellar and planetary parameters for the GJ 1214 system. In $\S$ \ref{sec-stellaractivity} we discuss the detection of a flare and a possible spot-crossing event. Finally, in $\S$ \ref{sec-conclusions} we summarize our findings.

\label{sec-conclusions} \N \textbf{A Transit Model Suited for Bayesian Analysis:} We show that fitting for the transit duration ($t_T$) and the ingress/egress duration ($t_G$) results in a parameter set with few mutual degeneracies (see Figure \ref{figure_mcmc_1}). This condition is suited very well for MCMC methods, which are regularly used to determine uncertainties on parameters derived from transit lightcurves. Our joint analysis of multi-wavelength data using this parameter set was able to reproduce previous estimates of system parameters for GJ 1214b (see Table \ref{table_pars2}, chain003a). We also find that milli-magnitude photometry may not be sufficient to constrain limb-darkening parameters using transit lightcurves. We show that MCMC runs where we fit for these parameters were slow to converge (see Table \ref{table_mcmcstats}), and posterior probability distributions for various parameters were plagued with degeneracies (see Figure \ref{figure_mcmc_2}). Estimates of system parameters from these runs were generally unreliable when compared to runs where the limb-darkening parameters were kept fixed (see chains `b' vs. `a' in Tables \ref{table_pars1} and \ref{table_pars2}). \N \textbf{Transit Timing Variations:} Data gathered so far do not indicate significant variations in the times of transit for GJ 1214b (see Figure \ref{figure_ttv}). APOSTLE will continue making observations of GJ 1214b and a more detailed analysis of timing data will follow in a future paper. \N \textbf{System Parameters for GJ 1214:} From fitting SEDs to photometry, we constrained GJ 1214's observed flux and luminosity ($\S$ \ref{sec-mr}). The luminosity allowed us to constrain GJ 1214's mass and since we obtained stellar density from transit lightcurves it allowed us to estimate GJ 1214's radius. We find the derived values of mass and radius to be in agreement with previous estimates, however we find GJ 1214 deviates from well-known mass-radius relations for low-mass stars (see Figure \ref{figure_mrplot}). Simple calculations using the formalism presented in \citet{chabrier07} show that GJ 1214's position on the mass-radius plot can be explained by the presence of cool regions on its surface. From RV, transit data \citep{charbonneau09} and absolute stellar properties we determined various properties of GJ 1214b (see Table \ref{table_pars3}). The planetary mass and radius (6.37$\pm$0.87 M$_{\earth}$, 2.74$_{-0.05}^{+0.06}$ R$_{\earth}$) places GJ 1214b between the terrestrial and ice-giant regime of planets (2M$_{\earth} <$ M$_{p}$ $<$ 10M$_{\earth}$). Its classification as a ``Super-Earth'' remains and the planetary density confirms it is not like the rocky bodies of our solar system (see Table \ref{table_pars3}). \citet{rogersseager10} propose 3 scenarios for the origin of its gaseous envelope: i) primordial H/He, ii) sublimated ices (H$_2$O,CO$_2$) or iii) volcanic outgassing. \citet{millerrfortney10} propose that space-based observations of the transmission spectra of GJ 1214b's atmosphere should be able to tell us how Hydrogen-rich its atmosphere is. The largest source of uncertainty in our estimate of planetary mass was the velocity semi-amplitude ($K$). Errors in the planetary radius follow from our uncertainty in measuring the absolute size of the star, which ultimately hinges on our luminosity estimate (see $\S$\ref{sec-mr}). Improved precision on radial velocity, flux and distance would tighten our constraints on the absolute mass and radius of GJ 1214b. \N \textbf{Evidence for Stellar Activity:} The detection of a low energy stellar flare and the possible transit of the planet over a star-spot (see Figures \ref{figure_flare} and \ref{figure_spot}) indicate that GJ 1214 is active. However, considering its age and comparing the flare energy to flares on the younger AD Leo confirms that GJ 1214 is a quiet star for its spectral type \citep{hawley96,hawley03}. We find a fast-rise exponential decay profile fits the flare signal (UTD 2010-04-21) quite well. A symmetric rise in the normalized flux ratio during the transit on UTD 2010-06-06 could indicate the planet occulted a star-spot on the surface of GJ 1214. The signal is weak, and our attempt at fitting a simplified spot model shows that spot properties are difficult to constrain from a single spot-crossing observation. Our results on spot-properties are inconclusive due to degeneracies between spot-size, planet-to-spot impact parameter and spot-to-star contrast ratio in our model. Detections of this signal from successive transits would have confirmed it as a star-spot and provided interesting constraints on the properties of an active stellar surface region \citep{dittmann09}. The stellar rotation rate might have also been estimated with such data.

1012.1180

We present three transits of GJ 1214b, observed as part of the Apache Point Observatory Survey of Transit Light Curves of Exoplanets. By applying Markov Chain Monte Carlo techniques to a multi-wavelength data set which included our r-band light curves and previously gathered data of GJ 1214b, we confirm earlier estimates of system parameters. Using spectral energy distribution fitting, mass-luminosity relations, and light curve data, we derived absolute parameters for the star and planet, improving uncertainties by a factor of two for the stellar mass (M <SUB>sstarf</SUB> = 0.153<SUP>+0.010</SUP> <SUB>-0.009</SUB> M <SUB>sun</SUB>), stellar radius (R <SUB>sstarf</SUB> = 0.210<SUP>+0.005</SUP> <SUB>-0.004</SUB> R <SUB>sun</SUB>), planetary radius (R<SUB>p</SUB> = 2.74<SUP>+0.06</SUP> <SUB>-0.05</SUB> R <SUB>⊕</SUB>), and planetary density (ρ<SUB> p </SUB> = 1.68 ± 0.23 g cm<SUP>-3</SUP>). Transit times derived from our study show no evidence for strong transit timing variations. We also report the detection of two features in our light curves which we believe are evidence for a low-energy stellar flare and a spot-crossing event.

12168117

2011ApJ...733...31H

1012

1012.0075.txt

\label{sec:intro} In order to explain observations of massive galaxy evolution in the universe, current models of galaxy formation require a form of energetic feedback that is thought to result from the effects of a central active galactic nucleus (AGN) \citep{croton2006,somerville2008,dimatteo2008}. Energy and momentum input from the AGN into the galaxy's interstellar medium (ISM) can serve to heat or remove gas so that it is no longer available for star formation. ``AGN feedback'' has been presented as one of the major factors giving rise to the red sequence in massive galaxies \citep{silk1998,kauffmann2003b}. Furthermore, every galaxy bulge appears to contain a supermassive black hole \citep{kormendy1995}, whose mass is correlated with bulge properties such as stellar velocity dispersion \citep{merritt2001,gultekin2009}. These correlations offer evidence of coupling between the formation of the black hole and bulge, which may result from the effects of AGN feedback \citep{silk1998,murray2005,hopkins2008}. Outflows have been observed in strongly star-forming galaxies over a range of redshifts \citep{franx1997,shapley2003,martin2005,steidel2010,rupke2005}, and post-starburst galaxies at $z \sim 0.6$ \citep{tremonti2007}. For AGNs, outflows have been observed in local Seyfert galaxies \citep{crenshaw2003,krug2010}, and at higher redshifts in broad absorption line quasars \citep{korista2008,ganguly2008}, radio galaxies \citep{nesvadba2006,nesvadba2008}, and ULIRGs \citep{alexander2010}. However, at early times, outflows have thus far not been fully examined in a sample of active galaxies that can be quantitatively compared to a non-active sample with similar host galaxy properties. This paper examines the outflow properties of such a sample at $2 \leq z \leq 3$, when both star-formation density and black hole (BH) accretion were at their peak \citep{madau1996, ueda2003, richards2006, reddy2008, silverman2008}. The rest-frame UV portion of a galaxy spectrum is ideally suited for the study of the ISM. In star-forming galaxies, this spectral region contains emission or absorption from \ion{H}{1} Lyman $\alpha$, as well as low- and high-ionization metal absorption lines that have been used to infer the presence of outflows \citep{pettini2000,pettini2001,pettini2002,shapley2003}. At $z \sim 2 - 3$, the rest-frame UV part of the spectrum is shifted into the observed optical, and is accessible using ground-based facilities. At these redshifts, individual galaxy spectra have low continuum signal-to-noise (S/N), which makes robust absorption line measurements challenging. With a large enough sample, however, a higher S/N composite spectrum can be created, allowing measurements of the global properties and spectral trends within the sample. \citet{shapley2003} and \citet{steidel2010} have used such composite spectra to explore the outflow properties of UV-selected star-forming galaxies at $z \sim 2 - 3$. Because the black hole accretion disk and broad-line region are obscured from view, the light from a narrow-lined AGN is not dominated by emission from the central source, but rather that of the host galaxy. The ability to study the host galaxy allows for a comparison between the galaxy-scale properties of a sample of narrow-lined AGNs and those of a similar non-AGN sample. In order to undertake such a study, we augment the sample of narrow-lined UV-selected AGNs at $z \sim 3$ presented in \citet{steidel2002}, extending it to include objects at $z \sim 2$. The original sample enabled, for the first time, an estimation of the fraction of star forming galaxies within the Lyman Break Galaxy (LBG) survey that showed evidence for AGN activity on the basis of their rest frame UV spectroscopic properties. With the expanded sample of AGNs, we construct composite spectra that reveal the properties of outflowing gas in these objects. The galaxies that harbor these narrow-lined AGN were selected on the basis of their broadband rest-frame UV colors, and should have host galaxies similar to those of the non-AGN LBGs. As this AGN sample appears to be hosted by galaxies drawn from the same parent population as the non-AGN LBG sample \citep{steidel2002,adelberger2005d}, we can conduct a controlled experiment to understand how the AGN impacts the gas properties of the host galaxy. The sample of UV-selected AGNs is presented in \S \ref{sec:sample}, while in \S \ref{sec:generating} we describe the creation of the AGN composite spectrum. This spectrum and its basic properties are shown in \S \ref{sec:features}, including the detection of blueshifted high-ionization absorption features. In \S \ref{sec:trends}, we examine spectral trends within the AGN sample that are highlighted by separating objects according to Ly$\alpha$ equivalent width (EW), UV magnitude, and redshift. We conclude in \S \ref{sec:discussion} with a discussion placing the results from the composite spectrum analysis into the context of our understanding of AGNs. Throughout our analysis, we assume $\Omega_\mathrm{M} = 0.27$, $\Omega_\Lambda = 0.73$, and $H_0 = 71$ km s$^{-1}$ Mpc$^{-1}$.

