Datasets:

JSALT2024-Astro-LLMs

/

astro_paper_corpus

like 0

Follow

JSALT2024 Evaluating LLMs for Research Astronomy 7

3105324

2011A&A...526A..11A

1012

1012.1951_arXiv.txt

Many efforts have been made in past years to constrain the epoch of formation of massive elliptical galaxies. Recent observational results suggest that massive red-sequence galaxies were already assembled in cluster cores at $z\sim1.2$ (e.g. De Propris et al. 1999; Andreon 2006a; Lidman 2008) and possibly up to $z\sim2$ (Andreon 2010). Their colors and morphologies, as revealed from observational studies in mass selected samples, remain unchanged from $z\sim1$ (e.g. Holden et al. 2007; Huertas-Company et al. 2009), which again suggests that these galaxies were assembled at even higher redshifts. These observations may contradict the predictions of the $\Lambda$CDM model, in which mergers of gas-rich galaxies are the main driver of spheroid formation (e.g. Toomre \& Toomre 1972) and the main explanation for the over density of red galaxies in cluster cores (e.g. Dressler 1984). Finding the cluster of galaxies at the highest redshifts allows the epoch of formation of elliptical galaxies to be pushed back, making it a very intense activity in past years (e.g. Papovich et al. 2010). Recently, Andreon et al. (2009) detected what seems to be the most distant cluster of galaxies to date (JKCS\,041 at $z_{phot}=1.9$) by looking at over densities in red galaxies. JKCS\,041 lies in the SWIRE/CFHTLS field at RA=2h, Dec=-4 $deg$ and benefits from an extensive multi-wavelength follow-up, spanning 1.4 Ghz to X rays and including SWIRE/Spitzer imaging in 7 IR bands (3.6, 4.5, 5.8, 8.0, 24, 70, 160 microns), as well as optical CFHTLS ugriz, WIRDS JHK, and UKIDSS JK data. The cluster has not been spectroscopically confirmed yet. However, in the Chandra follow-up, it is detected as an extended X-ray source with $T\sim7.4$ keV, which confirms there is a hot intracluster medium, present in formed clusters and lacking in protoclusters. In this paper 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 $z'-J$ color. Second, by comparing the red sequence color 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 a much higher redshift, and we quantify its photometric redshift ($z\sim 2.2$). 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=15610f1t.ps,width=5truecm,clip=}} \centerline{\psfig{figure=15610f1b.ps,width=5truecm,clip=}} \caption{Stellar locus from photometry as published (top panel) and after minor corrections to match locii determined from different catalogs (bottom panel). Red triangles and open circles indicate stars in the direction of JKCS\,041 and IRC0218A, respectively. The solid line shows the mean fitted locus. The shading marks the 68 \% error (highest posterior interval) of the model.} \label{fig:stellarlocus} \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. This implies $z\gg 1.62$ and rules out $z \le 1.49$ the latter claimed by a recent paper. The 0.32 mag color difference of the two red sequences 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 color of both JKCS\,041 and IRC0218A clusters. We can thus 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 be 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 $z\sim2$ cluster candidates. Getting spectroscopic redshifts for all of them, or even a small part, is too time-consuming with current telescopes. Therefore, photometric redshifts based on the red sequence color will necessarily become very popular in the next years, so we need to get used to them.

1012.1951

This paper aims at robustly determining the redshift of the cluster of galaxies JKCS 041 and at putting constraints on the formation epoch of the color-magnitude sequence in two very high redshift clusters. New deep z' - J data show a clear narrow red sequence that is co-centered with, and similary concentrated on, 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 and ruling out z ≤ 1.49 the latter claimed by a recent paper. The color 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 color of both JKCS 041 and IRC0218A clusters. We do not observe any sign of truncation of the red sequence for both clusters down to J = 23 mag (1 × 10<SUP>11</SUP> solar masses), which suggests that it is already in place in clusters rich and massive enough to heat and retain hot gas at these high redshifts.

12224511

2011PhRvD..83d3528S

1012

1012.1998_arXiv.txt

\label{sec:introduction} Cosmic strings are possible remnants from the early universe (for a review see \cite{VilenkinShellard}) and there is significant effort to try and detect them. A positive detection of cosmic strings will open up a window to very high energy fundamental physics and can potentially have strong implications for astrophysical processes. Hence it is of great interest to continue to discover new observational signatures of cosmic strings, as well as to refine features of known signatures. In this paper we address the radiation of photons by cosmic strings. There is an extensive literature on gravitational radiation from cosmic strings, particularly motivated by upcoming and future gravitational wave detectors. More relevant to the work presented here, however, is the analysis in \cite{JonesSmith:2009ti} and \cite{Chu:2010zzb} of the emission of particles due to the time-dependent metric of cosmic strings from the viewpoint of Aharonov-Bohm radiation. The case of photon emission -- which we treat in the present paper -- was not explicitly discussed there. A crucial feature which emerged in these calculations is that cusps and kinks on cosmic strings emit radiation with a flat spectrum all the way up to the string scale. However, those results were based on studying two rather specific loop configurations with cusps and kinks. As we show here in more generality (namely for any loop configuration) the flat spectrum also applies to the emission of photons from cusps and kinks, as well as kink-kink collisions. Thus light emitted from cosmic strings in this way leads to a new and observable signature of cosmic strings that is completely independent of the details of the underlying particle physics model. As we shall see, the effect is small, however, being proportional to $(G\mu)^2$ where $G$ is Newton's constant and $\mu$ the string tension. Despite that, since photons are being emitted, it may be more easily measurable than, say, the gravitational wave (GW) bursts also emitted by cusps and kinks. The total power emitted in scalar particles from cosmic strings due to their gravitational coupling was first considered in \cite{Garriga:1989bx}, using formalism developed in \cite{Frieman:1985fr}. In this paper we calculate the differential power emitted in photons from cosmic strings due to the gravitational coupling. We call this the ``gravitational Aharonov-Bohm'' effect because the metric is flat everywhere except at the location of the string, and is closely analogous to the case of the electromagnetic Aharonov-Bohm effect due to a thin solenoid. In Sec.~\ref{sec:gAB} we set up the calculation and evaluate the invariant matrix element for the production of two photons. The emission is dominant in three cases -- from cusps, kinks and kink-kink collisions. Integrals relevant to these cases are evaluated in Sec.~\ref{sec:integrals}. In Sec.~\ref{power} we find the power emitted from cusps, kinks and kink-kink collisions on strings. Our results are summarized in Sec.~\ref{conclusions}, where we also consider observational signatures. Our numerical estimate in Eq.~(\ref{calN}) indicates that light from cosmic strings may potentially be detectable by current detectors for a range of string tensions.

\label{conclusions} In this paper we have calculated the flux of photons from cosmic strings. In general this falls off exponentially with harmonic number $n$. However, as we have shown, the power emitted from cusps, kinks and kink-kink collisions does not fall off with $n$ -- rather, it is $n$-{\it independent}. Thus the emission from these features on the string dominates over the emission from the rest of the string, at least at high frequencies. If we denote the differential energy flux at frequency $\omega_0$ by $F$ {\it i.e.} \be F \equiv \frac{d^3E}{{\rm d} t \; {\rm d}\omega_0 \; {\rm d}\Omega_{\rm s}} \ee where $\Omega_{\rm s}$ denotes solid-angle, then our results can be summarized as follows: \ba && \hskip -0.8 cm F_{\rm cusp} \approx \frac{(G\mu)^2}{L} (\omega_0 L)^{2/3}, \ \ \Omega_{\rm s} < (\omega_0 L)^{-2/3} \\ && \hskip -0.5 cm F_{\rm kink} \approx \frac{(G\mu)^2}{L} (\omega_0 L)^{1/3}, \ \ \theta < (\omega_0 L)^{-1/3} \\ && F_{\rm k-k} \approx \frac{(G\mu)^2}{L} \ . \ea The cusp emits a beam within a solid angle, the kink emits along a curve, while the kink-kink emission is in all directions. The duration of the cusp and kink beams is given by Eq.~(\ref{beamduration}), while the duration of the kink-kink radiation is given by the wavelength at which the emission is observed. Also, the cusp radiates at frequencies $\omega_0 < M\sqrt{ML}$ whereas the kink and kink-kink collisions radiate for $\omega_0 < M$, where $M$ is the string scale. Our results so far provide the emission characteristics from certain features on strings. Now we briefly discuss the cumulative effect of having many such features on a given loop of string. The cusp and kink emissions are beamed and this makes the analysis more involved. However, the emission due to kink-kink collisions is not beamed and is easier to estimate. Eq.~(\ref{onekk}) gives the energy emitted from a single kink-kink collision. To get the energy emitted from a string segment, we need to sum over all the kink-kink collisions occurring on that string segment of length $\Delta l$ in an interval of time $\Delta t$ \begin{equation} \frac{dE}{d\omega} \sim (G\mu)^2 \int d\psi_a ~ d\psi_b ~ \psi_a \psi_b \frac{dn_a}{d\psi_a} \frac{dn_b}{d\psi_b} ~ \Delta l ~ \Delta t \label{Eall} \end{equation} where $n_a (\psi_a, t)$ is the number of kinks of sharpness $\psi_a$ at time $t$ per unit length of string, and similarly for $n_b$. We will consider emission from a loop of string that formed at time $t_f$ by breaking off a long string. The loop inherits a large number of (shallow) kinks from the long string and from Ref.~\cite{Copeland:2009dk} we can write \begin{equation} \int d\psi_a \psi_a \frac{dn_a}{d\psi_a} \sim \frac{1}{t_f} \left ( \frac{t_f}{t_*} \right )^\alpha \ . \label{avgpsi} \end{equation} The exponent $\alpha$ is $\sim 0.7$ in the radiation-dominated epoch. In the matter-dominated epoch, strings contain all the kinks accumulated until the epoch of matter-radiation equality, $t_{\rm eq}$, and from then on the scaling in (\ref{avgpsi}) has $\alpha \sim 0.4$ \cite{CopelandKibble}. The time $t_*$ denotes the epoch at which frictional effects on strings became unimportant. Hence Eq.~(\ref{Eall}) can be written as \begin{equation} \frac{dP_\omega}{dl d\omega} \sim \frac{(G\mu)^2}{t_f^2 } \left ( \frac{t_f}{t_*} \right )^{2 \alpha} \label{Pall} \end{equation} where $P_\omega$ denotes the power emitted at frequency $\omega$. We now obtain some numerical estimates, leaving a detailed analysis for future work. The photons emitted from loops deep into the radiation epoch will get thermalized. Hence the emission from loops in the post-recombination era is most relevant for direct observation. Loops at the epoch of recombination could have been produced in the radiation epoch and for simplicity we consider a loop that was formed at the epoch of radiation-matter equality, $t_f = t_{\rm eq} \approx 10^{11} {\rm s}$. With $G\mu \sim 10^{-8}$, $t_* \sim t_P/(G\mu)^2 \sim 10^{-27} {\rm s}$ \cite{VilenkinShellard}, where $t_P \approx 10^{-43}$ s is the Planck time, and $\alpha = 0.7$, we get \begin{equation} \frac{dP_\omega}{dl d\omega} = \frac{(G\mu)^{2+4\alpha}}{t_f^{2-2\alpha} t_P^{2\alpha}} \approx 10^{-22} ~\frac{\rm ergs}{\rm cm} \ . \end{equation} Detectors on Earth observe a flux of photons and it is more relevant to calculate the number of photons arriving at the detector. This follows from $E = N_\omega \omega$ where $N_\omega$ is the number of photons of frequency $\omega$ emitted by the string, \begin{equation} \frac{dN_\omega}{dt dl} \approx 10^{-22} ~\frac{\rm ergs}{\rm cm} ~ \frac{d\omega}{\omega} \ . \label{Ntl} \end{equation} If the loop of length $t_{\rm eq} \sim 10^{21} {\rm cm}$ is at a distance comparable to the present horizon, $r \sim 10^{27}{\rm cm}$, then the flux of photons at the detector is obtained by multiplying (\ref{Ntl}) by $t_{\rm eq}/r^2$, \begin{equation} \frac{d{\cal N}_\omega}{dt dA} \biggl |_{\rm 1~loop} \approx \frac{10^{-10}}{\rm km^2-yr} \frac{d\omega}{\omega} \left ( \frac{G\mu}{10^{-8}} \right )^{2(1+2\alpha)} \label{calN1loop} \end{equation} where ${\cal N}_\omega$ denotes the number of photons arriving at the detector with collecting area $dA$. If at $t_{\rm eq}$ there was one loop of length $t_{\rm eq}$ per horizon, the number of loops that can contribute to the flux at the detector is given by the number of horizons at $t_{\rm eq}$ that fit within a comoving volume equal to our present horizon volume: $t_0^3/(t_{\rm eq}z_{\rm eq})^3$, where $t_0 \sim 10^{17}{\rm s}$ and $z_{\rm eq} \approx (t_0/t_{\rm eq})^{2/3}$. Therefore the number of contributing loops is $\sim t_0/t_{\rm eq} \sim 10^6$ and the photon flux due to all of these loops is \begin{equation} \frac{d{\cal N}_\omega}{dt dA} \biggl |_{\rm loops~at~rec.} \approx \frac{10^{-4}}{\rm km^2-yr} \frac{d\omega}{\omega} \left ( \frac{G\mu}{10^{-8}} \right )^{2(1+2\alpha)} \label{calN} \end{equation} This estimate suggests that a detector with collecting area $~ (100 ~{\rm km})^2$ -- comparable to the Auger observatory -- will detect $\sim 1$ photon/year in every logarithmic frequency interval emitted by string loops from the recombination era. The estimate (\ref{calN}) holds for frequencies all the way up to the string scale $\sim 10^{15} ~{\rm GeV}$ but it does not take into account any propagation effects. Neither does it take into account the network of long strings and the spatial and length distribution of loops. Note that the emission falls steeply with decreasing string tension. The effect can only be useful for small $G\mu$ if the amount of string in a horizon volume is inversely proportional to some high power of $G\mu$. The pattern of photon emission from a string is lineal and this may be helpful to distinguish it from conventional sources. It is also possible that the emission from the string will be polarized (though our analysis so far has summed over polarizations and hence erases the polarization information). The beamed emission from cusps and kinks may provide distinctive events that can signal the presence of strings. We plan to explore these signatures in future work. We would like to close with a cautionary note. The emission rate from kink-kink collisions is greatly enhanced by the factor $(t_f/t_*)^\alpha$ in Eq.~(\ref{avgpsi}). This factor is due to the accumulation of kinks on strings from the time, $t_*$, when their dynamics became undamped. The exponent, $\alpha$, depends on dynamical factors, such as Hubble expansion, that tend to straighten out the kinks, but the estimate does not take radiation backreaction into account and this will have a tendency to reduce the emission rate. However, it is possible that emission at frequencies much lower than the string scale remain relatively unaffected by the backreaction. We can be more confident of our estimates of light from cosmic strings only once this issue is satisfactorily resolved.

1012.1998

The time-dependent metric of a cosmic string leads to an effective interaction between the string and photons—the “gravitational Aharonov-Bohm” effect—and causes cosmic strings to emit light. We evaluate the radiation of pairs of photons from cosmic strings and find that the emission from cusps, kinks and kink-kink collisions occurs with a flat spectrum at all frequencies up to the string scale. Further, cusps emit a beam of photons, kinks emit along a curve, and the emission at a kink-kink collision is in all directions. The emission of light from cosmic strings could provide an important new observational signature of cosmic strings that is within reach of current experiments for a range of string tensions.

12168356

2011ApJ...734...56P

1012

1012.5269_arXiv.txt

Protostellar disks likely undergo one or more phases of gravitational instability during the first $10^5$ to $10^6$ years of their life \citep{vorobyov10b, vorobyov10a, hayfield10}. Two factors determining the susceptibility to gravitational instability, as well as its ultimate outcome, are the mass loading from the molecular cloud envelope \citep{boley09,rafikov09,boleyetal10} and how fast the gas is able to cool and dissipate the energy generated by compression and shocks in the spiral arms \citep{riceetal05,boleyetal07,mayer07,cossins10,merubate10}, either by radiative cooling alone, or by a combination of radiative cooling and convective cooling \citep{boss03, durisetal07, mayer10}. Whether gravitational instability results in a self-regulated, marginally unstable state with long-lived spiral structure, or whether self-regulation is broken and fragmentation takes place within spiral arms, producing sub-stellar objects on brown dwarf or giant planet scales \citep{stamatellos09, boleyetal10,mayer10}, depends on the ability to attain a gas cooling timescale comparable to, or shorter than, the local orbital timescale. Both radiative and convective cooling ultimately depend on the opacity in the gas. Therefore it is of pivotal importance to investigate, in detail, how the opacity might change within spiral shocks. At the relatively low temperatures of protoplanetary disks reached at ten or more AU from the star the major source of the opacity is dust grains, and different prescriptions are used to calculate their contribution. Some of these prescriptions are normally used in the 3-D hydrodynamical simulations of gravitationally unstable protostellar disks \citep{d'alessio01} but it is unclear whether they account for all the important effects on the grains as they undergo rapid temperature and density variations within the spiral shocks. In this work we explore, in detail, the effect of a varying background temperature and gas density on the grain size distribution and its resultant opacity. In particular we consider the time-dependent density and temperature evolution in strong spiral shocks appearing in gravitationally unstable protoplanetary disks that attain a minimum Toomre $Q$ parameter just above the value needed for fragmentation [the necessary condition for fragmentation is that the minimum $Q$ be $1.4$ or lower - see e.g \citep{durisetal07}]. We extract such information from one of the 3-D SPH simulations presented in \cite{mayer07}, which included radiative transfer with the flux-limited diffusion approximation. The paper is organized as follows: First we describe the model adopted for the three-dimensional spiral shocks and for dust grains, then we describe our results on the evolution of dust grain opacity in our reference dust model, analyzing different locations throughout the spiral shock. We show results for the disk midplane first, and then for regions well above the midplane, including the effect of irradiation from the star. Furthermore, we investigate the effect of changing parameters in our model for dust, such as the dust-to-gas ratio in the spiral shocks, the water content of grains, and the distribution of grain sizes. Subsequently, we explore a model in which dust grains are composite. We conclude the paper with a discussion of the implications of our results for disk cooling and fragmentation into sub-stellar objects.

Using the time-dependent spiral shock profiles from a 3-D hydrodynamical simulation of a gravitationally unstable disk we have calculated the effects that such shocks have on the opacity of the gas. Our results show that opacity variations of about an order of magnitude can take place on timescales of several months to years. Such timescales are much shorter than the orbital time, which is larger than 10 years at the distances of the shocks in the simulations (10\,-\,20\,AU from the star). This means that the cooling time in the disk will change on timescales much shorter than the orbital time. In particular, when the opacity drops by nearly an order of magnitude on such short timescales it means that the cooling time will also decrease proportionally fast on such short timescales \citep[see also][]{cossins10}. The cooling disk may not be able to self-regulate on such short timescales and therefore the likelihood of fragmentation should be substantially increased. Likewise, the temperature increase caused by the shock provides the right conditions to produce a rapid increase in molecular weight within the spiral shock, supporting the claim of \cite{mayer07}, as long as rapid migration of solids toward the arms can occur simultaneously and increase the dust-to-gas ratio. The water vapor that is released by the grains will increase the mean molecular weight of the gas. As \cite{mayer07} point out, this will make the disk more unstable. If the molecular weight of the background gas is $\mu_x$ and this gas is a mass fraction $X$ of the total, then introducing a gas of molecular weight $\mu_z$ and mass fraction $1-X$ will increase the mean molecular weight, $\overline{\mu}$, to \begin{equation}\label{meanmu} \overline{\mu}=\mu_x\left[1+\frac{\mu_z-\mu_x}{\mu_z}\left(\frac{1-X}{X}\right)\right] \end{equation} For solar composition gas, $\mu_x\approx 2.4$, while $\mu_z=18$ for water vapor. For the baseline model we investigated the water vapor mass fraction is $9\times 10^{-3}$ so that $\overline{\mu}$ would increase by only 1\%. But if grains are concentrated by the spiral arms and the mass fraction of water vapor increase by an order of magnitude, $\overline{\mu}$ can increase from 2.4 to 2.6. This is large enough to have noticeable effects on the stability of the disk \citep{mayer07}. No published hydrodynamical calculations exist that can show the simultaneous effect of both variations. By combining the results of \cite{mayer07} and \cite{merubate10}, one can expect that the formation of giant planets in the inner disk would be much more likely with a concurrent rapid drop of opacity and increase in molecular weight, probably irrespective of whether convection occurs or not [see also \cite{jongam03}]. \cite{rudpol91} have pointed out that if the opacity, $\kappa$, has a temperature dependence of the form $$\kappa\sim T^{\beta}$$ then a gas with an adiabatic exponent of 1.45-1.5 (which is a good approximation in the temperature range of the shock at peak strength, between 150 and 240 K, for solar metallicity, see \cite{boleyetal07}) will satisfy the Schwarzschild criterion for convective instability if $\beta>0.5-1$. The rising parts of the dashed curves in Fig.\,\ref{opactemp} have slopes corresponding to $\beta=0.5$. As it can be seen from the figure, when the shock ramps up, in the region with T= 150\,K where the ice quickly migrates to the largest cores, the value of $\beta$ is temporarily $ > 1$, namely it is above the critical value expected at this temperature, for which $\gamma = 1.5$ according to \citep{boleyetal07}. The exact value of $\beta$ depends on the location within the shock (i.e. it varies across different annuli because the opacity-temperature relation varies), but it is always high enough to allow, in principle, the development of convection. Later on, as the temperature increases and then decreases again, the adiabatic index also varies, thus changing the critical threshold, and $\beta$ becomes sub-critical. The same is true for a second passage of the shock since the opacity has also changed in the meantime. Therefore, it appears that convection could be triggered initially but would be short-lived. However, since convection itself could change the density and temperature distribution through the shocking region, only a dynamical simulation which includes our coated grain model will be able to ascertain the thermodynamical evolution of the gas phase. Finally, in this discussion we have not considered how the conditions might change as a function of the altitude from the midplane since we have used azimuthally averaged quantities, but going beyond this will once again require a dynamical (3-D) simulations since convection, if it starts, is expected to change the physical conditions particularly across the vertical extent of the disk. Another possible effect could arise because molecular weight variations can in principle damp the convective instability. Indeed, since vertical convection requires $ds/dz < 0$, where $s$ is the specific entropy, and, ignoring adiabatic index gradients, $ds/dz = d \ln T/dz - (\gamma-1) d ln \rho/dz - d ln \mu/dz$ for an ideal gas, one sees that a vertical gradient in the molecular weight can have a stabilizing influence on convection. Therefore, it will be important to understand whether or not significant $\mu$ gradients would develop once $\mu$ can be self-consistently calculated in simulations including both gas and dust. The important point to remember is that, for convection to be relevant for disk fragmentation, what matters is that it affects the cooling of the disk at midplane, and not necessarily across the entire vertical extent of the disk. Therefore, as long as molecular weight gradients are negligible near the disk midplane convection will not be affected. In practice, one will need to show that there is a large enough region around the midplane in which heat can be redistributed rapidly by convection but within which the molecular weight is increased nearly uniformly following dust grain collection and sublimation through the shock. It is tempting to speculate that the gas becomes convectively unstable at this point, and this convection allows the gas to cool more quickly. If this is indeed the case, then the increase in opacity would be offset by convective cooling. Note that this high $\beta$ occurs only during the heating of the gas, and $\beta$ is negative during the cooling stage. Certainly the influence of ice coatings on grains needs to be investigated directly with numerical simulations since these effects are not necessarily additive; for example, one can imagine that an increasing molecular weight would affect the temperature evolution. In addition, the opacity variation across the shock would also be affected if the timescales of the two processes are even slightly uneven. In general, the present work shows the importance of modeling grain chemistry within the spiral shocks of self-gravitating protoplanetary disks. While at $50\,{\rm AU} < R < 200$\,AU fragmentation in disks loaded by high accretion rates from the molecular envelope is an almost unavoidable event which depends little on the details of the gas and dust chemistry because the cooling time is expected to be short \citep{boley09,boleyetal10,rafikov09,vorobyov10b}, fragmentation in the inner disk ($< 30$\,AU) is going to depend very strongly on thermodynamics and therefore on the properties and response of dust grains that affect the cooling time via the opacity. Other thermodynamical factors, such as the effective adiabatic index \citep{boleydur08} and the molecular weight \citep{mayer07}, are also bound to play an important role because they will change the temperature evolution across the shock. The models that we describe in this paper are not dynamical; the grains do not move in response to the gas, nor do they coagulate and change their size distribution. Recent calculations \citep{boleydur10} show that grains significantly above interstellar size, can migrate toward spiral arms, increasing the dust-to-gas ratio and even affecting fragmentation itself. In our calculations we have assumed a size distribution appropriate for interstellar grains, and no grain growth (other than the growth of an ice shell), so that the assumption of coupling between gas and dust is satisfied. In the spiral shock the dust-to-gas ratio might go up by a factor of 10 if most grains have already grown to cm-size grains. This could open a new avenue for forming planetesimals rapidly via gravitational instability of the dust layer \citep{riceetal06}. Dust accumulation in spiral shocks might also be an additional effect promoting fragmentation in the gas phase \citep{boleydur10}, concurrent with those that we have just described. The results of our paper show that the detailed modeling of the response of the dust grains to spiral shocks has an important effect on the evolution of the opacity, one of the crucial parameters controlling the development and outcome of gravitational instability in massive disks. In their present form our results should be regarded as a "proof of concept" since the temperature evolution of the shocks that we have assumed in the paper has not been recomputed using our dust model, but rather stems from the \citep{d'alessio01} model adopted in the simulations. Incorporating the dust model that we have described in this paper directly in our three-dimensional hydrodynamical simulations of disks will be the first and most important next step. The stronger and faster opacity changes that we have found might produce fragmentation in a disk model that does not fragment with the standard opacity models used in \cite{mayer07} as well as in other published work. Likewise, since grain growth might have already taken place by the time the disk becomes spirally unstable, we will need to extend our model to larger grains, with sizes comparable to those that can accumulate efficiently within spiral shocks. \begin{figure}[ht] \centerline{\includegraphics [width=12cm]{fig1newshockb.eps}} \caption{Color-coded temperature maps of the hydrodynamical disk simulation. The temperature ranges from 30\,K (dark blue) to 250\,K (yellow). The $m=3$ mode is evident, with three strong spiral shocks developing in the region located between 12 and 15 AU from the central star (shown as a dark spot at the center). The disk is approximately 35 AU in radius, with the outer boundary naturally changing in time as a result of outward angular momentum transport from spiral density waves. The box is identified in a Lagrangian fashion by tracking the particles initially within the shock,and therefore its size changes in time (see text for a detailed description). Initially the box has an azimuthal extent of about 4\,AU and a height of 0.15\, AU. The four panels, equally spaced in time, cover about 20 years of evolution (from upper left to bottom right) before and after one of the spiral shocks reaches maximum amplitude (which happens in the bottom left panel), approximately corresponding to the time spanned by the profiles in Fig.\,\ref{shock}. The maximum amplitude corresponds to the temperature maximum shown in Fig.\,\ref{shock}. The azimuthally averaged profiles in Fig.\,\ref{shock} have been computed setting the location of the overlaid circle to $r=0$ and then considering increasingly larger annuli (with widths indicated as in Fig.\,\ref{shock}) within the boundaries of the box.} \end{figure}\label{shockfig1} \begin{figure} \centerline{\includegraphics[width=12cm]{shock.eps}} \caption{Azimuthally averaged gas temperature (blue curves) and gas density ($\times 10^5$) (red curves) for three regions in the spiral shock. measured at the disk midplane. The three regions correspond to annuli at $r= 0.2 AU$ (solid curve), 0.54 AU (dashed curve) and 1.3 AU (dotted curve), where $r=0$ is defined as the radius of the reference circle that approximately follows the shock peak amplitude, as shown in Figure 1.} \label{shock} \end{figure} \begin{figure}[p] \centerline{\includegraphics[width=10cm]{opac_t0.2.eps}\includegraphics[width=10cm]{opactime.54.eps}} \centerline{\includegraphics[width=10cm]{optime.72.eps}\includegraphics[width=10cm]{optime1.3.eps}} \caption{Opacity as a function of time for the annuli at 0.2 AU (upper left), 0.54AU (upper right), 0.72AU (lower left) and 1.3AU (lower right) in the density wave.} \label{opact} \end{figure} \begin{figure} \centerline{\includegraphics[width=20cm]{eflux.eps}} \caption{Energy flux from the grain to the gas as a function of time for the 5 grain sizes considered. Negative values indicate that the grain is colder than the gas. The numbers next to each curve refer to the grain size bin number corresponding to that curve. The gas temperature is shown as a dashed curve. The peak flux near $4\times 10^8$\,s is due to water vapor recondensing on the grains and heating them. This corresponds to a gas temperature of 165\,K, with the grains about 0.3\,K hotter.} \label{eflux} \end{figure} \begin{figure}[p] \centerline{\includegraphics[width=15cm]{shlthk.eps}} \caption{Shell thickness as a function of time for the 0.2\,AU annulus in the density wave. The numbers next to the curves give the core radius in microns. Note that the thickest shell does not surround the largest core.} \label{shlthck} \end{figure} \begin{figure}[p] \centerline{\includegraphics[width=8cm]{optemp.2.eps}\includegraphics[width=8cm]{optemp.54.eps}} \centerline{\includegraphics[width=8cm]{optemp.72.eps}\includegraphics[width=8cm]{optemp1.3.eps}} \caption{Opacity as a function of temperature for the 0.2 AU (upper left), 0.54AU (upper right), 0.72AU (lower left) and 1.3AU (lower right) annuli in the density wave. The black arrows show the direction of the opacity change as a function of time. The dotted curve in the upper left hand panel shows the opacity of \cite{pollack85} while the dashed curves in all the panels show the opacities of \cite{d'alessio01}.} \label{opactemp} \end{figure} \begin{figure} \centerline{\includegraphics[width=20cm]{repshock.eps}} \caption{Opacity as a function of temperature for the 0.2\,AU case. The green triangles show the case discussed earlier (no ice on the olivine cores at $t=0$). The blue diamonds show the same case but with an ice to olivine ratio of 2:1 at $t=0$. Because the original ice arrangement is different the opacity is different at the starting point ($t=0$). Once the shock heats the grains enough for the ice to migrate, it quickly rearranges itself in the same way, independent of the original distribution. After a second shock (red squares) the opacity - temperature relation remains unchanged.} \label{repshock} \end{figure} \begin{figure} \centerline{\includegraphics[width=8cm]{10rwoptime.2.eps}\includegraphics[width=8cm]{10wopact.2.eps}} \centerline{\includegraphics[width=8cm]{pwropact.eps}\includegraphics[width=8cm]{pwropacttemp.eps}} \caption{Opacity as a function of temperature for the 0.2\,AU annulus in the shock but where the abundances are 10 times the solar value (upper left), and for where only the water is 10 times the solar value(upper right). Also shown are solar abundances with 3 values of $\xi$: 1.5 (blue curve), 3.5 (red curve), and 5.5 (green curve). The opacity as a function of time is shown in the lower left hand panel and the opacity as a function of temperature is shown in the lower right hand panel.} \label{nonstopac} \end{figure} \begin{figure} \centerline{\includegraphics[width=14cm]{lowopt-t.eps}} \caption{Opacity as a function of temperature for the case of 0.2 AU in the shock but where the optical depth to the star is essentially zero. The blue curve is for a gas density equal to the gas density at midplane, and the red curve is for $10^{-4}$ of the midplane density.} \label{lowoptt} \end{figure}

1012.5269

We investigate the evolution of grains composed of an ice shell surrounding an olivine core as they pass through a spiral shock in a protoplanetary disk. We use published three-dimensional radiation-hydrodynamics simulations of massive self-gravitating protoplanetary disks to extract the thermodynamics of spiral shocks in the region between 10 and 20 AU from the central star. As the density wave passes, it heats the grains, causing them to lose their ice shell and resulting in a lowering of the grain opacity. In addition, since grains of different sizes will have slightly different temperatures, there will be a migration of ice from hotter grains to cooler ones. The rate of migration depends on the temperature of the background gas, so a grain distribution that is effectively stable for low temperatures can undergo an irreversible change in opacity if the gas is temporarily heated to above ~150 K. We find that the opacity can drop more and at a significantly faster rate throughout the spiral shocks relative to the prediction of the standard dust grain model adopted in hydrodynamical calculations of protoplanetary disks. This would lead to faster gas cooling within spiral arms. We discuss the implications of our results on the susceptibility of disks to fragment into sub-stellar objects at distances of a few tens of astronomical units.

12205274

2011IAUS..276..478G

1012

1012.0811_arXiv.txt

1012.0811

We present a re-analysis of archival HST/NICMOS transmission spectroscopy of the exoplanet system, HD 189733, from which detections of several molecules have been claimed. As expected, we can replicate the transmission spectrum previously published when we use an identical model for the systematic effects, although the uncertainties are larger as we use a residual permutation algorithm in an effort to account for instrumental systematics. We also find that the transmission spectrum is considerably altered when slightly changing the instrument model, and conclude that the NICMOS transmission spectrum is too dependent on the method used to remove systematics to be considered a robust detection of molecular species, given that there is no physical reason to believe that the baseline flux should be modelled as a linear function of any chosen set of parameters.

12132809

2010arXiv1012.2213R

1012

1012.2213_arXiv.txt

\noindent We have come a long way since \cite{b2fh,agw} but the mystery of the origin of the p-nuclides is still with us. Although there are considerable uncertainties in the astrophysical (hydrodynamical) modeling of the sites possibly producing p-nuclei, a sound base of nuclear reaction rates is essential for all such investigations. As long as an experimental determination of the rates around the deflection points is impossible, measurements of low energy cross sections of stable nuclides are essential to test and improve the theoretical calculations. This has been underlined by the recent results regarding optical potentials for the interaction of charged nuclei. Further measurements (at even lower energy) are highly desireable and, along with improved, self-consistent hydrodynamic simulations of the possible production sites, will gradually improve our understanding of p-nucleosynthesis.

1012.2213

A number of naturally occurring, proton-rich nuclides (the p-nuclei) cannot be made in the s- and r-process. It has been found that massive stars can produce p-nuclei through photodisintegration of pre-existing intermediate and heavy nuclei. This so-called gamma-process requires sufficiently high temperatures and occurs in pre-explosive or explosive O/Ne burning, depending on the mass of the star. Although the gamma-process has been successful in producing a large range of p-nuclei, two mass regions remain problematic, A&lt;110 and 150&lt;A&lt;165, where a number of p-nuclei are severely underproduced. The origin of the problems is yet to be identified. A large number of unstable nuclei with only theoretically predicted reaction rates are included in the reaction network and thus the nuclear input may involve uncertainties. Deficiencies in charged-particle optical potentials at gamma-process temperatures have been found for nuclei at stability. On the other hand, the gamma-process conditions (temperature profiles, entropy of the O shell, seed composition) also sensitively depend on details of the stellar structure and evolution, as well as on the initial metallicity. Nevertheless, especially the deficient low-mass p-nuclei may call for an additional production process or site, such as production in (subChandrasekhar) type Ia supernovae. Also the (im)possibility of a synthesis in the rp- and nu-p-processes is discussed. Were this the case, the production of p-nuclei would be realized as a superposition of several different processes, not necessarily in the same site.

12213692

2011MNRAS.413..573B

1012

1012.1606_arXiv.txt

Within the hierarchical structure--formation scenario, galaxy clusters are key targets that allow us to study both the dynamics on the gravity--dominated scale and the complexity of astrophysical processes dominating on the small scale. In such studies their mass is one of the most crucial quantities to be evaluated, and the bulk properties measured from X--ray observations still provide the best way to estimate the mass, primarily on the assumption of hydrostatic equilibrium \cite[][]{sarazin1988}. Mass estimates rely then on the assumptions made about the cluster dynamical state, since the Hydrostatic Equilibrium Hypothesis (HEH) implies that only the thermal pressure of the hot ICM is taken into account \cite[][]{rasia04}. Lately, it has been claimed in particular that non--thermal motions, as rotation, could play a significant role in supporting the ICM in the innermost region \cite[e.g. ][]{lau2009,fang2009} biasing the mass measurements based on the HEH. The analysis of simulated cluster--like objects provides a promising approach to get a better understanding of the intrinsic structure of galaxy clusters and the role of gas dynamics, which can be eventually compared to X--ray observations. Because of the improvement of numerical simulations, the possibility to study in detail the physics of clusters has enormously increased \cite[see][for a recent comprehensive review]{borgani2009} and future satellites dedicated to high--precision X--ray spectroscopy, such as ASTRO--H and IXO, will allow to detect these ordered motions of the ICM. With this perspective, we perform a preliminary study on the ICM structure for some clusters extracted from a large cosmological hydrodynamical simulation, investigating in particular the presence of rotational motion in the ICM velocity field.\\ \indent The paper is organized as follows. We describe the numerical simulations from which the samples of cluster--like haloes have been selected in \sec\ref{secSims}. In \sec\ref{secVrot} we consider a first set of simulated clusters and present results on build--up of rotation in the halo core for single cases of study (\sec\ref{subsecRot} and \sec\ref{subsecg51}), while results about the contribution to the mass estimations are given in \sec\ref{secMrot}. A second sample of clusters is then statistically investigated in \sec\ref{secPadme}. We discuss our results and conclude in \sec\ref{secConclusion}.\\ \indent Appendix A is devoted to comment on the effects of artificial viscosity, while in Appendix B we briefely comment on the ellipticity profiles of the simulated clusters. \section[]{Numerical Simulations} \label{secSims} We consider two sets of cluster--like haloes selected from two different parent cosmological boxes. In both cases the cosmological simulations were performed with the TreePM/SPH code GADGET--2 \cite[][]{springel2001,springel2005}, assuming a slightly different cosmological model (a standard \lcdm model and a WMAP3 cosmology, respectively) but including the same physical processes governing the gas component \cite[][]{springel2003}, i.e. radiative cooling, star formation, and supernova feedback (\textit{csf} simulation, see \cite{dolag2009} and references therein for a more detailed overview on different runs of the parent hydrodynamical simulations we refer to in our work). Additionally, we refer to simulations of the same objects without including radiative processes as \textit{ovisc}.\\ \indent \textbf{Set 1.} The first data set considered consists of 9 cluster--size haloes, re--simulated with higher resolution using the \openquote zoomed initial condition\closequote (ZIC) technique \cite[][]{tormen1997}. The clusters have been originally extracted from a large--size cosmological simulation of a \lcdm universe with $\om = 0.3$, $h = 0.7$, $\sigma_8 = 0.9$ and $\omb = 0.039$, within a box of $479 h^{-1} \mpc$ a side. The final mass--resolution of these simulations is $m_{DM}=1.13 \times 10^{9}h^{-1}\msun$ and $m_{gas}=1.69 \times 10^{8}h^{-1}\msun,$ for the DM and gas particles, respectively. The spatial resolution for Set 1 reaches $5 h^{-1}\kpc$ in the central parts and for the most massive clusters we typically resolve up to 1000 self--bound sub--structures within $R_{vir}$ \cite[as shown in][]{dolag2009}. The main haloes have masses larger than $\sim1.1 \times 10^{14}h^{-1}\msun$ and they have all been selected in a way that they are quite well--behaved spherically--shaped objects at present epoch, although a fair range from isolated and potentially relaxed objects to more disturbed systems embedded within larger structures is available. Having a reasonable dense sample within the time domain (e.g. 50 outputs between $z=1$ and today) and high resolution, we can study how common and significant the rotational support of the ICM is. In particular we focus on the detailed evolution of such rotational motions. \\ \indent \textbf{Set 2.} In the second data set we analyzed a volume limited sample of cluster--size haloes, where we computed the distribution of the ICM rotational velocity and compared their distribution at different redshift. This second set of clusters has been extracted from a large size cosmological simulation with a box--size of $300 h^{-1} \mpc,$ simulated with $2\times768^3$ particles, assuming the 3--year WMAP values for the cosmological parameters \cite[][]{spergel2007}, i.e. $\om = 0.268$, $\omb = 0.044$, $\sigma_8 = 0.776$ and $h = 0.704$. The final mass--resolution for this second set of simulations is $m_{DM}=3.71 \times 10^{9}h^{-1}\msun$ and $m_{gas}=7.28 \times 10^{8}h^{-1}\msun$. Given the larger sample, a fair investigation of the amount of rotational support within the ICM from a statistical point of view is then possible.\\ \indent The two sets of simulations analyzed were performed with a different value of $\sigma_8,$ meaning that differences in the merging histories can be introduced. Nevertheless, we stress that the purpose of the second set is only to enlarge the statistics on the build--up of rotational motions in simulated galaxy clusters, and not a direct comparison of single objects to the objects of Set 1. Therefore, we are confident that our conclusions do not depend on these differences in the parameters of the two simulations.