\label{sec:discussion} \begin{deluxetable*}{lrrrrrr} \tabletypesize{\scriptsize} %\rotate \tablecaption{Spectroscopic Properties of Composite Spectrum Subsamples\label{tab:binning}} \tablewidth{0pt} \tablehead{ & \colhead{$W_{Ly\alpha} > 63$} & \colhead{$W_{Ly\alpha} < 63$} & \colhead{$M_{UV,bright}$} & \colhead{$M_{UV,faint}$} & \colhead{$z > 2.7$} & \colhead{$z < 2.7$} } \startdata $\mathrm{N_{gal}}$ & 16\phm{,0000} & 17\phm{,0000} & 16\phm{,0000} & 17\phm{,0000} & 9\phm{,00000} & 24\phm{,0000} \\ $\mathrm{\langle z_{Ly \alpha}\rangle \tablenotemark{a}}$ & 2.60$\pm$0.08\phm{0} & 2.49$\pm$0.06\phm{0} & 2.65$\pm$0.08\phm{0} & 2.43$\pm$0.05\phm{0} & 2.94$\pm$0.07\phm{0} & 2.37$\pm$0.02\phm{0} \\ $\mathrm{\langle M_{UV} \rangle \tablenotemark{a}}$ & $-$20.6$\pm$0.2\phm{00} & $-$20.7$\pm$0.2\phm{00} & $-$21.2$\pm$0.1\phm{00} & $-$20.2$\pm$0.1\phm{00} & $-$21.1$\pm$0.2\phm{00} & $-$20.5$\pm$0.1\phm{00} \\ $\mathrm{\langle} W_{\mathrm{Ly\alpha}} \mathrm{\rangle \tablenotemark{a}}$ & 141$\pm$15\phm{.00} & 33$\pm$4\phm{.000} & 78$\pm$14\phm{.00} & 82$\pm$17\phm{.00} & 124$\pm$28\phm{.00} & 62$\pm$8\phm{.000} \\ $\mathrm{\langle \beta_{G-{\cal R}} \rangle \tablenotemark{a}}$ & $-$0.8$\pm$0.2\phm{00} & $-$1.2$\pm$0.1\phm{00} & $-$1.1$\pm$0.2\phm{00} & $-$0.9$\pm$0.2\phm{00} & $-$1.1$\pm$0.4\phm{00} & $-$1.0$\pm$0.1\phm{00} \\ $\mathrm{\langle \beta_{spec} \rangle \tablenotemark{a}}$ & $-$0.5$\pm$0.3\phm{00} & $-$0.6$\pm$0.3\phm{00} & $-$0.9$\pm$0.3\phm{00} & $-$0.3$\pm$0.3\phm{00} & $-$0.9$\pm$0.4\phm{00} & $-$0.4$\pm$0.2\phm{00} \\ $\mathrm{\beta_{spec}\tablenotemark{b}}$ & $-$0.1$\pm$0.4\phm{00} & $-$0.5$\pm$0.3\phm{00} & $-$0.7$\pm$0.3\phm{00} & $-$0.1$\pm$0.4\phm{00} & $-$0.7$\pm$0.5\phm{00} & $-$0.2$\pm$0.3\phm{00} \\ \\ $W_{\mathrm{Ly\alpha}}\tablenotemark{c}$ & 123$\pm$14\phm{.00} & 28$\pm$5\phm{.000} & 62$\pm$18\phm{.00} & 56$\pm$16\phm{.00} & 103$\pm$37\phm{.00} & 52$\pm$8\phm{.000} \\ $W_{\mathrm{NV,1240}}\tablenotemark{c}$ & 6.23$\pm$1.54\phm{0} & 4.42$\pm$1.01\phm{0} & 3.22$\pm$0.84\phm{0} & 8.21$\pm$1.80\phm{0} & 4.15$\pm$1.66\phm{0} & 5.83$\pm$1.12\phm{0} \\ $W_{\mathrm{NIV],1484}}\tablenotemark{c}$ & 2.13$\pm$0.86\phm{0} & 1.68$\pm$0.53\phm{0} & 1.05$\pm$0.52\phm{0} & 2.51$\pm$1.17\phm{0} & 3.07$\pm$1.36\phm{0} & 1.57$\pm$0.58\phm{0} \\ $W_{\mathrm{CIV,1549}}\tablenotemark{c}$ & 25.21$\pm$4.15\phm{0} & 7.99$\pm$2.56\phm{0} & 10.17$\pm$3.84\phm{0} & 16.40$\pm$4.15\phm{0} & 40.43$\pm$8.06\phm{0} & 12.01$\pm$2.89\phm{0} \\ $W_{\mathrm{HeII,1640}}\tablenotemark{c}$ & 9.16$\pm$1.88\phm{0} & 5.48$\pm$1.71\phm{0} & 6.01$\pm$1.23\phm{0} & 8.13$\pm$3.09\phm{0} & 14.03$\pm$5.29\phm{0} & 6.27$\pm$1.41\phm{0} \\ $W_{\mathrm{CIII],1909}}\tablenotemark{c}$ & 20.85$\pm$8.00\phm{0} & 6.06$\pm$2.66\phm{0} & 14.08$\pm$5.90\phm{0} & 15.62$\pm$6.74\phm{0} & 36.75$\pm$15.58 & 8.75$\pm$2.90\phm{0} \\ \\ $W_{\mathrm{SiII,1260}}\tablenotemark{c}$ & -\phm{00000} & $-$1.81$\pm$0.49\phm{0} & $-$0.87$\pm$0.34\phm{0} & -\phm{00000} & -\phm{00000} & $-$1.80$\pm$0.39\phm{0} \\ $W_{\mathrm{OI+SiII,1303}}\tablenotemark{c}$ & $-$2.00$\pm$0.82\phm{0} & $-$1.96$\pm$0.54\phm{0} & $-$2.36$\pm$0.47\phm{0} & -\phm{00000} & -\phm{00000} & $-$2.27$\pm$0.54\phm{0} \\ $W_{\mathrm{SiIV,1393}}\tablenotemark{c}$ & -\phm{00000} & $-$1.72$\pm$0.74\phm{0} & $-$1.66$\pm$0.55\phm{0} & -\phm{00000} & -\phm{00000} & $-$1.34$\pm$0.67\phm{0} \\ $W_{\mathrm{SiII,1527}}\tablenotemark{c}$ & -\phm{00000} & $-$1.03$\pm$0.53\phm{0} & $-$0.85$\pm$0.35\phm{0} & -\phm{00000} & -\phm{00000} & $-$0.79$\pm$0.35\phm{0} \\ \enddata \tablenotetext{a}{Sample average values for the composite spectra.} \tablenotetext{b}{UV-continuum slope, measured from the composite spectra.} \tablenotetext{c}{Rest-frame EW in \AA, measured from the composite spectra. Positive values indicate emission, while negative values indicate absorption. Uncertainties are calculated as described in \S \ref{sec:obsdata}. This table includes absorption and emission measurements with greater than 2$\sigma$ significance.} \end{deluxetable*} Analysis of the composite spectrum of the UV-selected AGNs at $z \sim 2 - 3$ reveals a number of results regarding the nature of AGN activity. We report the detection of weak absorption lines from both low- and high-ionization species. Most strikingly, the high-ionization \ion{Si}{4} absorption feature exhibits a significant blueshift of $\Delta v = -845\pm178$ km s$^{-1}$. The precise value of this blueshift is referenced to our estimate of the rest frame based on \ion{He}{2} $\lambda$1640. As discussed in \S \ref{sec:absfeat}, if we adopt a rest frame in which the low-ionization lines have the same blueshift as those in the non-AGN LBG composite spectrum of \citet{shapley2003}, then the magnitude of the inferred blueshift for \ion{Si}{4} would be even greater. While contamination from \ion{Si}{4} emission prevents us from tracing out the full \ion{Si}{4} velocity profile in absorption, the most strongly blueshifted material appears to be outflowing more rapidly than the associated gas in star-forming galaxies at $z \sim 2 - 3$ \citep{shapley2003,steidel2010}. \citet{thacker2006} presents results of a simulation that indicate that pure star formation, even when the full kinetic energy from each supernovae is applied to the outflowing material, can only produce maximum outflow velocities of roughly $v = 600$ km s$^{-1}$. Their modeling shows that only AGNs and quasars can produce high speed outflows with velocities greater than $10^{3}$ km s$^{-1}$. Outflows of this magnitude have also been seen in poststarburst galaxies at $z = 0.6$, using \ion{Mg}{2} $\lambda \lambda$2796, 2803, which are claimed to result from the effects of an AGN \citep{tremonti2007}. Previous studies of outflows observed in galaxies with AGNs can be used to place these UV-selected AGN results in context. \citet{krug2010} presented a study of outflows from a sample of local IR-faint AGN. For the narrow-lined objects, the outflow velocities (as calculated from the \ion{Na}{1} D interstellar absorption line doublet) are on the order of those from starburst galaxies, offering a conclusion that star formation was the process driving the outflows in Seyfert 2 systems. Based on a sample of local infrared-luminous starburst galaxies exhibiting AGN activity, \citet{rupke2005} show evidence for high velocity superwinds, which they compare to those from a non-AGN ULIRG sample presented in \citet{rupke2005b,rupke2005c}. Both have comparable outflow velocities, leading to the conclusion that the momentum and energy required for the outflow could have come equally from a starburst or the AGN. The current work represents the same type of differential comparison between AGNs and their non-active counterparts, but at high redshift, when the black hole and bulge are both actively forming. This analysis highlights the specific effects of the AGN on the outflowing ISM. Relative to work on AGN outflows at high redshift, which has focused on the extended line emission in individual systems or small samples of AGNs alone \citep{alexander2010, nesvadba2008}, the benefit of our analysis lies in our comparison to a control sample of non-AGN star-forming galaxies. In order to understand how the outflows observed in our sample of AGNs will ultimately affect the galactic gas content, a calculation of the mass outflow rate of the gas is needed. Such a calculation requires knowledge of the outflowing gas metallicity, column density, covering fraction and physical location with respect to the illuminating source. Specifically, it is necessary to determine whether the outflowing gas extends over the scale of the entire galaxy or is confined to the scale of the central engine. Based on current data, we cannot obtain a precise estimate of the location of the gas, metallicity, column density, or covering fraction. Without determinations of these properties, a full comparison to AGN feedback models \citep[e.g.][]{thacker2006} cannot be made. The \ion{N}{4}] $\lambda$1486 emission line is detected in our AGN composite. This feature is not detected in the non-AGN LBG composite \citep{shapley2003}, and is observed only rarely in broad-lined quasars in the Sloan Digital Sky Survey \citep{bentz2004,jiang2008}. At the same time, \ion{N}{4}] has been detected in the spectra of high redshift radio galaxies \citep{vernet2001,humphrey2008}. In order to determine the origin of the \ion{N}{4}] emission, \citet{humphrey2008} consider both photoionization and shock models. A comparison of model predictions for line ratios such as \ion{N}{4}] / \ion{C}{4} and \ion{N}{4}] / \ion{He}{2} with those measured in our composite spectrum suggests that the observed \ion{N}{4}] originates in photoionized gas. Our observed line ratios indicate that the gas is of solar or supersolar metallicity and subjected to a hard ionizing spectrum ($f_\nu \propto \nu^{\alpha}$, where $\alpha \geq -1.0$) with ionization parameter $U \geq 0.05$, where $U$ is defined here as the ratio of ionizing photons to H atoms at the surface of the model photoionized gas slab. On the other hand, shock excitation models cannot explain our high observed values of \ion{N}{4}] / \ion{C}{4}. The EW of Ly$\alpha$ in the UV spectra of our AGNs is indicative of both the strength of the AGN as well as the properties of star-forming regions. The observed Ly$\alpha$ EW is further modulated by various radiative transfer effects, due to its high scattering cross section. In non-AGN LBGs, the strength of Ly$\alpha$ emission has been shown to correlate with the EW of low-ionization interstellar absorption lines, such that strong Ly$\alpha$ emission is accompanied by weaker interstellar absorption \citep{shapley2003}. This result can be understood if the escape of Ly$\alpha$ photons is at least partially modified by the covering fraction of neutral gas in the ISM. Additionally, \citet{kornei2010} show that stronger Ly$\alpha$ emission is coupled with smaller dust obscuration as traced by the slope of the UV continuum, indicating that interaction with dust preferentially destroys Ly$\alpha$ photons, a result consistent with previous work by \citet{shapley2003}, \citet{pentericci2007}, and \citet{verhamme2008}. When separating our objects according to Ly$\alpha$ EW to create composite spectra, we find that the composite spectrum created from objects with large Ly$\alpha$ EW shows stronger \ion{C}{3}], \ion{C}{4} and \ion{He}{2} emission than the composite spectrum created from objects with smaller Ly$\alpha$ EWs. As shown in Table \ref{tab:binning}, the \ion{C}{3}], \ion{C}{4}, and \ion{He}{2} lines are two to three times weaker in the $W_{Ly\alpha} < 63$ composite spectrum than in the $W_{Ly\alpha} > 63$ composite spectrum. This result suggests that the strength of Ly$\alpha$ emission is modulated at least partially by the level of AGN activity, which is traced by the strength of these other emission lines. However, the fact that the Ly$\alpha$ EW is almost 5 times weaker in the weak Ly$\alpha$ EW composite spectrum indicates additional suppresion of Ly$\alpha$ photons beyond the reduced level of AGN activity. At the same time, the low-ionization interstellar absorption lines that indicate the covering fraction of cool gas are significantly stronger in the weak Ly$\alpha$ composite. Therefore, it is not only the strength of the underlying AGN that separates the objects by Ly$\alpha$ EW, but also the covering fraction of gas that might absorb and reradiate the Ly$\alpha$ emission. This trend agrees with the results of \citet{shapley2003} for non-AGN LBGs. In our AGN sample, the UV continuum of the strong Ly$\alpha$ composite is redder than that of the weak Ly$\alpha$ composite, though the difference is not significant. This discrepancy with the trends among non-AGN LBGs from \citet{shapley2003} and \citet{kornei2010} may be a result of our small sample size, or potentially because the Ly$\alpha$ flux we observe originates from both the general star-forming ISM as well as the nuclear region. These two sources of Ly$\alpha$ photons may have disjoint properties with respect to the geometry of dust extinction, suppressing the trend observed among the non-AGN LBGs. Future modeling of the spectral energy distributions of the AGN host galaxies (Hainline et al. in prep.) will allow for analysis of the UV spectra separated by stellar mass and E(B-V), as well as uncover spectral trends as a function of galaxy evolutionary state. We also will probe the origin of the strikingly red UV continuum slopes found in the narrow-line UV-selected AGN spectra.

1012.0075

We present new results for a sample of 33 narrow-lined UV-selected active galactic nuclei (AGNs), identified in the course of a spectroscopic survey for star-forming galaxies at z ~ 2-3. The rest-frame UV composite spectrum for our AGN sample shows several emission lines characteristic of AGNs, as well as interstellar absorption features detected in star-forming Lyman break galaxies (LBGs). We report a detection of N IV] λ1486, which has been observed in high-redshift radio galaxies, as well as in rare optically selected quasars. The UV continuum slope of the composite spectrum is significantly redder than that of a sample of non-AGN UV-selected star-forming galaxies. Blueshifted Si IV absorption provides evidence for outflowing highly ionized gas in these objects at speeds of ~10<SUP>3</SUP> km s<SUP>-1</SUP>, quantitatively different from what is seen in the outflows of non-AGN LBGs. Grouping the individual AGNs by parameters such as the Lyα equivalent width, redshift, and UV continuum magnitude allows for an analysis of the major spectroscopic trends within the sample. Stronger Lyα emission is coupled with weaker low-ionization absorption, which is similar to what is seen in the non-AGN LBGs, and highlights the role that cool interstellar gas plays in the escape of Lyα photons. However, the AGN composite does not show the same trends between Lyα strength and extinction seen in the non-AGN LBGs. These results represent the first such comparison at high redshift between star-forming galaxies and similar galaxies that host AGN activity. <P />Based, in part, on data obtained at the W. M. Keck Observatory, which is operated as a scientific partnership among the California Institute of Technology, the University of California, and NASA, and was made possible by the generous financial support of the W. M. Keck Foundation.

12204334

2011EPJWC..1104004H

1012

1012.2183_arXiv.txt

In searching for potentially habitable planets, M stars present the most promising targets. Because of their small masses, these stars have the greatest reflex acceleration due to an orbiting planet. The low surface temperatures of these stars place their (liquid water) habitable zones at distances of approximately 0.1 to 0.2 AU (corresponding to orbital periods of $\sim 20$ to 50 days) where the precision radial velocity surveys are normally at their optimal sensitivity. Given that within the Sun's immediate neighborhood, more 70\% of stars are of spectral type M, it is not surprising that for more than a decade, these stars have been the subject of research by many authors (Joshi et al. 1997; Segura et al. 2005; Boss 2006; Scalo et al. 2007; Grenfell et al. 2007; Tarter et al. 2007). In the past few years, such research has resulted in the detection of 25 extrasolar planets around 17 M stars. Slightly more than half of these planets are Neptune-mass or smaller, consistent with the fact that M stars have smaller circumstellar disks and their planets are less massive compared to those of G stars. Among these planets are the first Neptune-mass object around the star GJ 436 (Butler et al. 2004), the first Earth-size planet around the star GJ 876 (Rivera et al. 2005), and the recently discovered Earth-like planet in the habitable zone of the star GL 581 (Vogt et al. 2010). Although majority of currently known planets around M stars have been detected using the radial velocity technique, these stars have also been targets of transit photometry searches. The MEarth project, a robotically controlled set of eight 40 cm telescopes at Whipple observatory on Mt. Hopkins in Arizona, is a transit photometry survey that is dedicated to detecting M stars. This program has been successful in discovering a 6.6 Earth-mass planet around M star GJ 1214 (Charbonneau et al. 2009). The transit timing variation method has also been considered as a mechanism for detecting small planets around M stars. As shown by Kirste \& Haghighipour (2009, 2011), the variations in the transit timing of a transiting giant planet due to the perturbation of an Earth-size body or a super-Earth can be large enough to match the temporal sensitivity of {\it Kepler} space telescope. Figures 1 and 2 show samples of the results by these authors. As shown in figure 1, an Earth-size planet in a 10-day orbit around a 0.32 solar-mass star produces strong TTVs on a transiting Jupiter-mass planet when the two objects are in (1:2), (2:3), (5:2), and (2:1) mean-motion resonances. Figure 2 shows the mean-motion resonances for which an Earth-like planet in the habitable zone of an M star will produce TTVs of the order of 10 s or larger on a transiting Jupiter-like body. Although the calculations by Kirste \& Haghighipour (2009, 2011) point to the detectability of terrestrial planets in systems studied by these authors, the low masses of circumstellar disks around M stars cast doubt in the existence of their assumed planetary configurations. Computational simulations have indicated that circumstellar disk around M stars are not massive enough to accommodate the formation of giant planets, even in orbits as large as that of Jupiter around the Sun (Laughlin et al. 2004). The fact that observational surveys have been able to detected many Jovian-type planets around M dwarfs [e.g. GJ 876 with two Jupiter-like planets and a Uranus-mass body in approximately 30, 60, and 120 days orbits (Rivera et al. 2005, 2010), or HIP 57050 with a Saturn-mass planet in a 42 days orbit (Haghighipour et al. 2010)] suggests that these giant planets were probably formed at larger distances, where the disk contained more material, and migrated to their current short-period orbits. It would therefore be necessary to study how such a migration affects the formation of terrestrial planets around M stars and their final orbital configuration as the giant planet approaches short-period orbits. \vskip -3pt \begin{figure}[ht] \begin{center} \epsfig{width=12cm,file=fig1.eps} \caption{Transit timing variations of a 1 Jupiter-mass planet around a 0.32 solar-mass M star. The perturber is an Earth-sized planet in a 10-day orbit. The graph shows the values of TTVs for different ratios of the orbital periods of the two planets. As shown here, when the two planets are in (1:2), (2:1), (5:3), and (2:3) resonances, the TTVs have values larger than 100 sec. Figure from Kirste \& Haghighipour (2011).} \end{center} \end{figure} \begin{figure}[ht] \begin{center} \epsfig{width=12cm,file=fig2.eps} \vskip -10pt \caption{Graph of transit timing variation for different resonances between an Earth-size perturber and a Jupiter-mass transiting planet. The central star is a 0.36 Solar-mass M star. The habitable zone (HZ) and continuous habitable zone (CHZ) of the star are shown.The shaded area corresponds to TTVs larger than 10 sec. A shown here, (1:3), (1:2), and (3:5) resonances produce large TTVs when the perturber is in the continuous habitable zone.} \end{center} \end{figure}

Simulations results indicate that although for both disks models and all migration rates, terrestrial planets were formed in the protoplanetary disk, they did not maintain stability and were ejected from the system. The time of the ejection is inversely proportional to the rate of giant planet migration. Our study suggests that if a terrestrial planet is detected in resonance with a transiting giant planet around an M star, 1) the terrestrial planet is unlikely to have formed in-situ, 2) formation at far distances followed by resonance capture and migration while in resonance seems to be more viable, 3) the capture probability varies with the migration rate which itself depends on the mass of the protoplanetary disk. The latter suggests that slow migration rates and small protoplanetary disks may in fact facilitate the formation and subsequent resonance capture of a terrestrial planet with a close-in giant planet around M stars.

1012.2183

It has been shown that an Earth-size planet or a super-Earth, in resonance with a transiting Jupiter-like body around an M star, can create detectable TTV signals (Kirste &amp; Haghighipour, 2011). Given the low masses of M stars and their circumstellar disks, it is expected that the transiting giant planet to have formed at large distances and migrated to its close-in orbit. That implies, the terrestrial planet has to form during the migration of the giant planet, be captured in resonances, and migrate with the giant body to short-period orbits. To determine the possibility of this scenario, we have studied the dynamics of a disk of protoplanetary embryos and the formation of terrestrial planets during the migration of a Jupiter-like planet around an M star. Results suggest that unless the terrestrial planet was also formed at large distances and carried to its close-in resonant orbit by the giant planet, it is unlikely for this object to form in small orbits. We present the details of our simulations and discuss the implication of the results for the origin of the terrestrial planet.

12163003

2011A&A...526L..11C

1012

1012.5076_arXiv.txt

The extreme carbon star IRC+10216 is one of the best studied evolved low-mass stars on the Asymptotic Giant Branch (AGB) owing to its proximity (d $\sim$ 180 pc) and the rich chemistry of its wind. Indeed, more than 60 molecular species have been detected at millimetre (mm), submillimetre (submm) and infrared (IR) wavelengths, probing the entire gas conditions pertaining to the outflow (Ziurys \cite{zu06}). Because the star has already experienced third dredge-up, it is 'carbon-rich', i.e., its photosphere is characterised by a C/O ratio greater than 1. Many carbon-bearing species have been identified in the inner envelope extending from 1 to $\sim$ 10 \rstar, including CO, HCN, C$_2$H$_2$, CS, and SiC$_2$ at mid-IR and submm wavelengths (Keady \& Rigdway \cite{kea93}, Sch{\"o}ier et al. \cite{sho07}, Fonfr{\`i}a et al. 2008, Patel et al. \cite{pa09}, Cernicharo et al. \cite{cerni10}). Oxygen-bearing species other than CO have also long been detected in IRC+10216. Submm observations of silicon monoxide, SiO, confirmed a formation locus close to the star (Sch{\"o}ier et al. 2006, Decin et al. \cite{dec10}). The water molecule, H$_2$O, was first detected with the submm satellite SWAS by Melnick et al. (2001), and its presence was confirmed by submm observations with the ODIN satellite (Hasegawa et al. \cite{hag06}) and supported by the detection of hydroxyl, OH, by Ford et al. (2004). Recently, water was detected by the SPIRE, PACS, and HIFI spectrometers onboard the {\it Herschel} submm telescope in the carbon star VY Cyg (Neufeld et al. \cite{neu10}), the S star $\chi$ Cyg (characterised by a C/O ratio $\sim$ 1) (Justtanont et al. \cite{jus10}), and IRC+10216 (Decin et al. \cite{dec110}, hereafter DAB10). These various detections of high-excitation rotational lines probe relatively high gas temperatures, thus implying the presence of H$_2$O fairly close to the star. Derived abundances with respect to H$_2$ span values from 10$^{-8}$ to 10$^{-7}$ for IRC+10216, $\sim$ 10$^{-6}$ for VY Cyg, and 10$^{-5}$ for $\chi$ Cyg. The derived abundances seem to logically decrease as the carbon content of the star increases, because CO traps most of the oxygen available for water formation. \begin{table*} \caption{Gas parameters in the inner wind of IRC+10216. The shocks form at 1.2 \rstar. For each radius, the temperature and number density are given for the pre-shock gas, the gas after the shock front (in the collision-induced H$_2$ dissociation zone), and the gas at the beginning and the end of the adiabatic cooling zone ($\equiv$ ballistic trajectory).} % \label{tab1} % \centering % \begin{tabular}{c c c c c c c c c c } % \hline\hline % r & Shock velocity &\multicolumn{2}{c}{Preshock gas} & \multicolumn{2}{c}{Shock Front} & \multicolumn{2}{c}{Adiabatic expansion - start} & \multicolumn{2}{c}{Adiabatic expansion - end} \\ % \hline % (\rstar) & (km s$^{-1})$&T & n$_{gas}$ & T & n$_{gas}$ & T & n$_{gas}$& T & n$_{gas}$ \\ \hline 1.2 & 20.0 &2062 & 3.63(13)& 19725 &1.98(14) & 4409& 5.97(14)&1480 & 3.63(13)\\ % 1.5 & 17.9 &1803 & 8.24(12)& 15922 &4.40(13) & 3870& 1.29(14)&1290 & 8.24(12)\\ 2 & 15.5 &1517 & 1.44(12)& 12081 &7.59(12) & 3200& 2.14(13)&1080 & 1.44(13)\\ 2.5 & 13.9 &1327& 4.24(11)& 9779&2.21(12) &2750& 6.08(12)&951& 4.24(11)\\ 3& 12.6& 1190 & 1.69(11)& 8245 &8.73(11) & 2430& 2.35(12)&848 & 1.69(11)\\ 4& 11.0 &1001 & 4.51(10)& 6284 &2.29(11) & 1790& 4.48(11)& 711& 4.51(10)\\ 5 & 9.8&876 & 1.79(10)&5096 &8.94(10)&1550&1.71(11)&621 &1.79(10)\\ \hline \end{tabular} \tablefoot{ Temperatures T are in Kelvin and gas number densities n$_{gas}$ are in cm$^{-3}$.} % \end{table*} Several formation mechanisms were advocated to explain the presence of H$_2$O in IRC+10216. Melnick et al. (\cite{mel01}) proposed that icy comet bodies orbiting the carbon star could be vaporised in the stellar outflow, providing a source of oxygen to the gas for water formation in the intermediate envelope. Willacy (\cite{wil04}) suggested that water could form on the surface of iron dust grains in the intermediate envelope by Fischer-Tropsch catalysis. Recently, DAB10 and Ag{\'u}ndez et al. (\cite{ag10}, hereafter ACG10) proposed that partial penetration of the interstellar ultraviolet (UV) radiation field occurs deep in the outflow owing to the clumpy nature of the wind. $^{13}$CO and SiO can thus photodissociate, providing atomic oxygen to the gas, which then leads to the formation of H$_2$O. All the proposed explanations for the presence of water have their drawbacks. The first mechanism is somehow extreme because it implies that orbiting icy cometary bodies should be a characteristic of all carbon and S stars. The second proposition requires that iron grains form in the wind of carbon stars, an assumption that still needs confirmation by theoretical models or observations. Finally, according to the third model, all stellar winds are sufficiently clumpy to allow for some penetration of UV photons as deep as 2 \rstar, that is, in the dust formation and wind acceleration region. If so, partial photodissociation of molecular dust precursors, among which the radical propargyl (C$_3$H$_3$) and its precursors CH and CH$_2$, should occur, which would hamper the dust-formation process to a certain degree. The formation of the oxygen-bearing species SiO in the inner wind of IRC+12016 was investigated theoretically by Willacy \& Cherchneff (\cite{wil98}, hereafter WC98), who showed that collisional dissociation of CO occurred in the gas layers that experienced the passage of periodic shocks induced by stellar pulsation. The released atomic oxygen then formed a population of OH radicals that triggered the synthesis of SiO via their reaction with atomic Si. Shock-induced chemistry could also explain the formation of CO$_2$ in the O-rich Mira star, IK Tau (Duari et al. \cite{dua99}). In a later study, Cherchneff (\cite{cher06}) modelled the non-equilibrium chemistry of the inner wind of AGB stars as a function of C/O ratios. Of importance was the finding that a few molecules, namely CO, SiO, HCN, and CS, efficiently formed in large amounts in the dust-formation zone whatever the C/O ratio of the star. The author concluded that this group of molecules was ejected as parent species in the intermediate and outer envelopes in O-rich Miras, S stars, and carbon stars. This hypothesis was additionally supported by the observations of high-excitation rotational lines of CO, HCN, SiO, and CS in O-rich, C-rich, and S AGB stars (Decin et al. \cite{dec08}). These results strongly support the assumption that the non-equilibrium chemistry induced by the passage of periodic shocks is responsible for the formation of C-bearing species in the inner envelope of O-rich Miras, and of O-bearing species in the inner wind of carbon stars. Here, we revisit the non-equilibirum chemistry of the inner wind of IRC+10216. The updated chemistry includes new processes such as the thermal fragmentation that is active at the high postshock gas temperatures and the formation and destruction of O-bearing and C-bearing species. The latter include acetylene C$_2$H$_2$, hydrocarbons, carbon chains, and the benzene and phenyl aromatic rings, C$_6$H$_6$ and C$_6$H$_5$, respectively. The chemistry of metal hydrides, chlorides, and sulphides, and phosphorous-bearing compounds is also considered. The complete results of this study will be presented in a forthcoming publication, and we report here on our results for water. Section 2 presents the physical and chemical model considered for the inner wind of IRC+10216, while our results for H$_2$O and other important species are summarised in Section 3, and a discussion is presented in Section 4.

The models of ACG10 and DAB10 allow for the penetration of the interstellar UV radiation field in the very deep layers of the inner wind owing to clumping in the outflow. The models use that of Ag{\'u}ndez \& Cernicharo (\cite{ag06}), i.e., a Eulerian description of a steady outflow penetrated by some UV radiation (in the form of a minor UV-illuminated wind component), and do not take into account the periodic shocks pervading the inner wind from 1 to $\sim 5$ \rstar. The region exposed to UV radiation is 10 \% of the stellar envelope mass, and the models are highly dependent on the clumping factor and stellar mass-loss. The free oxygen necessary for water synthesis is provided by the partial photodissociation of $^{13}$CO and SiO. This mechanism also triggers the formation of other species of interest, like ammonia. Indeed, NH$_3$ is predicted to form at radius r $\geq$ 3 \rstar~with high abundances peaking at $7\times 10^{-7}$ with respect to H$_2$ at $r \sim 4 $ \rstar. Keady \& Ridgway (\cite{kea93}) observed mid-IR vibrational transitions of ammonia in IRC+10216 and deduced that a NH$_3$ abundance distribution peaking at 10-20 \rstar~could better reproduce their data. Later, Monnier et al. (\cite{mon00}) carried out interferometric observations of mid-IR molecular absorption bands of NH$_3$ in IRC+10216 with very high spectral resolution. They found that NH$_3$ was located in a region of decaying gas turbulence at radii well beyond the inner wind (r $\geq$ 20 \rstar). These large radii were further used to model NH$_3$ submm lines observed with the satellite ODIN (Hagesawa et al. \cite{hag06}). The above-mentioned observations suggest that the formation locus of NH$_3$ is located well beyond the dust-formation zone (r $\geq$ 5 \rstar), while ammonia is formed as early as 3 \rstar with a high abundance extending to $\sim$100 \rstar ~according to DAB10 and ACG10. Our present model does not form NH$_3$ in the inner wind ($x$(NH$_3$) $\sim 4\times 10^{-13}$ with respect to H$_2$ at r = 5 \rstar), supporting larger formation radii for this species in accordance with mid-IR and submm observations. Decin et al. ({\cite{dec110}, DAB10) report on newly observed high-excitation lines of cyanoacetylene HC$_3$N with the IRAM telescope. These high-J lines present a flat-topped profile, and DAB10 and ACG10 construe this shape as evidence for a formation locus deeper in the envelope. They claim that the dissociation occurring in their minor UV-illuminated wind component explains the formation of HC$_3$N in the intermediate envelope. An enhancement (up to $\sim 3 \times 10^{-7}$) in the HC$_3$N abundance peaking at $r= 120$ \rstar~is required to reproduce the high-J IRAM lines. The radial distribution of HC$_3$N in IRC+10216 was mapped by Audinos et al. (\cite{au94}) with the IRAM telescope, and the line modelling for a steady homogeneous outflow already required an enhanced HC$_3$N abundance at these radii. Cherchneff \& Glassgold (\cite{cher93}) modelled the cyanopolyyne chemistry in IRC+10216 and found a shoulder in the HC$_3$N distribution at radii $<$ 200 \rstar~resulting from the synthesis of HC$_3$N by neutral-neutral channels. The predicted HC$_3$N abundances in the shoulder were a factor $\sim 7$ lower compared to values derived by Audinos et al. We notice that the HC$_3$N abundance distribution of DAB10 for their major UV-shielded component does not include this predicted shoulder, which points to a different chemistry used by DAB10 and ACG10. High angular resolution observations by Trung \& Lim (\cite{tru08}) of the J = 5 - 4 HC$_3$N line with the VLA clearly indicate a very clumpy shell distribution as far as 20'' from the star. Therefore, radiative transfer models taking into account the outflow clumpiness are required to accurately assess the HC$_3$N abundance distribution across the envelope. An updated chemical model is also necessary to clearly identify the prevalent formation pathways to HC$_3$N, including neutral-neutral and UV dissociation-induced processes. A shock-induced chemistry may also be a viable source of HC$_3$N at intermediate radii and must be considered in subsequent models. The present results support the hypothesis that H$_2$O forms very close to the star with high abundances with respect to H$_2$ between $1 \times 10^{-6}$ and $1 \times 10^{-4}$ and that the abundance gradually chemically freezes out to a value of $1.4 \times 10^{-7}$ at r $\geq$ 5 \rstar. The contribution to the H$_2$O line intensities from the high-abundance region comprised between 1.2 \rstar~and 2.5 \rstar~is not observable with {\it Herschel} because of its small extent and beam filling factor. Furthermore, the H$_2$O formation chemistry appears to be coupled to that of SiO and highlights the importance of the hydroxyl radical OH in the inner wind. A major advantage of the shock-induced chemistry hypothesis is that any star on the AGB pulsates and experiences the passage of shocks in its dust-formation zone. Therefore, it provides a non-restrictive, genuine mechanism for forming water and other molecules very close to the star for a variety of objects in different evolutionary stages, potentially explaining the detection of H$_2$O in O-rich, S and carbon stars with the {\it Herschel} telescope. The shock-induced chemistry is also successful in explaining the presence of O-bearing molecules in carbon stars and C-bearing species in O-rich Miras. Foremost, this hypothesis does not exclude other formation mechanisms for water at larger radii (comets, chemistry on grain surfaces, partial photodissociation in a clumpy outflow)\footnotemark. Because of its presence in AGB stars, H$_2$O can be added to the list of species (CO, SiO, HCN, CS) proposed by Cherchneff (\cite{cher06}) to efficiently form in the dust-formation zone from shock chemistry and be ejected as 'parent' species in the intermediate and outer envelopes. This hypothesis awaits testing by observations of the deep layers of AGB envelopes at mid-IR and submm wavelengths. \footnotetext{New data on water in IRC+10216 from {\it Herschel}/HIFI have just been released (Neufeld et al. \cite{neu101}) and confirm the formation of H$_2$O at radii $\leq$ 6 \rstar. They indicate a common locus for H$_2$O and SiO formation in agreement with the present results and rule out the comet and Fischer-Tropsch catalysis scenarios for H$_2$O formation. This is further supported by new {\it Herschel}/HIFI data for several carbon stars, confirming the widespread presence of water very close to the star (Neufeld et al. \cite{neu102}).}

1012.5076

Context. The presence of water in the wind of the extreme carbon star IRC+10216 has been confirmed by the Herschel telescope. The regions where the high-J H<SUB>2</SUB>O lines have been detected are close to the star at radii r ≤ 15 R<SUB>star</SUB>. <BR /> Aims: We investigate the formation of water and related molecules in the periodically-shocked inner layers of IRC+10216 where dust also forms and accelerates the wind. <BR /> Methods: We describe the molecular formation by a chemical kinetic network involving carbon-and oxygen-based molecules. We then apply this network to the physical conditions pertaining to the dust-formation zone which experiences the passage of pulsation-driven shocks between 1 and 5 R<SUB>star</SUB>. We solve for a system of stiff, coupled, ordinary, and differential equations. <BR /> Results: Non-equilibrium chemistry prevails in the dust-formation zone. H<SUB>2</SUB>O forms quickly above the photosphere from the synthesis of hydroxyl OH induced by the thermal fragmentation of CO in the hot post-shock gas. The derived abundance with respect to H<SUB>2</SUB> at 5 R<SUB>star</SUB> is 1.4 × 10<SUP>-7</SUP>, which excellently agrees with the values derived from Herschel observations. The non-equilibrium formation process of water will be active whatever the stellar C/O ratio, and H<SUB>2</SUB>O should then be present in the wind acceleration zone of all stars on the Asymptotic Giant Branch. <P />Appendix is only available in electronic form at <A href="http://www.aanda.org">http://www.aanda.org</A>