\label{secConclusion} In this work, we have presented the result of a study over two sets of hydrodynamical simulations performed with the TreePM/SPH code GADGET--2. The simulations include radiative cooling, star formation, and supernova feedback and assume slightly different cosmological models, a \lcdm one and a WMAP3 one. The main target of this analysis has been the importance of rotational gas motions in the central regions of simulated cluster--like haloes, as it is thought to be a crucial issue while weighing galaxy clusters and identifying them as relaxed systems. The objects selected from our samples guarantee a wide variety of virial masses and dynamical structures, so that a reliable investigation of this phenomenon is allowed.\\ \indent Our main results can be summarized as follows: \begin{itemize} \item As main conclusion, we notice that the occurrence of rotational patterns in the simulated ICM is strictly related to the internal dynamics of \textit{gas--rich} substructures in a complicated way, so that it is definitely important to take it into account as contribution to the pressure support, but it's not directly nor simply connected to the global dynamical state of the halo. \item In the first part of our analysis we focused on g51, a simulated cluster with a very smooth late accretion history, isolated and characterized by few substructures in comparison to the other massive objects within Set 1. Also, we compare it with a highly disturbed system (g1). Even in the radiative simulation of this cluster, likely to be considered relaxed in a global sense, \textit{no clear rotation} shows up \textit{at low redshift because of some minor merging events} occurring \textit{close to the innermost region}: the rotation of the core is found to be an intermittent phenomenon that can be easily destroyed by the passage of gas--rich subhaloes through the equatorial plane. Gas particles stripped from the subhalo passing close to the main--halo innermost region ($<0.1\rfive$), are likely to get mixed to the gas already settled and contribute over few orbits to change the inclination of the best equatorial plane, suppressing any pre--existing rotational pattern. \item The velocity maps plotted in \fig\ref{fig_maps} show several DM--only subhaloes moving close to g51 central core. In our study, they have been found not to disturb in any significant way the ordered rotational gas motions created in the innermost region. The central gas sloshing is mainly set off by gas--rich subhaloes, especially if they retain their gas during the early passages through the core. Interesting work on numerical simulations have been found to be relevant for the result presented here, as the study from \cite{ascasibar2006} on the origin of cold fronts and core sloshing in galaxy clusters. \item Mass measurements based on HEH are likely to misestimate the total mass of galaxy clusters because of contributions by non--thermal gas motions that have to be considered. In agreement with previous works, we also find that significant rotation of the ICM can contribute to the pressure support. While several studies have been carried out on turbulent motions in the ICM and on their effect on the cluster mass estimates \cite[e.g.][]{rasia04,fang2009,lau2009,zhuravleva2010}, only lately the work by \cite{fang2009} and \cite{lau2009} have been addressing the ordered rotational patterns that could establish in the innermost region ICM as the result of the cluster collapse. Therefore, a comprehensive analysis of the details of rotation build--up and suppression both in single high--resolution case--studies and in larger, statistically significant samples is extremely interesting, especially for relaxed objects where this should be more important than turbulence. Focusing on rotation specifically, we calculated the corresponding mass term, $\mrot,$ for the two clusters g51 and g1. As expected from the tangential velocity profiles at redshift $z=0,$ the mass term coming from ICM rotational motions contributes more in the case of g1 than in g51, providing evidence that rotational support of gas in the innermost region is more significant in the former than in the latter. While $\mrot$ accounts for few percents at radii close to $\rfive$ in both cases, in the central regions up to $\sim17\%$ of the total true mass in g1 is due to rotational motions of the ICM. As regards g51, this contribution is less important, as no strong rotation has been found at $z=0,$ but it still reaches a value of $\sim10\%$ for the pressure support in the cluster core. \item Extending the analysis to a larger sample, we have investigated the statistical distribution of rotational velocity over dynamically--different clusters, isolated in a limited--volume simulated box such that their virial mass ($\mtwom$) is above a chosen threshold. At $z=0$ as well as at higher redshift up to $\sim0.5,$ a fair sample of cluster--size haloes let us infer that, on average, no high--velocity rotational patterns show up in the halo cores (i.e. in the region $<0.1\rfive$). Also for the clusters of Set 2, we find typical values of $\sim 200-300 \kms$ for the rotational velocity in the innermost region. \item We do not find any increasing trend of the rotational velocity distribution peak with decreasing redshift, that can correspond to the smooth mass assembly of the cluster--like halo through collapse. Although such trend is generally expected, it must be easily suppressed by internal minor events disturbing the halo central region. \end{itemize} \indent We conclude that the build--up of rotational patterns in the innermost region of galaxy clusters is mainly related to the physical processes included in the \textit{csf} run to describe the intracluster gas. On the contrary, numerical effects such as different implementations of artificial viscosity \cite[][]{dolag2005} do not affect in any significant way our results (see Appendix~\ref{appendix}, for a detailed discussion).\\ \indent An analogous conclusion can be drawn with respect to the differences between the two samples introduced by cosmology and resolution. For both Set 1 and Set 2 the build--up and suppression of rotational patterns in the halo central part is found to be mainly related to the physics included in the radiative run. In fact, comparable subsamples of the two sets in the \textit{csf} simulations show very similar distributions of rotational velocities for the ICM component in the halo innermost region, meaning that the shape of the distribution is essentially dominated by the physics of the gas. \indent Usually, relaxed clusters are assumed to have little gas motions. Therefore they are likely to be the best candidates for the validity of the HEH, on which mass estimations are based. Nevertheless, rotational motions should establish preferentially in relaxed clusters with respect to disturbed systems as a consequence of the assembling process, potentially representing a danger for relaxed cluster masses. Here, however, we find that the processes described in the paper save the reliability of the HEH--based mass determinations in most of the cases. In fact, rotational motions are not significant enough to compromise dramatically mass determinations with the exception of few outliers. In our simulation, the identification of relaxed or non--relaxed clusters according to the presence of gas rotation in the central region is not straightforward, since it has been shown to appear and disappear periodically. Its contribution has to be considered whenever is present, but it is not directly related to the global state of the simulated halo. Also, it is likely to be strongly influenced by the overcooling problem affecting hydrodynamical simulations, which has the effect to enhance the process of building up rotational patterns in the ICM in the innermost regions of simulated clusters.\\ \indent Although various theoretical and numerical studies in addition to the present work have been investigating the existence of gas bulk, non--thermal motions and the possible ways to detect them in galaxy clusters \cite[e.g.][]{fang2009,lau2009,zhuravleva2010}, little is known from observations. In a recent study, \cite{lagana2009} have made use of assumptions from theoretical models and numerical simulations about cosmic rays, turbulence and magnetic pressure to consider these non--thermal contributions to the total mass measurement for five Abell clusters. From a pure observational point of view, previous work has been able to confirm only indirect indications of bulk gas motions associated to merging events in galaxy clusters \cite[see][for a review]{markev2007} or evidences for turbulent gas motions, like the ones found in the Coma cluster in \cite{schuecker2004} or those inferred, on the scale of smaller--mass systems, from the effects of resonant scattering in the X-ray emitting gaseous haloes of large elliptical galaxies \cite[][]{werner2009}. Although not possible so far, the most direct way to measure gas motions in galaxy clusters would be via the broadening of the line profile of heavy ions (like the iron line at $\sim6.7\kev$ in X--rays) for which the expected linewidth due to impact of gas motion is much larger than the width due to pure thermal broadening. The possibility to use the shape of the emission lines as a source of information on the ICM velocity field as been discussed in detail in \cite{inogamov2003} and \cite{sunyaev2003}, and lately in \cite{rebusco2008}. Though, the investigation of the imprint of ICM motions on the iron line profile requires high--resolution spectroscopy, which will become possible in the near future with the next--generation X--ray instruments such as ASTRO--H and IXO. This will allow us to directly detect non--thermal contributions to the cluster pressure support, such as rotational patterns in the ICM, and enable us to take this correctly into account as contribution to the total mass estimate. Ultimately, this is likely to be an important issue to handle in order to better understand deviations from the HEH, on which scaling laws are usually based.

1012.1606

Gas motions in the hot intracluster medium (ICM) of galaxy clusters have an important effect on the mass determination of the clusters through X-ray observations. The corresponding dynamical pressure has to be accounted for in addition to the hydrostatic pressure support to achieve a precise mass measurement. An analysis of the velocity structure of the ICM for simulated cluster-size haloes, especially focusing on rotational patterns, has been performed, demonstrating them to be an intermittent phenomenon, strongly related to the internal dynamics of substructures. We find that the expected build-up of rotation due to mass assembly gets easily destroyed by passages of gas-rich substructures close to the central region. Though, if a typical rotation pattern is established, then the corresponding mass contribution is estimated to be up to ∼17 per cent of the total mass in the innermost region and one has to account for it. Extending the analysis to a larger sample of simulated haloes, we statistically observe that (i) the distribution of the rotational component of the gas velocity in the innermost region has typical values of ∼200-300 km s<SUP>-1</SUP>; and (ii) except for few outliers, there is no monotonic increase in the rotational velocity with decreasing redshift, as we would expect from approaching a relaxed configuration. Therefore, the hypothesis that the build-up of rotation is strongly influenced by internal dynamics is confirmed and minor events like gas-rich substructures passing close to the equatorial plane can easily destroy any ordered rotational pattern.

12131653

2010arXiv1012.3090D

1012

1012.3090_arXiv.txt

\label{s-intro} Black holes are one of the most well studied objects in general relativity. When combined with theory, several astronomical observations exhibit strong evidence for existence of stellar mass (few times solar mass $M_{\odot}$)~\cite{Casares:2006vx} as well as supermassive (mass $10^5-10^9 M_{\odot}$)~\cite{2005SSRv..116..523F} black holes. The evidence of stellar mass black holes typically comes from observations of binary system, while by recent observations the supermassive black holes are expected to exist at the center of almost every galaxy. In recent times there are some observational evidences~\cite{2009Natur.460...73F} for intermediate mass black holes. Existence of primordial black holes~\cite{Carr:2005zd} in the early universe is also speculated. Most of the primordial black holes are expected to have very small mass. Since the black holes have a trapped surface, singularity theorems~\cite{HawkEllis} guarantee a singularity inside their horizon, although, they are silent about their structural details. Singularities also form without or before formation of a trapped surface or an event horizon. If such a singularity forms, then non-spacelike geodesics come out of it and in principle the singularity can be visible to an outside i.e., nonsingular observer; therefore, it is called a naked singularity. There are speculations~\cite{Witten:1992fp} that gamma ray bursts originate from such naked singularities. We cannot foretell laws of physics at a singularity, hence existence of a naked singularity can lead to breakdown of predictability. Normally, we do not want such a situation to arise in nature. With this in mind Penrose~\cite{Pen1} proposed the Cosmic Censorship Hypothesis (CCH). The CCH asserts that, naked singularities should not/do not form from collapse of a reasonable matter field when we start with a generic nonsingular initial data. There are many investigations supporting the CCH as well as many examples of existence of naked singularities~\cite{HaradaICgc,EardlySmarr, ChirstoTBL, NewmanTBL, PsjTP, SanjTBL1, SarayGen, Dwivedi:1998ts, PsjPrd93, ShrirLum1, ShrirTime, NolanNrng, Ori:1987hg, Ori:1989ps, HaradaPfluid, HaradaMena, MagliTan2, TangSanjMagGov, BarveTan, SanjCluster, PsjMahaGo, HaradaTimeFam, Lake, DwiPsjVaidya, MagliExact, SanjNeg, GhoshDadhich2, PsjCmp94, ChrisScalar}. Normally, these investigations put several restrictions on the model. Almost all these studies assume a very specific and very restrictive form of matter. They also assume only a certain kind of metric, typically, either with spherical or axial symmetry. Therefore, they may not satisfy the genericity of initial data or reasonable matter field criteria. Hence, they do not prove or disprove the CCH. The main difficulty arises as in general we cannot establish the relationship of initial data to the formation and structure of the singularity even after assuming lot of symmetry and taking a very simplistic equation for matter. Therefore, proving or disproving the CCH remains one of the most important open problems in classical general relativity. In such a situation we need to investigate, whether some other effects safeguard the physics. It may happen that, even if naked singularities exist geometrically, they have no physical or observational consequences. i.e. they cannot affect the world outside them and other physical effects make them benign. Naked singularities and the CCH get lot of attention and is important as general relativity is a physical theory which describes nature. Unpredictability is a serious issue for a physical theory and needs to be fixed. Therefore, it is important to find out whether a naked singularity can affect the physics outside it, i.e., whether a nonsingular observer can distinguish it observationally. With this in mind we studied the spherically symmetric dust collapse model~\cite{CanWe1} and calculated the redshift and luminosity of light rays coming out of the singularity. Subsequently, we also studied extremely wide class of (any type II matter) spherically symmetric collapse models for the same purpose. In both the cases we showed that at the most along one singular null geodesic coming out of a null (naked) singularity the redshift is finite while redshift diverges along all other outgoing singular null geodesics. Hence, we concluded that the null naked singularities will not be physically troublesome, as no energy can come out of them. Here, we generalize the proof for any null singularity. In this work, we consider a completely general spacetime metric without any symmetries. Apart from weak energy condition~\cite{HawkEllis}, no restrictions on form of matter are assumed. We show that if a null singularity forms in the collapse, then at the most for one singular null geodesic the redshift is finite while it is infinite for all other (infinite family) of singular geodesics. Hence no energy or information can come out of a null naked singularity. We need a wavepacket to carry energy, thereupon, redshift should be finite for a finite (though it can be very small) duration to get out energy from the singularity. As no energy can come out of the null singularity, it cannot affect the physics outside and there should not be any danger of breakdown of predictability. This means, physically the null naked singularities are not important or they are not dangerous. We also show that the redshift is always finite for null geodesics coming out of a timelike naked singularity. This means that, in principle such a singularity can be observed and they can be more problematic. However, as such we expect them to be very rare~\cite{CanWe1}.

\label{s-summary} In summary, in this work, writing the metric in the Bondi-Sachs form we have shown that no energy can come out of any null naked singularity. The advantage of using null formulation is that it is designed to study the null geodesics which we want to explore. Hence no extra work is needed to get out the quantities of our interest. We have not imposed any symmetries on the spacetime nor we have assumed any specific form of matter. Essentially the assumption that the singularity which forms is null and just the geometry of spacetime is enough for us to reach this conclusion. We also showed that a timelike singularity is in principle likely to be visible to an outside/nonsingular observer as redshift is always finite for rays coming out from it. However, we expect formation of a timelike singularity to be rare when energy conditions hold, as it seems~\cite{CanWe1} creation of such a singularity needs fine balance of factors; we require the collapse to stop as soon as the singularity is formed. That indicates they will be non-generic. The known examples of naked singularities are also mostly null, or else, they are formed in a very unusual matter fields and need lot of fine tuning. We have also given the fundamental principle/logic which drive the result. As we have explained above, the result is essentially very similar to the special relativistic result for a null (or timelike) source. For the null singularity (surface) the redshift basically diverges as the proper time goes to zero on null surface. So if a ray has finite frequency outside the singularity (surface) and it originates at the singularity, then it has to have infinite frequency at the singularity. Similarly for timelike singularity as the proper time is always finite (nonzero) at the singularity so the redshift is finite, though it can be very large. Our result is valid for any form of matter as it is purely based on the geometry at the source i.e., geometry at the singularity. That means our result(s) are most general and should be valid in any theory of gravity as well as, in higher dimensions. One can generalize these results easily along the timelike geodesics. The results basically means that, though the null singularity is geometrically naked (i.e., null geodesics can come out of it) essentially physically it is not visible (naked), as no energy can come out of it due to infinite redshift. That implies we cannot get any information from the null naked singularity and it will not have any undesirable physical effect outside. Therefore, null naked singularity cannot cause breakdown of predictability and they have no special physical significance; i.e. the fact that non-spacelike geodesics come out of null naked singularity does not have any significance. Our conclusion is extremely general and is valid for all possible spacetimes. This strongly supports/preserves the essence of comic censorship for null singularities. However, we need to study the timelike naked singularities in more detail for this purpose. I would liked to thank IUCAA, Pune and CTP, Jamia Millia Islamia, Delhi for hospitality where early part of this work was done.

1012.3090

In this work we study collapse of a general matter in a most general spacetime i.e., a spacetime with any matter and without (assuming) any symmetry. We show that the energy is completely trapped inside the null singularity and therefore they cannot be experimentally observed. This most general result implies, there is no physical significance of the null naked singularities irrespective of their existence. This conclusion strongly supports the essence of cosmic censorship hypothesis.

2775998

2011ApJ...728L..36P

1012

1012.0570_arXiv.txt

Recently a novel type of faint supernovae (SNe) with peculiar properties has been discovered. A group of eight such events has been identified, all spectroscopically similar to type Ib SNe, but faint (typical absolute peak magnitude of $\sim-15$) and calcium rich \citep{per+10}. Although type Ib SNe are generally thought to result from core-collapse of massive stars (e.g. \citealt{fil97}), a large fraction of the host galaxies of these faint, Ca-rich SNe are early type galaxies. Additionally, the ejecta mass of SNe in this subclass appear to be very low (e.g. $\sim0.3\,M_{\odot}$ found for SN 2005E, \citealp{per+10}; and $\lesssim1\,M_{\odot}$ found for SN 2005cz, \citealp{kaw+10}), less than expected and observed for core-collapse SNe of any type. These SNe were therefore suggested to originate from a different process involving the thermonuclear explosion of a helium-shell on a white dwarf \citep{per+10,she+10,wal+10}. Nevertheless, an alternative scenario involving a core-collapse of a $10-12\,M_{\odot}$ star, which is a part of a binary, was suggested by \cite{kaw+10} for the origin of one of the Ca-rich SNe; SN 2005cz. Here we study this possibility, and look for any evidence for SF or young stellar population near the location of SN 2005cz. In the following we discuss our results from observations of various SF tracers including $H\alpha$ emission, HST photometry, host galaxy spectra and UV emission. In addition we shortly discuss the star formation and merger history of the host galaxy.

In this letter we studied the local and global environment of SN 2005cz in the elliptical galaxy NGC 4589. We used various SF tracers including optical spectroscopy, $H\alpha$ emission, UV emission and HST photometry. We also reviewed the the merger history of the host galaxy. We found that although some H$\alpha$ emission (which in principle could trace SF activity) exists in the host galaxy, it is far (>1.5 kpc away) from, and unrelated to the close environment of SN 2005cz. Moreover, this emission is more likely to be associated with an AGN in the nucleus, rather than trace SF activity. Other star formation tracers (HST imaging of young massive stars, UV, R-band imaging and host galaxy spectrum) show no evidence for SF in the galaxy, and particularly close to the reported location of the SN. The UV emission data could trace stars down to lower mass than H$\alpha$ emission \cite[e.g.][]{gog+09}. Therefore, while H$\alpha$ emission may not be detected in older SF regions, in which core-collapse SNe from progenitors of $10-12\,M_{\odot}$ may explode, they should still present significant UV emission. The lack of such UV emission therefore suggests that recent SF activity has not occurred in this galaxy, and in particular close to the location of SN 2005cz. In addition, the overall structure and colors of the host galaxy show no evidence of recent SF in the last Gyr. The HST data exclude very massive progenitors, suggested to be the progenitors of type Ib SNe, and show no evidence for massive young clusters or $>15\,M_{\odot}$ supergiants near the SN location. In principle, a massive progenitor could have formed far from the observed SN location, and later have been ejected at high velocity to explode far from its birth place. Since we find no evidence for SF even up to 1 kpc from the SN location, such a star should have been a runaway star to form so far (velocities of 30-100 km s$^{-1}$, for a lifetime of 10-40 Myr). According to \cite{kaw+10} the progenitor is suggested to be a binary star. However, runaway (or hypervelocity) binary stars, especially massive ones as required for a SN progenitor, are rare \citep{leo+90,per09b,per+10,per+10b}. Therefore, although possible in principle, such a scenario would require fine tuned conditions. Taken together, the analysis of UV and H$\alpha$ emission, the spectrum of the host galaxy NGC 4589, the 2MASS and RCS photometry as well as out HST data, show no evidence for recent SF near the location of SN 2005cz or even at large distances from it. We conclude that our results strongly disfavor a young massive-star progenitor for SN 2005cz. These results are consistent with the host galaxy type of other Ca-rich faint type Ib SNe, found to be biased towards early type galaxies \citep{per+10}. Moreover, the only SNe of types II or Ib/c ever to be found in elliptical galaxies (SNe 2000ds and 2005cz) are both faint Ca-rich type Ib SNe similar to SN 2005E \citep{hak+08,per+10}. This provides additional support to the suggested origin of these SNe from a helium detonation in a WD-WD binary system, i.e. from a low mass old progenitor rather than a core-collapse of a young massive star.

1012.0570

The supernova SN 2005cz has recently attracted some attention due to the fact that it was spectroscopically similar to type Ib supernovae (SNe Ib), a class that is presumed to result from the core collapse of massive stars, yet it occurred in an elliptical galaxy, where one expects very few massive stars to exist. Two explanations for this remarkable event were put forward. Perets et al. associate SN 2005cz with the class of Ca-rich, faint SNe Ib, which likely result from old double-white-dwarf systems with an He-rich secondary. On the other hand, Kawabata et al. suggest that SN 2005cz is indeed a core-collapse event (in a binary system), albeit of a star at the lower end of the mass range, 10-12 M <SUB>sun</SUB>. The existence of this star in its elliptical host is explained as resulting from low-level star formation (SF) activity in that galaxy. Here we present extensive observations of the location of SN 2005cz, sensitive to a variety of SF tracers, including optical spectroscopy, Hα emission, UV emission, and Hubble Space Telescope photometry. We show that NGC 4589, the host galaxy of SN 2005cz, does not show any signatures of a young stellar population or recent SF activity either close to or far from the location of SN 2005cz.

12163199

2011A&A...528A.121B

1012

1012.0436_arXiv.txt

} During 2009 the space missions CoRoT and \kepler\ have generated a high level of activity in the groups specialized in the asteroseismic analysis of photometric time series of stars. In both cases the very long, continuous observing and the very low noise data has opened up a completely new world of possibilities for the asteroseismic investigation of stellar interiors. In particular, the seismic investigation of K~giants has taken a huge leap forward. The results from a time series analysis of 150 days of measurements obtained by the CoRoT space telescope increased the number of known pulsating giants from a handful to nearly 800 \citep{deridder}. The \kepler\ mission is observing the flux continuously of thousands of stars for at least 3 years and has increased both the number, the range in luminosity, and the length of the time series compared to CoRoT. The high-precision light curves from \kepler\ constitute important data for detailed asteroseismic investigations of red giants due to the long temporal coverage and low noise levels of the observations. This has extended the range of giants with detected oscillations to lower luminosities \citep{bedding,stello3,mosser}. Before we can hope to make a successful analysis of individual red giant stars observed by \kepler\, we need to measure accurate atmospheric parameters as discussed by \citet{brown}, \citet{creevey1}, and \citet{creevey2}. We have started the observations of about 100 \kepler\ red giant stars with the FIbre-fed Echelle Spectrograph (FIES) spectrograph at the Nordic Optical Telescope (NOT). % From the spectral analysis we can determine accurate atmospheric parameters, which are essential for constraining the stellar models when comparing asteroseismic observations and theory. We concentrate in particular on old, metal poor stars, which are important for the understanding the early history of the Galaxy. We will also obtain important insight about the observed variation of the pulsational behaviour with metallicity. At present only photometric determinations of the metallicity are available, based on the Kepler Input Catalogue \citep{kic}. The KIC is a photometric catalogue with estimated parameters of all stars down to $V \simeq 18$ in the \kepler\ field of view. The values of \feh\ have been shown to be inaccurate \citep{molenda}, and we here confirm this based on the present analyses. \begin{figure*}[t] \centering \includegraphics[width=16cm]{fig01.eps} \caption{Examples of diagnostic plots of \feone\ abundance vs.\ excitation potential and equivalent width for three different giants with KIC-IDs 11342694, 4157282, and 8017159 (top to bottom).} \label{eqwexi} \end{figure*}

We have determined accurate atmospheric parameters for a sample of 14 K~giant targets that are being observed with the NASA \kepler\ satellite. The parameters are mandatory to put constraints on asteroseismic models when comparing observations and theory. We confirm the results by \cite{molenda} that there are serious discrepancies for \feh\ when comparing with the photometric KIC catalogue (RMS scatter in \feh\ of 0.5 dex), while \teff\ and \logg\ values are in reasonable agreement. However, for \logg\ we find discrepancies of about $1$ dex for two stars with ${\rm [Fe/H]} < -1.0$, indicating that there may be a problem in the KIC catalogue at low metallicities. We have validated our method and evaluated systematic errors from the analysis of 6 bright giants with well-known parameters and compared with results in the literature, confirming that our analysis is reliable. Also we see good agreement between our parameters and the ones found from asteroseismology. The uncertainties on \logg\ and \feh\ in KIC are too large to match the quality of the data produced by \kepler, and this emphasizes the importance and need for further, detailed spectroscopic studies of the \kepler\ giant targets. This paper will be followed by a second paper presenting the results for an additional 50 K~giants. We have verified that one of the \kepler\ giants is a population II star (KIC 8017159), and we expect to find several more in our larger sample of stars. Until now, only one nearby population~II star, $\nu$~Ind, has been studied using asteroseismic techniques \citep{nuind}. \begin{table} \caption{Comparison of parameters from VWA and the PASTEL catalogue \citep{soubiran} for six bright giants.} \centering \setlength{\tabcolsep}{3pt} % \begin{tabular}{r|ccc|ccc} \hline \hline & \multicolumn{3}{c|}{VWA} & \multicolumn{3}{c}{PASTEL}\\ ID & \teff & \logg & \feh & \teff & \logg & \feh \\ \hline $\alpha$~Mon & $4850\pm70$ & $ 2.77\pm0.25$ & $ +0.08\pm0.08$ & $ 4794 $ & $ 2.62$ & $ -0.04 $ \\ $\mu$~Leo & $4660\pm90$ & $ 2.63\pm0.24$ & $ +0.53\pm0.11$ & $ 4509 $ & $ 2.29$ & $ +0.31 $ \\ $\alpha$~Boo & $4300\pm70$ & $ 1.43\pm0.27$ & $ -0.52\pm0.09$ & $ 4316 $ & $ 1.71$ & $ -0.55 $ \\ $\lambda$~Peg & $4830\pm90$ & $ 2.56\pm0.26$ & $ -0.08\pm0.08$ & $ 4775 $ & $ 2.47$ & $ -0.09 $ \\ $\mu$~Peg & $5100\pm70$ & $ 2.96\pm0.26$ & $ +0.05\pm0.08$ & $ 4986 $ & $ 2.74$ & $ -0.08 $ \\ $\psi$~UMa & $4600\pm70$ & $ 2.11\pm0.25$ & $ -0.04\pm0.10$ & $ 4605 $ & $ 2.38$ & $ -0.13 $ \\ % \hline \hline \multicolumn{7}{l}{Uncertainties on \teff, \logg, and \feh\ for PASTEL are 80\,K, 0.1 dex}\\ \multicolumn{7}{l}{and 0.1 dex, respectively.} \end{tabular} \label{soubtable} \end{table}

1012.0436

Context. Accurate fundamental parameters of stars are mandatory for the asteroseismic investigation of the Kepler mission to succeed. <BR /> Aims: We determine the atmospheric parameters for a sample of six well-studied bright K giants to confirm that our method produces reliable results. We then apply the same method to 14 K giants that are targets of the Kepler mission. <BR /> Methods: We used high-resolution, high signal-to-noise ratio spectra acquired using the FIES spectrograph on the Nordic Optical Telescope. We applied the iterative spectral synthesis method VWA to derive the fundamental parameters from carefully selected high-quality iron lines and pressure-sensitive Calcium lines. <BR /> Results: We find good agreement with parameters from the literature for the six bright giants. We compared the spectroscopic values with parameters based on photometric indices in the Kepler Input Catalogue (KIC). We identify serious problems with the KIC values for [Fe/H] and find a large RMS scatter of 0.5 dex. The log g values in KIC agree reasonably well with the spectroscopic values displaying a scatter of 0.25 dex after excluding two low-metallicity giants. The T<SUB>eff</SUB> values from VWA and KIC agree well with a scatter of about 85 K. We also find good agreement with log g and T<SUB>eff</SUB> derived from asteroseismic analyses for seven Kepler giant targets. <BR /> Conclusions: We determine accurate fundamental parameters of 14 giants using spectroscopic data. The large discrepancies between photometric and spectroscopic values of [Fe/H] emphasize the need for further detailed spectroscopic follow-up of the Kepler targets. This will be mandatory to be able to produce reliable constraints for detailed asteroseismic analyses and interpretation of possible exo-planet candidates found around giant stars. <P />Based on observations made with the Nordic Optical Telescope, operated on the island of La Palma jointly by Denmark, Finland, Iceland, Norway, and Sweden, in the Spanish Observatorio del Roque de los Muchachos of the Instituto de Astrofisica de Canarias.Reduced spectra are only available in electronic form at the CDS via anonymous ftp to cdsarc.u-strasbg.fr (130.79.128.5) or via <A href="http://cdsarc.u-strasbg.fr/viz-bin/qcat?J/A+A/528/A121">http://cdsarc.u-strasbg.fr/viz-bin/qcat?J/A+A/528/A121</A>

12224090

2011PhRvD..83b4036B

1012

1012.2433_arXiv.txt

The concepts of dark energy (DE) and dark matter (DM) \cite{DE1}-\cite{DM3} considered as two manifestations of one unified dark fluid \cite{DF1}-\cite{DF5} are basic elements of modern cosmology and astrophysics. Because the dark fluid is estimated to accumulate about $95\%$ of the Universe's energy, the coupling of DE with DM predestine the main features of the Universe evolution. In the first part of the work \cite{antigauss} we introduced the so-called Archimedean-type force, which is linear in the four-gradient of the DE pressure and acts on the DM particles. This force is a relativistic generalization of the classical Archimedean force and belongs to the class of effective forces described in \cite{BZ1}-\cite{BZ6}. From mathematical point of view the presented model gives us a new self-consistent nonlinear scheme of interaction between two constituents of the cosmic medium. From physical point of view the Archimedean-type force is an effective redistributor of the total energy of the Universe between the DE and DM constituents. The Archimedean-type model can be also considered in terms of two-fluid representation of the cosmic medium \cite{M1}-\cite{M6}, in which the interaction terms $\pm Q$ appear in the right-hand sides of separate balance equations for the DE and DM with opposite signs and which disappear in a sum, when one deals with the total balance equation. We formulated in \cite{antigauss} the theory of interaction between DE and DM by using the relativistic hydrodynamics for dark energy and relativistic kinetics for dark matter. In the first paper \cite{antigauss} we focused on the submodels with some special values of guiding parameters, which admit exact analytical solutions of the total self-consistent system of master equations; in particular, we discussed the so-called anti-Gaussian solution. Now we consider the results of numerical and qualitative analysis for all admissible sets of guiding parameters and initial data. This paper is organized as follows. In Sec.II we recall briefly the key equations and basic formulas of the model with Archimedean-type coupling between dark energy and dark matter. In Sec.III the numerical results are presented: in Sec.III A we explain the scheme of numerical analysis and details of results representation; in Sec.III B we classify the models using the number of transition points and discuss details of seven submodels, the perpetually accelerated and periodic universes being among them. In Sec.IV we analyze the model of Archimedean-type coupling in terms of dynamic system associated with nonlinear key equation for the DE pressure; in Sec.IV A we study the toy model, which relates to the autonomous dynamic system, find critical points and discuss two basic phase portraits of this system; in Sec.IV B we analyze instantaneous phase portraits of a general nonautonomous dynamic system in the context of numerical results presented in Sec.III; in Sec.IV C we consider asymptotical behavior of the models. Sec.V contains discussions: in Sec.V A we attract an attention to the inflationary behavior of the model in the early Universe; in Sec.V B we show that all the models under discussion manifest the late-time accelerated expansion; in Sec.V.C we touch on the coincidence problem.

The Archimedean-type model can describe qualitatively two cornerstones of the Universe's evolution: the inflation in the early Universe and the late-time accelerated expansion. Also, this model offers a variant explanation of the so-called coincidence problem, and it gives some motivation for the specific division of the Universe's history into epochs and eras. Let us consider these questions in more detail. \subsection{Inflation in the early Universe} When $\tau \leq 1$ Fig.1(a), 2(a),3(a),4(a),5(a),6(a), and 7(a) display the first inflationary-type epoch: one can reveal a quick growth of the DE energy density $\rho(\tau)$. During this period of the Universe evolution the scale factor $a(t)$ increases essentially: in order to visualize this inflationary effect we presented additionally the fragment of the plot $a(t)$ in the natural scale $t$ instead of the logarithmic one [see Fig.11(a)]. The steepness of the plot of $a(t)$ is predetermined by the combination of the guiding parameters; there are a lot of possibilities to fit the observed curve by the curve predicted in the framework of the Archimedean-type model. In the context of expansion the growth of the DE energy density $\rho(\tau)$ in the early Universe relates to the heating of the DE component of the dark fluid. Clearly [compare, e.g., the panels (a) and (b) in Figs.1-7] the sum $\rho{+}E$ remains non-negative for the whole interval of time. This feature provides the $H^2$ to be non-negative, which guarantees that $H$ is a real function. At the same time the DM energy density $E(\tau)$ [see the panels (b)] decreases monotonically; nevertheless, the rate of its effective cooling differs essentially from the rate given by the standard powerlaw function ($1/a^4$ for the ultrarelativistic DM or $1/a^3$ for the cold DM). Principally new detail can be found in Fig.4(d) and Fig.5(d): the DM pressure is described by a nonmonotonic function; it grows quickly, reaches a maximum, and then decreases rapidly. \begin{figure} \centerline{\includegraphics{n_fig011.eps}} \caption {{\small In the left panel a fragment of the plot is presented, which illustrates the inflationary-type growth of the scale factor $a(t)$ for the early Universe, given in terms of cosmological time $t$. In the right panel one can see a typical late-time behavior of $E(\tau)$ and $P(\tau)$; this rescaled fragment illustrates the features of the DM evolution at $\tau >1$, which cannot be recognized on the plots given by Fig.1(b,d)-Fig.7(b,d). }} \end{figure} \subsection{Late-time accelerated expansion} When $\tau \to \infty$, the curves ${-}q(\tau)$, which present the evolution of the acceleration parameter in Fig.1(f),2(f),3(f),4(f),5(f),6(f), and 7(f) tend to the horizontal asymptote, the value ${-}q(\infty)$ being positive. Taking into account the behavior of the plots of the Hubble functions displayed in Fig.1(e), Fig.3(e)-Fig.7(e), one can conclude that $H(\tau \to \infty) \to H_{\infty}$, i.e., the Hubble function also tends to the positive constant value, when $\tau \to \infty$. This means, first, that a typical behavior of the Universe with Archimedean-type interaction between DE and DM is characterized by the late-time accelerated expansion; second, that asymptotically the Archimedean-type model converts into the quasi-de Sitter one. The DM state functions $E(\tau)$ and $P(\tau)$ decrease rapidly, and for $\tau > 1$ we should change the scale in order to visualize the behavior of these functions. A typical late-time behavior of $E(\tau)$ and $P(\tau)$ is presented in the right panel of Fig.11. In order to describe the effective equation of state of dark matter in the framework of the Archimedean-type model, we calculated numerically the ratio $w(\tau)\equiv \frac{P(\tau)}{E(\tau)}$; the plot of the function $w(\tau)$ is presented in the left panel of Fig.12. Clearly, this function is constrained: $0<w(\tau)<1$; maximal value of this ratio is about $1/3$ in the first epoch of the Universe's evolution, when the DM can be considered as an effectively ultrarelativistic substrate. In the early Universe the DM state evolves quasiperiodically, i.e., the eras with effective cooling were changed by the eras of effective heating. Starting from $\tau \simeq 4$ the function $w(\tau)$ tends to zero linearly in $\tau {=} \log{\frac{a(t)}{a(t_0)}}$, i.e., the DM behaves effectively as a cold gas (dust). \begin{figure} \centerline{\includegraphics{n_fig012.eps}} \caption {{\small The left panel contains the plot of the ratio $w(\tau){=}\frac{P(\tau)}{E(\tau)}$, which presents the effective equation of state for the DM component of the dark fluid. Maximal value of this ratio is about $1/3$ in the first epoch of the Universe's evolution, when the DM can be considered as an effectively ultrarelativistic substrate. In the early Universe the DM state evolves quasiperiodically; starting from $\tau \simeq 4$ the function $w(\tau)$ tends to zero linearly in $\tau$, i.e., the DM behaves effectively as a cold gas (dust). The right panel displays the cross-points of the $\rho$ and $E$ plots. This example illustrates the so-called coincidence problem: why the energy densities of the DE and DM components of the dark fluid are of the same order today ($72\%$ and $23\%$, respectively). }} \end{figure} \subsection{Coincidence problem} Although the DE and DM components of the dark fluid evolve at different rates throughout the history of the Universe, their magnitudes are of the same order today ($72\%$ and $23\%$, respectively). This is known as the coincidence problem \cite{coi1}-\cite{coi4}. The Archimedean-type model could help us to make a step toward solving this problem. In the right panel of Fig.12 we placed the example of the plots describing the DE and DM evolution with time $\tau$. The guiding parameter $\rho^{*}_0$ defines the asymptotical ratio between $\rho$ and $E$; this ratio can in principle be chosen so that $\frac{\rho}{E}\to \frac{72}{23}$. Other guiding parameters define how many cross-points of the $\rho$ and $E$ plots can exist. For instance, in Fig.12 there are examples with one and three cross-points. Using the guiding parameters we can remove one of the cross-points away from the point $\tau \simeq 1$, but we hope to discuss in detail the fitting problem in a special work. \subsection{On the partition of the Universe's history into epochs and eras and its multi-inflationary behavior} There are a few versions of dividing the history of the Universe into self-sufficient parts that are interesting from a physical point of view. For instance, one can consider the Hubble function $H(t)$, find its zeros and extrema and separate the admissible time interval in line with them (see, e.g., \cite{Izq}. We attached such a partition to the function ${-}q(t)$, i.e., based the classification of models on the zeros and extrema of the acceleration parameter. This seems to be motivated, since we are interested in picking out the epochs of accelerated and decelerated expansion of the Universe, thus the transition points appeared as points, in which the function $-q(t)$ changes the sign. We have shown above that there are a lot of models in which the Universe's history is multistage. In particular, the Universe's evolution can be quasiperiodic, and one can use the term {\it multi-inflationary} evolution, the first inflation being the sharpest, others being more and more smoothed. Following this line, we divide the epochs into eras by using the maximums and minimums of this function: this is equivalent to the method of statefinders proposed in \cite{statefinder}. We hope to consider such fine details of the Universe's history partition in a special work. \subsection{Conclusions} The model of Archimedean-type coupling between dark energy and dark matter makes it possible to explain the principal cornerstones of the Universe's evolution: the early-time inflationary expansion, the late-time accelerated expansion, and the coincidence phenomenon. The model of Archimedean-type coupling between dark energy and dark matter possesses a wide set of guiding parameters suitable for fitting of the model predictions to the observational data; we hope to devote a special work to this important question. The model of Archimedean-type coupling offers a natural approach to the partition of the Universe's history into epochs and eras using the number of transition points and number of extrema of the function ${-}q(t)$, the acceleration parameter of the Universe's evolution. The model allows us to speak about multistage evolution and about the multi-inflationary behavior of the Universe.