12133276

2010arXiv1012.5133A

1012

1012.5133.txt

A classical theory, in the description of a physical system, assumes that any underlying characteristic of the physical system can undergo only (deterministic) continuous changes. Noncontinuous changes (which can be nondeterministic) are attributed to statistically averaged characteristics, of a given physical system placed, in an ensemble (ie. a large collection) of physical systems. Quantum theory involves extensions, of the classical theoretical description of a physical system, in which some of the underlying characteristics of the physical system instead undergo noncontinuous changes (which may be deterministic, nondeterministic or partially deterministic). The classical description can be obtained from the quantum description in the limit where the noncontinuous changes are small enough to be approximately considered as continuous changes. The effects of noncontinuous changes are expected to be observed when the system is involved in high energy interactions, where dissociations are most likely to occur. \subsection{Quantum mechanics} Mechanics describes the characteristic changes of a given mechanical system (any physical system involved in mostly nondestructive interactions). Quantum mechanics focuses on an extension of the classical mechanical description to include also those underlying characteristics (electrical charge, radiative energy, angular momentum, etc) of the mechanical system that undergo noncontinuous changes. Quantum mechanics resulted from efforts that either predicted or explained observed phenomena such as the energy distribution in a black body's spectrum, the photoelectric effect, the Compton effect, electron diffraction, atomic spectra, etc. Early quantization ideas were presented by Planck, Einstein, Bohr, De Broglie, Hiesenberg and Schrodinger. Planck had to assume that the \emph{blackbody} consisted of oscillators that could emit or absorb energy only in fixed amounts $\vep$ that needed to depend linearly on the frequency only. That is $\vep=hf$, where $h$ is a constant. Similarly Einstein in order to explain the \emph{photoelectric effect} (the ejection of electrons from the light-illuminated surface of a metal, with the kinetic energy of the electrons depending linearly on frequency but not on the intensity of the light) assumed that the energy of light was quantized (distributed in space as localized lumps each of which can be produced, transported or absorbed only as a whole) so that the energy of a particle of light may be written as $E=hf$ and hence deduced a corresponding momentum with $|\vec{p}|={h\over c}f={h\over\ld}$. Thus the wave phase of light could then be rewritten as $e^{2\pi i(f t-\vec{k}\cdot \vec{x})}=e^{i{2\pi\over h}(Et-\vec{p}\cdot \vec{x})}$, in terms of its particles' states $(\vec{x},\vec{p}),\\ ~~\vec{p}=h\vec{k}={h\over\ld}\hat{v}={h\over\ld}{\dot{\vec{x}}\over |\dot{\vec{x}}|}$,~~ or~~ $(x^\mu,p^\mu)\eqv (ct,\vec{x},{E\over c},\vec{p})$. Since the energy and momentum of a massive free particle are related by $E^2=\vec{p}^2c^2+m^2c^4$ the particle of light is therefore a massless free particle. De Broglie postulated that the wave phase relation be applied also to massive free particles $E^2=\vec{p}^2c^2+m^2c^4$ in which case these particles should also display wave-like properties with \\ ~~$f={E\over h},~~\ld={h\over |\vec{p}|},~~\vec{p}={E\over c^2}\vec{v}$. This was confirmed in \emph{electron diffraction} experiments. It was then straightforward to write down ``wave'' or ``field'' equations (eg. the nonrelativistic Schrodinger equation $(i\del_t+{1\over 2m}\vec{\del}^2-V(\vec{x},t))\psi(\vec{x},t)=0$) for massive particles in an external potential $V(\vec{x},t)$, where a ``field'' $\psi(\vec{x},t)$ is a superposition or linear sum of ``waves''. Light quantization also explains the \emph{Compton effect}: the observed shift in wavelength of light when it scatters off free electrons. On the other hand, it was realized by Bohr and others that it is not possible to map out a clear path or orbit for the electron in an atom. In the continuum theory, the Fourier transform of the electron's electric dipole moment $eq$ predicted a continuous frequency spectrum for radiation with the Fourier coefficients of $eq$ giving the intensities associated with each radiated frequency. However, the observed frequencies were discrete implying that the Fourier representation was not an appropriate way to represent $eq$. A matrix representation was finally chosen by Heisenberg and others as an appropriate representation for $eq$, where the components of the matrix may be interpreted as ``transition probabilities'' among the discrete frequencies in analogy to the classical Fourier coefficients which were normally interpreted as radiation intensities associated with the continuous frequencies. Empirical results in \emph{atomic spectroscopy}, eg. Rydberg's wavelength formula ~${1\over\ld_{ij}} ={R\over n_i}-{R\over n_j}$ where $n_i,n_j$ are integers and $R$ a constant, indicate that the energy levels of an electron in a physical atom may be represented by the eigenvalues of a matrix called Hamiltonian $H$. The Hamiltonian $H$ is a matrix-valued ``function'' of equally matrix-valued \footnote{instead of a Fourier sequence} observable \footnote{In ``observable'' or ``measurable'', measurement of a quantity $U$ refers to an assignment of a number to the quantity $U$. An observable will randomly take on one value of its spectrum each time it is measurement.} quantities $q,p$ that represent the canonical position and momentum from classical Hamiltonian mechanics. The relations may be expressed as follows %{\footnotesize $\P$} \bea \label{fn:hamiltonian1}&&H=H({p},q,t)=(H_{mn}),\\ &&\label{heisenbergeq}{dF\over dt}={i\over \hbar}(HF-FH)+{\del F\over\del t},~~~~F=F({p},q,t)=(F_{mn}),\\ &&{q}p-p{q}=i\hbar,~~~~q=q(t)=(q_{mn}),~~{p}={p}(t)=({p}_{mn}),\\ && H\psi_\nu =\hbar\nu ~\psi_\nu~~~(\txt{Eigenvalue problem for the matrix $H$}),\\ \label{fn:hamiltonian2}&&H=S\Ld S^{-1},~~\Ld_{mn}=\hbar\omega_m\delta_{mn},~~\omega_m=2\pi\nu_m, \eea where the commutator $[H,F]=HF-FH$ may be interpreted as a quantum mechanical analogue of the classical Poisson bracket $\{h,f\}=\del_qh\del_pf-\del_ph\del_qf$. The equations above come from an empirically deduced form for the coordinate $q$ given by \bea &&q_{mn}(t)=q^0_{mn}~e^{i\omega_{mn}t},~~\omega_{mn}=2\pi(\nu_m-\nu_n),\nn\\ &&{d q_{mn}(t)\over dt}={i\over \hbar}(\Ld q-q\Ld)_{mn}=i\omega_{mn}~q_{mn}(t). \eea where $\nu_{mn}=\nu_m-\nu_n$ is the frequency of a photon emitted by an electron that ``drops'' from a higher energy level $m$ to a lower energy level $n$ (the energy of the photon is $h\nu$). The canonical quantization conditions \\ ~~$[{q}_i,p_j]=i\hbar\delta_{ij},~[{p}_i,{p}_j]=[q_i,q_j]=0$~ are an extension (see eqn (\ref{BScondition2}) of appendix \ref{cano-quant}) of the Bohr-Sommerfeld quantization\footnote{This model considers (planar) elliptical rather than (planar) circular orbits of Bohr's model for the electronic orbits of Hydrogen. This quantization condition is merely an additional constraint (to the usual classical equations of motion) imposed in order to obtain a discrete rather than a continuous set of orbits, energies, angular momenta and related quantities. It may also be written as ~$\oint_C(p_jdq^j-q_jdp^j)=2nh$~ or as ~$\oint_C\bar{z}_jdz^j=2nhi,~~z_j=q_j+ip_j$.} condition \bea \label{BScondition1}\oint_C {p}_idq^i=nh. \eea The time evolution equation (\ref{heisenbergeq}) generates a one parameter time translation group $\{e^{iHt}\}$ with the Hamiltonian $H$ as the sole generator. The spectrum (from the eigenvalue problem $H\psi_\nu=h\nu\psi_\nu$ for $H$) of $H$ is preserved by this time translation symmetry and consequently each atom has a unique emission or absorption spectrum that characterizes (or serves as a thumbprint for) the type of chemical element the atoms of that type produce. The eigenvalue problem for $H=H(\vec{q},\vec{p})$ may be seen as the problem of finding the irreducible representations of the one parameter time translation group and so each frequency represents an irreducible or elementary attributes (a single excitation, or energy, level of an electron of the atom) of a non-rotating atomic electron system. Naturally, the electron system can be free to rotate around or relative to the nucleus in which case we have invariance under the time translation plus rotation group whose irreducible representations would give the elementary attributes of the system. The canonical quantization condition for a system with several canonical degrees of freedom is $[{q}_i,p_j]=i\delta_{ij}\hbar,~[q_i,q_j]=[{p}_i,{p}_j]=0$. For a system with Hamiltonian $H=H(\vec{{p}}^2,\vec {{p}}\cdot\vec{q},\vec{q}^2)$ and angular momentum $L_{ij}={1\over 2}(q_i{p}_j-q_j{p}_i)$, $H$ commutes with $L_{ij}$ and $\{H,L^2=L_{ij}L_{ij}\}$ generate the center of the algebra of the symmetry group. \bea [L_{ij},L_{{k}{l}}]=-{1\over 2}(\delta_{i{k}}L_{j{l}}+\delta_{j{l}}L_{i{k}} )+{1\over 2}(\delta_{i{l}}L_{j{k}}+\delta_{j{k}}L_{i{l}} ). \eea All parts of the atomic system can also be displaced by the same amount in ``free'' space without disturbing the spectrum of the atomic system. Thus one needs to consider a Hamiltonian of the form $H=\sum_{ab}h(\vec{p}_a\cdot\vec{p}_b,\vec{p}_a(\vec{q}_a-\vec{q}_b),(\vec{q}_a-\vec{q}_b)^2)$ where $a,b$ label the various pieces or particles of the system. Then $H$ also commutes with the total momentum operator ~$\vec{P}=\sum_a\vec{p}_a$ ~ which is the generator of spatial translations. The canonical commutation relations are \bea &&[{q}^i_a,p^j_b]=i\delta^{ij}\delta_{ab}\hbar,~[q^i_a,q^j_b]=[p^i_a,p^j_b]=0\nn\\ \eea and the angular momentum operator will be the sum of the individual ones: \bea &&L^{ij}=\sum_aL^{ij}_a=\sum_a{1\over{2}}(q^i_aP^j-q^j_aP^i)={1\over{2}}(Q^iP^j-Q^jP^i),\nn\\ &&\vec{Q}=\sum_a\vec{q}_a. \eea The center of the algebra of the symmetry group of the atomic system is now generated by $(H,L^2,\vec{P})$. At this point one realizes that the problem of quantizing the atomic system includes the problem of finding the irreducible representations of its symmetry group (or equivalently of the algebra of the symmetry group) generated by $H,P^i,L^{ij}$. To include relativistic effects, one needs to replace the (spatial rotation plus spatial translation plus time translation) group with the Poincare group (spacetime rotation plus translation group). Then relativistic quantum mechanics involves the problem of finding the spectrum of the center of the group generated by the operators $\P^\mu,J^{\mu\nu}$ which have the canonical representation \bea &&\P^\mu=(\P^0(H),\vec{\P}(\vec{P})),~\Q^\mu=(\Q^0(Q^0),\vec{\Q}(\vec{Q})),\nn\\ &&\P^0(H)=H=H(\gamma^0,\vec{\gamma},t,\vec{Q},\vec{P}),~~\vec{\P}(\vec{P})=\vec{P},\nn\\ &&\Q^0(Q^0)=Q^0,~~\vec{\Q}(\vec{Q})=\vec{Q},\nn\\ &&J^{\mu\nu}={1\over{2}}(\Q^\mu \P^\nu-\Q^\nu \P^\mu)+{i\over 4}(\gamma^\mu\gamma^\nu-\gamma^\nu\gamma^\mu)+...\eqv L^{\mu\nu}\otimes 1_{S}+1_L\otimes S^{\mu\nu}+...,\nn\\ &&~~~~=J^{\mu\nu}_L+J^{\mu\nu}_S,\nn\\ &&[J_{{\mu}{\nu}},J_{\al\beta}]=-{1\over 2}(\eta_{{\mu}\al}J_{{\nu}\beta}+\eta_{{\nu}\beta}J_{{\mu}\al} )+{1\over 2}(\eta_{{\mu}\beta}J_{{\nu}\al}+\eta_{{\nu}\al}J_{{\mu}\beta} ),\nn\\ &&[J_L^{{\mu}{\nu}},J_S^{\al\beta}]=0,\nn\\ %&&[P^\mu,\gamma^\nu]=0,~[Q^\mu,\gamma^\nu]=0,~~\gamma^\mu=(\gamma^0,\vec{\gamma}),\nn\\ &&[P^\mu,Q^\nu]=i\eta^{\mu\nu},~~Q^\mu=(Q^0,\vec{Q}),\nn\\ &&\{\gamma^\mu,\gamma^\nu\}=2\eta^{\mu\nu}. \eea In the Schrodinger representation~~ $Q^\mu\ra \mu_{x^\mu},~P_\mu\ra i{\del\over\del x^\mu}$ (here $\mu_{x^\mu}$ denotes ordinary multiplication by the spacetime coordinates $x^\mu$), one then has the consistency condition \bea \label{schrodinger-eqn} i{\del\over\del t}=H(\gamma^0,\vec{\gamma},t,\vec{x},i{\del\over\del\vec{ x}}) \eea on the space of sections $E/\mathbb{\mathbb{R}}^{d+1}=\{\psi:\mathbb{R}^{d+1}\ra E\simeq O(\mathbb{C}^M\ot\mathbb{C}^N)\times(\mathbb{C}^M\ot\mathbb{C}^N)\}$ of a vector bundle $E$ over $\mathbb{R}^{d+1}$ where $\psi=\psi_L\otimes \psi_S$ is the product of the orbital and spin angular momentum wavefunctions and $O(\mathbb{C}^M\ot\mathbb{C}^N)$ is the space of linear operators on $\mathbb{C}^M\ot\mathbb{C}^N$. \begin{comment} One can always return to a lower-dimensional theory by restricting integration to a lower-dimensional subspace. \bea &&\int d^nx~\L(x,\vphi(x),\del\vphi(x))\ral \int d^pu\sqrt{h(u)}\L(x(u),\vphi(x(u)),\del\vphi(x(u))) ,\nn\\ &&~~~~h(u)=\det_{ab} (\eta_{\mu\nu}{\del x^\mu\over\del u^a}{\del x^\nu\over\del u^b}). \eea \end{comment} Even though it is not possible to say precisely where the atomic electron's orbit is, it is however possible to say that it is mostly around the nucleus of the atom; that is, the electron's orbit is localized in the region around the nucleus. A basic quantity introduced for the study of localization\footnote{A system is localized in a certain region $\D$ at a particular time if the total probability of finding it in $\D$ at that time is $1$. Alternatively, the region $\D$ is dense in the support of the probability density function of the system.} was Schrodinger's wavefunction in wave mechanics which is any function satisfying the consistency condition (\ref{schrodinger-eqn}). Schrodinger's wave mechanics is equivalent to Heisenberg's matrix mechanics which was discussed earlier. In general the wave function is a complex-valued function(al) of the quantized configuration variables such as canonical coordinate in quantum mechanics or fields in quantum field theory, whose absolute value can be interpreted as a joint probability density function for the quantized canonical variables on which it depends. When the quantum, ie. quantized classical, configuration variables are represented as elements of an algebra $\O(\H)$ of operators on a Hilbert space\footnote{A Hilbert space is a vector space completed into a metric space by a norm that is induced by an inner product measure defined on the vector space.} $\H$ then the wavefunction would be the value of a chosen linear functional\footnote{Wavefunctions of physical systems and probability amplitudes for various physical processes are examples of (values of) linear functionals on $\O(\H)$. The wavefunction for a physical system is a time-dependent linear functional whose value on a given quantum configuration is the probability amplitude for finding the system in that quantum configuration and it satisfies Schrodinger's equation.} on the quantum configuration variable in question. Thus the time evolution equation may also be written either in terms of the wavefunction or in terms of a corresponding vector in the Hilbert space $\H$. The time evolution equation in terms of the wavefunction is known as Schrodinger's equation. More specifically the sole irreducible representation, up to unitary equivalence, of the relations (\ref{fn:hamiltonian1}) through (\ref{fn:hamiltonian2}) on a Hilbert space is known as Schrodinger's representation. \subsection{Quantum field theory} Quantum field theory is a relativistic quantum theory of systems with arbitrary numbers and types of degrees of freedom. Quantum mechanics treats a system of $N$ (interacting) particles using a fixed number and type of $N$ (coupled) equations. However not all interacting systems have a fixed number and species of particles. Particle transformations and relativistic quantum effects such as particle creation and annihilation may occur. Particles of a kind are now regarded as localizable disturbances (ie. perturbations or fluctuations) in a field of that kind. In particular the field description treats elementary particles as (Fourier) modes of the oscillatory part of an associated field in direct analogy to the electromagnetic field, the modes of whose oscillatory part correspond to the various frequencies of the electromagnetic spectrum. One has an analog of the canonical quantization condition; \bea &&\vec{q}_n(t)\ra q_p^\al(t)=\sum_{\vec{x}}\psi^\al(\vec{x},t)u_p(\vec{x},t),\nn\\ &&\vec{p}_n(t)\ra \pi_p^\al(t)=\sum_{\vec{x}}\Pi^\al(\vec{x},t)u_p(\vec{x},t),\nn\\ &&q_p^\al(t)\pi_{p'}^\beta(t)-(-1)^{2s}\pi_{p'}^\beta(t)q_p^\al(t)=i\hbar\delta^{\al\beta}\delta_{pp'},\nn\\ &&\sum_{p}u^\ast_p(\vec{x},t)u_p(\vec{y},t)=\delta^3(\vec{x}-\vec{y}),\nn\\ &&\Pi^\al(\vec{x},t)\psi^\beta(\vec{y},t)-(-1)^{2s}\psi^\beta(\vec{y},t)\Pi^\al(\vec{x},t)=i\hbar\delta^{\al\beta}\delta(\vec{x}-\vec{y}),\nn\\ &&\Pi(\vec{x},t)={\del\L\over\del\del_t\psi }(\vec{x},t),\nn\\ &&S[\psi]=\int\L(x,dx,\psi,d\psi), \eea where $n$ is a discrete label for a collection of particles and the value of $x$ needs to be chosen in such a way as to obtain a consistent theory for the field $\psi$. For example Pauli's exclusion principle\footnote{The exclusion principle associates the shell structure of atomic electron systems, space occupying/shape forming properties of matter, stability of astronomical objects such as neutron stars, etc to the difficulty for two elementary matter systems to have exactly the same set of fundamental quantum labels. Electromagnetic fields for example and other force fields do not appear to exhibit these properties. The exclusion principle is connected to the idea of spin angular momentum by the requirement that the probability amplitude of a composite physical process must be a rotationally invariant/covariant functional of the probability amplitudes for the individual elementary processes of which it is composed.