1012.2433

In this (second) part of the work we present the results of numerical and qualitative analysis, based on a new model of the Archimedean-type interaction between dark matter and dark energy. The Archimedean-type force is linear in the four-gradient of the dark energy pressure and plays a role of self-regulator of the energy redistribution in a cosmic dark fluid. Because of the Archimedean-type interaction the cosmological evolution is shown to have a multistage character. Depending on the choice of the values of the model-guiding parameters, the Universe expansion is shown to be perpetually accelerated, periodic or quasiperiodic with a finite number of deceleration/acceleration epochs. We distinguished the models, which can be definitely characterized by the inflation in the early Universe, by the late-time accelerated expansion and nonsingular behavior in intermediate epochs, and classified them with respect to a number of transition points. Transition points appear, when the acceleration parameter changes the sign, providing the natural partition of the Universe’s history into epochs of accelerated and decelerated expansion. The strategy and results of numerical calculations are advocated by the qualitative analysis of the instantaneous phase portraits of the dynamic system associated with the key equation for the dark energy pressure evolution.

12135931

2010EAS....43..115P

1012

1012.0600_arXiv.txt

Astrophysical plasmas like stellar atmospheres, gaseous nebulae or accretion disks are not in any sense closed systems, as they emit photons into interstellar space. Therefore, the thermodynamic state of such plasmas is in general not described well by the equilibrium relations of statistical mechanics and thermodynamics for local values of temperature and density, i.e. by local thermodynamic equilibrium (LTE). The presence of an intense radiation field, which in character is very different from the equilibrium Planck distribution, results in deviations from LTE (non-LTE) because of strong interactions between photons and particles. The thermodynamic state is then determined by the principle of statistical equilibrium. All microscopic processes that produce transitions from one atomic state to another need to be considered in detail via the rate equations. A fundamental complication is that the distribution of the particles over all available energy states -- the level populations or occupation numbers -- in turn affect the radiation field via the effects of absorptivity and emissivity on the radiation transport. What is required is a self-consistent simultaneous solution of the radiative transfer and statistical equilibrium equations. A {\it model atom} is a collections of atomic input data required for the numerical solution of a given non-LTE problem. It is a {\it mathematical-physical approximation} to the quantum-mechanical system of a real atom, and its interaction with radiation and with other particles in a plasma. A model atom comprises, on one hand, data to specify the structure of the atom/ion like energy levels, statistical weights and ionization potentials. On the other hand, the transitions among the individual states need to be described, requiring oscillator strengths, cross-sections for photoionization and collisional excitation/ionization, etc. The number of levels in a model atom amounts typically to several tens to several hundred in modern work, and the number of transitions from hundreds to many (ten-)thousands. As only a very limited amount of atomic data have been determined experimentally up to now -- mostly energy levels, wavelengths and oscillator strengths --, most of the data have to be provided by theory. Large collaborative efforts have been made to compute the data required in astrophysical applications via {\it ab-initio} methods. The Opacity Project (OP; Seaton~\cite{seaton87}; Seaton et al.~\cite{seaton94}) and its successor the IRON Project (IP; Hummer et al.~\cite{hummer93}) provided enormous databases of transition probabilities and cross-sections for photoionization and excitation via electron impact. Many smaller groups and individuals have contributed additional data, most notably Kurucz~(see e.g.~Kurucz~\cite{kurucz06}) in a tremendous effort lasting already for about three decades. {\it Ab-initio} data for radiative processes between levels of principal quantum number $n \le 10$ are available for most of the ions of the lighter elements up to calcium, and for iron. {\it Ab-initio} data for excitation via electron collisions are by far less complete, typically covering transitions up to $n \le 3$~or~4 for selected ions of the lighter elements and for iron. Reliable data for other members of the iron group and for the heavier elements are only selectively available. The remainder of data -- still the bulk by number -- has to be approximated for practical applications. Consequently, the starting point for the construction of model atoms for many elements of astrophysical interest has improved tremendously since the mid-1980s. Nevertheless, building realistic model atoms is neither an easy nor a straightforward task. It is a common misconception that non-LTE {\it per se} brings improvements over LTE modelling. A careful LTE analysis of well-selected lines {\it can} be more reliable than a non-LTE study of the `wrong' lines with an inadequate model atom. On the other hand, computations using a realistic model atom {\it will} improve over LTE -- provided that the other ingredients of the modelling are also~realistic. The independence of microscopic processes from environment -- at least under not too extreme conditions -- provides a tool to assess the quality of model atoms by comparison with observation. Comprehensiveness and robustness of a model atom are given when it reproduces the observed line spectra over a wide range of plasma conditions. Standard stars should serve as test `laboratories', covering a wide range of effective temperature and surface gravity (particle density). In the following we discuss practical aspects of the construction of model atoms for non-LTE line-formation calculations of trace elements in stellar atmospheres. Guidelines and suggestions are given how to build up robust and comprehensive model atoms, and how to test them thoroughly.

1012.0600

Model atoms are an integral part in the solution of non-LTE problems. They comprise the atomic input data that are used to specify the statistical equilibrium equations and the opacities and emissivities of radiative transfer. A realistic implementation of the structure and the processes governing the quantum-mechanical system of an atom is decisive for the successful modelling of observed spectra. We provide guidelines and suggestions for the construction of robust and comprehensive model atoms as required in non-LTE line-formation computations for stellar atmospheres. Emphasis is given on the use of standard stars for testing model atoms under a wide range of plasma conditions.

12205189

2011hsa6.conf..339S

1012

1012.3002_arXiv.txt

} Much of our recently acquired understanding of the architecture of the Universe and its constituents derives from large surveys (e.g., 2dFGRS, SDSS, GEMS, VVDS, COSMOS to name but a few). Such surveys have not only constrained the evolution of global quantities such as the cosmic star formation rate, but also enabled us to link this with the properties of individual galaxies -- morphological types, stellar masses, metallicities, etc.. Compared to previous possibilities, the major advantages of this recent generation of surveys are: (1) the large number of objects sampled, allowing for meaningful statistical analysis to be performed on an unprecedented scale; (2) the possibility to construct large comparison/control samples for each subset of galaxies; (3) a broad coverage of galaxy subtypes and environmental conditions, allowing for the derivation of universal conclusions; and (4) the homogeneity of the data acquisition, reduction and (in some cases) analysis. \begin{figure}[t] \center \includegraphics[width=15.5cm,trim=20 252 56 228,clip=true]{sanchezsfF1.pdf} ~ \caption{\label{fig1} Postage stamp (90''$\times$90'') true-color images of a subset of galaxies within the CALIFA mother sample , extracted from the SDSS dataset, ordered following the $u-r$ vs. $r$ color-magnitude diagram. The figure spans from M$_r\sim$-23 mag from the left end, to M$_r\sim$-18 mag to the right end, and from $u-r\sim$3.5 mag from the top end, to $u-r\sim$1.5 mag to bottom end. The figure illustrates the large variety of galaxy types covered by the survey. } \end{figure} An observational technique combining the advantages of imaging and spectroscopy (albeit with usually quite small field of view) is Integral Field Spectroscopy (IFS). However, so far this technique has rarely been used in a `survey mode' to investigate large samples, with a few notable exceptions (e.g., SAURON, de Zeeuw et al. 2002). On the other hand, the large single fiber surveys mentioned above have limitations that can only be overcome by a statistical sample of nearby galaxies with spatially resolved spectroscopic information. In order to address this requirement, we proposed the CALIFA survey. This survey has been granted with 210 dark nights of the 3.5m telescope at Calar Alto Observatory (Spain), homogeneously distributed along 6 semesters, officially starting the 1st of July 2010. CALIFA will observe a well-defined sample of $\sim$600 galaxies in the local universe with the PMAS/PPAK integral field spectrophotometer (Roth et al. 2005; Kelz et al. 2006), mounted on the 3.5 m telescope at the Calar Alto Observatory (Spain). The sample to be observed was selected to comprise most galaxy types, covering the full color-magnitude diagram down to M$_B<-$18 mags. The observations will cover the optical wavelength range between 3700 and 7000\AA, using two overlapping setups, with resolutions of R$\sim$1650 and R$\sim$850. Considering this spectral coverage, and the large field-of-view of PPAK ($>$1 arcmin$^2$), CALIFA is thus the largest and the most comprehensive wide-field IFU survey of galaxies carried out to date. \begin{figure} \begin{center} \includegraphics[width=12cm]{sanchezsfF2.pdf} \caption{\label{fig:fig_gas_NII} Similar color-magnitude diagram as presented in Figure 1, with the color-maps showing the distribution of the emission line ratios between [NII]$\lambda$6583 and H$\alpha$, derived by the fitting procedure. In this case the solid-contours show the intensity of the H$\alpha$ emission. The dashed-blue-contours show 3 intensity levels of the continuum emission at $\sim$6550\AA\ (starting at 3\Funits). They have been included to indicate the physical extension of the continuum emission in the galaxies.} \end{center} \end{figure}

We have presented the CALIFA survey, the largest IFU survey currently being implemented. The results of this survey will allow us to make significant progress in many areas of galaxy evolution where large, single-fiber surveys are limited through a lack of spatial resolution and aperture biases. We will progressively report on the development of the survey in subsequent articles, and through its webpage {\sc http://www.caha.es/CALIFA}.

1012.3002

We present here the Calar Alto Legacy Integral Field spectroscopy Area survey (CALIFA). CALIFA is large collaboration comprising more than 50 astronomers of 8 different countries, which main aim is to obtain detailed spatially resolved spectroscopic information of ∼ 600 galaxies of any type at the Local Universe (0.005&lt;z&lt;0.03). The defining science drivers for the project are: (a) model the stellar population and constrain the star formation histories; (b) trace the distribution of ionized gas and estimate chemical abundances for the gas phase; and (c) measure the kinematic properties of the galaxies, both from emission and from absorption lines; all these quantities will be recovered from maps covering the entire luminous extent of the galaxies in the sample. To achieve these goals CALIFA will sample the full optical extension of the selected galaxies (to a 3 σ depth of μ ∼ 23 mag arcsec^{-2}), covering the wavelength range between 3700-7000 Å with two setups (R ∼ 850, and ∼ 1650), by using PPAK/PMAS at the 3.5 m telescope. The proposal was recently approved by the observatory, allocating ∼ 210 dark nights, distributed in 6 semesters and starting in July 2010. As a legacy survey, the fully reduced data will be delivered freely, once their quality has been verified. We report here on the results of the early analysis performed on the data taken so far (21 galaxies).

2775911

2011PhRvD..83f4038S

1012

1012.3144_arXiv.txt

Feynman has argued that no matter how beautiful or elegant a certain theory is, or how authoritative its proponents, if it does not agree with experiments, then it must be wrong. For the past 40 years, this philosophy has been applied to gravitational theories with great success. Many modified gravity theories that were prominent in the 1970's, have now been essentially discarded, as they were found to disagree with Solar System experiments or binary pulsar observations~\cite{lrr-2006-3}. Similarly, this decade is beginning to bring a wealth of astrophysical information that will be used to constrain new modified gravity theories. In fact, precision double binary pulsar observations~\cite{Burgay:2003jj,Lyne:2004cj,Kramer:2006nb} have already allowed us to constrain modified theories to exciting new levels~\cite{Yunes:2008ua,2010PhRvD..82h2002Y}. Future gravitational wave (GW) observations on Earth, with the Advanced Laser Interferometer Gravitational Observatory (aLIGO)~\cite{ligo,Abramovici:1992ah,:2007kva}, aVIRGO~\cite{virgo} and its collaborators, and in space, through the Laser Interferometer Space Antenna (LISA)~\cite{lisa,Prince:2003aa,Danzmann:2003tv,Danzmann:2003ad}, will allow new precision tests of strong field gravity~\cite{Schutz:2009tz}. Such tests of alternative theories of gravity will be very sensitive to the motion of compact bodies in a regime of spacetime where gravitational fields and velocities are large, i.e.~the so-called strong-field. Gralla~\cite{Gralla:2010cd} has shown that motion in classical field theories that satisfy certain conditions (the existence of a Bianchi-like identity and field equations no higher than second-order) is ``universally'' geodesic to leading-order in the binary system's mass-ratio, with possible deviations from geodesicity due to the bodies' internal structure. He also argues that one might be able to relax the second condition, as it does not seem necessary. In fact, motion in certain higher-order theories, such as Chern-Simons modified gravity~\cite{Alexander:2009tp}, is already known to be purely geodesic to leading-order in the mass-ratio, without influence of internal structure due to additional symmetries in the theory. Tests of modified gravity theories in the strong-field, however, not only require a prescription for the conservative sector of motion, but also of the dissipative sector, that which describes how the objects inspiral. Geodesic motion must thus be naturally corrected by a radiation-reaction force that drives non-geodesic motion toward an ultimate plunge and merger~\cite{Barack:2009ux}. Similarly, one can think of such motion as geodesic, but with {\emph{varying}} orbital elements~\cite{Pound:2007th,Gralla:2008fg,Pound:2009sm,Pound:2010pj,Pound:2010wa} (energy, angular momentum and Carter constant). The rate of change of such orbital elements is governed by the rate at which all degrees of freedom (gravitational and non-gravitational) radiate. In the gravitational sector and to leading order in the metric perturbation, such a rate of change is controlled by an effective stress-energy tensor for GWs, first computed by Isaacson in General Relativity (GR)~\cite{Isaacson:1968ra,Isaacson:1968gw}. In his approach, Isaacson expanded the Einstein equations to second order in the metric perturbation about an arbitrary background. The first-order equations describe the evolution of gravitational radiation. The second-order equation serves as a source to the zeroth-order field equations, just like a stress-energy tensor, and it depends on the square of the first-order perturbation. This tensor can then be averaged over several gravitational wavelengths, assuming the background length scale is much longer than the GW wavelength (the short-wavelength approximation). In this approximation, Isaacson found that the effective GW energy-momentum tensor is proportional to the square of first partial derivatives of the metric perturbation, i.e.~proportional to the square of the gravitational frequency. Components of this stress-energy then provide the rate of change of orbital elements, leading to the well-known quadrupole formula. Alternative theories of gravity generically lead to a modified effective GW stress-energy tensor. It is sometimes assumed that this stress-energy tensor will take the same form as in GR~\cite{Nelson:2010rt,Nelson:2010ru}, but this need not be the case. In GR, the scaling of this tensor with the GW frequency squared can be traced to the Einstein-Hilbert action's dependence on second-derivatives of the metric perturbation through the Ricci scalar. If the action is modified through the introduction of higher-powers of the curvature tensor, then the stress-energy tensor will be proportional to higher powers of the frequency. Therefore, the consistent calculation of the modified Isaacson tensor needs to be carried out until terms similar to the GR contribution (proportional to frequency squared) are obtained. This in turn implies that calculations of effective energy-momentum tensors in modified gravity theories to {\emph{leading-order in the GW frequency}} can sometimes be insufficient for determining the rate of change of orbital elements. In this paper, we present a formalism to compute the energy-momentum tensor consistently in generic classical field theories. We employ a scheme where the action itself is first expanded in the metric perturbation to second order, and the background metric and metric perturbation are treated as independent fields. Varying with respect to the background metric leads to an effective GW stress-energy tensor that can then be averaged over several wavelengths. This produces results equivalent to Isaacson's calculation. We exemplify this formulation by first considering CS gravity~\cite{Alexander:2009tp}. This theory modifies the Einstein-Hilbert action through the addition of the product of a scalar field with the contraction of the Riemann tensor and its dual. This scalar field is also given dynamics through a kinetic term in the action. The leading-order contribution to the CS-modified GW stress-energy tensor should appear at order frequency to the fourth-power, but Sopuerta and Yunes~\cite{Sopuerta:2009iy} have shown that this contribution vanishes at future null infinity. We here continue this calculation through order frequency cubed and frequency squared and find that such CS modifications still vanish at future null infinity. This is because the background scalar field must decay at a certain rate for it to have a finite amount of energy in an asymptotically-flat spacetime. If one insists on ignoring such a requirement, such as in the case of the non-dynamical theory, then frequency-cubed CS modifications to the energy-momentum do not vanish. We explicitly calculate such modifications for a canonical embedding, where the scalar field is a linear function of time in inertial coordinates. This is similar to previous work~\cite{Guarrera:2007tu} that calculated another effective stress-energy tensor for the non-dynamical version of Chern-Simons. In this case, the dominant modification to the radiation-reaction force is in the rate of change of radiated momentum, which leads to so-called recoil velocities after binary coalescence. In GR, such recoil is proportional to the product of the (mass) quadrupole and octopole when multipolarly decomposing the radiation field. In non-dynamical CS gravity with a canonical embedding, the recoil is proportional to the square of the mass quadrupole, which dominates over the GR term. We then construct a wide class of alternative theories that differ from GR through higher order curvature terms in the action coupled to a scalar field. We compute the GW stress-energy-momentum tensor in such theories and find that corrections to the Isaacson tensor vanish at future null infinity provided the following conditions are satisfied: (i) the curvature invariants in the modification are quadratic or higher order; (ii) the non-minimally coupled scalar field is dynamical; (iii) the modification may be modeled as a weak deformation away from GR; (iv) the spacetime is asymptotically flat at future null infinity. These results prove that the effective stress-energy tensor assumed in~\cite{Nelson:2010rt,Nelson:2010ru} is indeed correct\footnote{The authors of~\cite{Nelson:2010rt,Nelson:2010ru} presented an energy loss formula which was not evaluated at $\scri^+$. In the limit of $r\to\infty$, their energy loss formula reduces to the Isaacson formula.}. Even if the effective GW energy-momentum tensor is identical to that in GR, in terms of contractions of first derivatives of the metric perturbation, this does not imply that GWs will not be modified. First, background solutions could be modified. For example, in dynamical CS gravity, the Kerr metric is not a solution to the modified field equations for a rotating black hole (BH)~\cite{Grumiller:2007rv}, but it is instead modified in the shift sector~\cite{2009PhRvD..79h4043Y}. Second, the solution to the GW evolution equation could also be modified. For example, in non-dynamical CS gravity, GWs become amplitude birefringent as they propagate~\cite{jackiw:2003:cmo,Alexander:2007kv,Yunes:2010yf}. Third, additional degrees of freedom may also be present and radiate, thus changing the orbital evolution. All of these facts imply that even if the Isaacson tensor correctly describes the effective GW energy-momentum tensor, GWs themselves can and generically will be modified in such alternative theories. In the remainder of this paper we use the following conventions. Background quantities are always denoted with an overhead bar, while perturbed quantities of first-order with an overhead tilde. We employ decompositions of the type $\met_{\mu \nu} = \bar{\met}_{\mu \nu} + \epsilon \; \tilde{h}_{\mu \nu} + \scO(\epsilon^2)$, where $\met_{\mu \nu}$ is the full metric, $\bar{\met}_{\mu \nu}$ is the background metric and $\tilde{h}_{\mu \nu}$ is a small perturbation ($\epsilon \ll 1$ is a book-keeping parameter). Covariant differentiation with respect to the background metric is denoted via $\bar\nabla^{}_\mu B^{}_{\nu}$, while covariant differentiation with respect to the full metric is denoted via $\nabla^{}_\mu B^{}_{\nu}$. Symmetrization and antisymmetrization are denoted with parentheses and square brackets around the indices respectively, such as $A^{}_{(\mu\nu)}\equiv [A^{}_{\mu\nu} + A^{}_{\nu\mu}]/2$ and $A^{}_{[\mu\nu]} \equiv [A^{}_{\mu\nu} - A^{}_{\nu\mu}]/2$. We use the metric signature $(-,+,+,+)$ and geometric units, such that $G = c = 1$. This paper is organized as follows: Section~\ref{sec:pert-Lag} describes the perturbed Lagrangian approach used in this paper to compute the effective GW stress-energy tensor. Section~\ref{sec:GRTab} applies this framework to GR. Section~\ref{sec:CS} discusses dynamical CS gravity. Section~\ref{sec:CSTab} computes the full effective stress-energy tensor in this theory. Section~\ref{sec:gen-alt-theories} generalizes the calculation to a wider class of alternative theories. Section~\ref{sec:conc} concludes and points to future research.

\label{sec:conc} We have here addressed the energy content of GWs in a wide class of modified gravity theories. We focused on theories that are weak deformations away from GR and calculated the effective stress-energy tensor where GWs are extracted: in the asymptotically-flat region of spacetime. The main calculation tool we employed was the perturbed Lagrangian approach. We demonstrated the calculation explicitly for GR, recovering the Isaacson effective stress-energy tensor. We also explicitly calculated this effective tensor in dynamical modified CS gravity, where again the result at $\scri^+$ reduces to the Isaacson tensor. The features of CS gravity that lead to the effective stress-energy tensor being identical to the one in GR are the dynamical nature of the scalar field and the topological nature of the curvature correction to the action. We then generalized this finding to all action modifications of a similar nature: a dynamical scalar field coupled to a scalar curvature invariant of rank 2 or higher in a spacetime that is asymptotically flat. For scalar curvature invariants of rank 3 or higher, we showed that there is no modification to the stress-energy tensor or the equations of motion at $\scri^+$. For rank 2, we calculated the contribution to the effective stress-energy tensor and to the first-order equation of motion. In the weak coupling limit, the only solutions to the first-order equations of motion satisfy the GR first-order equations of motion at $\scri^+$, namely $\bar\square\barh_{\mu\nu} = 0$. Evaluating the effective stress-energy tensor on-shell with these solutions leads, again, to the Isaacson stress-energy tensor. A few caveats are in order. As we have stressed before, this result is evaluated at asymptotically-flat, future null infinity, so it does not apply to cosmological spacetimes, e.g.~de~Sitter spacetime. Not all of the energy that is lost by a system is carried away by GWs to $\scri^+$: there is also radiation in the scalar field (which is calculated straightforwardly from $T_{\mu\nu}^{(\vartheta)}$), and both GWs and the scalar field radiation are lost to trapped surfaces. All of these effects must be accounted for in calculating the radiation-reaction of a system. Finally, we did not address modifications to the action of the form $f(\vartheta) R$, which reduce to a classical scalar-tensor theory. There are several avenues open for future work. Considering classical scalar-tensor modifications is one possible extension. The work should also be extended to the next simplest spacetimes, those that are asymptotically de~Sitter. This is appropriate for calculating GWs from inflation, for example. Extending this approach to calculating energy lost to trapped surfaces is another possibility. The most natural application of this work is in tests of GR with pulsar binaries and with GWs emitted by EMRIs. The former problem requires performing a post-Keplerian expansion of the motion of bodies orbiting each other. The latter requires knowing the BH spacetime (background) solution in the class of modified gravity theories and the geodesic or non-geodesic motion on that spacetime. Both of these programs require knowledge of radiation-reaction in GWs at $\scri^+$, which we have here computed for a large class of modified gravity theories.

1012.3144

The inspiral of binary systems in vacuum is controlled by the stress-energy of gravitational radiation and any other propagating degrees of freedom. For gravitational waves, the dominant contribution is characterized by an effective stress-energy tensor at future null infinity. We employ perturbation theory and the short-wavelength approximation to compute this stress-energy tensor in a wide class of alternative theories. We find that this tensor is generally a modification of that first computed by Isaacson, where the corrections can dominate over the general relativistic term. In a wide class of theories, however, these corrections identically vanish at asymptotically flat, future, null infinity, reducing the stress-energy tensor to Isaacson’s. We exemplify this phenomenon by first considering dynamical Chern-Simons modified gravity, which corrects the action via a scalar field and the contraction of the Riemann tensor and its dual. We then consider a wide class of theories with dynamical scalar fields coupled to higher-order curvature invariants and show that the gravitational wave stress-energy tensor still reduces to Isaacson’s. The calculations presented in this paper are crucial to perform systematic tests of such modified gravity theories through the orbital decay of binary pulsars or through gravitational wave observations.

12137736

2010idm..confE.122B

1012

1012.4764_arXiv.txt

\label{introduction} One of the most exciting topics in physics today is the nature of Dark Matter in the Universe. Although indirect evidence for cold dark matter is well established, its true nature is not yet known. The most promising explanation is Weakly Interacting Massive Particles (WIMPs), for they would naturally lead to the observed abundance and they arise in many of the potential extensions of the Standard Model of particle physics. WIMPs could be detected directly by their collisions with nuclei in underground experiments, such a discovery would be a milestone in physics. However, since the predicted signal rates are much lower than one interaction per kg of target material and day, large detector masses and ultra-low backgrounds are necessary ingredients of any experiment aiming to discover WIMPs. Results from noble liquid detectors have recently shown that these detectors are among the most promising technology to push the sensitivity of direct WIMP searches far beyond existing limits into the regime of favored theoretical predictions. Liquid argon (LAr) and xenon (LXe), having high charge and light yields for nuclear recoils expected from WIMP-nucleus scattering, are excellent WIMP targets. A noble liquid Time Projection Chamber (TPC) can offer a scalable, large, self-shielding, homogeneous and position sensitive WIMP detector. The relative size of the charge and light signals, as well as their timing allows efficient discrimination against electron recoil events, and good spatial resolution allows the identification of the neutron background. The XENON10 \cite{xenon10} and WARP (2.3 liter) \cite{warp2.3} experiments at the Gran Sasso Underground Laboratory (LNGS) and the ZEPLIN-III experiment \cite{zeplin3} at the Boulby laboratory were successful demonstrators, reaching competitive limits on both spin-independent and spin-dependent WIMP-nucleon cross sections. The XENON100 experiment is taking science data at LNGS, and has published first results from a commissioning run in late 2009 \cite{xenon_PRL105}. The WARP-140 experiment, a 100kg-scale LAr TPCs, is under commissioning at LNGS, while the ArDM-1t \cite{ardm} detector, with a LAr WIMP target of 850 kg, is under commissioning at CERN, with the goal of underground installation in the Canfranc Underground Laboratory in 2011. While the aimed sensitivities are around 10$^{-45}$cm$^2$ for the spin-independent WIMP-nucleon cross section, a few events per year would be detected for a cross section at the level of 10$^{-44}$cm$^2$. The proposed XENON1T detector \cite{xenon1t_tdr}, with a total of 1 ton of LXe in the fiducial volume (2.4\,t total liquid xenon mass), would reach another order of magnitude in sensitivity improvement by 2015.

1012.4764

DARWIN (DARk matter WImp search with Noble liquids) is an R&amp;D and design study towards the realization of a multi-ton scale dark matter search facility in Europe, based on the liquid argon and liquid xenon time projection chamber techniques. Approved by ASPERA in late 2009, DARWIN brings together several European and US groups working on the existing ArDM, XENON and WARP experiments with the goal of providing a technical design report for the facility by early 2013. DARWIN will be designed to probe the spin-independent WIMP-nucleon cross section region below 10-47cm^2 and to provide a high-statistics measurement of WIMP interactions in case of a positive detection in the intervening years. After a brief introduction, the DARWIN goals, components, as well as its expected physics reach will be presented.

12213936

2011MNRAS.416.1629M

1012

1012.3658_arXiv.txt

In this paper, we consider a set of new statistics designed to encapsulate much of the information content of third-order and higher statistics for spinorial fields defined in three dimensions. In cosmological applications such higher-order statistics can be very noisy, and the dimensionality of the space may also lead to a very large number of data to consider. Thus some form of data compression is attractive, but preferably in a way which does not reduce the cosmological information content inherent in the original statistics. In this paper, we build on the ideas originally presented in \cite{MuHe09}, where it was shown that one statistic, the power spectrum associated with the bispectrum, could be used very effectively to estimate non-gaussianity in the microwave background radiation, as it was a lossless compression for this purpose, and also had the important added benefit of being able to provide evidence that a non-gaussianity is primordial. In this paper, we extend the ideas to cover spin-weighted fields which are defined in three dimensions, with particular emphasis on weak lensing fields convergence, shear and flexion. Weak gravitational lensing of background source galaxies is caused by fluctuations in the intervening mass distribution. It manifests itself in a number of ways, most notably as distortions in their images. This effect arises due to the fluctuations of the gravitational potential and consequent deflection of light by gravity. Despite being a relatively young subject weak gravitational lensing \citep{MuPhysRep08} has made major progress within the last decade, since the first measurements were published \citep{BRE00,Wittman00,KWL00,Waerbeke00}. There has been considerable progress in analytical modelling, technical specification and the control of systematics. By its dependence on the mass power spectrum at lower redshifts, weak lensing surveys play a complementary role to the studies based on large-scale galaxy surveys and Cosmic Microwave Background (CMB) observations. Ongoing and future weak lensing surveys such as the CFHT legacy survey{\footnote{http://www.cfht.hawaii.edu/Sciences/CFHTLS/}}, Pan-STARRS {\footnote{http://pan-starrs.ifa.hawaii.edu/}}, the Dark Energy Survey, and further in the future, the Large Synoptic Survey Telescope {\footnote{http://www.lsst.org/lsst\_home.shtml}}, WFIRST {\footnote{http://wfirst.gsfc.nasa.gov/}} and Euclid {\footnote{http://sci.esa.int/euclid}} will provide a wealth of information in terms of mapping the distribution of mass and energy in the universe. Owing to the lack of photometric redshift information the traditional approach to weak lensing has largely adopted a 2D approach, analysing correlations of the shapes of galaxy images on the sky only. However the availability of photometric redshifts allows a 3D weak lensing analysis, which was introduced by \citet{Heav03}. Later developments by various authors \citep{HRH00, HKT06, HKV07, Castro05} have shown that it can play a vital role in constraining the dark energy equation of state \citep{HKT06} and the neutrino mass \citep{Kit08}. This has lead to recent progress in modelling weak lensing observables in 3D extending results previously obtained in projection or using tomographic techniques \citep{MuHeCo_wl1_10}. Early results on analytical modelling typically assumed a small survey size and adopted a 2D approach that uses a flat-sky formalism. This is related to the fact that first generation of surveys typically covered a small portion of the sky and lacked any redshift information \citep{JSW00}. Indeed such analytical modelling was very successful in predicting lower-order statistical properties of weak lensing convergence and shear very accurately \citep{MuJai01,Mu00,MuJa00}. These results depends on analytical modelling of underlying density perturbations using perturbative and empirical methods. \citep{MuJai01,Valageas00, MuVa05,VaMuBa04,VaMuBa05}. A tomographic step was next advocated to tighten the cosmological constraints. The tomographic studies typically divides the sources into a few redshift slices \citep{Hu99,TakadaWhite03,TakadaJain04,Massey07,Schrabback09}. These slices are then analyzed essentially using a two-dimensional approach but including the correlation between different redshift slices. A notable exception to the 2D analysis was \citet{Stebbins96} who developed an all-sky formalism for weak lensing surveys. The techniques developed in \citet{Stebbins96} rely on a tensorial formalism, whereas we will be using an equivalent treatment based on spin weight spherical harmonics. Extending previous studies by \citet{Heav03} and \citet{Castro05}, \cite{Mu_wl2_10} extended the all-sky formalism to 3D to take into account the photometric redshift information, as well as extending to higher-order statistics. However they focused on the convergence field which, being a spin-$0$ field, is relatively easier to analyze. The main motivation behind this work is to extend previous results to arbitrary spinorial fields such as shear and their derivatives flexion. Weak lensing at small angular scales probes the nonlinear regime of gravitational clustering, and the extra modes there can lift degeneracies about background cosmology present in studies involving the power spectrum alone see, e.g., \cite{BerVanMell97,JainSeljak97, Hui99,Schneider98,TakadaJain03}. The nonlinear regime is characterized by gravity-induced non-Gaussianity, and detailed studies that employ the Fisher matrix formalism have already demonstrated the potential of using higher-order non-Gaussianity information to lift cosmological degeneracies. Higher-order studies are also important in evaluating the variance of lower-order statistics, e.g. a proper knowledge of the trispectrum is essential for computing the error bars in the power spectrum\citep{TakadaJain09}. The modelling of higher-order statistics typically involves either perturbative techniques or empirical modelling of the underlying matter clustering \citep{Fry84,Schaeffer84, BerSch92,SzaSza93, SzaSza97, MuBaMeSch99, MuCoMe99a, MuCoMe99b, MuMeCo99, CMM99, MuCo00, MuCo02, MuCo03}. Using such prescriptions and their extensions, studies involving non-Gaussianity, have also been performed in projection (2D) as well as using tomographic information \citep{Hu99,TakadaJain04,TakadaJain03,Semboloni08} with remarkable success. Studies involving higher-order correlation functions have been performed using observational data \citep{BerVanMell97,BerVanMell02,Pen03,JBJ}. Most of these studies involve one-point moments (cumulants) which collapse the entire correlation function into a single number. Mode-by-mode estimates of higher-order correlation functions or multispectra though far more interesting is difficult given the low signal-to-noise of current observational data. Current studies by \cite{MuHe09} defined power spectra associated with each multispectrum that uses an intermediate option in data compression. While initially this concept was applied to CMB studies, recent work by \citep{MuHeCo_wl1_10} extended this concept to weak lensing. This initial work focused on convergence $\kappa$. Being a spin-$0$ (scalar) object, the analysis of convergence statistics is relatively simple. In their analysis \cite{Mu_wl2_10} used the similar statistics for shear and flexion fields but in projection (2D). Th main motivation for the present study is to use the full 3D information (available from photometric redshift surveys) in analyzing the non-Gaussianity not only in the convergence field but also in shear and flexion. This is particularly interesting as current photometric redshift surveys with good image quality will provide a wealth of data for the analysis of weak lensing which can be used to probe cosmological information. For our study, we combine well-motivated ansatz\'e in modelling the gravitational clustering with the Limber approximation. The results that we derive here are generic and will be useful in other areas of cosmology where integration along line of sight is involved. To keep the results simpler we will ignore the fact, that in a realistic survey, the average density of sources will decline with distance, and the distance estimated from photometry will also include error, but these are evidently important ingredients in a practical implementation of these statistics. The expressions for higher-order multispectra generically include multidimensional integrals involving multiple spherical Bessel functions. We will be using Limber approximation to simplify these results. We will show that, at each order, we can reduce the dimensionality of these integrals to unity by using Limber approximation. This will simplify the numerical evaluations of such integrals considerably. This paper is arranged as follows. In \textsection2 we discuss the basic formalism of 3D weak lensing. The formalism presented here is a generalization of \citep{MuHeCo_wl1_10} and \citep{Mu_wl2_10} and can analyze higher-order statistics of spinorial fields in 3D. In \citep{MuHeCo_wl1_10} results were derived for higher-order statistics for the convergence and in \citep{Mu_wl2_10} the focus was on higher-order statistics of spinorial objects but in projection (2D). The notations for 3D harmonic decomposition, which will be used in the following sections are also introduced here. In \textsection3 we introduce the models describing higher-order clustering of underlying matter which are then used to construct models for the bispectrum and trispectrum in the nonlinear regime. The results obtained are generic and can describe higher-order statistics of weak lensing convergence, shear and flexions. In \textsection4 we focus on power spectra associated with higher-order multispectra. Results presented in this section correspond to both all-sky and patch-sky coverage. In \textsection5 we focus on error analysis and derive results for scatter (or variance) of various estimators in the presence of observational noise and mask. Finally \textsection6 is devoted to discussion of the results. Though we have mainly focused on weak-lensing, the general formalism developed in the paper will have wider applicability. We will use the Hierarchical ansatz to model clustering of underlying mass distribution, but the treatment can also be adopted in the context of more elaborate scenarios of clustering e.g. halo model.

It is now well accepted that the next generation of weak lensing surveys will play an important part in further reducing the uncertainty in fundamental cosmological parameters, including those that describe the evolution of equation of state of dark energy \cite{Euclid}. They will also be instrumental in testing various alternative gravity models (e.g. \cite{HKV07,Amendola08,Benyon09,Schrabback09, Kilbinger09}). The power of weak lensing surveys largely depends on the fact that they can exploit both the angular diameter distance and the growth of structure to constrain cosmological parameters. It is therefore very important to develop analytical techniques and statistical tools that can fully exploit the potential of future weak lensing surveys. Typically without the redshift information the data from weak lensing surveys are analyzed in projection for the entire source distribution. However it was found that by binning sources in a few photometric redshift bins the constraints improve \citep{Hu99}. In recent years a full 3D formalism which exploits the photometric redshifts of individual sources were developed. Such an approach does not involve any binning; see \citet{Heav03,Castro05,HKT06}. Further studies along these lines demonstrate that 3D lensing can provide more powerful and tighter constraints on the dark energy equation of state parameter, and on neutrino masses \citep{deBernardis09, Jimenez10}, as well as testing braneworld and other alternative gravity models. These constitute the main science drivers for the future weak lensing surveys. Initial studies in weak lensing focused on two-point correlation functions or the power spectrum mainly due to the low signal-to-noise available for higher-order studies from most first generation surveys. With the availability of modern surveys it is useful to include the non--Gaussianity information in the data analysis pipeline \citep{TakadaJain04, Semboloni09} that can help to lift some of the degeneracies in estimation of cosmological parameters involving power spectrum alone. In their recent work \cite{MuHeCo_wl1_10} have explored the possibility of extending the higher-order statistics of convergence to 3D. The main motivation of this work is to generalize those results to spinorial objects and perform a full 3D analysis for the higher-order statistics. In this sense this is also an extension of results derived in \cite{Mu_wl2_10} which analyzed higher-order statistics of spinorial fields but only in projection (2D). The results here are valid for all-sky surveys. It depends on full 3D spherical harmonic decomposition on the surface of the sky as well as along the radial directions. Such an approach in analyzing the data from future surveys which will cover a large fraction of the sky. The higher-order statistics of convergence $\kappa$, shear $\gamma_{\pm}$ or flexions $\cal F$ and $\cal G$ depend on accurate modelling of the underlying density contrast $\delta$. Various models are used, such as the hierarchical ansatz which we use here, known to be valid in the highly nonlinear regime. However the techniques developed here are generic and can also be used in association with other models such as the halo model. The higher-order multispectra contain invaluable information. Some of these information is however degenerate because of symmetries associated with higher-order correlation functions. It is difficult to estimate the higher-order multispectra mode by mode because of the associated scatter involved in such estimation especially from a noisy data set. In \cite{MuHe09}, various power spectra (skew spectrum, kurt spectra) were introduced, that are associated with a multispectra of a given order and can be estimated in the presence of mask and noise. These spectra carry some of the information contents of the multispectra from which they are constructed. In our present study we express the skew spectrum and two degenerate kurt spectra of generic spinorial fields in terms of the bi- and trispectrum. This extends earlier results for the convergence (spin-$0$) field. Extending the previously introduced 3D power spectrum $C_l(k_1,k_2)$ to higher-order, we introduce a series of power spectra related to multispectra at each order. We have introduced the {\em 3D skew spectrum} $C_l^{\sg\sgp,\sgpp}(k_1,k_2)$ associated with the bispectrum of arbitrary triplets of spinorial fields $\sg,\sg,\sgpp$. Analogously, at the level of trispectrum we have introduced two 3D kurt spectra $C_l^{\sg\sgp\sgpp,\sgppp}(k_1,k_2)$ and $C_l^{\sg\sgp,\sgpp\sgppp}(k_1,k_2)$ for arbitrary choice of spinorial fields. These extends the skew- and kurt spectra defined in \cite{MuHeCo_wl1_10} where harmonic decomposition was performed only on the surface of the celestial sphere and a real space analysis was performed on the radial direction leading to a mixed representation of skew spectrum $C_l^{(2,1)}(r_2,r_1)$ as well as their higher-order counterparts i.e. the two kurt spectra $C_l^{(2,2)}(r_2,r_1)$ and $C_l^{(3,1)}(r_2,r_1)$. The generic expression for the skew- and kurt spectra involve spherical Bessel functions. We simplified these radial integrals by using the Limber approximation, whose accuracy scales as ${\cal O}({1 \over l^4})$. We show that at each order the Limber approximation can reduce the dimensionality of the integrals to unity which dramatically reduces computational cost. Both the Limber approximation and the hierarchical ansatz are accurate at smaller scales and their joint use can help us to compute the skew- and kurt- spectra very efficiently with reasonable accuracy, but the method can accommodate different models for nonlinear clustering. We also present analytical results for dealing with a mask, via a pseudo-$C_l$ approach, encapsulated in a mode-mixing matrix. The estimation of unbiased skew- or kurt spectra are done by simple inversion of the mixing matrix $M$, which depends on the spins associated with the spinorial fields. Some regularisation will normally be required. The presence of an observational mask typically only induces mode-mixing on the celestial sphere and not on the radial direction. We have also showed how our formalism presented here can be used also for the computation of scatter under certain simplifying assumptions in the presence of an observational mask, and we have identified individual terms that correspond to contributions from noise, partial sky coverage (cosmic variance) and cross terms. The results presented here will be relevant for the study of cosmic magnification studies in 3D as well as in many other contexts where integrated radial information is used. The estimators for skew or kurt spectra that we have described here can be improved by inverse variance weighting of 3D harmonics. Finally, to summarize: \begin{itemize} \item{We have studied higher-order multispectra in the context of 3D weak lensing surveys.} \item{We use a full 3D Fourier decomposition which employ spin-weight spherical harmonics.} \item{Our generic results are valid for arbitrary 3D spinorial objects.} \item{The results are relevant for convergence $\kappa$, magnification $\mu$, shear $\gamma_{\pm}$ as well as flexions $\cal F$ and $\cal G$ or an arbitrary scalar tracer field $\Phi$.} \item{In our analysis we define power spectra $C_l(k_1,k_2)$ that are related to the bispectrum (skew spectra) and to the trispectrum (kurt spectra).} \item{We provide both all-sky exact results and corresponding approximate results using the Limber approximation.} \item{Use of Limber's approximation reduces multidimensional integrations along the radial direction to one-dimensional integrals.} \item{We show how the multi-spectra can be recovered from a masked sky in the presence of noise, and show how the presence of masks mixes modes not only on the surface of the sky but in the radial direction.} \item{The modelling was done using the hierarchical ansatz but the formalism can work with any input underlying density multispectra.} \item{Under certain simplifying approximations, we also obtain expressions for the covariance of our power spectra and skew spectra estimators.} \item{We outline how inverse variance weights can be introduced and optimal estimators can be defined for the detection of a specific type of non--Gaussianity.} \item{The formalism can be relevant in many other contexts where line-of-sight integrations of non-Gaussianities are performed or in studies involving cross-spectra or mixed-bispectra.} \end{itemize} In this paper we have ignored many observational complexities for simplicity, such as that in a realistic survey the lensing potential can only be sampled at the discrete positions of galaxies, and the average number of source galaxies will decline with redshift. We also ignore photometric redshift errors.

1012.3658

We introduce a collection of statistics appropriate for the study of spinorial quantities defined in three dimensions, focusing on applications to cosmological weak gravitational lensing studies in three dimensions. In particular, we concentrate on power spectra associated with three- and four-point statistics, which have the advantage of compressing a large number of typically very noisy modes into a convenient data set. It has been shown previously by Munshi &amp; Heavens that, for non-Gaussianity studies in the microwave background, such compression can be lossless for certain purposes, so we expect the statistics we define here to capture the bulk of the cosmological information available in these higher order statistics. We consider the effects of a sky mask and noise, and use Limber's approximation to show how, for high-frequency angular modes, confrontation of the statistics with theory can be achieved efficiently and accurately. We focus on scalar and spinorial fields including convergence, shear and flexion of three-dimensional weak lensing, but many of the results apply for general spin fields.

12099174

2009P&SS...57..816M

1012

1012.0692_arXiv.txt

The four giant planets of our solar system have hydrogen and helium envelopes which are enriched in heavy elements with respect to the solar composition. In Jupiter, for which precise measurements from the Galileo probe are available, C, N, S, Ar, Kr, Xe are all found to be enriched compared to the solar value by factors 2 to 4~\citep{1999Natur.402..269O,2004Icar..171..153W} (Assuming solar abundances based on the compilation by \citet{2003ApJ...591.1220L}). In Saturn, the C/H ratio is found to be 7.4 $\pm$ 1.7 times solar \citep{2005Sci...307.1247F}. In Uranus and Neptune it is approximately $45 \pm 20$ times solar \citep{GuillotGautier2006} (corresponding to about $30$ times solar with the old solar abundances). Interior models fitting the measured gravitational fields constrain enrichments to be between 1.5 and 8 for Jupiter and between 1.5 and 7 times the solar value for Saturn \citep{2004ApJ...609.1170S}. For Uranus and Neptune, the envelopes are not massive enough (1 to 4 Earth masses) for interior models to provide global constraints on their compositions.\\ Enriching giant planets in heavy elements is not straightforward. \citet{2000ASPC..219..475G} have shown that once the planets have their final masses, the ability of Jupiter to eject planetesimals severely limits the fraction that can be accreted by any planet in the system. The explanations put forward then generally imply an early enrichment mechanism:\\ \begin{itemize} \item \citet{2005A&A...434..343A} show that migrating protoplanets can have access to a relatively large reservoir of planetesimals and accrete them in an early phase before they have reached their final masses and started their contraction. This requires the elements to be mixed upward efficiently, which is energetically possible, and may even lead to an erosion of Jupiter's central core \citep{2004jpsm.book...35G}. \item The forming giant planets may accrete a gas that has been enriched in heavy elements through the photoevaporation of the protoplanetary disk's atmosphere, mainly made of hydrogen and helium \citep{2006MNRAS.367L..47G}. This could explain the budget in noble gases seen in Jupiter's atmosphere but is not sufficient to explain the enrichment in elements such as C, N, O because small grains are prevented from reaching the planet due to the formation of a dust-free gap \citep[eg.][]{2007A&A...462..355P}. The photoevaporation model requires that the giant planets form late in the evolution of disks, which appears consistent with modern scenarios of planet formation (see \citet{2004ApJ...604..388I}; \citet{2008ApJ...686.1292I}). It also implies that a significant amount of solids are retained in the disk up to these late stages, as plausible from simulations of disk evolution (e.g. \citet{2007ApJ...671.2091G}). \end{itemize} It has been recently suggested that the Solar System underwent a major change of structure during the phase called `Late Heavy Bombardment' (LHB) \citep{2005Natur.435..459T, 2005Natur.435..466G}. This phase, which occurred $\sim 650$~My after planet formation, was characterized by a spike in the cratering history of the terrestrial planets.\\ The model that describes these structural changes, often called the `Nice model' because it was developed in the city of Nice, reproduces most of the current orbital characteristics of both planets and small bodies. This model provides relatively tight constraints on both the location of the planets and on the position and mass of the planetesimal disk at the time of the disappearance of the proto-planetary nebula. Thus, it is interesting to study the amount of mass accreted by the planets at the time of the LHB, in the framework of this model.\\ The article is organized as follows: we first describe the orbital evolution model at the base of the calculation. Physical radii of the giant planets at the time of the late heavy bombardment are also discussed. We then present results, both in terms of a global enrichment and in the unlikely case of an imperfect mixing of the giant planets envelopes.

In this work, we evaluated the extent to which the late heavy bombardment could explain the observed enrichments of giant planets.\\ We calculated the mass accreted by each planet during this period thanks to several dynamical simulations of the LHB within the so-called "Nice" model. The accreted masses were found to be much smaller than those of the envelopes of each giant planet. In the realistic hypothesis of a global mixing in these envelopes, we found the enrichments over the solar value to be approximately two orders of magnitude smaller than the observations for Jupiter and Saturn and one order of magnitude smaller than the observations for Uranus and Neptune.\\ We then tested the possibility of an incomplete mixing in the giant planets envelopes to account for the observed enrichments. With a size distribution of planetesimals inferred from observations of trans-neptunian bodies, we found that the enrichments were always at least a factor of 2 lower than observed. Given the efficient convection expected in the deep atmosphere, we expect however the mixing to be complete. \\ Therefore we conclude that the enriched atmospheres of the giant planets do not result from the Nice model of the LHB and probably from any model describing the LHB. In fact \citet{2000ASPC..219..475G}'s calculations showed that the mass needed to explain Jupiter's and Saturn's enrichments would be certainly much too large, in any late heavy bombardment model. Earlier events should thus be invoked in the explanation of the enriched atmopspheres of giant planets. On the other hand the enrichment process during the LHB may not be completely negligeable when considering fine measurements of the compositions of giant planets \citep[eg.][]{2008ExA...tmp...34M}. When present it may also have a role in enriching the envelopes of close-in extrasolar giant planets because of their radiative structure.\\ \\ We acknowledge support from the Programme National de Plan\'etologie. We thank one of the referees, Brett Gladman, for comments that improved the article.

1012.0692

The giant planets of our solar system possess envelopes consisting mainly of hydrogen and helium but are also significantly enriched in heavier elements relatively to our Sun. In order to better constrain how these heavy elements have been delivered, we quantify the amount accreted during the so-called "late heavy bombardment", at a time when planets were fully formed and planetesimals could not sink deep into the planets. On the basis of the "Nice model", we obtain accreted masses (in terrestrial units) equal to 0.15±0.04M<SUB>⊕</SUB> for Jupiter, and 0.08±0.01M<SUB>⊕</SUB> for Saturn. For the two other giant planets, the results are found to depend mostly on whether they switched position during the instability phase. For Uranus, the accreted mass is 0.051±0.003M<SUB>⊕</SUB> with an inversion and 0.030±0.001M<SUB>⊕</SUB> without an inversion. Neptune accretes 0.048±0.015M<SUB>⊕</SUB> in models in which it is initially closer to the Sun than Uranus, and 0.066±0.006M<SUB>⊕</SUB> otherwise. With well-mixed envelopes, this corresponds to an increase in the enrichment over the solar value of 0.033±0.001 and 0.074±0.007 for Jupiter and Saturn, respectively. For the two other planets, we find the enrichments to be 2.1±1.4 (w/ inversion) or 1.2±0.7 (w/o inversion) for Uranus, and 2.0±1.2 (w/ inversion) or 2.7±1.6 (w/o inversion) for Neptune. This is clearly insufficient to explain the inferred enrichments of ∼4 for Jupiter, ∼7 for Saturn and ∼45 for Uranus and Neptune.

12143953

2010JSMTE..11..029S

1012

1012.5624_arXiv.txt

Cosmological observations are usually interpreted within a theoretical framework based on the simplest conceivable class of solutions of Einstein's law of gravitation. Namely, by using the assumptions that the universe is homogeneous and isotropic, one is able to work out, from the Einstein's field equations, the dynamics of space-time. The Friedmann-Robertson-Walker (FRW) geometry is derived under these two assumptions and it describes the geometry of the universe in terms of a single function, the scale factor, which obeys to the Friedmann equations \cite{weinberg}. In this situation, the matter density is constant in a spatial hyper-surface. On the top of the constant matter density one can consider a statistically homogeneous and isotropic small-amplitude fluctuations field. These fluctuations furnish the seeds of gravitational clustering which will eventually give rise to the structures we observe in the present universe. Evolution of fluctuations is not considered to have a sensible effect on the evolution of the space-time which is instead driven by the uniform mean field. While the simplicity of this scenario has its own appeal, in the standard model of cosmology one has to conjecture the existence of two fundamental constituents, if observational constraints are met, that both have yet unknown origin: first, a {\it dominant repulsive} component is thought to exist that can be modeled, for instance, by a positive cosmological constant. The physical nature of this component, named Dark Energy, is yet unknown and its abundance cannot be inferred from a-priori principles, whilst it is widely believed that dark energy is the biggest puzzle in standard cosmology today; e.g., the value of the cosmological constant in cosmology seems absurdly small in the context of quantum physics \cite{review_weinberg}. There is, secondly, a {\it non-baryonic} component that should considerably exceed the contribution by luminous and dark baryons and massive neutrinos. This component, named non-baryonic Dark Matter, is thought to be provided by exotic forms of matter, not yet detected in laboratory experiments. The main peculiar property of this matter component is that of weakly interacting with radiation in order to met the observational constraints given by observations \cite{einastodarkmatter}. According to the concordance model of standard cosmology, the contribution of the former converges to about 3/4 and that for the latter to about 1/4 of the total source of the standard cosmological equations (Friedmann equations), while up to a few percent has to be attributed to what is instead directly measurable by observations, namely ordinary baryonic matter, radiation and neutrinos. If the underlying cosmological model is not a perturbation of an exact flat FRW solution, the conventional data analysis and their interpretation is not necessarily valid and thus the estimations of Dark Matter and Dark Energy can be questionable. The breaking of the FRW solution can be caused, for instance, by strong inhomogeneities of large spatial extension in the matter distribution. If this were the case, the theoretical problem would then concern of whether inhomogeneous properties of the Universe can be described by the strong FRW idealization and/or in which limit this would be so. The question of whether observations of galaxy structures satisfy, on some scales, the assumptions used to derive the FRW metric is thus central. Surprisingly enough, studies of the large scale distribution of matter in the universe, as sampled by galaxy structures, seem not to be trivially compatible with such a theoretical scenario. Indeed, more and more structures on larger and larger scales have been discovered in the course of the last two decades with the advent of the three dimensional maps of the large scale universe. There structures were unexpected both because two-dimensional (angular) surveys were rather uniform and because theoretical models were unable to predict the existence of them. In many cases it was concluded that the particular three-dimensional survey under consideration had picked up a particularly ``rare'' fluctuation: this was in respect to the Gaussian distribution of fluctuations predicted by theoretical models. The statistical characterization of these structures, determining the range of correlations and the amplitudes and sizes of inhomogeneities, has thus posed a fundamental challenge to the standard picture of cosmology. The key-problem would be then to include these large fluctuations in the cosmological dynamics in a coherent way. As long as structures are limited to small sizes, and fluctuations have low amplitude, one can just treat fluctuations as small-amplitude perturbations to the leading order FRW approximation. However if structures have ``large enough'' sizes and ``high enough'' amplitudes, a perturbation approach may loose its validity and a more general treatment of inhomogeneities needs to be developed. From the theoretical point of view, it is then necessary to understand how to treat inhomogeneities in the framework of General Relativity. In this context the first issue is whether inhomogeneities can be described by the FRW idealization at least {\it on average}, by postulating that on large enough scales uniformity is eventually reached. In other words, the key-question is: does an inhomogeneous model of the Universe at relatively small scales and, uniform at large scales, evolves on average like a homogeneous solution of Einstein's law of gravitation ? Currently there is a wide discussion in the literature on this issue because, in the framework of averaged cosmological equations that has been provided by Buchert \cite{buchert}, it was found that a potential way to explain Dark Energy (and possibly also Dark Matter) can be, partially or fully, given by an effect of structure formation in an inhomogeneous cosmology. Inhomogeneities may mimic the effect of Dark Energy \cite{wiltshire}. Thus, while observations of galaxy structures have given an impulse to the search for more general solution of Einstein's equations than the Friedmann one, it is now under an intense investigation whether such a more general framework may provide a different explanation to the various effects that, within the standard FRW model, have been {\it interpreted} as Dark Energy and Dark Matter. In these proceedings we first review in Sect.\ref{galaxies} the situation with respect to galaxy structures: their observations and the analysis of their statistical properties. We then discuss in Sect.\ref{assumptions} the implications of the results on galaxy distribution with respect to the standard theoretical assumptions of the FRW model. This allows us to properly frame the problem of inhomogeneities from the point of view of theoretical modeling. We then draw, in Sect.\ref{conclusions}, our main conclusions.

\label{conclusions} We discussed several results showing that galaxy distribution in the newest galaxy samples is characterized by large fluctuations. These are manifest in the scaling properties of the conditional density which shows scaling behaviors. Particularly at small scales this statistics presents a power-law behavior with an exponent close to minus one, corresponding to a fractal dimension $D\approx 2$. The difference with the different dimension $D=1.2$ reported by authors (see e.g. \cite{dp83,davis88,park}) is due to the finite size effects which perturb the estimation of the two-point correlation function $\xi(r)$ \cite{book}. On larger scales and up to $r\approx 80$ Mpc/h a smaller correlation exponent is found to fit the data: the density depends, for $20 \le r\le 80$ Mpc/h, only weakly (logarithmically) on the system size. Correspondingly, we find that the density fluctuations follow the Gumbel distribution of extreme value statistics. This distribution is clearly distinguishable from a Gaussian distribution, which would arise for a homogeneous spatial galaxy configuration. While in the Gaussian case the rapid decay of the tails of the distribution cut large fluctuations, in the Gumbel case the large value tail decays slower. In such a situation the density field is still inhomogeneous, although not as wild as in the case in which the PDF presents power-law tails. We discussed that in several samples it is found that self-averaging properties, at large scales, are not satisfied. This is due to the presence of large scale galaxy structures which correspond to density fluctuations of large amplitude and large spatial extension, whose size is limited only by the sample boundaries. Note that the lack of self-averaging does not allow one to characterize the nature of fluctuations; this is however a clear indication that the distribution has not reached spatial uniformity. The large scale inhomogeneities detected in the three dimensional galaxy samples are at odds with the predictions of standard models. In particular according to these models the density field should present on large scales sub-Poissonian fluctuations, or a super-homogeneous nature with negative correlations \cite{glass,book}. Forthcoming redshift surveys will allow us to clarify whether on such large scales galaxy distribution is still inhomogeneous but statistically stationary, or whether the evidences for the breaking of spatial translational invariance found in the galaxy samples considered were due to selection effects in the data. Finally we discussed that the galaxy distribution is found to be compatible with the assumptions that this is transitionally invariant, i.e. it satisfies the requirement of the Copernican Principle that there are no spacial points or directions., while because of lack of spatial homogeneity galaxy distribution is not compatible with the stronger assumption of spatial homogeneity, encoded in the Cosmological Principle. \subsection*

1012.5624

The large-scale distribution of galaxies in the universe displays a complex pattern of clusters, super-clusters, filaments and voids with sizes limited only by the boundaries of the available samples. A quantitative statistical characterization of these structures shows that galaxy distribution is inhomogeneous in these samples, being characterized by large amplitude fluctuations of large spatial extension. Over a large range of scales, both the average conditional density and its variance show a non-trivial scaling behavior: at small scales, r &lt; 20 Mpc/h, the average (conditional) density scales as r<SUP> - 1</SUP>. At larger scales, the density depends only weakly (logarithmically) on the system size and density fluctuations follow the Gumbel distribution of extreme value statistics. These complex behaviors are different from what is expected in a homogeneous distribution with Gaussian fluctuations. The observed density inhomogeneities pose a fundamental challenge to the standard picture of cosmology but they also represent an important opportunity which points to new directions with respect to many cosmological puzzles. Indeed, the fact that matter distribution is not uniform, in the limited range of scales sampled by observations, raises the question of understanding how inhomogeneities affect the large-scale dynamics of the universe. We discuss several attempts which try to model inhomogeneities in cosmology, considering their effects with respect to the role and abundance of dark energy and dark matter.

12213504

2011MNRAS.413.2031M

1012

1012.0001_arXiv.txt

\begin{figure*} \resizebox{\hsize}{!}{\includegraphics[]{diagram1.eps}} \caption{Schematic diagram of the regimes of neutron star versus black hole formation in core collapse SNe at sub-solar metallicities ({\it solid line}) in the space of main sequence mass and initial proto-NS spin period $P_{0}$, taking into account the possible effects of rapid rotation and strong magnetic fields. The dotted line denotes the rotation rate above which the NS rotational energy $E_{\rm rot}$ (eq.~[\ref{eq:erot}]) exceeds the gravitational binding energy of the progenitor envelope. The dashed line denotes the rotational energy $E_{\rm rot} = 10^{52}$ ergs sufficient to power a `hypernova'. The right axis shows the magnetic field strength $B_{\rm dip}$ that would be generated if the magnetic energy in the dipole field is $\sim 0.1\%$ of $E_{\rm rot}$ (eq.~[\ref{eq:beq}]). The dot-dashed line is the minimum rotation rate required for a magnetar with a field strength $B_{\rm dip}$ to produce a classical GRB with energy $E_{\gamma} > 10^{51}$ ergs, based on the model presented in $\S\ref{sec:GRB}$.} \label{fig:diagram1} \end{figure*} Soon following the discovery of Gamma-Ray Bursts (GRBs; \citealt{Klebesadel+73}), there were possibly more theories for their origin than theorists \citep{Ruderman75}. However, once GRBs were confirmed to originate from cosmological distances (e.g.~\citealt{Metzger+97}), the joint requirements of supernova-scale energies, short (millisecond) timescales, and relativistic speeds significantly narrowed the list of plausible central engines. It is now generally accepted that GRBs result from the formation or catastrophic rearrangement of stellar-mass black holes (BHs) or neutron stars (NSs). This conclusion has only been strengthened in recent years due to the much richer picture of the prompt and afterglow emission provided by the {\it Swift} and {\it Fermi} missions. However, despite a wealth of new data, the identity of the central engine remains elusive. At least some long duration GRBs originate from the deaths of very massive stars \citep{Woosley&Bloom06}, as confirmed by their observed association with energetic core collapse supernovae (SNe) (e.g.~\citealt{Galama+98}; \citealt{Bloom+99}; \citealt{Stanek+03}; \citealt{Chornock+10}; \citealt{Starling+10}). It nevertheless remains unsettled whether the central engine is a rapidly accreting BH (\citealt{Woosley93}; \citealt{MacFadyen&Woosley99}; \citealt{Nagataki+07}; \citealt{Barkov&Komissarov08}; \citealt{Lindner+10}) or a rapidly spinning, strongly magnetized NS (a `millisecond magnetar'; \citealt{Usov92}; \citealt{Thompson94}; \citealt{Blackman&Yi98}; \citealt{Wheeler+00}; \citealt{Zhang&Meszaros01}; \citealt{Thompson+04}; \citealt{Metzger+07}; \citealt{Bucciantini+07,Bucciantini+08,Bucciantini+09}). Although much less is known about the origin of short duration GRBs, the properties of their host galaxies and their notable lack of an accompanying SN are consistent with an origin associated with the merger of NS-NS and NS-BH binaries (\citealt{Hjorth+05}; \citealt{Bloom+06}; \citealt{Berger+05}; see e.g.~\citealt{Berger10} for a recent review). However, the unexpected discovery that many short GRBs are followed by an energetic X-ray `tail' lasting $\sim 100$ seconds has challenged basic predictions of the merger model (e.g.~\citealt{Gehrels+06}; \citealt{Gal-Yam+06}; \citealt{Perley+09}) and may hint at an alternative origin for some events, such as magnetar formation via the accretion-induced collapse (AIC) of a white dwarf \citep{Metzger+08b}. The large range in length scales and the complexity of the physics involved in producing a GRB have thus far prevented all steps in the phenomena from being studied in a single work. Any attempt to construct a `first principles' model is hindered by uncertain intermediate steps relating the physics of the central engine to the properties of the relativistic jet and the gamma-ray emission mechanism. Nevertheless, in this paper we argue that the magnetar model is uniquely predictive. This allows us to construct a self-consistent model which can in principle be compared directly with observations. Although we focus on magnetars formed via the core collapse of massive stars, we also apply our results to AIC ($\S\ref{sec:AIC}$). Our primary conclusion is that a remarkable fraction of GRB properties find natural explanations within the proto-magnetar model. \subsection{Black Hole vs.~Magnetar} \label{sec:BHvsNS} In the original collapsar model, \citet{Woosley93} envisioned a `failed supernova,' in which the energy released by core collapse is insufficient to unbind the majority of the star, such that a black hole necessarily forms. If the collapsing envelope has sufficient angular momentum, it accretes through a centrifugally-supported disk. Energy released by accretion, or via the accretion-mediated extraction of the black hole's spin \citep{Blandford&Znajek77}, then powers a relativistic jet, which burrows through the star and ultimately powers the GRB at larger radii (\citealt{MacFadyen&Woosley99}; \citealt{Proga&Begelman03}; \citealt{Matzner03}; \citealt{Morsony+07}). The discovery that long GRBs are accompanied by hyper-energetic ($\sim 10^{52}$ erg) SNe propelled the collapsar model to the theoretical forefront. However, it also proved, somewhat ironically, that GRB-SNe are far from the complete `failures' envisioned by \citet{Woosley93}. Indeed, if the collapsar scenario is correct, then either (1) the BH forms promptly following stellar collapse and the explosion mechanism associated with GRB-SNe is fundamentally different than that associated with the death of normal (slower rotating) stars, which are instead powered by NS formation; or (2) a BH forms only after several seconds delay, due to the `fall-back' of material that remains gravitationally bound despite a successful and energetic SN (e.g.~\citealt{Chevalier93}; \citealt{Fryer99}; \citealt{Zhang+08}; \citealt{Moriya+10}). Modern core collapse simulations find that the shock produced at core bounce initially stalls due to neutrino and photo-dissociation losses (e.g.~\citealt{Rampp&Janka00}; \citealt{Liebendorfer+01}; \citealt{Thompson+03}). It has long been thought that neutrino heating from the proto-NS may revive the shock, resulting in a successful explosion \citep{Bethe&Wilson85}. Recent simulations suggest that the neutrino mechanism may work for low mass progenitors (e.g.~\citealt{Scheck+06}), but higher mass stars appear more difficult to explode. Although multi-dimensional effects not captured by present simulations may be a crucial missing ingredient (e.g.~\citealt{Nordhaus+10}), neutrinos alone may well prove incapable of powering $\sim 10^{52}$ erg explosions. GRB progenitors are, however, far from typical. Essentially all central engine models require rapid rotation and a strong, large-scale magnetic field ($\gtrsim 10^{15}$ G; e.g.~\citealt{McKinney06}). These ingredients may go hand-in-hand in core collapse because differential rotation provides a source of free energy to power field growth, via e.g. an $\alpha-\Omega$ dynamo in the convective proto-NS \citep{Duncan&Thompson92} or the magneto-rotational instability (MRI; e.g.~\citealt{Akiyama+03}; \citealt{Thompson+05}). The crucial question then arises: {\it Do SNe indeed fail and lead to BH formation if the progenitor core is rapidly rotating?} or stated more directly: {\it Are the requisite initial conditions for the collapsar model self-consistent?} An additional energy reservoir (rotation) and means for extracting it (magnetic fields) make magneto-rotational effects a more promising way to produce hypernovae than neutrinos alone (e.g.~\citealt{LeBlanc&Wilson70}; \citealt{Symbalisty84}; \citealt{Ardeljan+05}). Only recently, however, have simulations begun to capture the combined effects of MHD and neutrino heating (e.g.~\citealt{Burrows+07}). \citet{Dessart+08}, hereafter D08, calculate the collapse of a rotating 35$M_{\sun}$ ZAMS collapsar progenitor of \citet{Woosley&Heger06}, which they endow with a pre-collapse magnetic field that results in a $\sim 10^{15}$ G field strength when compressed to NS densities. This reproduces the field strength, if not the field topology, expected from the saturated state of the MRI. Soon after core bounce, a bipolar MHD-powered outflow develops from the proto-NS. Although the explosion is not initially successful over all solid angles, matter continues to accrete through an equatorial disk. By accreting angular momentum, the NS remains rapidly spinning, which in turn enhances the mass loss from higher latitudes due to magneto-centrifugal slinging (e.g.~\citealt{Thompson+04}; \citealt{Metzger+07}; see eq.~[\ref{eq:fcentmax}]). Importantly, in the strongly magnetized model of D08, the wind mass loss rate eventually exceeds the accretion rate, such that for $t \gtrsim 300$ ms the NS mass begins {\it decreasing}. Although D08 cannot address the possibility of later fall-back, and a different progenitor angular momentum profile could change the conclusion, their result is nonetheless suggestive: a core self-consistently endowed with the properties required to produce a GRB may not leave a BH at all. The results of D08 highlight the fact that BH versus~NS formation may not be a function of progenitor mass and metallicity alone. Delineating this dichotomy more definitively will, however, require addressing challenging theoretical issues, such as the precise mechanism responsible for amplifying the magnetic field (see \citealt{Spruit08} for a discussion). Figure \ref{fig:diagram1} is a schematic diagram of the possible effects of rapid rotation and strong magnetic fields on the regimes of NS versus BH formation as a function of main-sequence stellar mass $M_{\star}$ and the initial NS rotation period $P_{0}$. The collapse of slowly rotating, low mass stars may result in a normal SN with kinetic energy $\sim 10^{51}$ ergs powered by neutrinos. For higher mass stars, however, neutrino-powered explosions are less likely (or are accompanied by significant `fall-back' accretion) due to more massive, compact iron cores and higher envelope binding energies $E_{\rm bind}$. For these reasons it has been argued that stars with $M_{\star} \gtrsim 25M_{\sun}$ leave BH remnants at the sub-solar metallicities that appear to characterize GRB progenitors (e.g.~\citealt{Fryer99}; \citealt{Heger+03}; \citealt{OConnor&Ott10}). Above the dashed line in Figure \ref{fig:diagram1}, however, the rotational energy $E_{\rm rot}$ of the proto-NS (eq.~[\ref{eq:erot}]) exceeds the binding energy of the stellar envelope, where \begin{eqnarray} E_{\rm rot} &\simeq& (1/2)I\Omega^{2} \nonumber \\ &\approx& 3\times 10^{52}{\rm ergs}\left(\frac{M_{\rm ns}}{1.4M_{\sun}}\right)\left(\frac{R_{\rm ns}}{12{\,\rm km}}\right)^{2}\left(\frac{P}{{\rm ms}}\right)^{-2}, \label{eq:erot} \end{eqnarray} and $I = (2/5)M_{\rm ns}R_{\rm ns}^{2}$, $M_{\rm ns}$, $R_{\rm ns}$, and $\Omega = 2\pi/P$ are the NS moment of inertia, mass, radius, and rotation rate, respectively. We have defined $E_{\rm bind}$ exterior to $1.8 M_{\sun}$, as calculated by \citet{Dessart+10} from the stellar profiles of \citet{Woosley+02}. Although the efficiency with which $E_{\rm rot}$ couples to the SN shock depends on uncertain details during the first few hundred milliseconds after core bounce, if $E_{\rm rot} > E_{\rm bind}$ then a NS remnant could in principle result, even for very massive stars. The hypothetical boundary between NS and BH formation based on the above discussion is shown with a solid line in Figure \ref{fig:diagram1}. We note that there is indeed evidence that some Galactic magnetars may have stellar progenitors with masses $\gtrsim 40M_{\sun}$ (\citealt{Muno+06}), although (consistent with Fig.~\ref{fig:diagram1}) this does not exclusively appear to be the case (\citealt{Davies+09}). If an MHD-powered SN does not leave a BH, then a rapidly spinning, strongly magnetized NS (a `proto-magnetar') may instead remain behind the outgoing SN shock. The rotational energy $E_{\rm rot} \gtrsim 10^{52}$ ergs of a magnetar with $P_{0} \sim 1$ ms is more than sufficient to power most long GRBs. However, not all of this energy is available to produce high energy emission; a fraction of $E_{\rm rot}$, for instance, is expended as the jet emerges from the star or is used to power an accompanying hypernova ({\it dashed line}; Fig.~\ref{fig:diagram1}). The right axis in Figure \ref{fig:diagram1} shows the magnetic field strength $B_{\rm eq}$ that would be generated if the magnetic energy in the dipole field is $\sim 0.1\%$ of $E_{\rm rot}$ (eq.~[\ref{eq:beq}]). A dot-dashed line shows the minimum rotation rate required to produce a classical GRB from a magnetar with a field strength $B_{\rm dip}$, based on the model presented in $\S\ref{sec:GRB}$. The conditions for a hypernova and a GRB from a proto-magnetar are thus remarkably similar. \begin{figure*} \resizebox{\hsize}{!}{\includegraphics[]{sig_edot.eps}} \caption{Wind power $\dot{E}$ ({\it right axis}) and magnetization $\sigma_{0}$ ({\it left axis}; eq.~[\ref{eq:sigma}]) of the proto-magnetar wind as a function of time since core bounce, calculated for a neutron star with mass $M_{\rm ns} = 1.4M_{\sun}$, initial spin period $P_{0} = 1.5$ ms, surface dipole field strength $B_{\rm dip} = 2\times 10^{15}$ G, and magnetic obliquity $\chi = \pi/2$. Stages denoted I$.-$V.~are described in detail in $\S\ref{sec:stages}$.} \label{fig:edotsig} \end{figure*} \subsection{Summary of the Magnetar Model and This Paper} In this section we summarize the organization of the paper and orient the reader with a brief description of the model timeline (more details and references are provided in subsequent sections). In $\S\ref{sec:windevo}$ we present calculations of the time-dependent properties of proto-magnetar winds and quantify the stages of the proto-magnetar model. The basic picture is summarized by Figure \ref{fig:edotsig}, which shows the wind power $\dot{E}$ and magnetization $\sigma_{0}$ (maximum Lorentz factor) as a function of time following core bounce, calculated for a proto-magnetar with a surface dipole magnetic field strength $B_{\rm dip} = 2\times 10^{15}$ G, initial spin period $P_{0} = 1.5$ ms, and magnetic obliquity $\chi = \pi/2$. Changes in the wind properties with time are driven largely by the increase in $\sigma_{0}(t)$ as the proto-NS cools. Within the first few hundred milliseconds following core bounce, a successful SN shock is launched by neutrino heating or MHD forces (Stage I). Soon after, a wind heated by neutrinos expands freely from the NS surface into the cavity evacuated by the outgoing shock. The wind is initially non-relativistic ($\sigma_{0} \lesssim 1$) because the neutrino-driven mass loss rate is high (Stage II). However, as the proto-NS cools, $\sigma_{0}$ increases to $\gtrsim$ 1 and the wind becomes relativistic (Stage III). The wind is collimated by its interaction with the star into a bipolar jet, which breaches the stellar surface after $\sim 10$ seconds. After jet break-out, the relativistic magnetar wind is directed through a relatively clear channel out of the star and the GRB commences (Stage IV; $\S\ref{sec:GRB}$). Averaging over variability imposed by e.g.~interaction with the jet walls ($\S\ref{sec:variability}$), the time evolution of the power and mass-loading of the jet match those set by the magnetar wind at much smaller radii. In $\S\ref{sec:stages}$ we provide a more quantitative description of the individual model stages described above using an extensive parameter study of wind models. Although the site and mechanism of prompt GRB emission remain uncertain, in $\S\ref{sec:GRB}$ we calculate the light curves and spectra within two emission models. Depending on the means and efficacy of the jet's acceleration ($\S\ref{sec:acceleration}$), GRB emission may be powered by the dissipation of the jet's Poynting flux directly (`magnetic dissipation'; $\S\ref{sec:magdiss}$) near or above the photosphere; and/or via `internal shocks' within the jet at larger radii\footnote{In this paper we define `internal shocks' as those resulting from the interaction between the magnetar jet and the accumulated (slower) shell of material released at earlier times. This is in contrast to the standard internal shock model (e.g.~\citealt{Rees&Meszaros94}), which invokes the singular interaction between shells with similar properties released immediately after one another. As we discuss in $\S\ref{sec:internalshocks}$, the former dominate the latter in the magnetar model because the mean Lorentz factor of the jet increases monotonically in time.} ($\S\ref{sec:internalshocks}$). As Figure \ref{fig:edotsig} makes clear, self-interaction in the jet is inevitable because $\sigma_{0}-$and hence the jet speed$-$increase monotonically as the proto-NS cools. After $t\sim 30-100$ seconds, $\sigma_{0}$ increases even more rapidly as the proto-NS becomes transparent to neutrino emission. Because magnetic dissipation and jet acceleration become ineffective when $\sigma_{0}$ is very large, this abrupt transition likely ends the prompt GRB. In $\S\ref{sec:highsig}$ we address the possibility that residual rotational or magnetic energy may continue to power late time flaring or afterglow emission, such as the X-ray plateau. In $\S\ref{sec:discussion}$ we discuss the implications of our results for the diversity of GRB-related phenomena, including very luminous GRBs ($\S\ref{sec:VLGRBs}$), low luminosity GRBs ($\S\ref{sec:LLGRBs}$), thermal-rich GRBs/X-ray Flashes ($\S\ref{sec:XRF}$), Galactic magnetars ($\S\ref{sec:galactic}$), very luminous supernova ($\S\ref{sec:choked}$), and magnetar formation via AIC ($\S\ref{sec:AIC}$). We summarize our conclusions in $\S\ref{sec:conclusions}$. \begin{figure} \resizebox{\hsize}{!}{\includegraphics[]{ns.eps}} \caption{Geometry of magnetized proto-neutron star winds. The neutron star radius $R_{\rm ns}$ is initially large ($\gtrsim 20$ km) following the launch of the supernova shock, but decreases to its final value $R_{\rm ns} \approx 12$ km in a few seconds (Fig.~\ref{fig:pons}). The neutron star rotates at an angular velocity $\Omega = 2\pi/P$ about the vertical axis, where $P$ is the rotational period; the light cylinder radius is $R_{\rm L} = c/\Omega \simeq 50(P/{\rm ms})$ km. The magnetic dipole moment $|\mu| = B_{\rm dip}R_{\rm ns}^{3}$ makes an angle $\chi$ with respect to the rotation axis. The angle $\theta_{\rm open}$ defines the size of the open magnetosphere on the neutron star surface. The magnetosphere is closed at angles $\theta > \theta_{\rm open}/2$ from the magnetic pole, while field lines with $\theta < \theta_{\rm open}/2$ form an `open' or `wind' zone along which matter may escape to infinity. The size of the open zone affects both the spin-down rate and the mass loss rate from magnetized proto-neutron star winds. The bundle of last closed field lines intersects the magnetic equator at the `Y' point radius $R_{\rm Y}$. Ultra-relativistic, force-free winds ($\sigma_{0} \gg 1$) have $R_{\rm Y} \sim R_{\rm L}$, while less magnetized winds in general have $R_{\rm Y} < R_{\rm L}$ (see $\S\ref{sec:mdot}$ and Fig.~\ref{fig:radii}). } \label{fig:nscartoon} \end{figure}