} requires that s be a half integer for matter fields and and integer for interaction mediation fields. $p$ is a characteristic or typical value (an eigenvalue for a corresponding momentum operator as a Noether charge associated with translational invariance) for the momentum of an individual mode. This is because $\vec{q}_1$ and $\vec{q}_2$ (corresponding to $q^\al_{p_1},~q^\al_{p_2}$) denote different positions in space. $\al$ is a ``spin'' index which is an extension of the spatial vector index. The differential action or Lagrangian $\L(x,dx,\psi,d\psi)$ is a differential form on spacetime. Thus an individual mode is described by the triple $(q_p^\al(t),\pi_p^\al(t),u_p(\vec{x},t))$, where $|u_p(\vec{x},t)|^2d^3x$ is the probability of finding the mode in an infinitesimal neighborhood of $\vec{x}$ of volume $d^3x$ at any given time $t$. This means that the role of the point $\vec{x}$ is now being played by the linear functional \\ ~~$u_p:(q^\al_p(t),\psi_\al(\vec{x},t))\mapsto u_p(\vec{x},t)=\langle q^\al_p(t)\psi_\al(\vec{x},t)\rangle $. The field $\psi^\al$ can also be directly interpreted as the particle coordinate, where the particle is constrained to move along a time-parametrized path $\vec{q}:[0,1]\ra \C$ of the configuration space \\ ~~$\C=\bigcup_n\C_n\eqv\bigcup_n\{\vec{q}_n\}$ in many particle quantum mechanics meanwhile the particle is constrained to move along a spacetime-parametrized hypersurface\\ ~~ $\psi^\al:\M\simeq([0,1]^4,g)\ra \U=\bigcup_{x\in\M}\{\psi^\al(x)\}\subset\C$ of the configuration space \\ ~~$\C=\bigcup_{p}\{q^\al_p\} $ in quantum field theory and similarly the particle is constrained to move along a $(\sigma,\tau)$-parametrized two dimensional surface \\ ~~$X^\al:([0,1]^2,h)\ra \C=\mathbb{R}^{d+1}$ in string theory. In quantum field theory the role originally played by the Hamiltonian $H$ alone in quantum mechanics is now played by the 4-momentum operator \\ ~~$P_\mu=(H,\vec{P})$~~(A component $T^{\mu 0}$ of the Energy-momentum tensor \\ ~~$T^{\mu\nu}=\int d^3x~\T^{\mu\nu},~~~~\del_\mu \T^{\mu\nu}=0 $, a Noether charge corresponding to spacetime translation symmetry). The eigen-value problem for $H$, and any other quantities that commute with $H$, is replaced by the problem of finding the solutions $U$ of the equation \\ ~~$U(\Ld_1,b_1)U(\Ld_2,b_2)=U(\Ld_1\Ld_2,b_1+\Ld_1b_2)$~ which is Wigner's method of classifying elementary particle states. That is, finding the irreducible representations of the Lorentz-Poincare transformation \\ ~~$LP:{{\mathbb{R}^{3+1}}}\ra{{\mathbb{R}^{3+1}}},~ x\mapsto\Ld x+b,~\Ld^T=\Ld^{-1} ,~~x=(x_\mu)=(x_0,\vec{x})$, the automorphism or symmetry group of the spacetime ${{\mathbb{R}^{3+1}}}$. The irreducible representations correspond to free elementary point particles that can be localized in ${{\mathbb{R}^{3+1}}}$. In addition to reparametrization symmetry the Lorentz-Poincare transformation is a symmetry and thus a canonical transformation\footnote{Section \ref{cano-quant}} of the relativistic point particle action \bea S[x,\Gamma]=m\int_\Gamma\sqrt{\eta_{\mu\nu}dx^\mu dx^\nu} \eea since this Lagrangian involves only the metric $ds^2=\eta_{\mu\nu}dx^\mu dx^\nu$ which is the defining structure of the Minkowski spacetime. Thus given any space $\S$, one can also consider the problem of finding the irreducible representations of the automorphism group $G(\S):\S\ra\S$ of $\S$ so as to be able to characterize/classify all the possible elementary physical systems that can be localized in $\S$. Examples of spaces include topological metric spaces, manifolds (which also include Lie groups), fiber bundles, etc and other spaces derived from these using various mathematical constructs. Here it is also important to note that the symmetry group of a topologically nontrivial\footnote{A space is topologically nontrivial if any two of its subspaces cannot always be continuously deformed into each other. Topology is the study of invariance under continuous shape change or deformation (ie. geometry) transformations. Physically interesting geometries would be the fixed points of these geometry transformations.} space (as compared to the flat spacetime ${{\mathbb{R}^{3+1}}}$) is ``enlarged'' mainly due to additional discrete transformation channels leading to various periodicity types and therefore one expects additional distinct physical properties induced on the elementary systems in $\S$ by its nontrivial topology. Conversely, if the elementary systems in $\S$ are observed to display unexpected additional properties, say through experimentation, that do not seem to depend on the geometry, ie. shape/size structures, on $\S$ then they may be investigated by introducing nontrivial topology. Ways of introducing nontrivial topology include employing nondynamical constraints (like quotienting of a [topologically trivial] space by the actions of [discrete] transformation groups ) as well as dynamical constraints such as postulating the presence of unknown forms of ``elementary'' systems that can couple to the known elementary systems in a way that can explain the additional properties and also gives possible explanations as to whether the unknown forms of elementary systems could be experimentally detectable or not. For example, consider the variational problem for an electron (considered as the less physically realistic case, a point particle, so that it can only trace 1-dimensional paths) with action $S[q]=\int_0^1 L(t,dt,q,dq),~~q:[0,1]\ra {{\mathbb{R}^{3+1}}}$. If there is a very strong magnetic field confined in a thin infinitely long tube through the space ${{R^{3+1}}}$, then since any electron (and hence its path) with insufficient energy cannot penetrate this tube, it means that for such an electron the variational problem will have more than one solution as a path on one side of the magnetic tube cannot be continuously varied to a path on the opposite side of the tube. For the same reason a path that wraps around the tube n times cannot be varied to a path that wraps around it any $m$ times in the opposite sense or $m\neq n$ times in the same sense. Hence to every path is associated an integer parameter labeling the number of times and sense in which its path winds around the tube. Therefore if identical electrons of insufficient energy are produced at some point and later interact then one expects to observe the effect of the difference in the topological charges (winding numbers) they gained during their individual journeys. This effect may be included in the action by adding a non trivial but smooth path deformation independent term \bea &&\nu[q]=\int_0^1B_idq^i,~~\nn\\ &&\delta\nu[q]=\delta (\int_0^1B_idq^i)= \int_{0}^{1}(\delta B_i~dq^i+B_i\delta dq^i)= \int_0^1(\delta q^i\del_i B_j~dq^j+B_i{\del_j(\delta q^i)dq^j})\nn\\ &&~~~~=\int_0^1\delta q^i(\del_iB_j-\del_jB_i)dq^j+\int_0^1\del_i(B_j\delta q^j)dq^i\nn\\ &&~~~~=\int_0^1\delta q^i(\del_iB_j-\del_jB_i)dq^j+[B_j\delta q^j]|_0^1=0, \eea where the path $q(t)$ can be smoothly deformed to the path $q(t)+\delta q(t)$. That is, smooth path deformation independence requires $dB=0$~ in the region between between any two paths, with common end points, that can be continuously deformed into each other, ~$\oint_\Gamma B=n(\Gamma)\in \mathbb{Z},~~B=B_idq^i$. Alternatively, let $\Gamma_+$ ($\Gamma_-$) be the path oriented from $t=0$ to $t=1$ ($t=1$ to $t=0$),~ $(\Gamma+\delta\Gamma)_+$ be the varied (with end fixed $\delta q(0)=0=\delta q(1)$) path oriented from $t=1$ to $t=0$ and $\Gamma$ be the closed path $(\Gamma+\delta\Gamma)_++\Gamma_-$. Then Stokes' theorem implies that \bea &&\delta\nu[q]=\nu_{(\Gamma+\delta\Gamma)_+}[q]-\nu_{\Gamma_+}[q]=\nu_{(\Gamma+\delta\Gamma)_+}[q]+\nu_{\Gamma_-}[q]=\nu_{(\Gamma+\delta\Gamma)_++\Gamma_-}[q]\nn\\ &&~~~~=\nu_{\Gamma}[q]=\oint_\Gamma B=\int_{\txt{int}(\Gamma)} dB. \eea Therefore, $\delta\nu[q]=0$ unless the variation takes the path across the tube since $dB|_{{{\mathbb{R}^{3+1}}}\backslash\txt{tube}}=0$. $B$ may be normalized so that any non-zero contribution from $\nu[q]$ is an integer. One notes obviously that $B$ (as well as the tube) can also be a dynamical field. The configuration space of the electron is $\S\simeq{{\mathbb{R}^{3+1}}}\backslash\txt{tube}$ instead of $\mathbb{R}^{3+1}$. This same analysis can be carried out for the physically more realistic systems such as strings, $p$-branes, and fields in general as well; which can be sensitive to several other kinds of topologies. The additional terms $\nu(q)$ are known as Wess-Zumino terms and their gauge non-invariance can be adapted to cancel gauge anomalies and so they may be used to define gauge invariant functional integrals in quantum field theory\footnote{See for example \cite{nair} for a review of quantum field theory.}. \begin{comment} Nontrivial topology of spaces is classified mathematically by homotopy, homology\footnote{Homology/homotopy characterizes a topological space $X$ using the global non-triviality of locally trivial structures $s_{ij}:R_i(X)\ra R_j(X)$ defined on various representations $\{R_i(X)\}$ of $X$. Distinction between homology and homotopy lies in the representations $R_i(X)$ and the corresponding notion of triviality. The space of homotopy or continuous deformation invariant (closed but not exact) forms $\H_n(\M)=\{\Omega:\M\ra T^\ast(\M)^{\wedge n},~d\Omega=0,~\int_{\M\backslash\del\M}\Omega\neq 0\}$ is related to the set of homotopy equivalence classes of continuous maps $\H^n(\M)=\{[f],~f:S^n\ra \M\}$.} and their extensions. Also localization on a topological space $X$ may be discussed using sections $\Gamma=\Gamma(X,E)=\{\psi:X\ra E\},~\Gamma^\ast=\{u:\Gamma\ra \F(X),~u(s+t)=u(s)+u(t)~\}$ of vector bundles $\pi:E\ra X$ over $X$, where $\F(X)=\{f:X\ra \mathbb{\mathbb{C}}\}$ is the space of complex functions on $X$. Gauge symmetry comes in when effects, on the radiation spectrum (spectrum of $H$), of the electromagnetic interactions between the radiating electron, other electrons in the atom, and the nucleus of the atom are taken into account. The theory in this case is known as electrodynamics. An extension (also called internal symmetry) of gauge symmetry onto other field theories helps to describe nuclear interactions (known as weak and strong force). Due to analogy between Einstein's geometrical theory of gravity and electromagnetism, gauge symmetry is also involved in the treatment of gravitational interactions. \end{comment} \label{ch:two} We give an introductory review of quantum physics on the noncommutative spacetime called the Groenewold-Moyal plane. Basic ideas like star products, twisted statistics, second quantized fields and discrete symmetries are discussed. We also outline some of the recent developments in these fields and mention where one can search for experimental signals. \label{sec:intro} The CMB radiation shows how the universe was like when it was only $400, 000$ years old. If photons and baryons were in equilibrium before they decoupled from each other, then the CMB radiation we observe today should have a black body spectrum indicating a smooth early universe. But in 1992, the Cosmic Background Explorer (COBE) satellite detected anisotropies in the CMB radiation, which led to the conclusion that the early universe was not smooth: There were small perturbations in the photon-baryon fluid. The theory of inflation was introduced \cite{guth, Linde, Albrecht} to resolve the fine tuning problems associated with the standard Big Bang cosmology. An important property of inflation is that it can generate irregularities in the universe, which may lead to the formation of structure. Inflation is assumed to be driven by a classical scalar field that accelerates the observed universe towards a perfect homogeneous state. But we live in a quantum world where perfect homogeneity is never attained. The classical scalar field has quantum fluctuations around it and these fluctuations act as seeds for the primordial perturbations over the smooth universe. Thus according to these ideas, the early universe had inhomogeneities and we observe them today in the distribution of large scale structure and anisotropies in the CMB radiation. Physics at Planck scale could be radically different. It is the regime of string theory and quantum gravity. Inflation stretches a region of Planck size into cosmological scales. So, at the end of inflation, physics at Planck region should leave its signature on the cosmological scales too. There are indications both from quantum gravity and string theory that spacetime is noncommutative with a length scale of the order of Planck length. In this paper we explore the consequences of such noncommutativity for CMB radiation in the light of recent developments in the field of noncommutative quantum field theories relating to deformed Poincar\'e symmetry. The early universe and CMB in the noncommutative framework have been addressed in many places \cite{Greene, Lizzi, Brandenberger1, Huang, Brandenberger2, queiroz1, Fatollahi2, Fatollahi1}. In \cite{Greene}, the noncommutative parameter $\theta_{\mu \nu} = -\theta_{\nu \mu} =\textrm{constants}$ with $\theta_{0i} =0$, ($\mu, \nu = 0, 1, 2, 3$, with $0$ denoting time direction), characterizing the Moyal plane is scale dependent, while \cite{Brandenberger1, Brandenberger2, Huang} have considered noncommutativity based on stringy space-time uncertainty relations. Our approach differs from these authors since our quantum fields obey twisted statistics, as implied by the deformed Poincar\'e symmetry in quantum theories. We organize the paper as follows: In section II, we discuss how noncommutativity breaks the usual Lorentz invariance and indicate how this breaking can be interpreted as invariance under a deformed Poincar\'e symmetry. In section III, we write down an expression for a scalar quantum field in the noncommutative framework and show how its two-point function is modified. We review the theory of cosmological perturbations and (direction-independent) power spectrum for $\theta_{\mu \nu}=0$ in section IV. In section V, we derive the power spectrum for the noncommutative Groenewold-Moyal plane ${\cal A}_{\theta}$ and show that it is direction-dependent and breaks statistical isotropy. In section VI, we compute the angular correlations using this power spectrum and show that there are nontrivial ${\cal O}( \theta^{2})$ corrections to the CMB temperature fluctuations. Next, in section VII, we discuss the modifications of the $n$-point functions for any $n$ brought about by a non-zero $\theta^{\mu \nu}$ and show in particular that the underlying probability distribution is not Gaussian. The paper concludes with section VIII. The Moyal plane is the algebra $\A_\theta(\mathbb{R}^d)$ of functions on $\mathbb{R}^d$ with the $\ast$-product given by \bea \label{moyal1}&&(f\ast g)(x)=f(x)e^{{i\over 2}\ola{\del}_\mu\theta^{\mu\nu}\ora{\del}_\nu}g(x)\eqv f(x)e^{{i\over 2}\ola{\del}\wedge\ora{\del}}g(x),~~f,g\in \A_\theta(\mathbb{R}^d),\nn\\ &&\theta_{\mu\nu}=-\theta_{\nu\mu}=\txt{constant}. \eea If $\hat{x}_\mu$ are coordinate functions, $\hat{x}_\mu(x)=x_\mu$, then (\ref{moyal1}) implies that \bea [\hat{x}_\mu,\hat{x}_\nu]=i\theta_{\mu\nu}. \eea Thus $\A_\theta(\mathbb{R}^d)$ is a deformation of $\A_0(\mathbb{R}^d)$ \cite{qft-us}. There is an action of a Poincar$\acute{\txt{e}}$-Hopf algebra with a "twisted" coproduct on $\A_\theta(\mathbb{R}^d)$. Its physical implication is that QFT's can be formulated on $\A_\theta(\mathbb{R}^d)$ compatibly with the Poincar$\acute{\txt{e}}$ invariance of Wightman functions \cite{drinfeld,qft-us}. There is also a map of untwisted to twisted fields corresponding to $\theta_{\mu\nu}=0$ and $\theta_{\mu\nu}\neq 0$ (``the dressing transformation'' \cite{Grosse,Faddeev-Zamolodchikov}). For matter fields, if these are $\vphi_0$ and $\vphi_\theta$, \bea &&\vphi_\theta(x)=\vphi_0(x)e^{{1\over 2}\ola{\del}_\mu\theta^{\mu\nu}P_\nu}\eqv \vphi_0(x)e^{{1\over 2}\ola{\del}\wedge P},\\ &&P_\mu=\txt{Total momentum operator}. \eea While there is no twist factor $e^{{1\over 2}\ola{\del}\wedge P}$ for gauge fields, the gauge field interactions of a matter current with a gauge field are twisted as well: \bea \H^\theta_I(x)=\H^0_I(x)e^{{1\over 2}\ola{\del}\wedge P}, \eea where $\H^0_I$ can be the standard interaction $J^{0\mu}A_\mu$ of an untwisted matter current to the untwisted gauge field $A_\mu$. The twisted fields $\vphi_\theta$ and $\H^\theta_I$ are not causal (local). Thus even if $\vphi_0$ and $\H^0_I$ are causal fields, \bea &&[\vphi_0(x),\vphi_0(y)]=0,\\ &&[\H^0_I(x),\H^0_I(y)]=0,\\ &&[\H^0_I(x),\vphi_0(y)]=0,~~x\times y \eea ($x\times y$ means that $x$ and $y$ are relatively spacelike), that is not the case for the corresponding twisted fields. For example, \bea &&[\vphi_\theta(x),\H^\theta_I(y)]=e^{-{i\over 2}{\del\over\del x^\mu}\theta^{\mu\nu}{\del\over\del y^\nu}}\vphi_0(x)\H^0_I(y)-e^{-{i\over 2}{\del\over\del y^\mu}\theta^{\mu\nu}{\del\over\del x^\nu}}\H^0_I(y)\vphi_0(x)\neq 0,\nn\\ &&~~~~ x\times y. \eea Thus acausality leads to correlation between events in spacelike regions. The study of these correlations at finite temperatures at the level of linear response theory (Kubo formula) is the central focus of this paper. We will also formulate the Lehmann representation for relativistic fields at finite temperature for $\theta_{\mu\nu}\neq 0$. It is possible that some of our results for $\theta_{\mu\nu}=0$ and $\theta_{\mu\nu}\neq 0$ are known \cite{fradkin}. In section 3, we review the standard linear response theory \cite{fradkin} and the striking work of Sinha and Sorkin \cite{sorkin-sinha}. We also discuss the linear response theory for relativistic QFT's at finite temperature for $\theta_{\mu\nu}=0$. It leads to a natural lower bound on relaxation time, a modification of the result ``$(\Delta r)^2\approx \txt{constant}\times\Delta t$'' of Einstein and its generalization ``$(\Delta r)^2\approx \txt{constant}\times \log\Delta t$'' to the `` quantum regime'' by Sinha and Sorkin \cite{sorkin-sinha}. Section 4 contains the linear response theory for the twisted QFT's for $\theta_{\mu\nu}\neq 0$. A striking result we find is the existence of correlations between spacelike events: A disturbance in a spacetime region $M_2$ evokes a fluctuation in a spacetime region $M_1$ spacelike with respect to $M_2$ $(M_1\times M_2)$. Noncommutative corrections in four-momentum space also have striking periodicity properties and zeros as a function of the four-momentum $k$. They are also direction-dependent and vanish in certain directions of the spatial momentum $\vec{k}$. All these results are discussed in this section. The results of this section have a bearing on the homogeneity problem in cosmology. It is a problem in causal theories \cite{trodden-vachaspati}. The noncommutative theories are not causal and hence can contribute to its resolution. In section 5, we derive the finite temperature Lehmann representation for \\ ~~$\theta_{\mu\nu}=0$ and generalize it to $\theta_{\mu\nu}\neq 0$. The Lehmann representation is known to be useful for the investigation of QFT's. The concluding remarks are in section 6.