\label{sec:conclusions} In this paper we take the first steps towards developing the millisecond proto-magnetar model into a quantitative theory for gamma-ray bursts. Using detailed evolutionary models of magnetar spin-down, we explore a wide range of magnetar properties and calculate the prompt emission predicted by magnetic dissipation and internal shock models. Although the picture we construct may not be accurate in all details, it serves as a `proof of principle' that the basic concepts can be constructed into a self-consistent model. Our work also provides a baseline for future improvements, as will be necessitated in particular by advances in our understanding of the origin of prompt GRB emission. Several theoretical uncertainties remain that should be addressed with future work. These include a more detailed understanding of the effects of rotation and convection on the cooling evolution of the proto-neutron star, and the effects of strong magnetic fields on the neutrino-driven mass loss rate. Although most of our results are at least qualitatively robust to these uncertainties, predictions for the GRB duration (and how it correlates with other observables) is in particular sensitive to the time of neutrino transparency. The mass loss rate from the proto-NS (and, hence, the wind magnetization) also depends on fraction of the magnetosphere open to outflows, which depends on the poorly-understood sources of dissipation near the Y-point. Future studies would also be aided by a more detailed understanding of the dependence of the jet properties (e.g. break-out time and opening angle) on the properties of the proto-magnetar and the stellar progenitor. The source of the rapid rotation and strong magnetic fields required to produce millisecond magnetars also remains a major uncertainty. However, we note that black hole models place similar, if not more extreme, constraints on the progenitor rotation and the large-scale magnetic field of the central engine (e.g.~\citealt{McKinney06}). Our primary conclusion is that a surprisingly large fraction of GRB properties can be explained by the magnetar model. These include: \begin{itemize} \item{{\bf Energy.} Magnetars with properties in the `classical GRB' regime in Figure~\ref{fig:regimes} radiate $E_{\gamma} \sim 10^{50}-10^{52}$ ergs during the GRB phase, consistent with the beaming-corrected gamma-ray energies inferred from afterglow modeling. Magnetars with stronger(weaker) magnetic fields and/or shorter(longer) initial periods may produce very luminous(low luminosity) GRBs.} \item{{\bf Lorentz Factor.} Magnetars in the `classical GRB' regime produce jets with average and instantaneous magnetizations $\sigma_{0} \gtrsim 10^{2}-10^{3}$ (Fig.~\ref{fig:meansig}) which are remarkably similar to the typical Lorentz factors inferred from GRB observations (cf.~Fig.~\ref{fig:meangs}). The baryon loading of the jet is not fine-tuned or put in by hand, but instead results naturally from the physics of neutrino heating above the proto-magnetar surface. This is contrast to black hole models, for which current predictions for $\Gamma$ depend on the uncertain rate at which baryons diffuse into an otherwise clean jet (e.g.~\citealt{Levinson&Eichler03}; \citealt{McKinney05}). The magnetar model predicts that $\sigma_{0}$ (and probably $\Gamma$) increases monotonically with time during the burst. Among other things, this implies that any thermal emission present will be strongest at early times and will decrease in relative strength as the outflow becomes cleaner with time (Fig.~\ref{fig:magdis2}). } \item{{\bf Duration.} The GRB begins once the jet breaks out of the star and becomes optically thin at the internal shock or magnetic dissipation radius. The GRB ends once the jet magnetization increases sufficiently that jet acceleration and dissipation become ineffective. Because the latter generally occurs when the NS becomes transparent to neutrinos at $t = t_{\rm\nu-thin} \sim 10-100$ s (Fig.~\ref{fig:pons}), the magnetar model naturally explains the typical durations of long GRBs.} \item{{\bf Steep Decay Phase.} The abrupt onset of the high-$\sigma_{0}$ transition at $t \approx t_{\rm \nu-thin}$ (Fig.~\ref{fig:edotsig}) explains why GRB prompt emission decreases rapidly after the prompt emission ends (the `steep decay' phase; e.g.~\citealt{Tagliaferri+05}).} \item{{\bf Association with Hypernova.} It is natural to associate energetic, MHD-powered supernovae with magnetar birth. If the magnetar model is correct, all long GRBs formed from the core collapse of massive stars should be accompanied by an energetic (and possibly hyper-energetic) supernova. Magnetars formed via AIC, by contrast, may produce long GRBs not accompanied by a bright SN. This is a promising explanation for the $\sim 100$ second X-ray tails observed following some short GRBs ($\S\ref{sec:AIC}$) and explains why they resemble long GRBs in many of their properties.} \item{{\bf High Lorentz Factors$\leftrightarrow$Energetic Bursts.} The magnetar model predicts a positive correlation (with significant scatter) between the (energy-weighted) average magnetization $\sigma_{\rm avg}$ of the jet and the (beaming-corrected) GRB luminosity/energy (Fig.~\ref{fig:meansiglum}). This is consistent with the fact that energetic {\it Fermi} bursts appear to have the largest Lorentz factors.} \item{{\bf High Radiative Efficiency.} Both magnetic dissipation and internal shocks may occur in proto-magnetar winds, resulting in the prompt high-energy emission. Both models predict maximum radiative efficiencies $\epsilon_{\rm r} \sim 30-50\%$, consistent with the high values of $\epsilon_{\rm r}$ inferred from afterglow modeling (e.g.~\citealt{Panaitescu&Kumar01}; \citealt{Zhang+07}; \citealt{Fan&Piran06}).} \item{{\bf Amati-Yonetoku Relation.} Our spectral modelling favors magnetic dissipation over internal shocks as the prompt emission mechanism, in part because magnetic dissipation predicts a relatively constant spectral energy peak $E_{\rm peak}$ as a function of time (Fig.~\ref{fig:epeak}). Strong internal shocks may be suppressed by the residual magnetization of the ejecta or if the toroidal field geometry is not conducive to particle acceleration (e.g.~\citealt{Sironi&Spitkovsky10}). In combination with the predicted $\sigma_{\rm avg}-L_{\gamma}$ correlation (Fig.~\ref{fig:meansiglum}), the magnetic dissipation model reproduces both the slope and normalization of the observed Amati-Yonetoku correlations.} \item{{\bf Late-Time Emission.} Although we expect that prompt internal emission becomes ineffective when $\sigma_{0}$ becomes very large at late times, the plateau X-ray afterglow phase may also be powered by magnetar spin-down, as proposed by previous authors and suggested by Figure \ref{fig:edotsig}. The predicted correlation between the plateau luminosity and duration (Fig.~\ref{fig:tauedot}) is consistent with the sample of `internal' plateaus studied by \citet{Lyons+10}. Late-time X-ray flaring may be powered by residual rotational or magnetic energy. } \end{itemize} \begin{table*} \begin{center} \vspace{0.05 in}\caption{Properties of Proto-Magnetar Winds} \label{table:winds} \resizebox{18cm}{!}{ \begin{tabular}{cccccccccccccc} \hline \hline \multicolumn{1}{c}{$B_{\rm dip}^{(a)}$} & \multicolumn{1}{c}{$P_{0}$[P$|_{t=0}$]$^{(b)}$} & \multicolumn{1}{c}{$M_{\rm ns}$} & \multicolumn{1}{c}{$\chi^{(c)}$} & \multicolumn{1}{c}{$\eta_{\rm s}^{(d)}$} & \multicolumn{1}{c}{$\sigma_{\rm bo}^{(e)}$} & \multicolumn{1}{c}{$\sigma_{\rm avg}^{(f)}$} & \multicolumn{1}{c}{$\Gamma_{\rm s,avg}^{(g)}$} & \multicolumn{1}{c}{$T_{90}^{(h)}$} & \multicolumn{1}{c}{$t_{\rm end}^{(i)}$} & \multicolumn{1}{c}{$E_{\rm th}^{(j)}$} & \multicolumn{1}{c}{$E_{\gamma}^{(k)}$} & \multicolumn{1}{c}{$\dot{E}|t_{\rm end}^{(l)}$} & \multicolumn{1}{c}{$\tau_{\rm s}|_{t_{\rm end}}^{(m)}$} \\ (G) & (ms) & ($M_{\sun}$) & (rad) & - & - & - & - & (s) & (s) & ($10^{50}$ ergs) & ($10^{50}$ ergs) & ($10^{50}$ erg s$^{-1}$) & (s) \\ \hline \\ $2\times 10^{15}$ & 1.5[4.3] & 1.4 & $\pi/2$ & 3 & 22 & 570 & 68 & 47 & 60 & 1.6 & 24 & 0.25 & 270 \\ $2\times 10^{15}$ & 1.5[4.3] & 1.2 & $\pi/2$ & 3 & 25 & 500 & 74 & 38 & 51 & 1.3 & 19 & 0.23 & 240 \\ $2\times 10^{15}$ & 1.5[4.3] & 2.0 & $\pi/2$ & 3 & 14 & 760 & 51 & 80 & 96 & 1.8 & 35 & 0.24 & 400 \\ $2\times 10^{15}$ & 1.5[4.3] & 1.4 & 0 & 3 & 40 & 1200 & 140 & 46 & 56 & 0 & 14 & 0.19 & 430 \\ $2\times 10^{15}$ & 1.5[4.3] & 1.2 & 0 & 3 & 54 & 1300 & 180 & 37 & 47 & 0 & 11 & 0.18 & 380 \\ $2\times 10^{15}$ & 1.5[4.3] & 2.0 & 0 & 3 & 17 & 1100 & 70 & 72 & 85 & 0.7 & 19 & 0.17 & 660 \\ $2\times 10^{15}$ & 1.5[4.3] & 1.4 & $\pi/2$ & 1 & 37 & 360 & 93 & 18 & 29 & 0.7 & 12 & 0.28 & 240 \\ $2\times 10^{15}$ & 1.5[4.3] & 1.2 & $\pi/2$ & 1 & 46 & 340 & 110 & 13 & 23 & 0 & 7.7 & 0.27 & 210 \\ $2\times 10^{15}$ & 1.5[4.3] & 2.0 & $\pi/2$ & 1 & 21 & 450 & 60 & 41 & 54 & 1.9 & 22 & 0.25 & 390 \\ $2\times 10^{15}$ & 1.5[4.3] & 1.4 & 0 & 1 & 70 & 750 & 200 & 18 & 28 & 0 & 6.5 & 0.20 & 420 \\ $2\times 10^{15}$ & 1.5[4.3] & 1.2 & 0 & 1 & 100 & 830 & 280 & 12 & 22 & 0 & 4.0 & 0.19 & 370 \\ $2\times 10^{15}$ & 1.5[4.3] & 2.0 & 0 & 1 & 28 & 670 & 87 & 40 & 52 & 0.69 & 14 & 0.17 & 660 \\ $ 10^{16}$ & 1[2.8] & 1.4 & $\pi/2$ & 3 & 340 & 5100 & 740 & 54 & 64 & 0 & 61 & 0.30 & 50 \\ $ 10^{16}$ & 1[2.8] & 1.4 & 0 & 3 & 890 & $2.0\times 10^{4}$ & 2200 & 61 & 71 & 0 & 110 & 0.50 & 53 \\ $ 10^{16}$ & 1[2.8] & 1.4 & $\pi/2$ & 1 & 490 & 4400 & 1100 & 19 & 29 & 0 & 22 & 0.60 & 32 \\ $ 10^{16}$ & 1[2.8] & 1.4 & 0 & 1 & 1200 & $1.4\times 10^{4}$& 3000 & 19 & 29 & 0 & 43 & 1.40 & 33 \\ $5\times 10^{14}$ & 2[5.7] & 1.4 & $\pi/2$ & 3 & 1.9 & 55 & 6.4 & 30 & 60 & 0.15 & 0.86 & 0.020 & 2600 \\ $5\times 10^{14}$ & 2[5.7] & 1.4 & 0 & 3 & 2.2 & 62 & 7.5 & 30 & 55 & 0.063 & 0.43 & 0.010 & 5100 \\ $5\times 10^{14}$ & 2[5.7] & 1.4 & $\pi/2$ & 1 & 3.3 & 31 & 7.8 & 7.6 & 28 & 0.083 & 0.24 & 0.023 & 2300 \\ $5\times 10^{14}$ & 2[5.7] & 1.4 & $0$ & 1 & 4.0 & 38 & 9.7 & 9.2 & 27 & 0.040 & 0.15 & 0.011 & 4700 \\ $^{\dagger}2\times 10^{15}$ & 1.5[4.3] & 1.28 & $\pi/2$ & 3 & 36 & 350 & 91 & 14 & 25 & 0.94 & 12 & 0.39 & 190 \\ $^{\dagger}2\times 10^{15}$ & 1.5[4.3] & 1.28 & 0 & 3 & 81 & 910 & 230 & 14 & 24 & 0 & 5.9 & 0.26 & 330 \\ $^{\dagger} 10^{16}$ & 1[2.8] & 1.28 & $\pi/2$ & 3 & 590 & 4300 & 1300 & 13 & 23 & 0 & 32 & 1.3 & 21 \\ $^{\dagger} 10^{16}$ & 1[2.8] & 1.28 & 0 & 3 & 1800 & $1.9\times 10^{4}$ & 4400 & 14 & 24 & 0 & 61 & 2.7 & 21 \\ $^{\dagger}5\times 10^{14}$ & 2[5.7] & 1.28 & $\pi/2$ & 3 & 3.4 & 34 & 8.1 & 5.7 & 25 & 0.072 & 0.21 & 0.027 & 1900 \\ $^{\dagger}5\times 10^{14}$ & 2[5.7] & 1.28 & 0 & 3 & 4.6 & 47 & 11 & 7.4 & 24 & 0.036 & 0.14 & 0.012 & 4200 \\ \hline \end{tabular} } \end{center} {\small $^{(a)}$Surface dipole magnetic field strength following NS contraction. $^{(b)}$Spin period if NS were to contract to final radius with fixed angular momentum [Actual initial spin period at t=0]. $^{(c)}$Magnetic obliquity (see Fig.~\ref{fig:nscartoon}). $^{(d)}$`Stretch' correction applied to neutrino luminosities and energies to account for the effects of rotation (see eq.~[\ref{eq:stretch}]). $^{(e)}$Wind magnetization when the jet breaks out of the stellar surface at $t = t_{\rm bo} = 10$ s (eq.~[\ref{eq:tbo}]; Fig.~\ref{fig:sigbo}). $^{(f)}$Energy-weighted average magnetization between jet break-out and the end of the prompt emission, $t_{\rm bo} \lesssim t \lesssim t_{\rm end}$ (Fig.~\ref{fig:meansig}). $^{(f)}$Energy-weighted average Lorentz factor of the bulk shell produced by internal shocks (Fig.~[\ref{fig:meangs}]). $^{(h)}$Duration of the prompt GRB emission $T_{90} \equiv t_{\rm end} - t_{\rm thin,is}$ (Fig.~\ref{fig:T90}). $^{(i)}$Time after core bounce when the prompt GRB emission ends, defined as the point when the `saturation radius' $r_{\rm mag}$ (eq.~[\ref{eq:rsat}]) exceeds the internal shock radius $r_{\rm sh}$ (eq.~[\ref{eq:rsh}]). This transition generally occurs simultaneous with the transition of the proto-NS to neutrino transparency (see Fig.~\ref{fig:edotsig}). $^{(j)}$Maximum `thermal' energy produced by the jet (eq.~[\ref{eq:etherm}]). $^{(k)}$Maximum GRB energy, defined as the rotational energy released in the time interval $min[t_{\rm bo},t_{\rm thin,is}] < t < t_{\rm end}$ (see Fig.~\ref{fig:eGRB}). $^{(l)}$Wind power at $t \simeq t_{\rm end}$. $^{(m)}$Dipole spin-down timescale at $t = t_{\rm end}$. $^{\dagger}$Calculated using the neutrino cooling calculations of \citet{Hudepohl+10}. \\ } \end{table*}

1012.0001

Long duration gamma-ray bursts (GRBs) originate from the core collapse of massive stars, but the identity of the central engine remains elusive. Previous work has shown that rapidly spinning, strongly magnetized protoneutron stars ('millisecond protomagnetars') produce outflows with energies, time-scales and magnetizations σ<SUB>0</SUB> (maximum Lorentz factor) that are consistent with those required to produce long duration GRBs. Here we extend this work in order to construct a self-consistent model that directly connects the properties of the central engine to the observed prompt emission. Just after the launch of the supernova shock, a wind heated by neutrinos is driven from the protomagnetar. The outflow is collimated into a bipolar jet by its interaction with the progenitor star. As the magnetar cools, the wind becomes ultrarelativistic and Poynting flux dominated (σ<SUB>0</SUB>≫ 1) on a time-scale comparable to that required for the jet to clear a cavity through the star. Although the site and mechanism of the prompt emission are debated, we calculate the emission predicted by two models: magnetic dissipation and shocks. <P />Magnetic reconnection may occur near the photosphere if the outflow develops an alternating field structure due to e.g. magnetic instabilities or a misalignment between the magnetic and rotation axes. Shocks may occur at larger radii because the Lorentz factor of the wind increases with time, such that the faster jet at late times collides with slower material released earlier. Our results favour magnetic dissipation as the prompt emission mechanism, in part because it predicts a relatively constant 'Band' spectral peak energy E<SUB>peak</SUB> with time during the GRB. The baryon loading of the jet decreases abruptly when the neutron star becomes transparent to neutrinos at ? s. Jets with ultrahigh magnetization cannot effectively accelerate and dissipate their energy, which suggests this transition ends the prompt emission. This correspondence may explain both the typical durations of long GRBs and the steep decay phase that follows. Residual rotational or magnetic energy may continue to power late time flaring or afterglow emission, such as the X-ray plateau. We quantify the emission predicted from protomagnetars with a wide range of physical properties (initial rotation period, surface dipole field strength and magnetic obliquity) and assess a variety of phenomena potentially related to magnetar birth, including low-luminosity GRBs, very luminous GRBs, thermal-rich GRBs/X-ray flashes, very luminous supernovae and short-duration GRBs with extended emission.

12213959

2011MNRAS.416.1703E

1012

1012.4833_arXiv.txt

Redshift surveys have shown that the clustering properties of galaxies strongly depend on their luminosity, color and morphological (or spectral) type (e.g. \citealt{b12}; \citealt{b13}). This indicates that galaxies do not perfectly trace the distribution of the underlying dark matter, a phenomenon commonly referred to as `galaxy biasing'. Its origin lies in the details of the galaxy formation process which is shaped by the interplay between complex hydrodynamical and radiative processes together with the dark-matter driven formation of the large-scale structure. Attempts to infer cosmological parameters from galaxy clustering studies are severely hampered by galaxy biasing. A number of theoretical arguments and the outcome of numerical simulations both suggest that, on sufficiently large scales, the power spectra of galaxies and matter should be proportional to each other: $P_{\rm g}=b_1^2\, P_{\rm m}$ where the linear bias factor $b_1$ depends on galaxy type but is generally scale independent (e.g. \citealt{b14}; \citealt{b15}). Similarly, to model higher-order statistics, such as the galaxy bispectrum, it is generally assumed that galaxy biasing is a local process such that $\delta_g=b_1\delta_m+b_2 \delta_m^2/2+\dots$ where $\delta_g$ and $\delta_m$ are the (smoothed) galaxy and dark-matter density contrast, respectively \citep{b16}. However, the reliance of these phenomenological approximations limits cosmological studies to very large scales whereas data with better signal-to-noise ratio are already available on much smaller scales. Moreover, future studies of baryonic acoustic oscillations (e.g. \citealt{b17}; \citealt{b18}) will require measurements of the matter power-spectrum with percent or even sub-percent accuracy in order to shed new light on the source of cosmic acceleration. Understanding and controlling the effects of galaxy biasing with this precision will be challenging. All this provides a very strong motivation for developing more accurate (and physically driven) models of galaxy biasing. A number of authors have used the power spectrum statistics to explore the scale dependence of galaxy biasing based on numerical simulations (\citealt{b19}; \citealt{b20}; \citealt{b21}; \citealt{b22}; \citealt{b23}; \citealt{b24}) or analytical calculations (\citealt{b25}; \citealt{b26}; \citealt{b27}; \citealt{b8}) stemming from either perturbation theory or the halo model for the large-scale structure (see \citealt{b28} for a review). The general picture is that galaxy biasing is expected to be scale dependent (i.e. $P_{\rm g}(k)=b(k)^2\, P_{\rm m}(k)$) and the functional form of $b(k)$ can sensibly depend on the selected tracer of the large-scale structure. Since galaxies are expected to form within dark-matter haloes, understanding the clustering properties of the haloes is a key step to accurately model galaxy biasing. This is a much simpler problem, considering that dark-matter haloes form under the sole action of gravity. It is in fact expected that long-wavelength density fluctuations modulate halo formation by modifying the collapse time of localized short-wavelength density peaks (\citealt{b29}; \citealt{b30}). This argument (known as the peak-background split) predicts that, on large scales, the halo overdensity $\delta_{\rm h}=b\, \delta_{\rm m}$ where the bias coefficient $b$ varies with the halo mass \citep{b31}. The numerical value of the bias coefficient is determined by two different occurrences: first, haloes form out of highly biased regions in the linear density field (\citealt{b32}; \citealt{b33}) and, second, they move over time as they are accelerated towards the densest regions of the large-scale structure by gravity \citep{b31}. These two phenomena generally go under the name of ``Lagrangian biasing'' and ``Lagrangian to Eulerian passage'', respectively. \citet{b31} dealt with the second problem by assuming that long-wavelength density perturbations evolve according to the spherical top-hat model. A more sophisticated generalization of the peak-background split has been presented by \citet{b34} who assumed that also the large-scale motion of the density ``peaks'' is fully determined by the long-wavelength component of the density field. Since the halo population and the matter feel the same large-scale gravitational potential, their density fluctuations are strongly coupled and their time evolution must be solved simultaneously. This makes the process of halo biasing non-linear and non-local even if one starts from a linear and local Lagrangian biasing scheme (\citealt{b34}; \citealt{b35}). The bispectrum can be used to test this model against the standard Eulerian local biasing scheme \citep{b36}. In this paper, we present a novel and very promising approach to model the clustering of dark matter haloes. Adopting the formalism by \citet{b34} combined with a non-local Lagrangian biasing scheme for the haloes \citep{b4}, we simultaneously follow the growth of perturbations in the matter and in the halo distribution over cosmic time. We present perturbative solutions for the corresponding overdensity and velocity fields and we are able to resum the perturbative series in the limit of large wavenumbers. Moreover, we write down a system of equations for the power spectra $P_{\rm m}$ and $P_{\rm h}$ using the time-renormalization-group (TRG) approach by \citet{b1} and numerically integrate them. Our results are in excellent agreement with the output of a high-resolution N-body simulation, showing an improvement over linear theory, and we are able to predict the matter-halo cross spectrum with a precision within $5$ per cent for $k< 0.15\ h$ Mpc$^{-1}$. Related work has been very recently presented by \citet{b42} who computed the two-point correlation function of linear density peaks and followed its time evolution assuming that peaks move according to the Zel'dovich approximation. For massive haloes this results in a scale-dependent bias (with variations of $\sim 5$ per cent) on the scales relevant for baryonic-oscillation studies. Contrary to their approach, we do not deal with a point process but describe large-scale fluctuations in the distribution of dark-matter haloes as perturbations in a continuous fluid. On the other hand, we account for the full gravitational motion of the haloes and do not rely on simplified dynamical models as like as the Zel'dovich approximation. The structure of our paper is as follows. In Section \ref{mod} we present our model for the joint evolution of the matter and halo power spectra. The initial conditions for our evolutionary equations are discussed in Section \ref{init}. The solution of the linearized equations is presented in Section \ref{lin} where we also quantify the importance of the halo velocity bias. Using a perturbative technique, in Section \ref{analit} we compute analytic solutions for the propagator of perturbations (the two-time correlator). We derive 1-loop corrections and, in the limit of large wavenumbers, the fully resummed propagator. The discussion in Sections \ref{comp1} and \ref{comp2} is very technical and the less experienced readers can safely skip it without compromising understanding of the remainder. In Section \ref{num}, we numerically integrate the full equations for the evolution of halo and matter power spectra in the TRG formalism. We then compare the results against the outcome of a high-resolution N-body simulation. Finally, in Section \ref{con} we conclude.

\label{con} We have presented a novel approach to modeling the clustering of dark-matter haloes on mildly non-linear scales. This follows the motion of the regions that will collapse to form haloes (that we dub proto-haloes). Since the number of proto-haloes is conserved over time, for sufficiently large scales ($k<0.2\,h\,\mathrm{Mpc}^{-1}$), we can write a set of fluid equations that govern their evolution under the effect of gravity, which couples perturbations in the halo and matter density fields. We provide analytical solutions for the linearized equations and 1-loop perturbative corrections for the halo and matter power spectra. For the propagator, quantifying the memory of the density and velocity fields to their initial conditions, we also perform a resummation of perturbative corrections. Finally, for the power spectrum we compute the non-linear evolution using a semi-analytical procedure, namely an extension of the time renormalization group. The initial conditions for our evolutionary equations are specified adopting a Lagrangian bias model, originally developed to describe the clustering and motion of linear density peaks. We fix the parameters of the model so that to reproduce the distribution of proto-haloes in a high-resolution N-body simulation at $z=50$. We use the same simulation to test the predictions of our model at $z=0$. Our main results can be summarized as follows: \begin{itemize} \item Independently of the initial conditions, in the linear solution the halo density and velocity fields asymptotically match the corresponding matter fields at late times. This 'debiasing' develops at a different rate for the density and the velocity, being faster for the latter. \item Even if there is no initial density bias, the presence of a velocity bias generates a transient density bias that vanishes at late times. \item Neglecting any initial velocity bias alters the linear predictions for the halo-matter cross spectrum at redshift $z=0$ only by less than 3 per cent, for $k<0.3\,h\,\mathrm{Mpc}^{-1}$. This provides us with the motivation to ignore the velocity bias in the non-linear analysis. \item Unlike its linear counterpart, the resummed propagator is in good agreement with the N-body simulation, independently of halo mass. \item The halo-matter cross spectrum predicted by the TRG is accurate to $5$ per cent up to $k\simeq 0.1\, h\,\mathrm{Mpc}^{-1}$ for a broad range of halo masses. This does not hold for very massive haloes ($M>10^{14} h^{-1} M_{\odot}$), that have low number density and high initial velocity bias, for which discreteness effects are more important. \item The TRG result improves upon both linear theory and 1-loop corrections. Its performance is slightly enhanced accounting for the initial non-Gaussianity of the halo distribution. \item For low halo masses our model accurately describes the scale-dependent bias measured in the simulation at $z=0$. \end{itemize}

1012.4833

We address the issue of the cosmological bias between matter and galaxy distributions, looking at dark matter haloes as a first step to characterize galaxy clustering. Starting from the linear density field at high redshift, we follow the centre-of-mass trajectory of the material that will form each halo at late times (protohalo). We adopt a fluid-like description for the evolution of perturbations in the protohalo distribution, which is coupled to the matter density field via gravity. We present analytical solutions for the density and velocity fields, in the context of renormalized perturbation theory. We start from the linear solution, then compute one-loop corrections for the propagator and the power spectrum. Finally, we analytically resum the propagator, and we use a suitable extension of the time-renormalization-group method to resum the power spectrum. For halo masses M &lt; 10<SUP>14</SUP> h<SUP>-1</SUP> M<SUB>⊙</SUB>, our results at z= 0 are in good agreement with N-body simulations. Our model is able to predict the halo-matter cross-spectrum with an accuracy of 5 per cent up to k≈ 0.1 h Mpc<SUP>-1</SUP> approaching the requirements of future galaxy redshift surveys.

12104465

2010ApJ...725L.234L

1012

1012.1411_arXiv.txt

It is now widely agreed that relativistic outflows are launched hydromagnetically. The commonly invoked scenario is that a strong magnetic field in a rapidly rotating compact object serves to convert the rotational energy into the energy of the outflow (see, e.g., the review by Spruit (2010)). The main advantage of the magnetic launch mechanism is that the magnetic field lines, like driving belts, could transfer the rotational energy to a low density periphery of the central engine thus forming a baryon pure outflow with the energy per particle significantly exceeding the rest mass energy. Such outflows potentially capable of reaching relativistic velocities. Since the energy is transported, at least initially, in the form of the Poynting flux, the question arises of how and where the electromagnetic energy is eventually transformed to the plasma energy. An important point is that the flow could be considered as true matter dominated only when the electro-magnetic energy falls well below the equipartition level because only then the efficient shock dissipation becomes possible (Kennel \& Coroniti 1984). In the scope of ideal magnetohydrodynamics (MHD), the energy could in principle be transferred to the plasma via gradual acceleration by electromagnetic stresses. The problem is that in relativistic flows, the magnetic and electric forces nearly compensate each other so that the flow is nearly ballistic. It turns out that unconfined flows stops accelerating when still being highly Poynting dominated. A more efficient acceleration regime is achieved if the flow is collimated by an external medium. Such a configuration arises naturally in gamma-ray bursts (GRBs), where the relativistic jet from the collapsing stellar core pushes its way through the stellar envelope. In accreting systems, the outflow from the rotating black hole could be collimated by the pressure of a slow (and generally magnetized) wind from the outer parts of the accretion disc. However, a systematic study shows (Komissarov et al. 2007, 2009; Tchekhovskoy et al. 2008, 2010; Lyubarsky 2009) that even though externally confined flows could in principle be accelerated till the equipartition between the kinetic and magnetic energy, the conditions of that are rather restrictive. Moreover, the true matter dominated stage could in any case be achieved only at exponentially large distances (Lyubarsky 2010). An efficient conversion of the electro-magnetic into the plasma energy was claimed for impulsive flows (Granot et al. 2010,; Lyutikov 2010a,b). In a freely expanding shell, the magnetic energy is indeed eventually transformed into the kinetic energy. However, Levinson (2010) demonstrated that in the real world, it is difficult to provide conditions for the free expansion because even a tenuous ambient medium hampers the expansion. Intermittent ejection of small sub-shells could hardly help because the shells merge while still highly magnetized unless the distance between them is much larger than their width. The problem of the energy conversion could be resolved if one abandons the assumption of no dissipation. However, one has to find a reliable dissipation mechanism in highly conductive space plasma. One of the ideas is that some sort of MHD instability destroys the regular structure of the magnetic field thus triggering the anomalous dissipation. The kink instability is the best candidate (Begelman 1998; Giannios \& Spruit 2006). However, any global MHD instability could develop only if the proper Alfven crossing time is less than the propagation time, which implies $\theta\gamma<1$, where $\gamma$ is the jet Lorentz factor, $\theta$ the opening angle. This condition is rather restrictive; for example it could hardly be fulfilled in GRBs (e.g. Tchekhovskoy et al. 2010). Even if the kink instability does develop, it is still not clear whether the flow is disrupted or just helically distorted. The impact of the instability on the jet structure could be studied only with three-dimensional numerical simulations. Till now only a few attempts was made, the results still being controversial (Moll \& Spruit 2008; (McKinney \& Blandford 2008). Magnetic dissipation could occur if a small-scale structure with oppositely directed fields preexisted in the flow. In pulsar winds, such a structure arises naturally because the magnetic field in the equatorial belt of the flow changes sign every half of period. It is the annihilation of the oppositely directed fields that provides the main energy conversion mechanisms in pulsar winds (Coroniti 1990; Lyubarsky \& Kirk 2001; Kirk \& Skj{\ae}raasen 2003; Petri \& Lyubarsky 2007; Zenitani \& Hoshino 2007; see also the review by Kirk et al. (2009)). The outflow with alternating fields could in principle arise also in accreting systems if the magnetic field in the central engine changes sign. Then independently of the exact field configuration at the launch site, the flow expansion ensures that in the far zone, the overall magnetic structure is that of the "striped wind" with stripes of oppositely directed azimuthal fields separated by current sheets. The width of the stripes is small as compared with the scale of the flow therefore locally the structure of the flow is very simple: plane current sheets separate domains with oppositely directed magnetic fields. Within the current sheets, the magnetic field changes sign and the pressure balance is maintained by the thermal pressure of the plasma. Efficient annihilation of oppositely directed magnetic fields across the current sheet implies some sort of strong anomalous resistivity. The necessary condition is that the drift velocity of the current carriers becomes large enough, which is equivalent to the condition that the particle Larmor radius is not small as compared with the thickness of the sheet. Since the current in the sheets decays as $r^{-1}$ whereas the plasma density drops down as $r^{-2}$, the anomalous dissipation due to the charge starvation should inevitably occur at a large enough distance from the source. In pulsar winds, the plasma density is extremely low so that the alternating fields do dissipate in the far zone of the wind or at the wind termination shock (Kirk \& Skj{\ae}raasen 2003; Petri \& Lyubarsky 2007; Zenitani \& Hoshino 2007). However in accreting systems, the outflow is expected to be heavily loaded by the plasma from the accretion disk (and in GRBs, also by the plasma from the progenitor star (Levinson \& Eichler 2003)) therefore the charge starvation conditions are achieved only at unreasonably large distances. Spruit et al. (2001) postulated that in GRB outflows, the alternating fields annihilate with a rate $\sim 0.1$ of the speed of light. The model based on this assumption is capable of explaining many essential features of GRBs (Drenkhahn 2002; Drenkhahn \& Spruit 2002; Giannios 2006, 2008; Giannios \& Spruit 2007). The problem is that the reconnection rate of the order of 10\% from the Alfven velocity is the maximal possible reconnection rate, which is achieved only at special conditions (see review by Yamada et al. 2010). For example, there is practically no reconnection across the heliospheric current sheet, which is a prototype of large-scale current sheets in any astrophysical outflow. In this Letter, we propose a mechanism for fast reconnection of the alternating fields in Poynting dominated outflows. The basic idea is the following. If the flow is accelerated (or decelerated), an effective gravity force appears in the proper plasma frame so that the hot plasma within the current sheet is supported against the gravity force by the magnetic pressure. Such a configuration is known to be unstable with respect to the Kruskal-Schwarzschild instability, which is a magnetic counterpart of the Rayleigh-Taylor instability. Under the influence of the effective gravity force, the plasma could drip down between the magnetic field lines (e.g. Infeld (1989)) so that bridges between trickles thin out until the current sheet tears (Fig. 1). Thus the instability facilitates the reconnection. An important point is that the magnetic dissipation forces the flow to accelerate (Lyubarsky \& Kirk 2001; Drenkhahn 2002; Drenkhahn \& Spruit 2002; Kirk \& Skj{\ae}raasen 2003) therefore the process is self-sustaining: the acceleration is an aid to the reconnection whereas the reconnection promotes the acceleration. In this Letter, we show that the alternating fields could be completely dissipated by this mechanism and find the characteristic dissipation scale. \begin{figure*} \includegraphics[scale=0.8]{fig.eps} \caption{The development of the Kruskal-Schwarzschild instability in accelerating flows. In the presence of an effective gravity force, the plasma drips out of the current sheet that separates the oppositely directed magnetic fields. This promotes the magnetic reconnection.} \end{figure*}

In this Letter, we proposed a novel mechanism for the energy release in Poynting dominated outflows. The mechanism operates if a significant fraction of the Poynting flux is transferred by alternating fields. Such a structure inevitably arises in the wind from obliquely rotating pulsars and magnetars. In accreting sources, one has to assume that the magnetic polarity changes sign near the black holes. Earlier only anomalous resistivity mechanisms for the field dissipation were considered; they are efficient only in the current starvation regime when the thickness of the current sheet is comparable with the Larmor radius of the particles (Coroniti 1990; Lyubarsky \& Kirk 2001; Kirk \& Skj{\ae}raasen 2003; Zenitani \& Hoshino 2007). We found a universal mechanism for dissipation of alternating fields, which works even if the conditions of ideal MHD are fulfilled and anomalous resistivity does not appear. Thus we justified the assumption of Spruit et al. (2001) that in Poynting dominated outflows, the annihilation rate of alternating fields comprises a good fraction of the speed of light. The mechanism is based on the MHD instability of the current sheet in the accelerating flow. Under the influence of the inertia force, the plasma drips out of the current sheet allowing the oppositely directed fields come together. The ensuing field annihilation leads to the flow acceleration so that the process is self-sustaining. Note that this instability is easily suppressed by magnetic shear therefore it is very important that in the far zone of the flow, the field becomes predominantly toroidal, i.e. practically shearless. We have shown that the complete dissipation occurs at the scale (\ref{diss_radius}). In pulsar winds, $\gamma_{\rm max}\sim 10^4-10^6$ and $l\sim r_L$ so that with this mechanism, the alternating fields are dissipated at the distance $\sim 10^9-10^{13}r_L$, which is comparable or even larger than the dissipation scale due to the current starvation mechanisms (Lyubarsky \& Kirk 2001; Kirk \& Skj{\ae}raasen 2003; Zenitani \& Hoshino 2007). However, this mechanism could efficiently work in GRBs and AGNs because $\gamma_{\rm max}$ in these systems is not very large whereas the current starvation distance is huge. In AGNs, $\gamma_{\rm max}\sim 10$ and $l\sim r_g$ therefore this mechanism provides the energy release at the scale of $\sim 1000r_g$. This is compatible with observations because fitting of the blazar spectra implies that jets are already matter dominated at $\sim 1000 r_g$ (Ghisellini et al. 2010). The proposed mechanism provides also a natural basis for the "jets-in-a-jet" model (Giannios et al. 2009, 2010) that accounts for the observed ultra-fast TeV variability of blazars. In GRBs, $\gamma_{\rm max}\sim 1000$ so that the energy release occurs at the distance of the order of $10^{13}$ cm. The magnetic dissipation at this scale could reproduce the observed properties of the prompt GRB emission (Giannios \& Spruit 2007; Giannios 2008). High efficiency of the conversion of the magnetic energy into the energy of radiating particles makes this model a viable alternative to the internal shock model.

1012.1411

Reconnection of alternating magnetic fields is an important energy transformation mechanism in Poynting-dominated outflows. We show that the reconnection is facilitated by the Kruskal-Schwarzschild instability of current sheets separating the oppositely directed fields. This instability, which is a magnetic counterpart of the Rayleigh-Taylor instability, develops if the flow is accelerated. Then the plasma drips out of the current sheet, providing conditions for rapid reconnection. Since the magnetic dissipation leads to the flow acceleration, the process is self-sustaining. In pulsar winds, this process could barely compete with the earlier proposed dissipation mechanisms. However, the novel mechanism turns out to be very efficient at active galactic nucleus and gamma-ray burst conditions.

2388922

2012JCoPh.231..759P

1012

1012.1885_arXiv.txt

\label{sec:intro} Smoothed Particle Hydrodynamics (SPH), originally formulated by \citet{lucy77} and \citet{gm77}, is by now very widely used for many diverse applications in astrophysics, geophysics, engineering and in the film and computer games industry. Whilst numerous excellent reviews already exist \citep[e.g.][]{monaghan92,monaghan05,price04,rosswog09}, there remain -- particularly in the astrophysical domain -- some widespread misconceptions about its use, and more importantly, its fundamental basis. Our aim in this -- mainly pedagogical -- review is therefore not to provide a comprehensive survey of SPH applications, nor the implementation of particular physical models, but to re-address the fundamentals about why and how the method works, and to give practical guidance on how to formulate general SPH algorithms and avoid some of the common pitfalls in using SPH. Since such an understanding is critical to the development of robust and accurate methods for Magnetohydrodynamics (MHD) in SPH (hereafter referred to as ``Smoothed Particle Magnetohydrodynamics'', SPMHD), this will lead us naturally on to review the background and current status in this area -- particularly relevant given the importance of MHD in most, if not all, astrophysical problems. Whilst the paper is written with an astrophysical flavour in mind (given the topical issue of JCP for which it is written), the principles are general and thus are applicable in any of the areas in which SPH is applied. Finally, alongside this article I have released a public version of my \textsc{ndspmhd} SPH/SPMHD code, along with a set of easy-to-follow numerical exercises -- consisting of setup and input files for the code and step-by-step instructions for each problem in 1, 2 and 3 dimensions -- the problems themselves having been chosen to illustrate many of the theoretical points in this paper. Indeed, $\textsc{ndspmhd}$ has been used to compute all of the test problems and examples shown. The hope is that this will become a useful resource\footnote{\textsc{ndspmhd} is available from \url{http://users.monash.edu.au/~dprice/SPH/}. Note that we do not advocate the use of \textsc{ndspmhd} as a ``performance'' code in 3D, since it is not designed for this purpose and excellent parallel 3D codes already exist (such as the GADGET code by \citealt{springel05}). Rather it is meant as a testbed for algorithmic experimentation and understanding.} not only for advanced researchers but also for students embarking on an SPH-based research topic.

In summary, we have given an overview of SPH methodology, starting with the density estimate as the basis of all SPH formulations (Sec.~\ref{sec:calculatingdensity}) and showing how the equations of motion and energy can be self-consistently derived from the density estimate using a variational principle (Sec.~\ref{sec:densitytoequationsofmotion}). Kernel interpolation theory has been introduced mainly as a way of interpreting the SPH equations, and we have discussed how linear error analysis can be used to construct more accurate and very general derivative operators (Sec.~\ref{sec:interpolation}). In Sec.~\ref{sec:localconservation} the importance of local conservation was highlighted with respect to maintaining a regular particle distribution and thus accurate derivatives, giving us an understanding of how the tensile instability arises in SPMHD and why one should be careful in setting the ratio of smoothing length to particle spacing. Second derivatives in SPH were discussed in Sec.~\ref{sec:2ndderivs}, mainly as a way of formulating dissipative terms necessary to capture shocks and other kinds of discontinuities. In the second half of the paper, we have shown how SPMHD, like SPH, can also be formulated from a variational principle (Sec.~\ref{sec:spmhdfromL}) and have addressed the three main issues with regards to the accuracy of SPMHD: The tensile instability (Sec.~\ref{sec:mhdtensile}), the formulation of shock-capturing dissipation terms (Sec.~\ref{sec:mhddiss}) and the enforcement of the divergence-free condition on the magnetic field (Sec.~\ref{sec:divB}). Finally, this paper marks the public release of the \textsc{ndspmhd} SPH/SPMHD code that can be used to test and verify all of the ideas and methods that have been discussed. \vspace{-0.3cm}

1012.1885

This paper presents an overview and introduction to smoothed particle hydrodynamics and magnetohydrodynamics in theory and in practice. Firstly, we give a basic grounding in the fundamentals of SPH, showing how the equations of motion and energy can be self-consistently derived from the density estimate. We then show how to interpret these equations using the basic SPH interpolation formulae and highlight the subtle difference in approach between SPH and other particle methods. In doing so, we also critique several 'urban myths' regarding SPH, in particular the idea that one can simply increase the 'neighbour number' more slowly than the total number of particles in order to obtain convergence. We also discuss the origin of numerical instabilities such as the pairing and tensile instabilities. Finally, we give practical advice on how to resolve three of the main issues with SPMHD: removing the tensile instability, formulating dissipative terms for MHD shocks and enforcing the divergence constraint on the particles, and we give the current status of developments in this area. Accompanying the paper is the first public release of the NDSPMHD SPH code, a 1, 2 and 3 dimensional code designed as a testbed for SPH/SPMHD algorithms that can be used to test many of the ideas and used to run all of the numerical examples contained in the paper.

12225839

2011PhRvD..83k3012B

1012

1012.1761_arXiv.txt

The observation of neutrino masses and mixings provides the first conclusive evidence for physics beyond the standard model (BSM), so that an understanding of this phenomenon will open up one clear direction for the new BSM physics. Knowledge of the nature of the neutrino mass is crucial in order to make progress in this search. Unlike quarks and charged leptons, for which the form of the mass term is unambiguous, neutrinos are electrically neutral and therefore allow several possibilities. The two kinds of mass terms widely discussed are: (i) a Dirac type mass, which requires the theory to have a lepton number symmetry as well as a right-handed (sterile) neutrino degree of freedom; or (ii) a Majorana type mass, which necessitates the breaking of lepton number. Implementing the first possibility in the standard model requires a minimal set of assumptions, i.e., simply adding three right-handed neutrinos. In order to understand the small masses one requires that the associated Yukawa couplings are of order $10^{-12}$ or less. The challenge then becomes one of understanding this tiny Yukawa coupling or at least connecting it to some other phenomenon that requires it. In the case of Majorana neutrinos, one can write effective dimension five operators of the form $LHLH/M$, where $M$ represents the effect of higher scale physics. While no small couplings need be invoked in this case, one must understand the origin of the high scale $M$ and explore what physics is associated with it. The most widely discussed theories of this type are the seesaw models \cite{seesaw}, where the higher scale could come from the breaking of new symmetries such as $B-L$ or possibly a grand unified theory such as $SO(10)$. An intermediate possibility that has also been discussed in the literature is the pseudo-Dirac scenario, where a tiny Majorana mass is added for either one or both of the two two-component neutrino states that make up the Dirac neutrino \cite{pseudo}. If one considers all three active neutrinos to be pseudo-Dirac, then current observations put very stringent constraints on the magnitude of the Majorana mass \cite{andre}, i.e.,~$\ls 10^{-10}$ eV, in order to have the pseudo-Dirac mass splitting small enough to remain undetected in solar neutrino oscillation experiments. Roughly speaking, for $\nu_1$ and $\nu_2$, which contain a large amount of $\nu_e$, the pseudo-Dirac mass-squared difference should be smaller than $E/L \sim 10^{-11}$ eV$^{2}$ for solar neutrinos, otherwise the associated oscillations would have been observed in solar neutrino data. Detailed explanations can be found in Ref.~\cite{andre}. This tiny splitting makes these neutrinos almost Dirac particles, hence the name pseudo-Dirac. In a recent paper \cite{adm}, a new possibility for neutrino masses was pointed out, where some neutrino mass eigenstates are Dirac while the others are Majorana. This is only phenomenologically viable if one defines the Dirac or Majorana nature of neutrinos in terms of the mass eigenstates, rather than the flavor eigenstates. In this case, all neutrino flavors have large admixtures of both Dirac and Majorana type mass, and can be called ``bimodal flavor neutrinos'' (or schizophrenic neutrinos, as in Ref.~\cite{adm}). One then needs to add as many sterile neutrino states to the standard model as there are Dirac mass eigenstates. This is different from the pseudo-Dirac case in the sense that the lepton number violating and conserving terms have comparable magnitude. Another interesting feature of bimodal neutrinos is that unlike the case of pseudo-Dirac flavor neutrinos, where there exist stringent constraints on the Majorana admixture, here the oscillations of solar neutrinos (as well as all other oscillation observations) remain unaffected. In other words, in conventional neutrino oscillation experiments the bimodal flavor case looks the same as the pure Dirac or pure Majorana case. An obvious place where the bimodal scenario leads to a different effect from both the pure Majorana and pseudo-Dirac possibilities is in the predictions for neutrinoless double beta decay. This was noted for a very specific model in Ref.~\cite{adm}. In this paper we consider the most general implementation of this idea and present the predictions for neutrinoless double beta decay for the cases of both normal and inverted mass ordering. It was also pointed out in Ref.~\cite{adm} that since there is no symmetry guaranteeing the bimodal possibility, one-loop corrections can induce a tiny ($\leq 10^{-14}$ eV) amount of Majorana mass to the mass eigenstate that had a tree level Dirac mass, effectively making this mass eigenstate pseudo-Dirac. Although this value is well within the constraints from solar neutrino observations, there are implications for astrophysical neutrinos. This is one of the new points explored in this paper. \\ In Sec.~\ref{sect:models} we provide a brief review of models that could lead to the bimodal flavor neutrino scenario. Secs.~\ref{sect:cosmonu} and \ref{sect:0nubb} contain a discussion of the phenomenology of the bimodal model as it pertains to astrophysical neutrino flux ratios and neutrinoless double beta decay, respectively. The summary and conclusions are given in Sec.~\ref{sect:conclusion}.

\label{sect:conclusion} In conclusion, we have studied two different ways to test the bimodal (schizophrenic) neutrino hypothesis that one or two of the neutrino mass eigenstates are Dirac particles and the others Majorana. There are in total six nontrivial possible combinations, and we have performed a mostly phenomenological analysis of these scenarios. We noted that (i) flux ratios of extra-galactic high energy neutrinos, and (ii) the effective mass for neutrinoless double beta decay are sensitive to the different possibilities, showing nonstandard behavior in many cases. Figs.~\ref{fig:ratio_1e2mu} to \ref{fig:mee_sum_1majorana} summarize our results. In brief, we found that flux ratios can differ significantly from their standard values and the effective mass can either lie only in certain regions or even completely outside of its standard parameter space. The examples given show that the many different experimental signatures provide good tests of whether neutrino masses have the bimodal (or pseudo-Dirac) character. We have also discussed simple beyond the standard model scenarios in which such bimodal features can arise. Evidence for bimodal nature of neutrino mass will require major changes in our thinking about the physics of neutrino mass. Indeed, the field of neutrino physics has provided many surprising results in the past, and the question of neutrino mass origin is far from settled. If the hypothesis of bimodal neutrinos is supported by the experiments outlined here, it will not only provide a major departure from our current thinking about the nature of neutrino masses but also its theoretical origin from physics beyond the standard model. As such it will have major impact on the physics at and beyond the TeV scale. \begin{center} {\it \bf Acknowledgements} \end{center} JB and WR are supported by the ERC under the Starting Grant MANITOP and by the DFG in the project RO 2516/4-1 as well as in the Transregio 27. R.N.M.~is supported by the NSF under Grant No.~PHY-0968854.

1012.1761

The standard assumption is that all three neutrino mass states are either Dirac or Majorana. However, it was recently suggested by Allaverdi, Dutta, and one of the authors (R.N.M.) that mixed, or bimodal, flavor neutrino scenarios are conceivable and are consistent with all known observations (these were called “schizophrenic” in the ADM paper). In that case each individual mass eigenstate can be either Dirac or Majorana, so that the flavor eigenstates are “large” admixtures of both. An example of this “bimodal” situation is to consider one mass state as a Dirac particle (with a sterile partner), while the other two are of Majorana type. Since only Majorana particles contribute to neutrinoless double beta decay, the usual dependence of this observable on the neutrino mass is modified within this scenario. We study this in detail and, in particular, generalize the idea for all possible bimodal combinations. Inevitably, radiative corrections will induce a pseudo-Dirac nature to the Dirac states at the one-loop level, and the effects of the pseudo-Dirac mass splitting will show up in the flavor ratios of neutrinos from distant cosmological sources. Comparison of the effective mass in neutrinoless double beta decay as well as flavor ratios at neutrino telescopes, for different pseudo-Dirac cases and with their usual phenomenology, can distinguish the different bimodal possibilities.

2186760

2011ApJ...728...77N

1012

1012.4658_arXiv.txt

Before the advent of the {\em Fermi} $\gamma$-ray Space Telescope (hereafter {\em Fermi}), the {\em Compton Gamma Ray Observatory} (CGRO) succeeded in detecting GeV emission from a handful of pulsars, while a much higher number of $\gamma$-ray sources remained unidentified (\nocite{hbb+99}Hartman et al. 1999). It is generally expected that the most energetic pulsars, i.e.~with spin-down luminosity $L_{\rm sd}=4\pi^2I\dot{P}/P^3> 10^{34}$ erg s$^{-1}$ ($I=10^{45}$ g cm$^2$ is the neutron star's moment of inertia; $P$ is the spin period and $\dot{P}$ its first derivative), are the best candidates for detectable $\gamma$-ray emission (\nocite{tbb+99}Thompson et al.~1999). They are typically young pulsars of characteristic spin-down age $\tau_{\rm c}=P/(2\dot{P})<100$ kyr, with $\dot{P}>10^{-15}$ and $P\lesssim 0.1$ s. Those expectations received support with the discovery of the six EGRET pulsars, all of which have $L_{\rm sd}>3\times10^{34}$ erg s$^{-1}$. In that sample, the two oldest pulsars were PSR B1055$-$52, with a spin-down age of $\tau_{\rm c}=535$ kyr, and the radio-quiet Geminga, with $\tau_{\rm c}=342$ kyr; both these pulsars are considered middle-aged amongst the known sample of non-recycled pulsars (e.g.~Fig.~2 in Abdo et al.~2010a\nocite{aaa+10c}). The {\em Fermi} satellite was successfully launched on 2008 June 11. During the first six months of the mission, data collected with the Large Area Telescope (LAT) --- the main instrument on-board {\em Fermi} --- were analysed for pulsed $\gamma$ rays from a pre-selected list of energetic radio pulsars (\nocite{sgc+08}Smith et al.~2008). The list of pulsars was selected based on the spin-down luminosity, so that $L_{\rm sd}>10^{34}$ erg s$^{-1}$. Not surprisingly, the vast majority of non-recycled pulsars on the candidate list have $\tau_{\rm c}<10^3$ kyr. One of the few exceptions is the 96-ms pulsar PSR J2043+2740 ($P=0.0961$ s, $\dot{P}=1.23\times10^{-15}$), which is much older than the rest in the sample, with $\tau_{\rm c}=1.2\times 10^3$ kyr. A recent effort to detect a high-energy signal from this pulsar was made with the {\em AGILE} space-telescope by \nocite{ppp+09}Pellizzoni et al.~(2009). They reported a pulsed signature above 50 MeV, at the level of 4.2$\sigma$ above the background. However, a detection was not claimed and the authors calculated a 2$\sigma$ $\gamma$-ray flux upper limit of $F(>100 \ {\rm MeV})<6\times 10^{-8}$ cm$^{-2}$ s$^{-1}$. In addition to the flux estimation, the authors placed an upper limit on the $\gamma$-ray efficiency: $\eta=L_{\gamma}/L_{\rm sd}<0.01$, under the assumption of a 1-sr beam and a spectral index of 2.0. Following the AGILE observations, Abdo et al.~(2010a)\nocite{aaa+10c} folded the first six months of {\em Fermi} data, from the direction of PSR J2043+2740, with the radio-timing ephemeris supplied by Jodrell Bank (\nocite{hlk+04}Hobbs et al. 2004). That analysis yielded the first confident detection (at nearly 5$\sigma$) of $\gamma$-ray pulsations from PSR J2043+2740. PSR J2043+2740 was discovered in radio, in the Arecibo millisecond-pulsar survey at 430 MHz (\nocite{trkp94}Thorsett et al. 1994). Based on its dispersion measure, ${\rm DM}=21.0\pm 0.1$ pc cm$^{-3}$ (\nocite{rtj+96}Ray et al.~1996), and the NE2001 free-electron density model of Cordes \& Lazio (2002)\nocite{cl02}, the distance estimate for this pulsar is $D\approx 1.8$ kpc (\nocite{mhth05}Manchester et al.~2005). PSR J2043+2740 lies near the south-western shell of the Cygnus Loop ($\sim 15$ pc outside the observable edge), perhaps suggesting an association with the remnant. However, the evidence so far suggests that such an association is unlikely: the distance to the Cygnus Loop has been estimated to $540^{\, +100}_{\, -80}$ pc (Blair et al.~2005\nocite{bsr05}; Blair et al.~2009\nocite{bstc+09}); in addition, assuming that the pulsar was born within the observable limits of the remnant, the age of the latter ($<12$ kyr; \nocite{sb02}Sankrit \& Blair~2002) suggests a transverse velocity of $>980$ km s$^{-1}$ for the pulsar, which is significantly higher than the average birth velocity of the known pulsar sample ($400\pm 40$ km s$^{-1}$; \nocite{hllk05}Hobbs et al.~2005). Last but not least, the pulsar's characteristic age, as calculated from its spin parameters, is two orders of magnitude higher than the remnant's. Therefore, these discrepancies need to be reconciled before an association can be claimed. There are indeed examples of age discrepancy between pulsars and their associated remnants, like in the case of PSR J0538+2817 and the supernova remnant S147 (\nocite{acj+96}Anderson et al.~1996). This pulsar has a characteristic spin-down age of 620 kyr but has been confidently associated with the 40-kyr remnant via pulsar timing and Very Long Baseline Interferometry (VLBI; \nocite{rn03}Romani \& Ng 2003; \nocite{klh+03}Kramer et al.~2003; \nocite{nrb+07}Ng et al.~2007). Those VLBI measurements, combined with previous X-ray observations of the pulsar's thermal profile (\nocite{mkz+03}McGowan et al.~2003) have revealed a hot neutron star with a transverse velocity of 400 km s$^{-1}$, which matches the observed average. The above studies concluded that PSR J0538+2817 must have been born with a slow initial spin period, very close to that observed today, thus invalidating the usual assumption of a short birth period for this pulsar. The properties of PSR J2043+2740 make it an intriguing pulsar for high-energy studies with {\em Fermi}: it is one of the shortest-period, non-recycled $\gamma$-ray pulsars without a known SNR association. Indeed, there are a number of recent {\em Fermi} detections of non-recycled pulsars that have shorter periods than PSR J2043+2740 and for which there is yet no association with a remnant: i.e.~PSRs J1028$-$5819, J1718$-$3825, J1420$-$6048, J1813$-$1246 (Abdo et al.~2010a\nocite{aaa+10c}). However, these pulsars are $\sim 10^2$ times younger than PSR J2043+2740 and are much more likely to be associated with a nearby remnant; on the contrary, if the characteristic age for PSR J2043+2740 corresponds to its true age (see section~\ref{subsec:grayeff}), then it is very unlikely that there would be any visible remnant left for an association to be possible. Moreover, the fast rotation of PSR J2043+2740 implies a relatively high spin-down luminosity ($L_{\rm sd}=5.6\times 10^{34}$ erg s$^{-1}$) compared to pulsars of similar characteristic age. This means that, in terms of energetics, PSR J2043+2740 is on a par with much younger pulsars, like PSR J1835$-$0643, which is an order of magnitude younger. Apart from arousing observational interest, this pulsar's physical properties are also attractive in theoretical investigations. In terms of $\gamma$-ray observability, $L_\gamma/D^2$, PSR J2043+2740 has been previously considered as a strong candidate for $\gamma$-ray emission, in the framework of both Polar Cap (PC) and Outer Gap (OG) models (\nocite{rd98}Rudak \& Dyks 1998; \nocite{hib02}Hibschman 2002; \nocite{cz98}Cheng \& Zhang 1998; \nocite{mc00}McLaughlin \& Cordes 2000). It should be noted, however, that the majority of {\em Fermi} pulsar observations to date have produced $\gamma$-ray spectra that disagree with the predicted super-exponential cutoffs of PC models (e.g.~\nocite{aaa+09c}Abdo et al.~2009; \nocite{aaa+10a}Abdo et al.~2010b). Nevertheless, even within the framework of outer-magnetospheric models, there exist alternative emission geometries to the traditional OG that describe the production of $\gamma$ rays at high altitudes above the pulsar surface: one such geometry is that of the Two-Pole Caustic Model (TPC; \nocite{dr03}Dyks \& Rudak 2003). The detection of PSR J2043+2740 makes it possible to test the predictions of the above models for this pulsar against the properties derived from observations. In section \ref{subsec:HEmodels}, we discuss the results from the derived emission geometry from radio polarization data and the $\gamma$-ray lightcurves, and what those suggest for the model describing this pulsar's emission. Furthermore, the measurement of the $\gamma$-ray efficiency, $\eta=L_{\gamma}/L_{\rm sd}$, for the old PSR J2043+2740 extends the studies of $\eta$ by a factor of 2 in characteristic age. This can help us confirm or reject previous claims for an increasing $\gamma$-ray efficiency with pulsar age (\nocite{buc80}Buccheri 1980; \nocite{har81a}Harding 1981; Zhang \& Cheng 1998\nocite{zc98}). A discussion on this subject can be found in section \ref{subsec:grayeff}. The present article reports on the results of our analysis of 14 months of {\em Fermi} data from the direction of PSR J2043+2740. The LAT instrument on-board {\em Fermi} is sensitive to $\gamma$ rays of energies from 0.02 to 300 GeV; its sensitivity is an order of magnitude higher than EGRET and more than 3 times higher compared to {\em AGILE} above 0.5 GeV (\nocite{aaa+09s}Atwood et al.~2009). Using the data collected with LAT during the first 14 months of operation (2008 Aug 4 -- 2009 Oct 17), we have detected $\gamma$-ray pulsations at a very high significance ($\approx 7\sigma$) from PSR J2043+2740. This work strengthens the previously published detection of this pulsar with {\em Fermi} data (\nocite{aaa+10c}Abdo et al.~2010a).

\subsection{High-Energy Models} \label{subsec:HEmodels} All recent detections by {\em Fermi} suggest that the observed $\gamma$-ray emission from pulsars is distributed over a large fraction of the celestial sphere: i.e.~$f_\Omega\gtrsim 1$, where $f_\Omega$ is a flux-correction factor for the fact that the actual phase-averaged flux of a pulsar, integrated over the whole celestial sphere, may be higher ($f_\Omega>1$) or lower ($f_\Omega<1$) than would be inferred from assuming isotropic emission based on the phase-averaged flux for the Earth line-of-sight (\nocite{wrwj09}Watters et al.~2009). The phase-averaged flux integrated over the whole sky is a measure of the size of the emission region in the pulsar magnetosphere. In general, values of $f_\Omega$ greater than 1 are consistent with the ``fan-beam'' emission from outer-magnetospheric gaps, whereas PC models tend to produce narrow beams, i.e.~$f_\Omega\ll 1$. The geometrical constraints from our radio polarization measurements can be combined with information derived from the observed $\gamma$-ray lightcurve, to place more stringent limits on the allowed values of $\alpha$ and $\zeta$. More specifically, we can use the measured P1--P2 peak separation of $\Delta \approx 0.35$ and the gap-thickness value for PSR J2043+2740, $w=(L_{\rm sd}/10^{33} \ {\rm erg \ s}^{-1})^{-1/2}\sim 0.1$, defined in Watters et al.~(2009), to explore the allowed geometries that are consistent with the TPC and the OG model, based on the ATLAS maps of Watters et al.~(2009). Fig.~\ref{fig:atlas} shows the $\alpha$ and $\zeta$ values allowed by the RVM fit of the PA in greyscale contours. The green and pink points show the allowed geometries from the ATLAS of Watters et al.~(2009), for a two-peaked gamma-ray profile and for a profile with a phase separation of 0.35--0.4 between the major peaks, respectively. The regions where all contours (from the radio data and from the $\gamma$-ray models) overlap are delineated with solid, black lines. For the TPC model, the overlap between the contours covers roughly the ranges $\alpha \sim$ $52^\circ$--~$57^\circ$ and $\zeta \sim 61^\circ$--~$68^\circ$. For the OG model, there are two separate regions of overlap between the radio contours and the ATLAS maps: these are roughly $\alpha \sim$ $62^\circ$--~$73^\circ$ and $\zeta \sim 74^\circ$--~$81^\circ$; and $\alpha \sim$ $72^\circ$--~$83^\circ$ and $\zeta \sim 60^\circ$--~$75^\circ$. The presence of overlapping regions in both TPC and OG contour plots does not allow us to exclude either of the two models. However, it is worth noting that the two outer peaks expected in the TPC geometry both have progressive leading edges and sharp trailing edges, whereas the OG model predicts symmetric peaks (see Appendix of \nocite{wrwj09}Watters et al.~2009). Therefore, the gamma-ray lightcurve shape we observe is consistent with the OG model prediction and inconsistent with the TPC. Furthermore, as was mentioned earlier, the P1/P2 ratio decreases with energy. There are already several investigations of the energy dependence of P1 and P2 for Fermi pulsars (\nocite{aaa+10e}Abdo et al.~2010e; \nocite{aaa+10f}Abdo et al.~2010f; \nocite{aaa+10g}Abdo et al.~2010g). In all those cases, a double-peaked profile is observed that reveals a harder spectral index for P2 compared to that of P1. This indicates, for curvature radiation-dominated models, a higher accelerating electric field and/or a smaller radius of field-line curvature in the P2 emission region; future modeling can use these spectral variations to probe the underlying physics. Interestingly, the highest-energy $\gamma$ ray is found nearer to P1 than to P2, which means that it is either part of the surviving photons from the respective cut-off process operating near the P1 region, or simply a background photon not associated with the pulsar. \subsubsection{Altitude of Emission} An interesting feature that emerges from the radio-polarization profile of PSR J2043+2740 is the shape of the circularly polarized flux, which seems to change handedness very near the peak of the total flux (it lags the latter by $\approx$ 0.6\% of the pulse period). Such features in pulsar profiles have been associated with emission from near the magnetic pole, viewing angles almost along the local magnetic field (Rankin 1986\nocite{ran86}). Interestingly, the RVM fit for this pulsar gives a $\phi_0$ that lags the V-swing by $\approx 2$\% of the pulse period. According to the original relativistic model of Blaskiewicz, Cordes \& Wasserman (1991)\nocite{bcw91}, for which Dyks (2008)\nocite{dyk08} provided recently a simple explanation in terms of the relative acceleration of the corotating pulsar magnetosphere to the observer's reference frame, such a difference between the fiducial phase of an RVM swing and that of the profile is expected if the emission is altitude-dependent. In that case, the radio emission would be generated from an altitude that is roughly 2\% the size of the light-cylinder ($R_{\rm LC}\sim 4.59\times 10^4$ km). On the other hand, one can place limits on the altitude of the $\gamma$-ray emission based on the highest observed photon energy from PSR J2043+2740. The most energetic $\gamma$ ray in the pulsar's lightcurve corresponds to an energy of $\epsilon_{\rm max}=4.9$ GeV. In the framework of a standard PC model, \nocite{bar04}Baring (2004) placed a lower limit on the radius of the high-energy emission, based on $\epsilon_{\rm max}$. The minimum radius (from the neutron star's center), based on the $\gamma$-$B$ absorption of $\gamma$ rays propagating through the magnetosphere, was estimated to be $r\gtrsim (\epsilon_{\rm max}B_{12}/1.76 \ {\rm GeV})^{2/7}P^{\: -1/7}R_{\star}$, where $R_{\star}$ is the neutron star radius. Substituting for $P=0.0961$ s and $B_{12}=B/(10^{12} \ {\rm G})=0.354$, the above inequality yields $r\gtrsim 1.39 R_{\star}$, which rules out emission models that produce $\gamma$ rays very near the polar caps. Recently, Lee et al.~(2010)\nocite{ldw+10} performed a largely similar alternative analysis of the altitude of gamma-ray emission, employing a well-known high-energy asymptotic form for the magnetic pair creation rate. They concluded that the minimum emission height can be approximated with $r\gtrsim 0.11\times (\epsilon_{\rm max}B_{12})^{2/5}P^{\: -1/5}R_{\star}$. Substituting the values of PSR J2043+2740, we get $r\gtrsim 3.5R_{\star}$, which again shows that the highest-energy emission from this pulsar is not produced near the polar caps. \subsection{$\gamma$-Ray Efficiency} \label{subsec:grayeff} As was mentioned in the introduction, the high characteristic age of PSR J2043+2740 gives us an opportunity to investigate the claims for a correlation between $\gamma$-ray conversion efficiency and pulsar age, over a wider range of characteristic ages. However, it should be cautioned that any conclusions from such a study should take into account the dependency of both efficiency and characteristic age on the spin parameters: i.e.~$\eta = L_\gamma/L_{\rm sd} \propto L_\gamma(P^3/\dot{P})$ and $\tau_{\rm c}\propto P/\dot{P}$. Therefore, other pulsar parameters that depend on the period and/or its time derivatives, for example the surface polar field $B$ or the field strength at the light cylinder radius, may well also exhibit a degree of correlation with the efficiency. The $\gamma$-ray luminosity of a pulsar can be written as $L_\gamma = 4 \pi f_\Omega G_{\mathrm{> 0.1\ GeV}} D^2$. Based on the measured $\gamma$-ray luminosities of non-recycled pulsars recently detected with {\em Fermi}, Abdo et al.~(2010a)\nocite{aaa+10c} calculated the corresponding $\gamma$-ray efficiencies under the assumption of $f_\Omega=1$ and using distance estimates based either on the pulsar DM or known kinematic properties and supernova-remnant associations. A scatter plot of the $\gamma$-ray efficiency versus the characteristic age of non-recycled $\gamma$-ray pulsars with available distances is shown in Fig.~\ref{fig:efficiency}. In addition, the plot includes an upper limit on $\eta$, denoted with a solid triangle, corresponding to the efficiency of PSR J1836+5925. It should be noted that the efficiency of PSR J1836+5925 was recently estimated by Abdo et al.~(2010d)\nocite{aaa+10b}, assuming $\eta\propto 1/\sqrt{L_{\rm sd}}$. However, we did not use their estimate in this work, because such an assumption introduces an artificial correlation between the pulsar's spin parameters and the $\gamma$-ray efficiency. Despite the large error bars on $L_{\gamma}$, Fig.~\ref{fig:efficiency} reveals a correlation between $\eta$ and $\tau_{\rm c}$, with older pulsars having on average higher $\gamma$-ray efficiencies compared to younger ones. The Spearman's rank coefficient for our data set is $r_{\rm s}=0.57^{\, +0.06}_{\, -0.07}$, which implies a positive correlation between efficiency and age. The probability of chance correlation given that value of $r_{\rm s}$ is only $0.11^{\, +0.46}_{\, -0.09}\%$. In the plot of Fig.~\ref{fig:efficiency}, the data point corresponding to the highest characteristic age --- without being an upper limit --- is that of PSR J2043+2740 and is denoted by a solid black square. It is notable that the efficiency of PSR J2043+2740 is more comparable with that of pulsars that are at least 10 times younger, while the next three youngest pulsars after PSR J2043+2740, being roughly half as young, have efficiencies that are 2--6 times higher. For a number of pulsars in Fig.~\ref{fig:efficiency}, we have two distance estimates: from the pulsar DM and the NE2001 model, and from and independent method; the latter is based either on the Doppler shift of the HI lines of objects associated with the pulsar, combined with a rotation model for the Galaxy, or on trigonometric parallax, or on estimates by other means (\nocite{aaa+10c}Abdo et al.~2010a). For those pulsars, we have assigned two markers connected with a dashed line, each corresponding to the distance estimate from either method. For PSR J2043+2740, we can suppose that its association with the Cygnus Loop is true: this in turn would imply that its distance is $\approx 540$ pc and, moreover, that its true age is roughly 12 kyr instead of 1.2 Myr. Therefore, if the association is true, then the marker for PSR J2043+2740 should be shifted to the position of the solid gray square in Fig.~\ref{fig:efficiency}, which would make the efficiency of this pulsar more consistent with the observed overall trend shown by the rest of the data points. There are indeed a handful of examples for which a pulsar--remnant association has led to a significant revision to the pulsar's age: the case of PSR J0538+2817 and S147 has already been mentioned in the introduction; other examples include the 65-ms X-ray pulsar PSR J1811$-$1925, whose association with the Galactic SNR G11.2$-$0.3 suggested an age that was a factor 12 smaller than the spin-down age (\nocite{krv+01}Kaspi et al.~2001), and PSR B1951+32 --- one of the original 6 $\gamma$-ray pulsars detected with EGRET --- with a characteristic age that is more than 1.5-times higher than the age of its birth site, the SNR CTB 80 (\nocite{mgb+02}Migliazzo et al.~2002). The reason for such discrepancies between $\tau_{\rm c}$ and the true age of a non-recycled pulsar is that the calculation of the characteristic spin-down age assumes in all cases that the pulsar was born with an initial period, $P_0$, that is negligible compared to that observed today. Evidently, in the abovementioned cases, such an assumption does not hold, and those pulsars must have been born with a $P_0$ very close to the period observed today. Hence, in the case where $P_0\ll P$ does not hold, the true age of the pulsar should be calculated from $\tau=P/[(n-1)\dot{P}]\times[1-(P_0/P)^{n-1}]$, where $n=2-P\ddot{P}/\dot{P}^2$ is the pulsar braking index. If $P_0\ll P$, $\tau$ can be approximated with $\tau_{\rm c}$. Under the usual assumption of $n=3$, for pure-dipole magnetic braking, one can calculate $\tau$ for a range of $P_0$ values. By simply equating the true age of PSR J2043+2740 with the upper limit on the age of the Cygnus Loop, one finds that the only possibility of association with the remnant is if the pulsar was born with $P_0\approx 95.6$ ms. A conclusive remark based on the above is that there is certainly some support for the pulsar--remnant association on the basis of $\gamma$-ray energetics. Nevertheless, it should be recognised that, as of yet, there exists no firm evidence for such an association and that a chance alignment between PSR J2043+2740 and the Cygnus Loop is as likely. Furthermore, it is important to note that a possible rejection of the above association does not provide indisputable support for the alternative distance to PSR J2043+2740, i.e.~the one derived from the NE2001 model: it is quite possible, given the inaccuracies often associated with distances based on DM (e.g.~Deller et al.~2009\nocite{dtb+09}), that this pulsar is at an even greater distance than predicted by its DM. In fact, PSR J2022+2854, which is $\approx 4\fdg5$ away from PSR J2043+2740 in th sky, is at a DM-distance of 2.7 kpc and has a comparable DM and RM to PSR J2043+2740 (${\rm RM}_{\rm J2022+2854}=-75$ rad m$^{-2}$); while other nearby pulsars, e.g. J2113+2754 at 2 kpc, have a much lower RM (${\rm RM}_{\rm J2113+2754}=-37$ rad m$^{-2}$). \subsection{Pulsar Orientation, Proper Motion and Birthplace} The Cygnus Loop region has been searched by a number of authors for the central compact object, but so far there has been no conclusive evidence that led to a positive association (\nocite{trkp94}Thorsett et al.~1994; \nocite{rtj+96}Ray et al.~1996; \nocite{mtt+98}Miyata et al.~1998; \nocite{mot+01}Miyata et al.~2001). Fig.~\ref{fig:cygloop} shows a radio map of the Cygnus Loop remnant relative to the position of PSR J2043+2740; the map was compiled from data from the Effelsberg 1.4 GHz Medium Latitude Survey (EMLS; \nocite{ufr+98}Uyan{\i}ker et al.~1998; \nocite{ufr+99}Uyan{\i}ker et al.~1999; \nocite{rfr+04}Reich et al.~2004). It has been proposed by Uyan{\i}ker et al.~(2002)\nocite{ury+02}, and was later supported with polarization observations at 6 cm (Sun et al.~2006\nocite{srh+06}), that the Cygnus Loop consists of two supernova remnnants: the northernmost G74.3$-$8.4, centered at ($\alpha$,$\delta$)=(20$^{\rm h}$ 51\fm36, +31$^{\circ}$ 3$^{\prime}$), and G72.9$-$9.0, a blowout region in the south-western rim of the Cygnus Loop, centered at ($\alpha$,$\delta$)=(20$^{\rm h}$ 49\fm56, +29$^{\circ}$ 33$^{\prime}$). Moreover, Miyata et al.~(1998)\nocite{mtt+98} have reported the discovery of a compact X-ray source, AX J2049.6+2939, near the center of G72.9$-$9.0 (see Fig.~\ref{fig:cygloop}). Although it may seem likely that AX J2049.6+2939 could be the central compact source associated with G72.9$-$9.0, follow-up observations with the {\em ASCA} and {\em RXTE} observatories showed significant X-ray variability, which makes the identification of AX J2049.6+2939 as an ordinary, rotation-powered pulsar unlikely (Miyata et al.~2001\nocite{mot+01}). Given the lack of alternative candidates, the proximity of PSR J2043+2740 ($\approx 1\fdg5$ outside the edge of G72.9$-$9.0) makes this pulsar the only possible prospect for an association. Ultimately, a conclusion to whether PSR J2043+2740 is associated with the Cygnus Loop can be drawn by measuring the pulsar's proper motion. One method of estimating pulsar proper motions is through pulsar timing. We examined 7 years of timing observations with the Lovell telescope, at 1.4 GHz. We fitted 234 TOAs for the pulsar period and its first and second time derivatives. The post-fit timing residuals displayed the typical, long-term (i.e.~``red''), correlated timing signature often seen in the TOAs of non-recycled pulsars: this phenomenon is commonly referred to as ``timing noise'' (see \nocite{hlk+04}Hobbs et al. 2004; \nocite{hlk10}2010). Younger and more energetic pulsars often exhibit a higher amount of timing noise than older ones. PSR J2043+2740 is ranked amongst the pulsars with the highest amount of timing noise, a fact which is demonstrated by this pulsar's high and significant value of $\ddot{\nu}=56.04(57)\times 10^{-24}$ s$^{-3}$ (compare with Table 1 of Hobbs et al.~2010). The fitting procedure that can be used to derive the pulsar proper motion from timing data only provides correct results when applied to statistically ``white'' data. Therefore, some form of pre-whitening needs to be applied to effectively remove the timing noise mentioned above (e.g.~\nocite{hem06}Hobbs et al.~2006). We attempted the standard technique of fitting harmonically related sinusoids, but given the strength of the timing noise in this pulsar we were unable to pre-whiten the timing residuals sufficiently to obtain a plausible fit for proper motion. Another method for inferring the pulsar's proper-motion direction is through our radio-polarization measurements (section~\ref{subsec:poldata}). It has been claimed that there is strong evidence for a correlation between the projected spin-axis orientations of pulsars and their respective velocity-vector directions (Johnston et al.~2005\nocite{jhv+05}; Johnston et al. 2007\nocite{jkk+07}). It has also been suggested that such a correlation --- if there is indeed a physical mechanism that aligns pulsar spin axes with their velocities --- should vanish for old pulsars, as the Galactic gravitational potential will have had enough time to alter the velocities of those pulsars. On the other hand, the young population of pulsars should display a stronger case for alignment. The position angle ${\rm PA}_0$ defines the orientation --- relative to our polarization dipoles --- of the plane of linear polarization of the pulsar's emission, at the closest approach of the magnetic pole to our line-of-sight. Assuming that the plane of polarization coincides with the plane defined by the magnetic-field lines and the magnetic axis, ${\rm PA}_0$ should also correspond to the orientation of the spin axis, projected on our field-of-view: i.e.~${\rm PA}_{\rm r}\equiv{\rm PA}_0$. Note that the calculation of the PA from Stokes $Q$ and $U$ can only provide a `headless' vector with $180^\circ$ ambiguity. In addition, as was shown by Backer, Rankin \& Campbell (1975)\nocite{brc75} and Manchester, Taylor \& Huguenin (1975)\nocite{mth75}, pulsar emission can occur in two orthogonal modes: i.e.~the polarization plane can either coincide with the field line--magnetic axis plane or be perpendicular to it. Therefore, there is an additional 90$^\circ$ ambiguity in the projected direction of the spin axis. If we assume that the scenario of alignment between the spin axis and velocity is true for young pulsars and that PSR J2043+2740 was indeed born within the bounds of the Cygnus Loop --- making it a young pulsar --- then we expect that the spin-axis orientation is, within the measurement errors, directed towards the SNR. Our polarization measurements show that the most favorable case to the above scenario, given the $90^\circ$ ambiguity, gives a position angle for the spin axis of $17^\circ$ North-through-East. The spin-axis orientation relative to the Cygnus loop is shown in Fig.~\ref{fig:cygloop}, drawn at the pulsar's position. From the figure, it is immediately evident that the spin-axis orientation is significanlty offset from the direction to the centre of G72.9$-$9.0, the offset being $\approx 19^\circ$. A direct implication of the above offset could be that our assumptions are not valid and that PSR J2043+2740 is not associated with the Cygnus Loop. Another possibility is that PSR J2043+2740 was born in the supernova but the velocity vector is not aligned with the pulsar's spin axis: it is true that in the work of Johnston et al.~(2005) the offsets between the velocity vectors and the spin axes show a significant spread. However, from the same work, Johnston et al.~(2005) derive that roughly 90\% of the pulsars have an offset of $<19^\circ$ between the spin-axis orientation and the velocity vector, which makes the measured offset for PSR J2043+2740 an unlikely product of any underlying alignment mechanism leading to the observed distribution of offsets. In conclusion, the above arguments for the pulsar--remnant association, based on the pulsar orientation and its inferred proper-motion direction, do not appear to favor the Cygnus Loop as the birthplace of PSR J2043+2740. However, all of these arguments are heavily based on uncertain assumptions; a proper-motion measurement is still required for an unequivocal verdict. Recently, it was proposed that a large part of the observed timing noise of PSR J2043+2740, as well as that of other pulsars, is due to pulse-shape variations (\nocite{lhk+10}Lyne et al.~2010); those variations appear to be correlated with changes in $\dot{P}$. Moreover, it was suggested that one can fit for such pulse-shape variations and mitigate the large variations in the timing residuals by a significant factor. In particular, PSR J2043+2740 shows major pulse variations that cause 100\% change in the pulse's FWHM on time scales of $\approx 200$ days. Our data set, covering nearly 500 days, is certainly affected by those changes; still, for a proper-motion fit that would use years of data --- as is needed to obtain a precise measurement --- such an effect dominates over the proper-motion signature: i.e.~for PSR J2043+2740, the RMS of the timing residuals over 4,000 days is $>1$ s (e.g.~see Fig.~1 of \nocite{lhk+10}Lyne et al.~2010). In conclusion, such a method is certain to assist proper-motion estimates through pulsar timing for a large number of `noisy' pulsars, such as PSR J2043+2740. Alternatively, a direct measurement of the pulsar's proper motion can be made with VLBI observations. If the pulsar was born 12 kyr ago in the supernova explosion that created G72.9$-$9.0, then the angular separation between the center of the SNR and the pulsar, i.e.~$\approx {\rm 2\fdg3}$, translates to a transverse velocity of $V_{\perp}\sim 1,770$ km s$^{-1}$, assuming both pulsar and remnant are at 540 pc distance. Indeed, this corresponds to a large value of proper motion ($\approx 690$ mas yr$^{-1}$), with the highest inferred transverse velocities published so far being $V_{\perp}\sim 1,600$ km s$^{-1}$ (\nocite{hllk05}Hobbs et al.~2005). We have performed TEMPO2 simulations that generate fake TOAs having a proper-motion signature corresponding to a pulsar at the position of PSR J2043+2740, with 690 mas y$^{-1}$. Fitting those TOAs with the pulsar's spin parameters but setting the proper motion to zero, reveals the characteristic periodic sine-wave, modulated with time difference from the reference position epoch. The maximum amplitude of the simulated wave over a few thousand days, due to the proper motion, was $\sim 10$ ms (see also Fig.~8.2d in \nocite{lk05}Lorimer \& Kramer 2005). This otherwise large amount of residual RMS is still swamped, for the case of PSR J2043+2740, by the much larger timing noise of this pulsar. However, if PSR J2043+2740 possesses such a large proper motion, it will be easily measurable with VLBI, which will conclusively confirm or reject the pulsar--remnant association. \vspace{0.5cm}

1012.4658

We report on the first year of Fermi γ-ray observations of pulsed high-energy emission from the old PSR J2043 + 2740. The study of the γ-ray efficiency of such old pulsars gives us an insight into the evolution of pulsars' ability to emit in γ rays as they age. The γ-ray light curve of this pulsar above 0.1 GeV is clearly defined by two sharp peaks, 0.353 ± 0.035 periods apart. We have combined the γ-ray profile characteristics of PSR J2043 + 2740 with the geometrical properties of the pulsar's radio emission, derived from radio-polarization data, and constrained the pulsar-beam geometry in the framework of a two-pole caustic (TPC) and an outer gap (OG) model. The ranges of magnetic inclination and viewing angle were determined to be {α, ζ} ~ {52°-57°, 61°-68°} for the TPC model, and {α, ζ} ~ {62°-73°, 74°-81°} and {α, ζ} ~ {72°-83°, 60°-75°} for the OG model. Based on this geometry, we assess possible birth locations for this pulsar and derive a likely proper motion, sufficiently high to be measurable with VLBI. At a characteristic age of 1.2 Myr, PSR J2043 + 2740 is the third oldest of all discovered, non-recycled, γ-ray pulsars: it is twice as old as the next oldest, PSR J0357 + 32, and younger only than the recently discovered PSR J1836 + 5925 and PSR J2055 + 25, both of which are at least five and ten times less energetic, respectively.