} In this chapter, we have shown that the introduction of spacetime noncommutativity gives rise to nontrivial contributions to the CMB temperature fluctuations. The two-point correlation function in momentum space, called the power spectrum, becomes direction-dependent. Thus spacetime noncommutativity breaks the rotational invariance of the CMB spectrum. That is, CMB radiation becomes statistically anisotropic. This can be measured experimentally to set bounds on the noncommutative parameter. The next chapter (see \cite{numerical}) presents numerical fits to the available CMB data to put bounds on $\theta$. We have also shown that the probability distribution governing correlations of fields on the Groenewold-Moyal algebra ${\cal A}_{\theta}$ are non-Gaussian. This affects the correlation functions of temperature fluctuations. By measuring the amount of non-Gaussianity from the four-point correlation function data for temperature fluctuations, we can thus set further limits on $\theta$. We have also discussed the signals of non-causality of non-commutative field theories in the temperature fluctuations of the CMB spectrum. It will be very interesting to test the data for such signals. %\end{comment} %%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%% %\include{chapters/cmb2.p} %****************************** %%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%% %\begin{comment} \begin{center} Summary of Chapter \ref{cmb2} \end{center} \begin{itemize} \item The noncommutativity parameter is not constrained by WMAP data, however ACBAR and CBI data restrict the lower bound of its energy scale to be around $10$ TeV \item Upper bound for the noncommutativity parameter: $\sqrt{\theta} < 1.36 \times 10^{-19}\textrm{m}$. This corresponds to a ~$10$ TeV~ lower bound for the energy scale. \item Amount of non-causality coming from spacetime noncommutativity for the fields of primordial scalar perturbations that are space-like separated \bea &&\Delta \varphi(\alpha, {\bf x}_{1}) \Delta \varphi(\alpha, {\bf x}_{2}) \geq \Big|\frac{1}{(2\pi)^{3}}\int d^{3}k~P_{\Phi_{0}}(k)~\textrm{sinh}(H\vec{\theta}^{0}\cdot {\bf k})~e^{-\frac{{\bf k}^{2}}{2 \alpha} -i{\bf k}\cdot({\bf x}_{1}-{\bf x}_{2})}\Big|.\nn \eea \end{itemize} %\end{comment} \chapter{Constraint from the CMB, Causality}\label{cmb2} We try to constrain the noncommutativity length scale of the theoretical model given in \cite{cmbpaper} using the observational data from ACBAR, CBI and five year WMAP. The noncommutativity parameter is not constrained by WMAP data, however ACBAR and CBI data restrict the lower bound of its energy scale to be around $10$ TeV. We also derive an expression for the amount of non-causality coming from spacetime noncommutativity for the fields of primordial scalar perturbations that are space-like separated. The amount of causality violation for these field fluctuations are direction dependent. % }A major result of this chapter is the derivation of acausal and noncommutative effects in finite temperature QFT's. They are new and are expected to have applications for instance in the homogeneity problem in cosmology. We have also treated the finite temperature Lehmann representation on the commutative and Moyal planes in detail. This representation succinctly expresses the spectral and positivity properties of the underlying QFT's in a transparent manner and are thus expected to be useful. %%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%% %\include{chapters/conclusion.p} %******************************########### %%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%% \chapter{Conclusions} \label{ch:conclus} We have given a brief review of quantum theory as well as an introduction to quantum field theory in noncommutative spacetime. The concept of deformed Lorentz invariance in noncommutative spacetime led to the following effects which may be susceptible to experimental tests. \begin{enumerate} \item Deformed statistics of quantum fields whose consequences include 1) modification of the statistical interparticle force and hence degeneracy pressure which determines the fate of galactic nuclei after fuel burning seizes, 2) the possibility of observing Pauli forbidden transitions, 3) observation of Lorentz, P, PT, CP, CPT and causality violations. \item The presence of noncommutativity dependent temperature fluctuations in the CMB radiation, through a noncomutativity dependent post inflation power spectrum; giving an estimated upper bound for the noncommutativity parameter and a corresponding lower bound for the energy scale. \item Encounter with noncommutativity-induced causality violation and a non-Gaussian probability distribution during cosmological inflation. \item Noncommutativity induces noncausal, and potentially periodic, corrections to the susceptibility in linear response theory. \end{enumerate} % things discussed % \& importance of work % Pending work To summarize we have investigated, in the context of quantum field theory, the scope of applicability of a new concept of Lorentz invariance. This new concept is a deformation of the usual concept of Lorentz invariance motivated by the form of invariance in Moyal's treatment of quantum mechanics. The investigations were based on certain available theoretical models and experimental data. Results of these investigations can point to alternative and hopefully simpler solutions to both expected and observed physical phenomena whose experimental energies fall within the range of validity of the noncommutativity models. %%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%% \appendix %\include{chapters/appenda.p} %******************************############ %%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%% % % sample appendix file % \chapter{Some physical concepts}

1012.5133

Noncommutative field theories are a class of theories beyond the standard model of elementary particle physics. Their importance may be summarized in two facts. Firstly as field theories on noncommutative spacetimes they come with natural regularization parameters. Secondly they are related in a natural way to theories of quantum gravity which typically give rise to noncommutative spacetimes. Therefore noncommutative field theories can shed light on the problem of quantizing gravity. An attractive aspect of noncommutative field theories is that they can be formulated so as to preserve spacetime symmetries and to avoid the introduction of irrelevant degrees freedom and so they provide models of consistent fundamental theories. In these notes we review the formulation of symmetry aspects of noncommutative field theories on the simplest type of noncommutative spacetime, the Moyal plane. We discuss violations of Lorentz, P, CP, PT and CPT symmetries as well as causality. Some experimentally detectable signatures of these violations involving Planck scale physics of the early universe and linear response finite temperature field theory are also presented.

12132941

2010arXiv1012.5811O

1012

1012.5811_arXiv.txt

The Cosmic Origins Spectrograph (COS), installed in the Hubble Space Telescope in May, 2009, was intended to provide high sensitivity, moderate to low resolution spectroscopy between 1130\AA~and 3200\AA~(Green 2003). In addition to meeting this goal, COS has demonstrated sensitivity down to wavelengths approaching 900\AA~(McCandliss 2010), providing coverage in the Far Ultraviolet Spectroscopic Explorer (FUSE) band at sensitivities comparable to individual FUSE channels. The nominal G140L/1230 mode provides coverage from the detector cutoff at 1850\AA~down to $<$910\AA, and two new modes, G130M/1096 and /1055, provide higher sensitivities and potentially much higher signal to noise at these wavelengths. \begin{figure} \plotone{osterman.fig.1.eps} \caption{COS Light Path. The FUV channels are distinguished by having only one reflection and a windowless, two segment detector. The holographically-ruled aspheric FUV diffraction gratings provide dispersion, reimaging, and astigmatism and aberration correction in a single reflection. The two segment detector allows for one segment to be disabled to prevent overlight if asymmetric illumination is anticipated. (Froning 2009)} \label{fig:lightpath} \end{figure} The COS light path is shown in figure \ref{fig:lightpath}. Light from the Optical Telescope Assembly (OTA) enters COS through an oversized aperture admitting the entire aberrated wavefront from a point-like source. The aperture is windowless so that any short wavelength ($>1130$\AA) light remaining after the two reflections in the OTA will travel unobstructed to the FUV grating. The gratings perform diffraction, aberration correction and focus in a single reflection in order to minimize reflections, maximizing short wavelength sensitivity. Light then travels to the windowless FUV detector. Laboratory testing indicated that the detector retains relatively high ($>$30\%) quantum efficiency down to at least 800\AA. While the COS diffraction gratings and the OTA mirrors are coated with MgF$_{2}$ protected aluminum (typically used for wavelengths longer than 1150\AA), these optics were expected to retain some first-surface reflectivity below the MgF$_{2}$ transmission cutoff at $\sim$1150\AA. This was verified for the G140L-C (flight spare) grating in laboratory testing (fig.\ref{fig:g140leff}) (Osterman 2002), and for the OTA during COS on-orbit testing. \begin{figure} \plotone{osterman.fig.2.eps} \caption{G140L-C (flight spare) laboratory test results showing $\sim$5\% total efficiency (groove efficiency$\times$reflectivity) at 1066\AA. } \label{fig:g140leff} \end{figure} \vspace{-0.4cm}

The Hubble Space Telescope has provided spectacular imaging and spectroscopy longwards of 1150\AA\ for over 20 years. The Cosmic Origins Spectrograph, while originally intended to provide spectroscopic coverage from 1130\AA\ to 3200\AA, has now extended the usable wavelength range of HST to the Lyman limit, providing spectroscopic access to wavelengths unobservable since the end of the FUSE mission in 2007. The capabilities of these new and proposed modes are summarized in table \ref{table:summary}. \begin{table} \caption{New and Proposed COS Observing Modes \label{table:summary}} \centering \begin{tabular}{|c|c|c|c|c|} \hline COS & Wavelength & Effective & Modeled & Background \\ Mode & Range & Area & Resolution & (cts/resl/ksec) \\ \hline \underline{G140L/1230} & $<$920-1160\AA & $\sim$8-10 cm$^{2}$ & $\sim$2100 & 0.3 \\ & 1230-1850\AA & (at 1000\AA) & (1000\AA) & \\ \hline {\it G140L/800} & $<$920-1850\AA & $\sim$8-10 cm$^{2}$ & $\sim$2100 & 0.2 \\ & (a segment) & (at 1000\AA) & (1000\AA) & \\ \hline \underline{G130M/1096} & 940-1081\AA & $\sim$15-25 cm$^{2}$ & $\sim$2300 & 1.3 \\ & 1096-1238\AA & (b segment) & (1000\AA) & \\ \hline \underline{G130M/1055} & 900-1041\AA & $\sim$15-25 cm$^{2}$ & $\sim$1800 & 1.8 \\ & 1055-1196\AA & (b segment) & (1000\AA) & \\ \hline {\it G130M/1215} & 1065-1205\AA & $\sim$30-2000 cm$^{2}$ & $\sim$16,000 & 0.1 \\ & 1220-1360\AA & (b segment) & (1100\AA) & \\ \hline \multicolumn{5}{|l|}{Underlined modes are available or will be made available to observers in}\\ \multicolumn{5}{|l|}{cycle 19. Modes in italic have not been tested and performance projections }\\ \multicolumn{5}{|l|}{are based on modeling and on similar modes.}\\ \hline \end{tabular} \end{table} By expanding coverage to these shorter wavelengths, we make possible a range of investigations not previously accessible to HST, including studies of the Lyman continuum escape fraction from low redshift galaxies, of the Gunn Peterson effect at redshifts between 2 and 2.8 along multiple lines of sight, and observations of the O VI $\lambda\lambda$ 1032, 1038 doublet. The higher resolution of the G130M/1215 mode could support observation of atomic and molecular diagnostics that could be used to study winds and atmospheres of massive stars, as well as the bulk of the mass in the translucent ISM out of which those stars form. \vspace{-0.2cm}

1012.5811

The far-ultraviolet (FUV) channel of the Cosmic Origins Spectrograph (COS) is designed to operate between 1130Å and 1850Å, limited at shorter wavelengths by the reflectivity of the MgF2 protected aluminum reflective surfaces on the Optical Telescope Assembly and on the COS FUV diffraction gratings. However, because the detector for the FUV channel is windowless, it was recognized early in the design phase that there was the possibility that COS would retain some sensitivity at shorter wavelengths due to the first surface reflection from the MgF2 coated optics. Preflight testing of the flight spare G140L grating revealed ~5% efficiency at 1066Å, and early on-orbit observations verified that the COS G140L/1230 mode was sensitive down to at least the Lyman limit with 10-20 cm^2 effective area between 912Å and 1070Å, and rising rapidly to over 1000 cm2 beyond 1150Å. Following this initial work we explored the possibility of using the G130M grating out of band to provide coverage down to 900Å. We present calibration results and ray trace simulations for these observing modes and explore additional configurations that have the potential to increase spectroscopic resolution, signal to noise, and observational efficiency below 1130Å.