12211755

2011JPhCS.312d2016N

1012

1012.5973_arXiv.txt

The discovery of a new class of neutron stars with very strong magnetic fields called magnetars has greatly enhanced the interest in the study of neutron star properties in the presence of strong magnetic fields \cite{lai}. Their surface magnetic fields could be $\geq 10^{15}$, as predicted by observations on soft gamma-ray repeaters and anomalous x-ray pulsars \cite{vas,kouvel98}. Such strong magnetic fields might be generated by dynamo processes in newly born neutron star \cite{thomp93}. In the core region the magnetic field may be even higher, the limiting value ($B_{max}$) is obtained by the scalar virial theorem \cite{lai91}. For a typical neutron star ($M=1.5 M_{\bigodot},R=15$ km) this value is $B_{max}\sim 10^{18}$ G. Such high magnetic fields can have significant effects on the equilibrium nuclear composition and equation of state in neutron star crust and interior \cite{lai91,ban}. Nonmagnetic equilibrium composition and equation of state for the outer crust was reported in a seminal paper by Baym, Pethick and Sutherland (BPS) \cite{bps}. Outer crust contains nuclei arranged in a body-centered (bcc) lattice immersed in a gas of free electrons which are relativistic above the density $\rho\sim 10^7 $g cm$^{-3}$. Though the lattice effect is small on the equation of state, it changes the equilibrium nucleus to a heavier one and lowers the total energy of the system by reducing the coulomb energy of the nucleus. At $\rho\sim 10^4$ g cm$^{-3}$, $^{56}$Fe is present as the equilibrium nucleus, but with increasing density equilibrium nuclei become more and more neutron rich through electron capture process. At a density $\rho \simeq 4\times10^{11}$ g cm$^{-3}$ neutrons begin to drip out of nuclei - this is called neutron drip point. The inner crust begins from here. In the inner crust nuclei are immersed in a neutron gas as well as a uniform background of electrons. Nuclei are also arranged in a lattice in the inner crust. The composition and the equation of state of the inner crust were earlier calculated by different groups \cite{bbp,neg}. The composition and equation of state of the outer crust of nonaccreting cold neutron stars in the presence of strong magnetic fields were first studied by Lai and Shapiro \cite{lai91}. In the presence of a magnetic field the motion of electrons perpendicular to the field get quantized into Landau orbitals. This causes the electron density to change which in turn modifies the coulomb energy. If the magnetic field is very strong then electrons occupy only the low-lying Landau levels and it may affect the sequence of nuclei and the equation of state as well as any nonequilibrium $\beta$-processes \cite{lai91}. However, there is no calculation of the inner crust composition and equation of state in the presence of magnetic fields. This paper is organised in the following way. We revisit the magnetic BPS \cite{lai91} adopting recent experimental and theoretical nuclear mass tables in Section II. The inner crust in strong magnetic fields is discussed in Section III. We conclude in Section IV.

We have revisited the BPS model of outer crust in the presence of strong magnetic fields $\sim 10^{16}$G or more using the recent experimental mass table. Further we have included the correction in the lattice energy due to the finite size of a nucleus. For zero magnetic field case, it is noted that maximum densities for heavier nuclei in this calculation are higher than those of the previous calculation \cite{lai91}. In the presence of a strong magnetic field, there are modifications in the sequence of nuclei compared with the zero field case. We have further investigated the inner crust in the presence of strong magnetic fields. For $B=10^{17}$G, it has been observed that proton fraction is enhanced at lower densities and mass numbers of nuclear clusters after subtracting the gas part are higher than those of the zero field. It would be worth investigating the implications of our results for the transport properties such as thermal and electrical conductivities and shear viscosity of the crust in magnetars.

1012.5973

We discuss the effects of strong magnetic fields through Landau quantization of electrons on the structure and stability of nuclei in neutron star crust. In strong magnetic fields, this leads to the enhancement of the electron number density with respect to the zero field case. We obtain the sequence of equilibrium nuclei of the outer crust in the presence of strong magnetic fields adopting most recent versions of the experimental and theoretical nuclear mass tables. For B ~ 10<SUP>16</SUP>G, it is found that some new nuclei appear in the sequence and some nuclei disappear from the sequence compared with the zero field case. <P />Further we investigate the stability of nuclei in the inner crust in the presence of strong magnetic fields using the Thomas-Fermi model. The coexistence of two phases of nuclear matter - liquid and gas, is considered in this case. The proton number density is significantly enhanced in strong magnetic fields B ~ 10<SUP>17</SUP>G through the charge neutrality. We find nuclei with larger mass number in the presence of strong magnetic fields than those of the zero field. These results might have important implications for the transport properties of the crust in magnetars.

12218304

2011PhLB..697..351H

1012

1012.0551_arXiv.txt

There have been recent suggestions \cite{s1,s2,s3} that in an expanding universe with Hubble parameter $H$, QCD can produce a dark energy density of order $H\Lambda_{\rm QCD}^3$. This is an intriguing claim since numerically this energy density is quite close to (although somewhat above) the observed value, and it leads to a cosmology not that dissimilar to the cosmological constant cosmology. But there are reasons to be skeptical. Given that QCD has a mass gap there would seem to be a clear decoupling between the vastly different QCD and Hubble length scales. One can use the Casimir effect as an analogy, where the effects of massive fields are exponentially suppressed. And even with a massless field the expected contribution to vacuum energy is only of order $H^4$, again in analogy with the Casimir effect. For a result different from the Casimir effect, one can consider an infinite volume and simply remove the contribution of the large wavelength modes to the zero-point vacuum energy. Then the change in vacuum energy density is \begin{equation} -\int^H \frac{d^3k}{(2\pi)^3}\frac{1}{2}\omega_k\sim -H^3m ,\label{e13}\end{equation} where we give the example of a massive field with $m\gg H$. This could be considered the maximum effect that Hubble scale physics could have on a Hubble scale contribution to vacuum energy. But this is still a miniscule contribution and it is of the wrong sign. The $H\Lambda_{\rm QCD}^3$ behavior is said to rely on various properties of QCD, including strong interactions, topological effects and ghosts that decouple in flat spacetime. But it is presently unclear to us how the arguments of \cite{s1,s2,s3} and the related work in lower dimensions \cite{s6,s7} overcome the naive expectations. We shall define $\Delta\rho_{\rm vac}$ to be the $H$ dependent contribution to the QCD vacuum energy density due to effects on Hubble scales. $\Delta\rho_{\rm vac}$ should be finite and its dependence on $H$ means that it can be discussed independently of the cosmological constant problem. From the previous remarks, if $\Delta\rho_{\rm vac}$ is to be large enough to be of any interest there must be some enhancement of the long wavelength contributions to the QCD vacuum energy. We shall argue that existing lattice studies can help to establish whether or not such an enhancement exists. We find particularly relevant a set of studies that are helping to elucidate the mechanism of confinement. The Coulomb gauge picture of confinement, originally due to Gribov \cite{s5} and developed over the years \cite{s15}, connects confinement to the nontrivial structure of the gauge field configuration space (as manifested by Gribov copies). It has received support from numerous lattice studies, with \cite{s4} being a recent example. Intrinsic to this picture of confinement is the nontrivial scaling behavior of various quantities in the far infrared. It is this type of effect that is of interest to $\Delta\rho_{\rm vac}$. On the lattice some finite volume analog of $\Delta\rho_{\rm vac}$ could in principle be measured without choosing a gauge. But this would be difficult in practice and so more useful is the fact that the lattice can provide information on the scaling behavior of 2-point functions in a particular gauge. This information can then be carried over to a continuum description where a gauge choice is necessary, and where a direct estimate of $\Delta\rho_{\rm vac}$ (a physical quantity) is more feasible.\footnote{The use of gauge-fixed lattice results in this way has recently been extolled in \cite{s16}.} The choice of the Coulomb gauge is also appropriate for another reason, its non-covariance matches the existence of a preferred frame in the cosmological setting. What is of particular interest in Coulomb gauge is the infrared behavior of the longitudinal component of the color electric field, the color Coulomb potential. The massless long wavelength modes of this field are enhanced, and this enhancement must be sufficiently large to be consistent with confinement. We shall follow a treatment \cite{s8} that derives a necessary condition for confinement in terms of the infrared divergence of the Coulomb self-energy of a color charge. In contrast the propagation of the transverse modes of the gauge field are suppressed in the infrared. This behavior is also related to confinement and it can be described by an effective mass function $m(k)\approx\Lambda_{\rm QCD}^2/k$ for these modes.\footnote{It was pointed out in \cite{s2} that this would enhance the contribution in (\ref{e13}), but not sufficiently and without correcting the sign.} We then turn to the vacuum energy and explore an interplay between the nonperturbatively enhanced long wavelength longitudinal modes and the modes that have developed a mass gap. The result is an infrared enhancement of $\Delta\rho_{\rm vac}$. We end section 2 by making use of a recent lattice study \cite{s9} to help determine the enhancement. In section 3 we return to the cosmological setting and discuss how $\Delta\rho_{\rm vac}$ feeds back and affects cosmological evolution. An interesting cosmology arises. $\Delta\rho_{\rm vac}$ depends on $H$ by definition and implicit in our picture is that there is no $H$ independent contribution to the vacuum energy; only then is the Minkowski vacuum consistent. Why this should be the case is the original cosmological constant problem. We only remark here on what we feel to be a necessary condition for a solution: the ultimate theory should contain no explicit dimensionful parameters and all masses should arise dynamically. The quantity $\Delta\rho_{\rm vac}$ of interest here is an example of a vacuum energy that appears due to the introduction of the dimensionful quantity $H$.

1012.0551

The infrared divergence of the self-energy of a color charge is due to an enhancement of the long wavelength modes of the color Coulomb potential field. There are also long wavelength contributions to the QCD vacuum energy that are similarly enhanced. Vacuum modes of Hubble scale wavelengths may be affected in a cosmological setting and this can lead to a residual positive energy density of the form H<SUP>d</SUP>ΛQCD4-d. Lattice studies constrain d. If the dark energy takes this form then the universe is driven towards de Sitter expansion, and we briefly study this cosmology when d is just slightly above unity.

12131713

2010arXiv1012.2886L

1012

1012.2886_arXiv.txt

\label{sec:intro} {\sc IGMtransfer} is a numerical code written in Fortran 90/95, intended for simulating the radiative transfer (RT) of light in the vicinity of the Ly$\alpha$ line through the intergalactic medium (IGM). Originally written with the purpose of investigating how the IGM close to Ly$\alpha$ emitting galaxies reshapes the emission line, it can also be used to probe the general transmission properties of the IGM, thus making it useful in simulations of cosmic reionization. Ly$\alpha$ is a resonant line, meaning that a photon born in the gas surrounding a hot star doesn't escape the galaxy until it has scattered millions of times on neutral hydrogen, constantly changing direction and frequency. To calculate the exact spectrum of Ly$\alpha$ photons escaping a galaxy, the full resonant scattering RT needs to be solved. Once a Ly$\alpha$ photon is the tenuous IGM, the probability of being scattered is much smaller, yet in general not negligible. Although the physics of scattering in galaxies and that of scattering in the IGM is not inherently different, the difference in physical conditions imposes a natural division of the two schemes: in the dense gas of galaxies, photons are continuously scattered in and out of the line of sight, whereas in the IGM, once a photon is scattered out of the line of sight, it is ``lost'', becoming part of the background radiation. The probability of a background photon being scattered \emph{into} the line of sight, on the other hand, is vanishingly small. An illustration of this is seen in \fig{r0}. If you haven't already done so, download all the necessary files from the URL \href{http://www.dark-cosmology.dk/~pela/IGMtransfer.html} {www.dark-cosmology.dk/\~{}pela/IGMtransfer.html}. Besides this documentation, the archive \verb+IGMtransfer-+$[$\emph{version}$]$\verb+.tar.gz+ contains the following files: \begin{center} \begin{tabular}{p{2.9cm}p{8cm}} \verb+IGMtransfer.f90+ & Main code.\\ \verb+ProcessIGM.f90+ & Processes the output from {\sc IGMtransfer}.\\ \verb+F_lam.pro+ & IDL code visualizing the output from {\sc ProcessIGM}.\\ \verb+fold[IPF].vim+ & Three Vim-scripts that structure the contents of the above three codes, if your editor is Vim.\\ \verb+test.in+ & Example input file for {\sc IGMtransfer}.\\ \verb+testdir/+ & A subdirectory for example files containing the following two files:\\ \ -- \verb+CellData.bin+ & Example input data file for {\sc IGMtransfer}, containing the physical parameters of the gas in a snapshot of a cosmological simulation at $z = 3.5$.\\ \ -- \verb+GalData.dat+ & Example input file containing physical parameters of the galaxies in the snapshot.\\ \verb+toymodel.in+ & Input file for a small toy model.\\ \verb+toymodel/+ & A subdirectory containing the following three files:\\ \ -- \verb+CellData.dat+ & ASCII data file with toy model gas parameters.\\ \ -- \verb+dat2bin.f90+ & Converts ASCII data to binary, ready for {\sc IGMtransfer}.\\ \ -- \verb+GalData.dat+ & Data for two toy galaxies.\\ \end{tabular} \end{center} After the following description of the basic principles and the physics of the main code, the individual programs are explained. A more thorough description is given in \citet{lau11}, which represents the work first employing {\sc IGMtransfer}, and in \citet[][my Ph.D. thesis]{lau10}. \begin{figure}[!t] \centering \ifnum\figtype=1 \includegraphics [width=0.90\textwidth] {r0.eps} \fi \ifnum\figtype=2 \includegraphics [width=0.90\textwidth] {r0.pdf} \fi \caption{{\cap Illustration of the difference between the galactic RT and the IGM RT. Close to the galaxy, photons are scattered both in and out of the line of sight. In the rarefied IGM, photons are mainly scattered out of the line of sight, obviating the need for a full Ly$\alpha$ RT. The exact value of the distance $r_0$ from a galaxy to begin the IGM RT is somewhat arbitrary, but is of the order of the virial radius of the galaxy.}} \label{fig:r0} \end{figure} \subsection{Underlying concepts} \label{sec:concept} {\sc IGMtransfer} performs the IGM RT in a ``computational box'' with a (possibly adaptively refined) cell-based structure. The final results are obtained by considering the average of the RT performed for many sightlines emerging in many directions from many galaxies. For a given sightline, a (normalized) spectrum is emitted, suffering random absorption lines (i.e.~the Ly$\alpha$ forest) as it is continuously redshifted when receding from the galaxy. When the edge of the computational volume is reached, the sightline continues in a random, inward angle, thus ``bouncing'' around until the bluest wavelength of the simulated spectrum has been redshifted into the Ly$\alpha$ resonance (\fig{bounce}). {\sc IGMtransfer} was originally applied in a non-periodic, spherical volume. The present version (v1.1) includes the possibility of using the full volume of the computational box, and a future version will include the possibility of utilizing periodic boundary conditions, i.e.~once the edge is reached, the sightline continues on the opposite side of the box. \begin{figure}[!t] \centering \ifnum\figtype=1 \includegraphics [width=0.50\textwidth] {bounce.eps} \fi \ifnum\figtype=2 \includegraphics [width=0.50\textwidth] {bounce.pdf} \fi \caption{{\cap Illustration of how sightlines are cast through the cosmological volume (in the case of a spherical subvolume). To sample sufficiently the full solid angle of $4\pi$ around the individual galaxies, {\tt n\_los} ($\sim$$10^3$) sightlines are cast from each galaxy (of which four are shown here). Each sightline is started at a distance $r_0$ --- which reflects the distance from the center at which the full scattering RT is no longer necessary --- from the center of a galaxy and traced until the bluest wavelength of the emitted spectrum has been redshifted into resonance. When the edge of the spherical volume is reached, the ray ``bounces'' back, i.e.~continues in a random inward angle.}} \label{fig:bounce} \end{figure} The spectrum of each individual sightline is written to an output file which can subsequently be processed by the program {\sc ProcessIGM}. This two-step process also enables the user to use the calculated spectra to perform other analyses, e.g.~investigate them for the relative abundance of different absorption systems. \subsection{Main output} \label{sec:out} The final, main output are the two related quantities: \begin{enumerate} \item the \emph{transmission function} $\Flam$, giving the fraction of light transmitted as a function of wavelength (as an average over many sightlines cast through the cosmological volume), and \item the average transmission $\T$ of the IGM in a wavelength interval blueward of the Ly$\alpha$ line. \end{enumerate} \subsubsection{Transmission function} \label{sec:Ftheo} The resulting value of $\Flam$ at wavelength $\lambda$ for a given sightline is \begin{equation} \label{eq:F} \Flam = e^{-\tau(\lambda)}. \end{equation} The optical depth $\tau$ is the sum of contributions from all the cells encountered along the line of sight: \begin{equation} \label{eq:tau} \tau(\lambda) = \sum_i^{\mathrm{cells}} r_i \,n_{\textrm{{\scriptsize H}{\tiny \hspace{.1mm}I}},i} \,\sigma(\lambda + \lambda v_{||,i}/c). \end{equation} Here, $n_{\textrm{{\scriptsize H}{\tiny \hspace{.1mm}I}},i}$ is the density of neutral hydrogen in the $i$'th cell, $r_i$ is the distance covered in that particular cell, $v_{||,i}$ is the velocity component of the cell along the line of sight, and $\sigma(\lambda)$ is the cross section of neutral hydrogen. Due to the resonant nature of the transition, the largest contribution at a given wavelength will arise from the cells the velocity of which corresponds to shifting the wavelength close to resonance. However, for sufficiently high column density absorbers (the so-called ``damped Ly$\alpha$ absorbers) the damping wing of the profile may cause absorption at velocities quite far from this. If dust is present in the simulation, $\nhi\sigma(\lambda)$ is replaced by $\nhi\sigma(\lambda) + \nd\sigma_{\mathrm{d}}(\lambda)$. Assuming a spectrum of light $\Iem$ escaping a galaxy (simulated, e.g., using Monte Carlo simulations, as in \citet{lau09a,lau09b}), the final, observed spectrum $\Iobs$ is then \begin{equation} \label{eq:Iobs} \Iobs = \Flam \, \times \, \Iem. \end{equation} \Fig{spXF} illustrates this effect. \begin{figure}[!t] \centering \ifnum\figtype=1 \includegraphics [width=1.00\textwidth] {spXF.eps} \fi \ifnum\figtype=2 \includegraphics [width=1.00\textwidth] {spXF.pdf} \fi \caption{{\cap Illustration of the effect of the IGM on the observed Ly$\alpha$ profile emerging from a galaxy at $z \sim 3.5$. Without taking into account the IGM, the two peaks are roughly equally high (\emph{left panel}). However, when the spectrum is transmitted through the IGM characterized by the transmission function $\Flam$ (\emph{middle panel}, with the \emph{cyangrayish} region representing the 68\% confidence interval), the blue peak is dimished, resulting in an observed spectrum with a higher red peak (\emph{right panel}). The figure is taken from \citet{lau11}}.} \label{fig:spXF} \end{figure} Due to the correlation of the IGM with the source, $\Flam$ is non-trivial close to the Ly$\alpha$ line, but far from the source, or, spectrally speaking, at very blue wavelengths, it becomes a constant function of wavelength (as long as one does not enter an appreciably different redshift epoch). \subsubsection{Average transmission} \label{sec:Ttheo} The transmission function is probably only accurate if you have very high resolution in your simulation; if your cell size is of the order of the virial radius of your galaxies, you will probably overestimate the absorption. However, {\sc IGMtransfer} may still sweeten your life by calculating the \emph{average}, or rather median, transmitted fraction $\T$ in a large wavelength interval bluward of the Ly$\alpha$ line. \Fig{Songaila} shows such calculated fractions at different redshifts, compared to the observations of \citet{son04}. \begin{figure}[!t] \centering \ifnum\figtype=1 \includegraphics [width=0.80\textwidth] {Songaila.eps} \fi \ifnum\figtype=2 \includegraphics [width=0.80\textwidth] {Songaila.pdf} \fi \caption{{\cap Comparison of observations (\emph{black data points}) and simulations (\emph{colored data points}) of the transmitted flux blueward of the Ly$\alpha$ line as a function of redshift. The figure is taken from \citet{lau11}, but the observed data points are from \citet{son04}.}} \label{fig:Songaila} \end{figure} \subsection{{\tt fold[IPF].vim}} \label{sec:fold} {\sc IGMtransfer} is written in one single file. If your editor is Vim, \emph{first} time you open the source code use \begin{verbatim} vim -s foldI.vim IGMtransfer.f90 \end{verbatim} This will fold distinct parts of the code into single lines, making it more manageable. To inspect the lines in a fold, go to that fold and press space. To close the fold again, press \verb+zc+ while standing in the fold. If you screw something up, exit and open one more time with \verb+-s foldI.vim+. Similarly, \verb+ProcessIGM.f90+ and \verb+F_lam.pro+ can be folded with \verb+foldP.vim+ and \verb+foldF.vim+, respectively. If the folds are gone when you open the files next time, try adding these lines in your \verb+.vimrc+: \begin{verbatim} au BufWinLeave * mkview au BufWinEnter * silent loadview \end{verbatim}

1012.2886

This document describes the publically available numerical code "IGMtransfer", capable of performing intergalactic radiative transfer (RT) of light in the vicinity of the Lyman alpha (Lya) line. Calculating the RT in a (possibly adaptively refined) grid of cells resulting from a cosmological simulation, the code returns 1) a "transmission function", showing how the intergalactic medium (IGM) affects the Lya line at a given redshift, and 2) the "average transmission" of the IGM, making it useful for studying the results of simulations of cosmic reionization.

12207316

2011JCAP...04..022W

1012

1012.0883_arXiv.txt

\label{sec1} Dark energy cosmology has been one of the most active fields in astronomy and physics, since the exciting discovery of current accelerated expansion of our universe~\cite{r1}. The first evidence came from the observation of Type~Ia supernovae (SNIa) in 1998~\cite{r2}. Five years later, cosmology entered the so-called ``precision era'' in 2003 when the first year Wilkinson Microwave Anisotropy Probe (WMAP) observations of the cosmic microwave background (CMB) had been released~\cite{r3}. Up to now, the CMB observation is still a very powerful probe for cosmology, and the CMB data from WMAP mission provide the most important basis. However, in the passed years, some unusual phenomena have been found from the CMB data released by WMAP team. Remarkably, among these unusual phenomena, it is claimed that there exists a preferred direction in the CMB temperature map (known as the ``Axis of Evil'' in the literature)~\cite{r4}. In fact, there are two different approaches to deal with this problem. The first one is to admit that this phenomenon is an observational fact and try to explain it in the cosmological theories (see e.g.~\cite{r5,r6,r7} and references therein). The second approach is to consider instead that this phenomenon might be an artifact due to the observational systematics. For example, in a series of works by Liu and Li (L\&L)~\cite{r8,r9,r10}, they claimed that there exists a timing asynchrony of $-25.6\,$ms between the spacecraft attitude and radiometer output timestamps in the original raw WMAP time-ordered data (TOD). If one fixed this problem, the most of CMB quadrupole component disappear. It is worth noting that recently their findings has been confirmed by several independent authors (see e.g.~\cite{r11,r12,r13}). In~\cite{r8}, L\&L reprocessed the WMAP data while the aforementioned timing asynchrony has been corrected, and they obtained an alternative CMB map in which the quadrupole dropped to nearly zero. In addition, using the alternative CMB data, they constrained the cosmological parameters of the standard flat $\Lambda$CDM model, and found that these cosmological parameters have been changed notably. For convenience, we reproduce their main results in Table~\ref{tab1}. It is easy to see that the fractional matter density $\Omega_{m0}=\Omega_{b0}+\Omega_{c0}$ of L\&L~\cite{r8} is considerably larger than the one of WMAP~\cite{r14,r15}.\\[-1mm] \begin{table}[htbp] \begin{center} \begin{tabular}{cccc} \hline\hline\\[-3mm] Description & Symbol & WMAP~\cite{r14,r15} & L\&L~\cite{r8} \\[1.2mm] \hline \\[-3mm] ~Hubble constant (km/s/Mpc)~~~~~ & $H_0$ & $71.9^{+2.6}_{-2.7}$ & $71.0\pm 2.7$ \\ Baryon density & $\Omega_{b0}$ & ~~~~~$0.0441\pm 0.0030$~~~~~~ & $0.052\pm 0.0030$~ \\ Cold dark matter density & $\Omega_{c0}$ & $0.214\pm 0.027$ & $0.270\pm 0.027$ \\ Dark energy density & $\Omega_{\Lambda 0}$ & $0.742\pm 0.030$ & $0.678\pm 0.030$ \\ Fluc. Ampl. at $8h^{-1}\,$Mpc & $\sigma_8$ & $0.796\pm 0.036$ & $0.921\pm 0.036$ \\ Scalar spectral index & $n_s$ & $0.963^{+0.014}_{-0.015}$ & $0.957\pm 0.015$ \\ Reionization optical depth & $\tau$ & $0.087\pm 0.017$ & $0.109\pm 0.017$ \\[1.2mm] \hline\hline \end{tabular} \end{center} \caption{\label{tab1} The cosmological constraints on the standard flat $\Lambda$CDM model, reproduced from~\cite{r8}.} \end{table} To our knowledge, there is still controversy about the findings of L\&L in the community. In the present work, we try to be neutral as much as possible. Instead of discussing the detailed data process of WMAP mission (as in the works by L\&L~\cite{r8,r9,r10} or other authors~\cite{r11,r12,r13}), in this work we would like to see the implications to dark energy cosmology if L\&L are right. While L\&L claimed that there is a bug in the WMAP pipeline which leads to significantly different cosmological parameters, an interesting question naturally arises, namely, how robust is the current dark energy cosmology with respect to systematic errors and bugs? So, in this work, we adopt the alternative CMB data of L\&L as a strawman to study the robustness of dark energy predictions. As is well known, using the full CMB data to perform a global fitting consumes a large amount of computation time and power. As a good alternative, one can instead use the shift parameter $R$ from the CMB data, which has been considered extensively in the literature. It is argued that the shift parameter $R$ is model-independent, and it contains the main information of the CMB data~\cite{r16}. So, in Sec.~\ref{sec2} we firstly derive the corresponding shift parameter $R$ from the alternative CMB data of L\&L~\cite{r8}. In addition, we briefly introduce the other observational data, such as SNIa and the baryon acoustic oscillation (BAO), which are also used in the present work. In Sec.~\ref{sec3}, we perform the data analysis by using the cosmological observations introduced in Sec.~\ref{sec2}. In particular, we discuss the tension between CMB and SNIa in Sec.~\ref{subsec3}; we study the age problem in dark energy models in Sec~\ref{subsec4}; and the cosmological constraints on dark energy models are considered in Sec.~\ref{subsec5}. Finally, the conclusion and discussions are given in Sec.~\ref{sec6}.

\label{sec6} Recently, in a series of works by L\&L~\cite{r8,r9,r10}, they claimed that there exists a timing asynchrony of $-25.6\,$ms between the spacecraft attitude and radiometer output timestamps in the original raw WMAP time-ordered data (TOD). In~\cite{r8}, L\&L reprocessed the WMAP data while the aforementioned timing asynchrony has been corrected, and they obtained an alternative CMB map in which the quadrupole dropped to nearly zero. In the present work, we try to see the implications to dark energy cosmology if L\&L are right. While L\&L claimed that there is a bug in the WMAP pipeline which leads to significantly different cosmological parameters, an interesting question naturally arises, namely, how robust is the current dark energy cosmology with respect to systematic errors and bugs? So, in this work, we adopt the alternative CMB data of L\&L as a strawman to study the robustness of dark energy predictions. In this work, we found that L\&L's alternative CMB data~\cite{r8} favor a larger $\Omega_{m0}$ in all the dark energy models. As a result, we found that the tension between CMB and SNIa can be alleviated to some extent, since SNIa dataset usually favors a large $\Omega_{m0}$. However, the age problem becomes even worse in the dark energy models, since a larger $\Omega_{m0}$ usually leads to a smaller age of our universe at any redshift $z$. On the other hand, we found that L\&L's alternative CMB data~\cite{r8} do not significantly change the rank of dark energy models from the perspective of model comparison. Of course, it is a big advantage that the quadrupole dropped to nearly zero in the alternative CMB map of L\&L (see~\cite{r8,r9,r10}). Altogether, we consider that it is better to keep neutral to L\&L's findings so far. While WMAP has been critical in establishing a highly successful cosmology, it is most important that these results are verified by other independent experiments, such as Planck. Finally, we stress that this work is a hypothetical study in fact, since the works of L\&L~\cite{r8,r9,r10} are highly controversial in the community. One should be aware of the risk that the present work is likely to be completely irrelevant if it turns out that the works of L\&L are indeed flawed. Here are some remarks to be taken serious (we thank the anonymous referee for pointing out these issues). Firstly, in fact there are many discussions on L\&L's claims in the community (see e.g.~\cite{r77,r78,r79}), and many people (including both WMAP and Planck experts) believe that L\&L are mistaken in their claims. Secondly, if the high-$\ell$ spectrum really is as discrepant as claimed by L\&L, one could expect that many ground-based experiments would also have seen this effect. Thirdly, it is worth asking whether the timing offset was applied correctly in the actual map making but not in the calibration step. Up to now, to our knowledge, the discussions on L\&L's claims in the community are still unsettled, whereas L\&L are still defending their claims persistently in more papers (see~\cite{r80} for instance) and in many conferential talks around the world. We hope that this ongoing controversy could be finished with a firm and reliable conclusion in the near future.

1012.0883

Recently, in a series of works by Liu and Li (L&amp;L), they claimed that there exists a timing asynchrony of -25.6 ms between the spacecraft attitude and radiometer output timestamps in the original raw WMAP time-ordered data (TOD). L&amp;L reprocessed the WMAP data while the aforementioned timing asynchrony has been corrected, and they obtained an alternative CMB map in which the quadrupole dropped to nearly zero. In the present work, we try to see the implications to dark energy cosmology if L&amp;L are right. While L&amp;L claimed that there is a bug in the WMAP pipeline which leads to significantly different cosmological parameters, an interesting question naturally arises, namely, how robust is the current dark energy cosmology with respect to systematic errors and bugs? So, in this work, we adopt the alternative CMB data of L&amp;L as a strawman to study the robustness of dark energy predictions.

12201831

2011ASPC..442..207H

1012

1012.2909_arXiv.txt

Upcoming radio observation facilities such as the Australian Square Kilometre Array Pathfinder (ASKAP), the MeerKAT Karoo Array Telescope, the Low Frequency Array (LOFAR), and ultimately the Square Kilometre Array (SKA) will pose a significant challenge for current astronomical data analysis and visualization tools. The expected data product sizes (e.g. 2.2 TB per 8 hour of observation for the proposed WALLABY all-sky HI survey\footnote{See http://www.atnf.csiro.au/research/WALLABY for details.}) are orders of magnitude larger than astronomers, and existing astronomy software, are accustomed to dealing with. In \citet{Hassan:2010a}, we presented a distributed GPU framework to interactively volume render larger-than-memory astronomical data cubes. Throughout this work, we demonstrated how volume rendering offers an alternative to standard 2D visualization techniques and provided a way to overcome the technological barrier caused by the computational requirement of volume rendering for large data cubes. The presented framework utilizes a heterogeneous CPU and GPU hardware infrastructure, combining shared- and distributed-memory architectures, to yield a scalable volume rendering solution, capable of volume rendering image cubes larger than a single machine memory limit, in real-time and at interactive frame rates. The usage of GPUs as the main processing backbone for the system, and the remote visualization architecture adopted in our design for this framework enables further enhancement for the astronomer's visualization experience. In this paper we present an extension to this framework to provide: better visualization output by integrating external quantitative information with the volume rendering output (see section \ref{sct:Duchamp}) , and high resolution output via multi-panel displays (see section \ref{sct:Opti}).

We demonstrate the ability to extend our GPU volume rendering framework to offer better visualization outcomes. Two main features were presented: integrating source finder output, and utilizing multi-panel displays. Both of these features demonstrate two main strong points in the framework design, namely the usage of GPU cluster as the main processing back-bone, and the separation between the rendering and the result display. We anticipate the ability to further enhance this for integrating more quantitative visualization tools and a better utilization of the GPU processing power.

1012.2909

The Australian SKA Pathfinder (ASKAP) will be producing 2.2 terabyte HI spectral-line cubes for each 8 hours of observation by 2013. Global views of spectral data cubes are vital for the detection of instrumentation errors, the identification of data artifacts and noise characteristics, and the discovery of strange phenomena, unexpected relations, or unknown patterns. We have previously presented the first framework that can render ASKAP-sized cubes at interactive frame rates. The framework provides the user with a real-time interactive volume rendering by combining shared and distributed memory architectures, distributed CPUs and graphics processing units (GPUs), using the ray-casting algorithm. In this paper we present two main extensions of this framework which are: using a multi-panel display system to provide a high resolution rendering output, and the ability to integrate automated data analysis tools into the visualization output and to interact with its output in place.

12166218

2011AnPhy.326.1998A

1012

1012.5821_arXiv.txt

1012.5821

Thermal leptogenesis explains the observed matter-antimatter asymmetry of the universe in terms of neutrino masses, consistent with neutrino oscillation experiments. We present a full quantum mechanical calculation of the generated lepton asymmetry based on Kadanoff-Baym equations. Origin of the asymmetry is the departure from equilibrium of the statistical propagator of the heavy Majorana neutrino, together with CP violating couplings. The lepton asymmetry is calculated directly in terms of Green's functions without referring to "number densities". Compared to Boltzmann and quantum Boltzmann equations, the crucial difference are memory effects, rapid oscillations much faster than the heavy neutrino equilibration time. These oscillations strongly suppress the generated lepton asymmetry, unless the standard model gauge interactions, which cause thermal damping, are properly taken into account. We find that these damping effects essentially compensate the enhancement due to quantum statistical factors, so that finally the conventional Boltzmann equations again provide rather accurate predictions for the lepton asymmetry.