12135933

2010EAS....43..199P

1012

1012.0605_arXiv.txt

Near-infrared (near-IR) spectroscopy of early-type stars is a relatively new field. It is motivated by the desire to study early phases of massive star formation (when the stars are deeply embedded in their parental clouds) and to investigate young stellar populations throughout the Galactic disk and at the Galactic centre in particular. Despite hot stars showing a steep decline in flux levels towards longer wavelengths, observations in the near-IR become highly competitive with optical region spectroscopy for regions of high extinction. In some cases near-IR observations are the only means to penetrate the dust. The most prominent example is the Galactic centre, where the extinction in the $K$-band amounts to 3\,mag, compared with more than 30\,mag~in~$V$. Near-IR spectroscopy became feasible only after substantial developments in detector technology over the past 15\,years. Systematic observational studies of (normal) hot stars in the near-IR are nonetheless scarce. A spectral catalogue at low resolution ($R$\,$=$\,$\lambda/\Delta\lambda$\,$\lesssim$\,3000) was compiled by Hanson et al.~(\cite{Hansonetal96}) for OB-type stars. Classification spectra of early-type stars in the $J$, $H$ and $K$ bands can also be found in the atlases of Wallace \& Hinkle~(\cite{WaHi97}), Meyer et al.~(\cite{Meyeretal98}) and Wallace et al.~(\cite{Wallaceetal00}). More recently, intermediate-resolution spectroscopy ($R$\,$\approx$\,10\,000) became the standard (Fullerton \& Najarro~\cite{FuNa98}; Hanson et al.~\cite{Hansonetal05}). Unguided by observations, the field attracted little interest on the theoretical side, except for an early prediction of photospheric Br$\alpha$ emission in B-type stars by Auer \& Mihalas~(\cite{AuMi69a}). Later, Zaal et al.~(\cite{Zaaletal99}) could explain the observed emission cores of many near-IR hydrogen lines in early B-type stars via non-LTE effects, though several discrepancies between models and observation remained. The systematic behaviour of near-IR hydrogen and helium lines with spectral type and luminosity were investigated on theoretical grounds by Lenorzer et al.~(\cite{Lenorzeretal04}) for O stars. Finally, the most comprehensive study of near-IR spectra of OB stars and a comparison with analyses in the optical was performed by Repolust et al.~(\cite{Repolustetal05}). Considerable work has also been done on the massive star population near the Galactic centre. For the most part, this comprises extreme stars like Luminous Blue Variables and Wolf-Rayet stars (e.g.~Najarro et al.~\cite{Najarroetal97}, \cite{Najarroetal09}; Martins et al.~\cite{Martinsetal07}). However, such objects are beyond the scope of the present discussion as their line spectra are formed in the stellar wind.

We have seen that amplification of non-LTE effects is ubiquitous for near-IR lines in hot stars. This makes non-LTE line-formation calculations challenging, as every aspect of the modelling has to be considered properly. On the other hand, the high sensitivity offers the possibility to put tight constraints on the atomic data selected for the construction of non-LTE model atoms. Deficits in the data will show up in the comparison of the model predictions with observation. The near-IR studies of hot stars so far comprise relatively strong lines of hydrogen and helium, based on medium-resolution spectra. The new generation of high-resolution near-IR spectrographs like {\sc Crires} on the VLT promise to revolutionise the field. In particular, a multitude of (weak) metal lines will become accessible for the first time (see e.g.~Przybilla et al.~\cite{Przybillaetal09}; Nieva et al.~\cite{Nievaetal09}). The development of proper non-LTE models will turn quantitative near-IR spectroscopy of hot stars into a crucial tool for Galactic studies, and with the upcoming extremely large telescopes (diffraction-limited observations using adaptive optics will be feasible only at near-IR wavelengths) also for extragalactic stellar astronomy.

1012.0605

Line-formation calculations in the Rayleigh-Jeans tail of the spectral energy distribution are complicated by an amplification of non-LTE effects. For hot stars this can make quantitative modelling of spectral lines in the near-IR challenging. An introduction to the modelling problems is given and several examples in the context of near-IR line formation for hydrogen and helium are discussed.

12168171

2011ApJ...733...39S

1012

1012.3606_arXiv.txt

Chaotic dynamics due to interaction of nonlinear resonances in Hamiltonian systems is studied in a broad range of application areas, from plasma physics to celestial mechanics (e.g., \citealt{C79,LL92,A06}). The characteristic time of predictability of any motion is nothing but the Lyapunov time (the inverse of the maximum Lyapunov exponent) of the motion. Generally, the estimation of the Lyapunov exponents is one of the most important tools in the study of chaotic motion \citep{LL92}, in particular in celestial mechanics. The Lyapunov exponents characterize the mean rate of exponential divergence of trajectories close to each other in phase space. Non-zero Lyapunov exponents indicate chaotic character of motion, while the maximum Lyapunov exponent equal to zero signifies regular (periodic or quasiperiodic) motion. The development of methods of numerical computation of the Lyapunov exponents has more than a 30 year history (e.g., \citealt{F84,LL92}). In contrast, methods of analytical estimation of the Lyapunov exponents started to be developed only recently \citep{HM96,MH97,S00a,S02,S07,S08b}. In studies of the dynamics of the Milky Way, analysis of the Lyapunov exponents has not yet been widely used, but nevertheless there are important achievements. \cite{F01} used Lyapunov exponents as a tool to find bar-induced manifestations of chaos in the local disk stellar kinematics. Taking into account the perturbations from the bar solely (for some particular bar strengths), he constructed ``Lyapunov diagrams'', presenting the Lyapunov timescales of the orbits in the $u$~--~$v$ velocity plane at fixed space positions, and identified regular and chaotic domains in velocity space as a function of space position with respect to the bar. The Lyapunov exponents were calculated on a Cartesian grid of planar velocities. The fraction of chaotic orbits was demonstrated to obviously increase with the bar strength. However, the diagrams in \cite{F01} can hardly be used to estimate real values of Lyapunov times, because the saturation of the computed values of the Lyapunov exponents was not controlled, while the adopted computation time, corresponding to three Hubble times (equivalently, $\sim 100$ Galactic years), might not be enough for the saturation. In other words, the obtained values characterize not the whole chaotic regions of the phase space, but only rather small vicinities of the initial data. Therefore, the computed values are the ``local'' values of the Lyapunov exponents. This is testified by the fact that the variation of the Lyapunov exponents in the diagrams of \cite{F01} is smooth, while there must be a sharp distinction between the chaotic regions (with non-zero Lyapunov exponents) and regular regions (with zero Lyapunov exponents) in the divided phase space. In connection with the problem of estimation of Lyapunov timescales in the solar neighborhood, \cite{Q03} noted that, according to \cite{HM96}, for a fully overlapped system, the chaotic zone should have a Lyapunov time $\sim 2 \pi / \nu$ (where $\nu$ is the frequency of perturbation), corresponding to the separatrix pulsation period. In what follows we shall consider the model set of \cite{Q03} for the stellar dynamics in the solar neighborhood and show that the heuristic estimate $\sim 2 \pi / \nu$ severely underestimates the real Lyapunov time. Besides, we shall see that it is rather seldom that the considered dynamical systems, modeling the dynamics in the solar neighborhood, can be called ``fully overlapped''. Besides obtaining the Lyapunov times, we shall estimate the diffusion times in the chaotic domain of the phase space, in the same model set. This will allow one to judge on the possibility for migration of the Sun from the inner regions of the Milky Way to its current location. Such a possibility, arising due to the overlapping of resonances in the phase space, was advocated and studied in detail by \cite{MF10} and \cite{MFC11}.

We have considered how the Lyapunov and diffusion timescales of the stellar dynamics in the solar neighborhood can be estimated. We have used Quillen's (2003) model to describe interaction of the ``spiral'' and ``bar'' nonlinear resonances in the phase space of the motion. A method of analytical estimation of the maximum Lyapunov exponents of the orbital motion has been applied to the solar neighborhood dynamics. The analytical treatment has been performed within a framework of the separatrix map theory \citep{S00a,S02,S07}, describing the motion near the separatrices of a perturbed nonlinear resonance. The Lyapunov times turn out to be basically in the range from 6 to 13 Galactic years. In comparison with the Lyapunov times of the solar system bodies (made in adequate time units), the Galactic dynamical chaos is rather strong in general terms of the loss of predictability of the motion. An interesting inference is that, as soon as the age of the Milky Way measured in its Lyapunov times is about 5--10, now it is already practically impossible to restore exact initial conditions for the stellar dynamics in the solar neighborhood from any observational data. We have estimated also the diffusion times, based on the approach developed initially by \cite{CV86,CV89} for the purposes of studies in cometary dynamics. We have found that, in a number of models, the diffusion times turn out to be small enough to permit radial chaotic migration of the Sun from the inner regions of the Milky Way to its current location. In other words, dynamically adequate models that permit large-scale radial chaotic migration do exist. This confirms the dynamical possibility of the migration concept advocated by \cite{MF10} and \cite{MFC11}. Due to the possibility of ballistic flights inside the chaotic layer, arising because $\lambda \sim 1$, the chaotic mixing might be even far more effective and quicker than in the case of normal diffusion. We have shown that only in a narrow range of possible values of the problem parameters $\epsilon$, $\beta$, $\nu$, and $\delta$ the Galactic chaos is adiabatic because the values of the separatrix map parameter $\lambda$, playing the role of the resonance overlap parameter, are typically greater than 1/2; in other words, adiabatic chaos ($\lambda < 1/2$) seems to be not characteristic for the dynamics in the solar neighborhood. The author is thankful to Alice Quillen for advice and comments. It is also a pleasure to thank the referee for helpful remarks.

1012.3606

We estimate the Lyapunov times (characteristic times of predictability of motion) in Quillen's models for the dynamics in the solar neighborhood. These models take into account perturbations due to the Galactic bar and spiral arms. For estimating the Lyapunov times, an approach based on the separatrix map theory is used. The Lyapunov times turn out to be typically of the order of 10 Galactic years. We show that only in a narrow range of possible values of the problem parameters the Galactic chaos is adiabatic; usually it is not slow. We also estimate the characteristic diffusion times in the chaotic domain. In a number of models, the diffusion times turn out to be small enough to permit migration of the Sun from the inner regions of the Milky Way to its current location. Moreover, due to the possibility of ballistic flights inside the chaotic layer, the chaotic mixing might be even far more effective and quicker than in the case of normal diffusion. This confirms the dynamical possibility of Minchev and Famaey's migration concept.

4519452

2011NuPhS.217..193G

1012

1012.4356_arXiv.txt

1012.4356

We analyzed the electron neutrino data of the Gallium radioactive source experiments and the electron antineutrino data of the reactor Bugey and Chooz experiments in terms of neutrino oscillations. We found a hint of a CPT-violating asymmetry of the effective neutrino and antineutrino mixing angles.

2763416

2011ApJ...728..130G

1012

1012.3789_arXiv.txt

The physics of angular momentum transport is at the core of accretion disk studies. Classical viscous thin disk theories (\cite{ss, lbp, nt73}) assume the existence of a local turbulent viscous stress, thus provide a simple local parameterization, i.e., ``anomalous viscosity" $\alpha$, for disk momentum transport and dissipation. Since the early 90's, the magnetorotational instability (MRI, \cite{bh91}) has been regarded as the best candidate to drive accretion disk turbulence, although gravitational torque or magnetic winds of a \cite{bp82} type can also enhance angular momentum transport. Classical thin disk theories are vertically integrated and azimuthally averaged, therefore essentially one dimensional. Currently, disk vertical structure can only be obtained from numerical simulations where turbulence is established from first-principle instabilities such as the MRI. Global disk simulations are just starting to investigate thin disks (\cite{rf08, shafee08, nkh09}), but they are computationally expensive and not yet able to fully resolve turbulent structures in the disk. Shearing box simulations, on the other hand, can concentrate resolution on disk dynamics at scales of order the disk scale height $H \equiv c_s/\Omega$, therefore are more suitable to study accretion flows in detail. Past studies of shearing box simulations with vertical gravity (e.g., \cite{bnst95, shgb96, ms00, hks06, bhk07, si09}) have revealed a rich set of structures and dynamics in stratified disks. However, all these stratified shearing box simulations were done with a box of limited radial extent $L_x \sim H$, therefore they were not able to explore any structure on scales larger than $H$. Recently, \cite{dsp10} have studied stratified shearing box of radial extent $L_x = 4H$, and \cite{jky09} have adopted models of box size up to $L_x \sim 10H$ in their zonal flow studies. However both these studies are limited to the small veritical extent ($\sim \pm 2H$) and physically unrealistic periodic vertical boundary conditions. In this paper we study the dynamics and structure in isothermal stratified disks using large shearing box with domain sizes $L_x \geq 10 H$ in all directions. We still do not know whether a magnetized turbulent disk is well modeled as a steady-state, locally dissipated disk model. It is possible, for example, that structures (gas and/or fields) develop at a scale large compared to $H$, and that these structures could be associated with nonlocal energy or angular momentum transport. Large scale structures might also develop in the magnetic field in the form of dynamo. The disk might also be secularly unstable (see the overview by \cite{pir78}), that could cause the disk to break up into rings. It is well known that a Navier-Stokes viscosity model for disk turbulence leads, for some opacity regimes, to both viscous \citep{le74} and thermal \cite{pir78} instability, although it is now believed that thermal instability can be removed by delays imposed through finite relaxation time effects in MRI-driven turbulence (\cite{hkb09}). >From an observational point of view, the level of fluctuations (inhomogeneity) at different locations in disks and how these different locations communicate with each other have important consequences for disk spectra modeling \citep{dbht05, bdhks06}. In these models observational diagnostics require integrating over the disk surface, so radially extended structure in the disk model may change the disk spectrum. Our disk model is isothermal (we do not solve an energy equation) and is therefore not capable of investigating dissipation and radiation. It is possible that larger fluctuations would appear in physically richer models where thermodynamics and radiative effects are taken into account (e.g., \cite{tsks03, turner04}). It would then be interesting in the future for spectral modelers to consider disk models with larger radial domains. A shearing box larger than $H$ is also essential to catch the field structure and dynamics in the accretion disk magnetic coronae (ADC; \cite{tp96}; also see a discussion in \cite{uzgoo}), where the field has a characteristic curvature $l \sim v_a/\Omega \geq H$, and $v_a$ is the characteristic Alfv\'en speed in the region. Recently, it has also been pointed out that a large box size may be important to study the saturation properties of the MRI-driven turbulence, either on the ground of resolving parasitic modes \citep{pg09}, or in a phenomenological model of an MRI driven dynamo \citep{v09}. Saturation mechanisms in stratified disk may be fundamentally different from those in unstratified disks. Recent numerical experiments on unstratified disks suggest that: (a) with a zero-net flux, the saturation is dependent on the microscopic Prandtl number $\pr_M$ in the disk, at least at the low Reynolds number \citep{fp07,ll07, sh09}; (b) with a net (toroidal or vertical) flux, the saturation increases with resolution \citep{hgb95, ggsj09}. Stratified disk models, which are closer to real disks, may well maintain a net (most likely, toroidal; see a discussion in \cite{gg09}) field in the disk region because of the magnetic buoyancy induced by stratification. Therefore we expect a saturation in stratified disk models to differ from unstratified models. It is worth enumerating the assumptions we adopt in this work: (1) we use an isothermal equation of state (EOS) in our models; (2) the vertical support comes from the gas and magnetic pressure rather than the radiation pressure; (3) there is no explicit viscosity or resistivity; (4) our initial conditions consist of a uniform toroidal field in a region near the disk midplane; (5) we use outflow boundary conditions for the vertical boundaries. The paper is organized as follows. In \S 2 we give a description of the local model and summarize our numerical algorithm. In \S 3 we present a fiducial model and analyze its structure in the saturated state. In \S 4 we describe how this structure depends on model parameters. In \S 5 we give a report on quasiperiodic oscillations (``butterfly diagrams'') and present a phenomenological model to describe them that is based on a mean-field dynamo model; \S 5 contains a summary of our results.

We have carried out stratified shearing box simulations with domain size $L_x = H$ to $L_x = 32H$ to study properties of isothermal accretion disks on a scale larger than the disk scale height $H$. Our numerical models have vertical extent $\geq 5H$ above and below the disk midplane with outflow boundary conditions. All models start from a net mean toroidal field in the central disk region and the mean fields are allowed to change in the evolution. We find the disk has an oscillating mean toroidal field and $\overline{\<\alpha\>} \sim 0.012 - 0.025$ in the parameter range we explored. We have not found a clear dependence of $\overline{\<\alpha\>}$ on $L_x$ in our models, although the temporal variances in volume averaged quantities decreases with $L_x$. The highest resolution used here is modest ($20-40$ zones per $H$), and we have observed $\overline{\<\alpha\>}$ increases with resolution. Recently, Stone el al. report a converged $\overline{\<\alpha\>} \sim 0.04$ in $L_x \sim H$ high resolution stratified disk simulations with zero-net-flux and periodic vertical boundary conditions (so that the volume-averaged field cannot change during the evolution). The sustained turbulence may be due to the presence of a mean toroidal field in the region close to the disk midplane, lending plausibility to the idea that the saturation mechanism of MRI in stratified disks near the midplane is similar to that in unstratified disks with a net toroidal field. In the saturated state the disk vertical structure consists of (a) a turbulent disk at $|z| \leq 2H$ and (b) a magnetically dominated upper region at $|z| > 2H$, confirming earlier small ($L_x \sim H$) box results. At $|z| \leq 2H$, the disk is mainly supported by gas pressure, and a Gaussian density profile is observed. The plane averaged magnetic energy density $[E_B](z)$ and Maxwell stress $[M_{xy}](z)$ are nearly uniform with vertical height $z$ in this region, where the disk is marginally stable to the Parker instability. At $|z| > 2H$, exponential dependences on $z$ are observed for both $[\rho]$ and $[E_B]$. Fitting formulae for $[\rho](z)$ and $[E_B](z)$ are given in Eqn (\ref{eqn:rhoz}) and Eqn (\ref{eqn:ebz}) respectively. Using a two-point correlation function analysis, we found that close to the midplane, the disk is dominated by small scale ($\leq H$) turbulence, very similar to what we have observed in unstratified disk models. In the corona, magnetic fields are correlated on scales of $\sim 10H$, implying the existence of meso-scale structures. Recently \cite{jky09} have also observed large scale pressure and zonal flow structures in their large shearing box simulations. We will give a detailed report of meso-scale structure in isothermal disks in a forthcoming paper. We have adopted a statistical approach to study the geometry of coronal magnetic fields. Only $\approx 4\%$ of coronal field lines are open. For closed field lines, we calculated the magnetic loop distribution function for the loop foot separation $\Delta \bx$ in the $x-y$ plane, loop maximum height $\Delta z_{\rm max}$, and loop orientation angle $\theta _{\rm foot}$. The loops are dominantly toroidal due to the differential shear. The loop foot distribution between $H-20H$ is a power law with an index $k \sim -1.25$. In the phenomenological model of UG, this corresponds to the limit where reconnection is slow compared to the shear. These comparisons are limited because our models are working in an ideal MHD regime and reconnection is purely numerical. In our models both vertical energy and momentum flux are negligible in the steady state. The mass loss rate from the disk surface is small and decreases with increasing $L_z$. The surface effects are therefore minimal and indicate a lack of disk winds in our stratified disk models. The weak winds are consistent with the constraint that we have a zero-net vertical magnetic flux in these models. A Blandford-Payne type wind requires the existence of a vertical net field (e.g., see \cite{si09}), although we note that in their models the most unstable wavelength for the extremely weak field are probably not resolved.) Initial investigations show that even a weak ($\beta _{0} \sim 1600$) net z field will induce very violent accretion in stratified shearing box models: at certain region of the disk accretion will run away, eventually causing the disk break into rings. Similar phenomena were reported in net vertical field models of \cite{ms00}. We have confirmed the ``butterfly" diagram seen in earlier stratified disk models of size $L_x \sim H$. The butterfly diagrams persist even in our largest runs with $L_x = 32H$. We also report the reversal of the mean fields (for both the dominant toroidal field and a weak radial field) in the disk on a timescale twice that of $[E_B]$. The periods for the butterfly diagram are close in all our models, $P \sim 5$ orbits for $[E_B]$ and $P \sim 10$ orbits for $[B_y]$. The mean field reversal and butterfly diagram may indicate the existence of a mean field dynamo in stratified disks, perhaps controlled by the MRI and magnetic buoyancy. We have presented a toy model for an $\tilde{\alpha}$ type mean field dynamo in stratified disks and found an $\tilde{\alpha} _{\rm imp} \sim 0.01$ will produce the reported period. Further exploration of parameter dependences would be useful for analytical modeling. In the future it would also be interesting to test whether the butterfly oscillations persist when averaging over a large range of radii in global disk simulations. The butterfly diagram may be associated with low frequency QPOs and therefore a good observational diagnostic for accretion flows. On the other hand, we also report a power-law index $k \sim -2.3$ in the temporal power spectrum for coronal magnetic energy fluctuations, consistent with results from recent GRMHD black hole accretion disk simulations. Our stratified disk models are primarily limited by the assumption that the disk is isothermal. Effects of thermodynamics and radiation therefore are neglected in this work. Our models are also limited by finite resolution, box size, evolution time, and the absense of explicit dissipation. Additional insights may also provided by the future explorations on magnetic field strength and geometry in disks.

1012.3789

We consider local, stratified, numerical models of isothermal accretion disks. The novel feature of our treatment is that radial extent L<SUB>x</SUB> and azimuthal extent L<SUB>y</SUB> satisfy H Lt L<SUB>x</SUB> , L<SUB>y</SUB> Lt R, where H is the scale height and R is the local radius. This enables us to probe mesoscale structure in stratified thin disks. We evolve the model at several resolutions, sizes, and initial magnetic field strengths. Consistent with earlier work, we find that the saturated, turbulent state consists of a weakly magnetized disk midplane coupled to a strongly magnetized corona, with a transition at |z| ~ 2H. The saturated α ~= 0.01-0.02. A two-point correlation function analysis reveals that the central 4H of the disk is dominated by small-scale turbulence that is statistically similar to unstratified disk models, while the coronal magnetic fields are correlated on scales ~10H. Nevertheless angular momentum transport through the corona is small. A study of magnetic field loops in the corona reveals few open field lines and predominantly toroidal loops with a characteristic distance between footpoints that is ~H. Finally, we find quasi-periodic oscillations with characteristic timescale ~30 Ω<SUP>-1</SUP> in the magnetic field energy density. These oscillations are correlated with oscillations in the mean azimuthal field; we present a phenomenological, alpha-dynamo model that captures most aspects of the oscillations.

1894320

2011ASSP...27...55M

1012

1012.0449_arXiv.txt

\label{intro} Our panchromatic study of LFs in the Shapley supercluster core (SSC), from UV to IR wavebands, aims to investigate the relative importance of the processes that may be responsible for the galaxy transformations examining in particular the effect of the environment through the comparison of LFs in regions with different local densities. The SSC is constituted by three Abell clusters: A\,3558, A\,3562 and A\,3556, and two poor clusters SC 1327-312 and SC 1329-313. The target has been chosen since the most dramatic effects of environment on galaxy evolution should occur in superclusters, where the infall and encounter velocities of galaxies are greatest ($>$1 000 km s$^{-1}$), groups and clusters are still merging, and significant numbers of galaxies will be encountering the dense intra-cluster medium (ICM) of the supercluster environment for the first time. This work is carried out in the framework of the joint research programme ACCESS \footnote{http://www.oacn.inaf.it/ACCESS} ({\it A Complete Census of Star-formation and nuclear activity in the Shapley supercluster}, \cite{MMH10}) aimed at determining the importance of cluster assembly processes in driving the evolution of galaxies as a function of galaxy mass and environment within the Shapley supercluster. We assume $\Omega_M$ =0.3, $\Omega_{\Lambda}$=0.7 and H$_0$ =70 km s$^{-1}$ Mpc$^{-1}$.

\label{sec:7} We find that optical and the $K$-band LF faint-end slope becomes steeper from high- to low-density environments, although the changes in slope are less dramatic at NIR wavebands indicating that the faint galaxy population increases in low-density environments. Differently from the LF no environmental effect is found in the slope of the SMFs. On the other hand, the $\mathcal{M}^\star$ increase from low- to high-density regions and the excess of galaxies at the bright-end is also dependent on the environment. These results seem indicate that the physical mechanism responsible for the transformation of galaxies properties in different environment are mainly related to the quenching of the SF. Moreover, the NUV and FUV LFs obtained have steeper faint-end slopes than the local field population, while the 24$\mu$m and 70$\mu$m galaxy LFs for the Shapley supercluster have shapes fully consistent with those obtained for the Coma cluster and for the local field galaxy population. This apparent lack of environmental dependence for the shape of IR luminosity functions suggests that the bulk of the star-forming galaxies that make up the observed cluster infrared LF have been recently accreted from the field and have yet to have their SF activity significantly affected by the cluster environment. \begin{acknowledgement} This work was carried out in the framework of the FP7-PEOPLE-IRSES-2008 project ACCESS. AM acknowledges financial support from INAF-OAC and the JENAM grant to attend the conference. \end{acknowledgement}

1012.0449

We present a panchromatic study of luminosity functions (LFs) and stellar mass functions (SMFs) of galaxies in the core of the Shapley supercluster at z = 0.048, in order to investigate how the dense environment affects the galaxy properties, such as star formation (SF) or stellar mass. We find that, while the faint-end slope of optical and NIR LFs steepens with decreasing density, no environment effect is found in the slope of the SMFs. This suggests that mechanisms transforming galaxies in different environments are mainly related to the quench of SF rather than to mass-loss. The Near-UV (NUV) and Far-UV (FUV) LFs obtained have steeper faint-end slopes than the local field population, while the 24 and 70 μm galaxy LFs for the Shapley supercluster have shapes fully consistent with those obtained for the local field galaxy population. This apparent lack of environmental dependence for the infrared (IR) LFs suggests that the bulk of the star-forming galaxies that make up the observed cluster IR LF have been recently accreted from the field and have yet to have their SF activity significantly affected by the cluster environment.

12163182

2011A&A...528A.114T

1012

1012.0163_arXiv.txt

} Stellar mergers have for a long time been recognized to play an important role in the evolution of stellar systems. High stellar densities in globular clusters can lead often to collisions and mergers of stars \citep{leon89}. In this way the origin of blue strugglers can be explained. In dense cores of young clusters multiple mergers of protostars have been suggested as a way of formation of the most massive stars \citep{bonn98}. Some binary stars, in particular contact binaries, are suggested to end their evolution as stellar mergers \citep{robeggl}. \begin{figure*} \includegraphics[height=\hsize]{l_curve.eps} \caption{Light curve of V1309 Sco from the OGLE-III and OGLE-IV projects: $I$ magnitude versus time of observations in Julian Dates. Time in years is marked on top of the figure. At maximum the object attained $I \simeq 6.8$. } \label{lightcurve} \end{figure*} The powerful outburst of V838 Mon in 2002 \citep{mun02}, accompanied by a spectacular light echo \citep{bond03}, raised interest in a class of stellar eruptions named "red novae", "optical transients" or "V838~Mon type eruptions". These objects, which typically reach a maximum luminosity of $\sim10^6~{\rm L}_\odot$, evolve to low effective temperatures and decline as very cool (super)giants. Apart form V838~Mon, in our Galaxy the class also includes V4332~Sgr, whose outburst was observed in 1994 \citep{martini}, and V1309~Sco, which erupted in 2008 \citep{mason10}. As extragalactic eruptions of this kind one can mention M31~RV \citep[eruption in 1989,][]{mould}, M85~OT2006 \citep{kulk07}, and NGC300~OT2008 \citep{berger09}. Several interpretations of the eruptions were proposed. They include an unusual classical nova \citep{tutu,shara}, a late He-shell flash \citep{lawlor}, or a thermonuclear shell flash in an evolved massive star \citep{muna05}. \citet{tylsok06} presented numerous arguments against these mechanisms. They showed that all the main observational characteristics of the V838~Mon-type eruptions can be consistently understood as resulting from stellar collisions and mergers, as originally proposed in \citet{soktyl03}. For these reasons \citet{st07} proposed to call these type of eruptions {\it mergebursts}. Recently, \citet{kfs10} and \citet{ks11} have suggested that some of the V838~Mon-type eruptions are of the same nature as the eruptions of luminous blue variables, and that they can be powered by mass-transfer events in binary systems. In the present paper, we show that the recent red nova, V1309~Sco, is a Rosetta stone in the studies of the nature of the V838~Mon type eruptions. Archive photometric data collected for the object in the OGLE project during about six years before the outburst allow us to conclude that the progenitor of V1309~Sco was a contact binary. The system quickly evolved towards its merger, which resulted in the eruption observed in 2008.

The principal conclusion of our study is that all observed properties of V1309~Sco, i.e. the light curve of the progenitor during six years before the 2008 eruption, as well as the outburst itself, can be consistently explained by a merger of a contact binary. This is the first case of a direct observational evidence showing that the contact binary systems indeed end their evolution by merging into single objects, as predicted in numerous theoretical studies of these systems. Our study also provides a conclusive evidence in favour of the hypothesis that the V838~Mon-type eruptions (red novae) result from stellar mergers, as originally proposed in \citet{soktyl03} and \citet{tylsok06}. In particular, long- and short-term variabilities of the progenitors, as those of V838~Mon and V4332~Sgr \citep{goran07,kimes07}, which were sometimes raised as evidence against the merger hypothesis, appear now natural in view of our data for the V1309~Sco progenitor. We do not claim that all observed eruptions of the V838~Mon-type are mergers of contact binaries. There can be different ways leading to stellar mergers. What the case of V1309~Sco evidently shows is that the observational appearances of a stellar merger are indeed the same as those observed in the V838~Mon-type eruptions. The outburst of V1309 Sco was shorter and less luminous than those of V838~Mon and the extragalactic red novae. The latter objects attained luminosities of about or above $10^6~{\rm L}_\odot$ and their eruptions lasted a few months. These differences are most likely caused by the masses of merging stars. For V838~Mon, an $\sim8~{\rm M}_\odot$ primary was probably involved \citep{tylsok06} instead of a $\sim1~{\rm M}_\odot$ one of V1309~Sco. As noted in Sect.~\ref{sect_cb}, the orbital period of the V1309~Sco progenitor of $\sim$1.4 day is long for the observed population of contact binaries. From a study of contact binaries discovered by the OGLE project in a sky region very close to the position of V1309~Sco, \citet{rucin98} concluded that the W~UMa type sequence sharply ends at the orbital period of 1.3--1.5~days \citep[see also][]{pacz06}, i.e. just at the orbital period of the V1309~Sco progenitor. This can be a pure coincidence (just one case observed), but can also indicate that binaries passing through contact at periods $\ga$1~day are not rare, but that the contact phase in their case is relatively short and quickly leads to a merger. V1309 Sco, an overlooked object \citep[only one research paper published so far, i.e.][]{mason10}, deserves much more attention of the observers and astrophysicists, as do the other V838~Mon-type objects. Apart from supernovae, they belong to the most powerful stellar cataclysms. As often happens in nature, cataclysms destroy old worlds, but also give birth to new ones. What will develop from the stellar mergers? Fast rotating giants, similar to FK~Com? Peculiar stars with circumstellar discs, when new generation planets can be formed? To answer these questions, we just have to follow the evolution of the V838~Mon-type objects, V1309~Sco in particular.

1012.0163

Context. Stellar mergers are expected to take place in numerous circumstences in the evolution of stellar systems. In particular, they are considered as a plausible origin of stellar eruptions of the V838 Mon type. V1309 Sco is the most recent eruption of this type in our Galaxy. The object was discovered in September 2008. <BR /> Aims: Our aim is to investigate the nature of V1309 Sco. <BR /> Methods: V1309 Sco has been photometrically observed in course of the OGLE project since August 2001. We analyse these observations in different ways. In particular, periodogram analyses were done to investigate the nature of the observed short-term variability of the progenitor. <BR /> Results: We find that the progenitor of V1309 Sco was a contact binary with an orbital period of ~1.4 day. This period was decreasing with time. The light curve of the binary was also evolving, indicating that the system evolved towards its merger. The violent phase of the merger, marked by the systematic brightenning of the object, began in March 2008, i.e. half a year before the outburst discovery. We also investigate the observations of V1309 Sco during the outburst and the decline and show that they can be fully accounted for within the merger hypothesis. <BR /> Conclusions: For the first time in the literature we show from direct observations that contact binaries indeed end up by merging into a single object, as was suggested in numerous theoretical studies of these systems. Our study also shows that stellar mergers indeed result in eruptions of the V838 Mon type. <P />Based on observations obtained with the 1.3-m Warsaw telescope at the Las Campanas Observatory of the Carnegie Institution of Washington. The photometric data analysed in the present paper are available from the OGLE Internet archive: ftp://ogle.astrouw.edu.pl/ogle/ogle3/V1309_SCO