12168673

2011ApJ...733...83B

1012

1012.2892_arXiv.txt

Observations of the cosmic microwave background (CMB) strongly suggest that the dominant source of perturbations in the early universe was adiabatic and nearly Gaussian in nature \citep{Komatsu:2010fb}, consistent with an inflationary scenario. Inflationary scenarios also induce a small amount of non-Gaussianity, typically characterised by a parameter $f_{\mathrm{NL}}$ \citep[for example]{Liguori:2010hx} which from the WMAP seven-year results is currently constrained to be $-10<f_{\mathrm{NL}}<74$ at 95$\%$ confidence \citep{Komatsu:2010fb}. It is expected that these bounds will tighten significantly with the forthcoming Planck data \citep{Liguori:2010hx}. A detection of a significant $f_{\mathrm{NL}}$ could provide vital information about the early universe and high-energy physics, as it could rule out a variety of inflationary models including the simplest single-field models. However, the inflationary model does not exclude the possibility that non-linear sources might also play a role in sourcing perturbations. Frequently-studied examples of such sources include cosmic defects \citep[for example]{Turok:1997gj} and cosmological magnetic fields \citep[for example]{Giovannini:2003yn,Tsagas:2004kv}. Magnetic fields are observed on many scales in the cosmos, including on cluster scales with a coherence length on the order of megaparsecs and field strengths on the order of between nano-Gauss and micro-Gauss \citep{Kronberg:1993vk,Grasso:2000wj,Giovannini:2003yn,Xu:2005rb,Kahniashvili:2010wm}. Fields on larger scales are notoriously difficult to detect but there are some suggestions that such fields might exist with field strengths up to the order of micro-Gauss. Recent observations of radiation from blazars implies a lower limit in the extra-galactic medium of $B\gtrsim\mathcal{O}(10^{-16})$-$\mathcal{O}(10^{-15})$, permeating extracluster voids \citep{Neronov:1900zz,Tavecchio:2010ja,Dolag:2010ni}. The presence of fields on such fields implies either a primordial origin or an efficient transfer of fields from within galaxies deep into the intergalactic medium. The precise origin of fields on such large scales remains debated but magnetogenesis scenarios exist which were viable in the extremely early universe. Such scenarios can be produced directly during inflation \citep{Turner:1987bw,Bassett:2000aw,Prokopec:2004au,Giovannini:2007rh,Campanelli:2007cg,Bamba:2008my} or at a phase transition \citep{Baym:1995fk,Martin:1995su,Hindmarsh:1997tj,Boyanovsky:2002wa,Kahniashvili:2009qi}. Fields can also be produced by the production of non-linear vorticity from linear density perturbations\citep{Gopal:2004ut,Matarrese:2004kq,Takahashi:2005nd,Ichiki:2006cd,Siegel:2006px,Kobayashi:2007wd,Maeda:2008dv}, but the impact these have on the CMB is complicated by their evolving, non-trivial nature. The impact primordial magnetic fields have on the CMB and its anisotropies have been well-studied, with \citep{Barrow:1997mj,Subramanian:1998fn,Durrer:1998ya,Koh:2000qw,Kahniashvili:2000vm,Mack:2001gc,Clarkson:2002dd,Lewis:2004ef,Kahniashvili:2006hy,Kahniashvili:2008hx,Yamazaki:2008gr,Finelli:2008xh,Paoletti:2008ck,Bonvin:2010nr,Giovannini:2009fu,Yamazaki:2010nf,Paoletti:2010rx,Kahniashvili:2010wm} being some instructive examples. While older literature tended to assume a homogeneous background component with an inhomogeneous perturbation, more recent work has typically focused on tangled configurations without a background component and I assume this throughout. These studies fairly consistently suggest that the field is constrained to be of at most nano-Gauss in magnitude. The spectral index is restricted to be approximately scale-invariant \citep{Yamazaki:2010nf,Paoletti:2010rx}, with limits growing extremely tight for a primordial magnetic field with index far from scale-invariance \citep{Caprini:2001nb}. A large-scale homogeneous field also introduces characteristic correlations between multipole moments with $\Delta l\in\{-2,0,2\}$ and $\Delta m\in\{0,\pm 1,\pm2\}$ which vanish in the standard scenario \citep{Kahniashvili:2008sh}. However, the magnetic 2-point signal is overwhelmed on large-scales by the standard perturbations, with the $B$-mode polarisation being perhaps the most realistic option if we are to detect it directly. The increasing accuracy of measurements of the CMB non-Gaussianity provides an alternative. The stress tensor of a magnetic field is non-linear, implying that the statistics induced on matter perturbations are intrinsically non-Gaussian, regardless of the nature of the underlying magnetic field. Since the standard scenario contains relatively few sources of primordial non-Gaussianity, it is possible that a magnetic signal is dominant. Viewed another way, predicted signals from a magnetic field are likely to be of a characteristic nature, and must be found in and cleaned from the CMB data before any conclusions on early-universe physics can be made. Aspects of the three-point moments have been studied in a series of papers in the last few years \citep{Brown:2005kr,Brown:2006wv,Seshadri:2009sy,Caprini:2009vk,Trivedi:2010gi,Cai:2010uw,Shiraishi:2010yk,Kahniashvili:2010us}. A bispectrum is set by three wavevectors, which we denote with $\mathbf{k}$, $\mathbf{p}$ and $\mathbf{q}$. Since to retain statistical isotropy these must form a closed triangle, this geometry can equivalently be expressed with the scalars $k,r,\phi$, where $p=rk$ and $\phi$ is the angle between $\mathbf{p}$ and $\mathbf{q}$. Employing these variables the bispectrum geometry can be written as a foliation of planes of constant $r$ and for each constant angle $\phi$ we then have a one-dimensional line through the bispectrum which in broad terms is expected to act in a similar manner to the power spectra. The magnetic bispectra studied thus far have typically been along only three such lines, all in the $r=1$ plane -- the ``colinear'' case where $k=p=q/2$ and so $\phi=0$ \citep[hereafter BC05, B06 and CFPR09]{Brown:2005kr,Brown:2006wv,Caprini:2009vk}, the ``equilateral'' case where $k=p=q$ and so $\phi=2\pi/3$ \citep[hereafter SS09, CFPR09 and TSS10]{Seshadri:2009sy,Caprini:2009vk,Trivedi:2010gi}, and the ``local'' or degenerate case where $k\approx p$, $q\approx 0$ and so $\phi\approx\pi$ (SS09, CFPR09, TSS10). TSS10 also considered configurations where $\phi=0$ but $k\neq p$. The recent studies have expanded the previous results considerably. SS09 considered the equilateral and degenerate lines of the bispectrum of the magnetic energy density for nearly scale-invariant magnetic fields, concluding that the degenerate line provides the greatest contribution to the integral and employing an approximation to this dominant term to estimate the CMB signal. Likewise, CFPR09 considered the bispectrum of the energy density and considered the colinear, equilateral and degenerate lines. The authors generally relied on approximations that neglect angular terms in the integrations, or apply only on large scales. Doing so recovers the scaling behaviour of the bispectrum at the expense of an accurate calculation of the relative amplitudes between lines. Since the degenerate line was found to diverge as $q^{2n+3}$ as $q\rightarrow 0$ this term is likely to dominate. CFPR09 also present exact solutions for the colinear case for both a causal field and a field relatively close to scale-invariance, which enable them to test their approximations. The approximations are certainly reasonable, but not ideal. In particular, since the bispectra are not positive-definite it is unclear whether there are strong cancellations to the degenerate line arising from other parts of the bispectrum. \citet{Cai:2010uw} employed the approximations of SS09 and CFPR09 and extended the treatment to full transfer functions. More recently, TSS10 considered the bispectrum of the anisotropic pressure of a near scale-invariant magnetic field. Unlike the previous papers they evaluated the bispectrum along the degenerate line in full, without neglecting any angular terms, finding it to be positive (in contrast with the approximate result, which would be negative). However, there are four components of the magnetic stress tensor comprising six degrees of freedom -- the isotropic pressure/energy density, the anisotropic pressure, a vector component, and a transverse-traceless tensor component. Each of these are of roughly equal magnitude, as evidenced by the two-point moments \citep[for example]{Paoletti:2008ck,Brown:2010ms}. There are many correlations that can arise between these at the three-point level, and in principle all of them must be considered for the CMB. The rotationally-invariant combinations are listed in BC05 and B06, and numerical solutions for the colinear line are given, although for $n_B<-1$ these took the form of noisy statistical realisations. \citet{Shiraishi:2010sm} recently presented a formalism capable of considering the vector and tensor auto-correlation in full generality; our previous work considered only a fully-contracted form of the tensor auto-correlation and neglected the vector auto-correlation entirely. In \citet{Shiraishi:2010yk} this formalism was applied to the 3-point vector auto-correlation. In this paper the authors employed a technique to integrate the full bispectrum and, therefore, evaluate an unambiguous CMB signal.\footnote{The authors followed this with a detailed presentation in \citet{Shiraishi:2011fi} while this manuscript was under review.} A similar paper was presented by \citet{Kahniashvili:2010us}. These authors focused on the symmetries of the intrinsic bispectrum and on the appearance of anisotropic and off-diagonal terms but in principle considered the full bispectrum. It is this consideration of the full bispectrum that is lacking for the other auto-correlations. This causes problems for causal fields and other fields far from scale-invariance in particular, where it is not necessarily to be expected that power is concentrated on the degenerate line. In this paper I address the full bispectrum across the range of $\{k,r,\phi\}$. For numerical study I consider a white noise field and a field near to scale-invariance, and evaluate the three auto-correlations $\av{\tau^3}$, $\av{\tau_S^3}$ and $\av{\tau_T^3}$ representing the bispectra of the magnetic energy density, anisotropic pressure, and gravitational wave source respectively. It is also possible to find exact solutions for the degenerate line of fields far from scale-invariance, and I present the large-scale limits of these. (Interested readers can find the full solutions using the techniques presented here, or else contact the author for details.) These solutions complement two particular solutions presented for the colinear line by CFPR09. Due the complexity of the integration volume, also discussed in that work, I focus otherwise on numerical techniques. In section \ref{MagneticFields} I present a brief overview of the model and in section \ref{GeneralConsiderations} set up the formulae necessary to find the bispectra. Sections \ref{Bulk}, \ref{Colinear} and \ref{Squeezed} consider the general case, the colinear case and the squeezed case respectively and then I present my results in section \ref{Results}. Section \ref{Conclusions} provides a brief conclusion. The appendices contain some additional formulae and some analytical solutions for the degenerate bispectrum on large scales.

\label{Conclusions} This paper presents numerical integrations of the auto-correlations $\BTrTrTr$, $\BTsTsTs$ and $\BTtTtTt$ which are required for an evaluation of the CMB angular bispectrum imprinted by magnetically-induced scalar or tensor modes. On large scales, quantified as lying below a coherence length $k_\mathrm{Coh}=k_\mathrm{Coh}(n_B,r,\phi)$, previous authors have found that the bispectra from ultra-violet fields with $n_B\geq -1$ act as white noise, while those from infra-red fields with $n_B<-1$ scale with $k^{3n_B+3}$, except along the degenerate line where they scale as $q^{2n_B+3}$. This behaviour is confirmed, ultra-violet fields having a coherence length $k_\mathrm{Coh}\sim k_c(\eta)/100$ and infra-red fields having a coherence scale $k\approx k_c(\eta)$, where the values are taken at $r=1$. In the region applicable to CMB studies the bispectra are well-modelled by a 2-dimensional plane $\mathcal{B}_\star(r,\phi)$ sampled at a pivot of $k=k_\star$. The ultra-violet bispectra typically have power up to a large $r$, with $k_\mathrm{Coh}$ decaying with increasing $r$, so if one were to put an $n_B>-1$ bispectrum onto the CMB in principle one should set a pivot on very large scales and allow it to run with $r$. In practice this would complicate the CMB integration significantly and, while it is not clear that the standard formalism directly applies to such fields, care should be taken. For $n_B<-1$ the evolution of $k_\mathrm{Coh}$ is less important since this remains approximately on the order of $k_\mathrm{Coh}\sim\mathcal{O}(1)$ except in regions when the bispectrum is negligible. These 2-dimensional planes are non-trivial but (relatively) quick to evaluate since evaluation is only needed at a constant $k_\star$. The sampling can be made arbitrarily fine as $r\rightarrow 1$ and $\phi\rightarrow\pi$ to capture the most important region, and can be relatively sparse outside of these regions. To constrain magnetic parameters from the CMB parameter estimation is required, which involves many multiple evaluations in a Monte-Carlo chain, so this increase in speed is necessary. There are two possibilities: evaluating $\mathcal{B}_\star(r,\phi)$ as requested for each value of $n_B$, or pre-evaluating a number of planes with discrete choices of $n_B$ and interpolating between them as required. The first case would be the ideal but might unfortunately remain prohibitively slow. The latter option rests on the assumption that the plane $\mathcal{B}_\star(r,\phi)$ evolves smoothly with varying $n_B$. In this case one would imagine sampling choices of $n_B\in(-1,-3)$ and stacking these planes into a cube in $\{r,\phi,n_B\}$-space. This issue is under current study. The CMB signal from the scalar modes is being increasingly studied (see for example SS09, CFPR09, TSS10 and \citet{Cai:2010uw}) and that from the vector auto-correlation has recently also been evaluated \citep{Shiraishi:2010yk,Kahniashvili:2010us}. An immediate consequence of this study is that we can evaluate the CMB angular bispectrum arising from magnetic gravitational waves and we will present the results in a forthcoming work. Necessary extensions to enable a full parameter estimation include evaluating the scalar cross-correlations $\BTrTrTs$, $\BTrTsTs$, $\BTsTsTs$, the scalar/vector and scalar/tensor cross-correlations $\BTrTvTv$, $\BTsTvTv$, $\BTrTtTt$, $\BTsTtTt$ and the vector/tensor cross-correlation $\BTvTtTv$, the colinear lines of all of which were considered in BC05 and B06. In the first instance we would want to select a constant spectral index $n_B\approx -5/2$ and evaluate the bispectra across a section of the coherent range to confirm the nature of the scalings and the behaviour of permutations of the wavenumbers. For more general $n_B$ it should suffice to evaluate instead the plane $\mathcal{B}_\star(r,\phi)$ at $k_\star\approx 10^{-4}$. We may however find for CMB integration that, given the form of the integration \citep{Ferreira:1998kt,Wang:1999vf,Shiraishi:2010sm} swapping back to $\{k,p,q\}$-space is more convenient. This would depend on the balance between the speed-increase from considering only 2D planes and the complications of remapping from the $\{k,r,\phi\}$ coordinate system, incorporating as well the permutations for cross-correlations, and is currently under study. Since BC05 and B06 demonstrate that each of these cross-correlations is non-vanishing and of the same order-of-magnitude along the colinear line each one will, in principle, leave traces on the CMB. Furthermore, the fact that the degenerate line possesses a different sign to the scalar signal strongly suggests that the \emph{total} magnetic signal -- the sum of all possible CMB bispectra, covering auto- and cross-correlations of all the components of the stress tensor -- will exhibit cancellations. In particular, since magnetic bispectra are not positive-definite, contributions from unexpected features such as the ``river'' in $\av{\tau_S^3}$, could introduce cancellations in the integrations particularly for indices further from scale-invariance where the degenerate line isn't so dominant. Perhaps more importantly, the complete set ($\BTrTrTr$, $\BTrTrTs$,\ldots,$\BTvTtTv$,\ldots) contributes to the magnetised CMB angular bispectrum and the nature of these terms -- particularly their sign near to the degenerate line -- is unknown. Whether these effects have a significant impact on the CMB bounds is uknown at present, but conceivably they could significantly tighten (or weaken) them. Until the complete set of intrinsic bispectra are known, and then wrapped onto the CMB with the corresponding transfer functions, CMB constraints on magnetic fields from bispectra should be treated with some caution.

1012.2892

Forthcoming data sets from the Planck experiment and others are in a position to probe the cosmic microwave background (CMB) non-Gaussianity with higher accuracy than has yet been possible, and potentially open a new window into the physics of the very early universe. However, a signal need not necessarily be inflationary in origin, and possible contaminants should be examined in detail. One such is provided by early universe magnetic fields, which can be produced by a variety of models including during an inflationary phase, at phase transitions, or seeded by cosmic defects. Should such fields have been extant in the early universe, they would provide a natural source of CMB non-Gaussianity. Knowledge of the CMB angular bispectrum requires the complete Fourier-space (or "intrinsic") bispectrum. In this paper, I consider in detail the intrinsic bispectra of an early-universe magnetic field for a range of power-law magnetic spectra.

12167664

2011ApJ...728...36R

1012

1012.1296_arXiv.txt

The $^{12}$C-to-$^{13}$C isotopic ratio in the interstellar medium (ISM) is an important diagnostic for probing the history of nucleosynthesis and chemical enrichment in the Galaxy. The abundance of $^{13}$C increases relative to $^{12}$C over Galactic timescales because $^{13}$C is a secondary product of stellar nucleosynthesis. Moreover, millimeter-wave emission studies of CO (Langer \& Penzias 1990), H$_{2}$CO (Henkel et al. 1982), and CN (Savage et al. 2002; Milam et al. 2005) in dense molecular clouds have demonstrated the existence of a Galactic gradient in the $^{12}$C/$^{13}$C ratio, with increasingly lower values found closer to the Galactic center (see also the reviews by Wilson \& Rood 1994 and Wilson 1999). The increase in $^{12}$C/$^{13}$C with Galactocentric distance presumably reflects the reduced rates of star formation, and subsequent stellar nucleosynthesis and mass loss, that are characteristic of the outer Galactic disk. The enrichment process, acting over the 4.6~Gyr since the formation of the solar system, also offers a natural explanation for the reduction in the $^{12}$C/$^{13}$C ratio in the vicinity of the Sun. The terrestrial value of $^{12}$C/$^{13}$C, which should approximate the interstellar value at the time of collapse of the solar nebula, is 89 (Lodders 2003), while values close to 70 are typical of local interstellar clouds. Estimates for the present-day, local ISM value of $^{12}$C/$^{13}$C are usually obtained from absorption or emission studies of carbon-bearing molecules, which serve as proxies when direct measurements of the ambient carbon isotopic ratio are lacking. Wilson (1999), for example, derived an average $^{12}$C/$^{13}$C ratio of $69\pm6$ from millimeter-wave data on CO and H$_{2}$CO emission for 13 local sources. More recently, Milam et al. (2005) added millimeter observations of CN in dense clouds to those of CO and H$_{2}$CO and found an average for the present-day, local $^{12}$C/$^{13}$C ratio of $68\pm15$. In this paper, we shall follow Sheffer et al. (2007) and take the ambient $^{12}$C/$^{13}$C ratio in the solar neighborhood to be $70\pm7$ (see also Sheffer et al. 2002). When molecular proxies are used to evaluate the interstellar abundance ratio of $^{12}$C to $^{13}$C, the values sometimes show evidence of chemical fractionation. While this is generally not the case in dense molecular clouds, where the processes responsible for fractionation are less effective (see Milam et al. 2005), fractionation can have a significant impact in diffuse clouds. Detailed models predict that the fractionation processes in diffuse environments will not affect each interstellar molecule in the same way. The molecular ion CH$^+$ is not subject to chemical fractionation because it must be formed at high effective temperatures. The production of CH$^+$ is linked to the endothermic reaction C$^+$~+~H$_2$~$\to$~CH$^+$~+~H (Elitzur \& Watson 1978, 1980), which has an activation energy of $\Delta E/k_{\mathrm{B}}=4640$ K. Since CH$^+$ is associated with cold diffuse clouds, its formation requires that nonthermal processes, such as magnetohydrodynamic shocks or propagating Alfv\'en waves, provide the additional heating. As a result, the $^{12}$CH$^+$/$^{13}$CH$^+$ ratio is believed to be equilibrated and thus to be the best measure of the ambient carbon isotopic ratio in the diffuse ISM. Many observations (e.g., Centuri\'on \& Vladilo 1991; Crane et al. 1991; Stahl \& Wilson 1992; Centuri\'on et al. 1995) have confirmed this expectation for CH$^+$, revealing $^{12}$CH$^+$/$^{13}$CH$^+$ ratios very near 70 for local diffuse clouds. Centuri\'on et al. (1995) give a weighted mean value of $^{12}$CH$^+$/$^{13}$CH$^+$ of $67\pm3$ for five sight lines to stars within approximately 500 pc of the Sun. However, some authors (e.g., Hawkins \& Jura 1987; Vladilo et al. 1993; Casassus et al. 2005; Stahl et al. 2008) have reported discrepant values of $^{12}$CH$^+$/$^{13}$CH$^+$. Vladilo et al. (1993) found ratios of $126\pm29$ and $98\pm19$ toward HD~152235 and HD~152424, respectively. Since both of these stars belong to the Sco OB1 association (at $d\simeq2$~kpc), Vladilo et al. (1993) suggest that $^{12}$CH$^+$/$^{13}$CH$^+$ may vary on scales of $\sim$1 kpc or smaller, depending on the distance to the clouds. Hawkins \& Jura (1987) found $^{12}$CH$^+$/$^{13}$CH$^+$ ratios of $40\pm9$ and $41\pm9$ toward 20~Tau and 23~Tau, respectively. Because these stars are members of the nearby Pleiades cluster ($d=110$ pc), such small values of $^{12}$CH$^+$/$^{13}$CH$^+$ seem to challenge the view that a ratio near 70 characterizes the solar neighborhood. Carbon-bearing molecules susceptible to fractionation, especially CO, may exhibit $^{12}$C-to-$^{13}$C ratios that are either enhanced or reduced with respect to the ambient value. Two competing processes are capable of altering the $^{12}$CO/$^{13}$CO ratio in diffuse molecular gas. Selective photodissociation (SPD) favors $^{12}$CO since it is the more abundant isotopologue and is thus protected to a greater extent via self shielding (e.g., van Dishoeck \& Black 1988; Visser et al. 2009). The process is effective in the case of CO because the photodissociation of this molecule is governed by line absorption. At lower gas kinetic temperatures, CO is influenced by the isotopic charge exchange (ICE) reaction $^{13}$C$^+$~+~$^{12}$CO~$\to$~$^{12}$C$^+$~+~$^{13}$CO~+~$\Delta E$, where the difference in zero-point energies $\Delta E/k_{\mathrm{B}}=35$~K (Watson et al. 1976) favors $^{13}$CO. Because CO is the most abundant carbon-bearing molecule in the ISM, an enhancement in $^{13}$CO, resulting from isotopic charge exchange, will deplete the carbon reservoir of $^{13}$C. Conversely, if $^{13}$CO is selectively destroyed through photodissociation, then the carbon reservoir will be enhanced in $^{13}$C. Any molecule arising from the remaining carbon in the reservoir should therefore possess a $^{12}$C-to-$^{13}$C ratio that is fractionated in the opposite sense compared to the ratio in CO. Such behavior is expected for CN because this molecule coexists with CO in diffuse molecular clouds (Pan et al. 2005). Existing data on isotopologic ratios in CO, CN, and CH$^+$ for the well-studied sight line to $\zeta$~Oph offer some evidence that the chemical predictions are borne out. The diffuse gas in this direction has $^{12}$CO/$^{13}$CO = $167\pm15$ (Lambert et al. 1994), $^{12}$CN/$^{13}$CN = $47.3^{+5.5}_{-4.4}$ (Crane \& Hegyi 1988), and $^{12}$CH$^+$/$^{13}$CH$^+$ = $67.5\pm4.5$ (Crane et al. 1991). Thus, the ratios in CO and CN are fractionated in opposing directions, while the ratio in CH$^+$ is consistent with the presumed ambient carbon isotopic ratio. In this investigation, we seek to measure $^{12}$CN/$^{13}$CN and $^{12}$CH$^+$/$^{13}$CH$^+$ ratios along lines of sight where $^{12}$CO/$^{13}$CO is either enhanced or reduced. Recent UV surveys of $^{12}$CO/$^{13}$CO along diffuse and translucent sight lines (Sonnentrucker et al. 2007; Burgh et al. 2007; Sheffer et al. 2007) have typically yielded ratios consistent with the average value of $^{12}$C/$^{13}$C for local interstellar clouds. Indeed, the weighted mean value of $^{12}$CO/$^{13}$CO for the 25 sight lines studied by Sheffer et al. (2007) is $70\pm2$. However, there are some notable exceptions. Enhanced ratios are found, not only in the direction of $\zeta$~Oph, but also toward $\rho$~Oph~A and $\chi$~Oph, where the respective $^{12}$CO/$^{13}$CO ratios are $125\pm23$ and $117\pm35$ (Federman et al. 2003), as well as toward $\zeta$~Per, where $^{12}$CO/$^{13}$CO = $108\pm5$ (Sheffer et al. 2007). Reduced ratios, with respect to the ambient value, are found in the directions of 20~Aql, where $^{12}$CO/$^{13}$CO = $50\pm15$ (Hanson et al. 1992), and HD~154368, where $^{12}$CO/$^{13}$CO = $37\pm8$ (Sheffer et al. 2007), among several others (see Sheffer et al. 2007). Sonnentrucker et al. (2007) find a low $^{12}$CO/$^{13}$CO ratio toward HD~73882 ($25\pm22$), but the relative uncertainties are large. Our goal is to obtain $^{12}$CN/$^{13}$CN and $^{12}$CH$^+$/$^{13}$CH$^+$ ratios for as many of these sight lines as possible, so that, when our results are combined with the existing data on $^{12}$CO/$^{13}$CO, the suite of measurements can be used to test the predictions of chemical models for diffuse clouds. We reexamine the sight line to $\zeta$~Oph as a check on our general methodology. We also obtain $^{12}$CH$^+$/$^{13}$CH$^+$ ratios toward the Pleiades stars, 20~Tau and 23~Tau, to investigate potential scatter in the carbon isotopic ratio within the solar neighborhood. Since high-quality data are needed to detect the weak features associated with $^{13}$CN and $^{13}$CH$^+$, the observations examined here also allow precise determinations of rotational excitation temperatures in CN. In some situations, these measurements can then be used to constrain the electron density, an important parameter for modeling the physical conditions within a given cloud (e.g., Black \& van Dishoeck 1991). It is widely understood that the CN molecule is maintained in radiative equilibrium with the cosmic microwave background (CMB) in interstellar space (see the review by Thaddeus 1972). The CMB is the primary source of radiation in the universe at 2.64 mm and 1.32 mm, the wavelengths of the two lowest rotational transitions in CN. As a result, CN excitation temperatures reflect the temperature of the CMB at these wavelengths, in the absence of any local sources of excitation. Observed excitation temperatures in CN, in fact, do exhibit an excess over the temperature of the CMB, as derived from the Far Infrared Absolute Spectrophotometer (FIRAS) onboard the \emph{COBE} satellite ($T_{\mathrm{CMB}}=2.725\pm0.002$ K; Mather et al. 1999; see also Fixsen 2009), but this excess is small ($<0.1$ K; e.g., Palazzi et al. 1992). Since electron impact should dominate any local contribution to CN excitation (Thaddeus 1972), an observed excess provides an estimate for the density of electrons in the portion of the cloud traced by CN. In order to investigate isotopologic ratios in CN and CH$^+$ and CN rotational excitation in diffuse molecular clouds, we examine high-resolution, very high signal-to-noise ratio observations of optical absorption lines arising from electronic transitions within the CN $B$~$^2\Sigma^+-$$X$~$^2\Sigma^+$ and CH$^+$ $A$ $^1\Pi-$$X$ $^1\Sigma^+$ systems. The observations and data reduction procedures are described in \S{} 2. In \S{}~3, we provide detailed information concerning the profile synthesis routine, with which we derive our final column densities and isotopologic ratios. The analysis and discussion of our results on $^{12}$CH$^+$/$^{13}$CH$^+$ and $^{12}$CN/$^{13}$CN ratios appears in \S\S{} 4.1 and 4.2, respectively. In \S{} 4.3, we explore the relationship between $^{12}$CN/$^{13}$CN and $^{12}$CO/$^{13}$CO in an effort to evaluate the effects of chemical fractionation in diffuse molecular gas. The topic of CN rotational excitation is examined in \S{} 5, and our main findings are summarized in \S{} 6. An appendix gives results on weak Ca~{\small I} and Ca~{\small II} absorption toward $\alpha$ Leo and $\alpha$ Vir, which were observed as unreddened standard stars to aid in the reduction of the McDonald Observatory data (see \S{} 2.1).

Isotopologic ratios in interstellar molecules containing carbon are effective probes of the chemical processes active in diffuse environments. In this investigation, we examined optical absorption lines of CN and CH$^+$ along a total of thirteen lines of sight through diffuse molecular clouds. The very high signal-to-noise ratio observations enabled us to extract precise $^{12}$CN/$^{13}$CN and $^{12}$CH$^+$/$^{13}$CH$^+$ ratios, which were used to assess various predictions of diffuse cloud chemistry. Our results on $^{12}$CH$^+$/$^{13}$CH$^+$ confirm that this ratio does not deviate from the ambient $^{12}$C/$^{13}$C ratio in local interstellar clouds, as expected if CH$^+$ is formed via nonthermal processes. Each of the McDonald sight lines in our sample with detectable absorption from $^{13}$CH$^+$ had at one time been found to have either a much lower or a much higher value of $^{12}$CH$^+$/$^{13}$CH$^+$. Such values are not confirmed here. In particular, the $^{12}$CH$^+$/$^{13}$CH$^+$ ratios that we derive toward the Pleiades stars, 20~Tau and 23~Tau, are both indistinguishable from 70, precluding a $^{12}$C/$^{13}$C ratio near 40, as advocated by Hawkins \& Jura (1987), for the gas associated with this nearby cluster. Considering these results, we find no evidence for variation in $^{12}$C/$^{13}$C within the solar neighborhood. In contrast, the studies by Casassus et al. (2005) and Stahl et al. (2008), based on UVES spectra, found variation in $^{12}$CH$^+$/$^{13}$CH$^+$ at about the 16\% level. These authors attributed the variation to intrinsic scatter in the local $^{12}$C/$^{13}$C ratio due to inefficient or incomplete mixing in the interstellar medium. From our analysis of some of the same UVES data, we conclude that the variation in $^{12}$CH$^+$/$^{13}$CH$^+$ is likely related to the difficulties of obtaining accurate $^{12}$CH$^+$/$^{13}$CH$^+$ ratios from these spectra. The major difficulties involve the placement of the rather poorly-defined continua, and the blending of absorption from $^{13}$CH$^+$ with weak $^{12}$CH$^+$ components blueward of the main components. It is probably not a coincidence that the $^{12}$CH$^+$/$^{13}$CH$^+$ ratios derived from our high-resolution McDonald spectra, which have flat, well-behaved continua, along lines of sight with simple CH$^+$ absorption profiles show little variation. While intrinsic scatter in the local $^{12}$C/$^{13}$C ratio may exist at some level, due to small differences in the amount of chemical enrichment and subsequent mixing from one site to another, it is probably too early to say that this variation has been detected. Unlike in the case of CH$^+$, the isotopologic ratios in CN and CO show evidence for significant fractionation away from the ambient value of $^{12}$C/$^{13}$C in some diffuse environments. We attribute these results to the competing influences of isotope-selective photodissociation and isotopic charge exchange reactions, which vary with the strength of the UV radiation field and the physical conditions within a cloud. The results on $^{12}$CN/$^{13}$CN and $^{12}$CO/$^{13}$CO in lines of sight where both ratios have been determined are suggestive of the inverse relationship between these quantities that is anticipated, assuming that CN and CO coexist in diffuse cloud cores. This relationship was first directly suggested by Sheffer et al. (2007), but the isotopologic results for CN that existed at the time were not extensive enough to adequately test the idea. With the results presented here, the situation is now significantly improved, but further observations of both $^{12}$CN/$^{13}$CN and $^{12}$CO/$^{13}$CO would help to clarify the precise relationship between the two quantities. Specifically, measurements of $^{12}$CO/$^{13}$CO toward HD~169454, where $^{12}$CN/$^{13}$CN is unfractionated, and toward HD~170740, where CN is fractionated upward (by a factor of $1.9\pm0.5$ over the ambient value), would prove useful. One aspect of the chemistry of diffuse clouds that future modeling efforts will need to examine is the role played by isotopic charge exchange with CN, which we have found to be important in cold diffuse cloud cores. An ICE reaction involving CN is the most logical explanation for why the enhancements in $^{12}$CN/$^{13}$CN are not as large as expected along lines of sight with low $^{12}$CO/$^{13}$CO ratios, such as HD~73882 and HD~154368. The effects of photochemical fractionation in carbon-bearing molecules are predicted to be the strongest in diffuse molecular clouds because dark clouds are better shielded from UV radiation and lack a significant abundance of carbon ions. This has been convincingly demonstrated by millimeter-wave emission studies of CN, CO, and H$_2$CO (see Milam et al. 2005). However, it is interesting to note that the expression developed in \S{} 4.3, which yields the fraction of carbon locked up in CO, may still apply in denser regions. Keene et al. (1998) detected 809 GHz emission from both $^{12}$C$^0$ and $^{13}$C$^0$ near the center of the Orion Nebula, finding a $^{12}$C$^0$/$^{13}$C$^0$ ratio of $58\pm12$. This was after Boreiko \& Betz (1996) observed 158 $\mu$m emission from both isotopic variants of C$^+$ in this region and found a $^{12}$C$^+$/$^{13}$C$^+$ ratio of $58\pm6$. The $^{12}$C$^{18}$O/$^{13}$C$^{18}$O ratio, also derived by Keene et al. (1998) from observations of the 2$-$1 and 3$-$2 transitions, is $75\pm9$. While all of these ratios are consistent with an ambient $^{12}$C/$^{13}$C ratio of 70, the exact correspondence between the ratios in neutral and ionized carbon suggests that there is a real difference between those ratios and that found in C$^{18}$O. Since these observations sample gas in a well-known photon-dominated region (PDR), it is reasonable to assume that the slight increase in the $^{12}$C$^{18}$O/$^{13}$C$^{18}$O ratio over the ambient value is the result of selective photodissociation. An application of equation (2) to these measurements, with the substitution of $f$($^{13}$C$^+$) for $f$($^{13}$CN), would yield a value of $f$(CO) = 0.76, which is consistent with expectations for gas in a dense molecular cloud like that in Orion. We contrast this result with the values of $f$(CO) we obtain toward the diffuse molecular clouds in Ophiuchus, namely 0.40 toward $\rho$~Oph~A and 0.45 toward $\zeta$~Oph. The Ophiuchus results are appropriate for sight lines that probe the transition region between diffuse and dense clouds, where the conversion of C$^+$ into CO is expected to occur. Beyond isotopologic ratios, the high quality optical spectra examined in this investigation allowed us to obtain precise rotational excitation temperatures for both $^{12}$CN and $^{13}$CN. Our weighted mean value of $T_{01}$($^{12}$CN) = $2.754\pm0.002$ K implies an excess over the temperature of the cosmic microwave background of only $29\pm3$~mK. This excess is considerably smaller than that suggested in the recent survey by S\l{}yk et al. (2008), but is also reduced compared to the excess found by Palazzi et al. (1992). Both of these previous studies indicated the need for an additional excitation mechanism beyond electron and neutral collisions (and the CMB) to account for the observed excitation of CN. The more modest excess found here eliminates the need for such a mechanism in the general ISM. Only three sight lines in our sample exhibit significant excesses in $T_{01}$ over the cosmic microwave background temperature. For the clouds in these directions, we derived upper limits to the electron densities in the range 0.3$-$0.9 cm$^{-3}$. The rotational excitation temperatures observed in $^{13}$CN, from strong detections of the $R$(0) and $R$(1) lines in four directions and the first interstellar detection of the $P$(1) line, show no excess over the temperature of the CMB.

1012.1296

We present very high signal-to-noise ratio absorption-line observations of CN and CH<SUP>+</SUP> along 13 lines of sight through diffuse molecular clouds. The data are examined to extract precise isotopologic ratios of <SUP>12</SUP>CN/<SUP>13</SUP>CN and <SUP>12</SUP>CH<SUP>+</SUP>/<SUP>13</SUP>CH<SUP>+</SUP> in order to assess predictions of diffuse cloud chemistry. Our results on <SUP>12</SUP>CH<SUP>+</SUP>/<SUP>13</SUP>CH<SUP>+</SUP> confirm that this ratio does not deviate from the ambient <SUP>12</SUP>C/<SUP>13</SUP>C ratio in local interstellar clouds, as expected if the formation of CH<SUP>+</SUP> involves nonthermal processes. We find that <SUP>12</SUP>CN/<SUP>13</SUP>CN, however, can be significantly fractionated away from the ambient value. The dispersion in our sample of <SUP>12</SUP>CN/<SUP>13</SUP>CN ratios is similar to that found in recent surveys of <SUP>12</SUP>CO/<SUP>13</SUP>CO. For sight lines where both ratios have been determined, the <SUP>12</SUP>CN/<SUP>13</SUP>CN ratios are generally fractionated in the opposite sense compared to <SUP>12</SUP>CO/<SUP>13</SUP>CO. Chemical fractionation in CO results from competition between selective photodissociation and isotopic charge exchange (ICE). An inverse relationship between <SUP>12</SUP>CN/<SUP>13</SUP>CN and <SUP>12</SUP>CO/<SUP>13</SUP>CO follows from the coexistence of CN and CO in diffuse cloud cores. However, an ICE reaction with CN may mitigate the enhancements in <SUP>12</SUP>CN/<SUP>13</SUP>CN for lines of sight with low <SUP>12</SUP>CO/<SUP>13</SUP>CO ratios. For two sight lines with high values of <SUP>12</SUP>CO/<SUP>13</SUP>CO, our results indicate that about 50% of the carbon is locked up in CO, which is consistent with the notion that these sight lines probe molecular cloud envelopes where the transition from C<SUP>+</SUP> to CO is expected to occur. An analysis of CN rotational excitation yields a weighted mean value for T <SUB>01</SUB>(<SUP>12</SUP>CN) of 2.754 ± 0.002 K, which implies an excess over the temperature of the cosmic microwave background (CMB) of only 29 ± 3 mK. This modest excess eliminates the need for a local excitation mechanism beyond electron and neutral collisions. The rotational excitation temperatures in <SUP>13</SUP>CN show no excess over the temperature of the CMB. <P />Based in part on observations made with the Very Large Telescope of the European Southern Observatory, Paranal, Chile, under programs 065.I-0526, 071.C-0367, 071.C-0513, and 076.C-0431.

4931782

2011ApJ...729..132C

1012

1012.4016_arXiv.txt

\label{sec:intro} Wavefront errors created by imperfections in the optical components (mirrors, lenses, beam splitters, ... etc.) of a high-contrast imaging instrument manifest as a complex pattern of quasi-static intensity variations, or so-called ``speckles", in the focal plane at the science camera. Nominally bright compared to the signal of a substellar companion or debris disk, speckle intensities must first be minimized by hardware, usually with an adaptive optics (AO) system, and then further suppressed through data processing. Since individual speckles have the same intrinsic width as the instrument point-spread-function (PSF), this latter step is crucial for distinguishing between stellar artifacts and true companions. Speckle suppression is accomplished through differential imaging by exploiting differences between the properties of residual starlight that reaches the detector and any incoherent radiation arriving from an off-axis source. To preserve companion flux, some form of ``diversity" must be introduced into the system, whereby a sequence of images is recorded such that individual frames are trivially different from one another. Once measured or calibrated, the image diversity is reversed using software to remove the speckles and reveal previously hidden companions. Image diversity can be achieved in a number of ways. Examples include modulation of companion position with respect to speckles through instrument rotation or field rotation \citep{krist_07,marois_08,marois_10}, modulation of input polarization state \citep{potter_03,oppenheimer_08,hinkley_09}, and modulation of the actual target being observed \citep{serabyn_10,crepp_10,mawet_09}. The level to which residual starlight is removed depends on the modulation timescale compared to the speckle decorrelation timescale. Quasi-static speckles\footnote{Other families of speckles related to the atmosphere and adaptive optics system can decorrelate on a much faster timescale \citep{macintosh_05}.} generally decorrelate in tens or hundreds of seconds as the result of changes in ambient conditions and optical system alignment (e.g., \cite{hinkley_07}). Ideally, science images and PSF reference images are recorded simultaneously to minimize the influence of speckle pattern evolution. An alternative approach to discriminate between companions and starlight is to generate diversity through the wavelength dependence of speckles. In this case, modulation of companion position with respect to speckles is achieved naturally through diffraction by acquiring images simultaneously in multiple different spectral channels. Chromatic differential imaging not only has the benefit of ``freezing" the speckle pattern in time, but it also yields the spectrum of any companions, thus enabling study of their chemical composition and bulk physical properties. First proposed by \cite{sparks_ford_02}, concurrent search and characterization may be accomplished in practice using an integral-field spectrograph (IFS). Compared to other speckle suppression techniques \citep{absil_mawet_09,oppenheimer_hinkley_09}, chromatic differential imaging has a high duty-cycle efficiency and can help maximize instrument scientific return by: accommodating a relatively wide bandpass\footnote{A special case of chromatic differential imaging, called simultaneous differential imaging (SDI), operates over a narrow bandpass and has shown promise for high-contrast applications involving methanated sources as demonstrated by the NICI campaign \citep{biller_10,liu_10}.}; allowing for the inter-changeable usage of science images and PSF references; maintaining an inner-working-angle that is independent of target coordinates relative to the observatory; avoiding smearing of the companion or dust disk PSF during an exposure; and obviating the need to observe a nearby PSF calibration star with similar brightness and colors. Chromatic differential imaging may also be implemented in combination with other techniques. AO-assisted integral field spectroscopy has been used previously to study the companions orbiting GQ Lupi \citep{mcelwain_07}, AB Doradus \citep{thatte_07}, AB Pictoris \citep{bonnefoy_10}, and recently the outer-planets of HR 8799 \citep{janson_10,bowler_10}. In this paper, we present the first on-sky experiments that combine an IFS with a coronagraph \citep{hinkley_10_instrument} and speckle suppression using chromatic differential imaging \citep{sparks_ford_02}. To this, we also implement the locally optimized combination of images (LOCI) algorithm which improves the signal-to-noise ratio of detections \citep{lafreniere_07}. Project 1640 (hereafter, P1640) is a ground-based high-contrast imaging instrument that was recently installed and tested on the Hale 200-inch telescope at Palomar. The hardware incorporates a near-infrared coronagraph and IFS and will soon receive a corrected beam from the PALM-3000 ``extreme" AO system \citep{hinkley_10_instrument,bouchez_09,soummer_05}. P1640 has made several discoveries to date, including Alcor B and Zeta Virginis B \citep{zimmerman_10,hinkley_10}. Though faint compared to their host stars, these companions have masses of $\approx0.25M_{\odot}$ and $\approx0.17M_{\odot}$ respectively, and were noticed in raw data. In the following, we describe our technique for detecting companions having brightness and angular separation that places them beneath the noise floor set by speckles prior to data processing. This paper is accompanied by a companion paper by Pueyo et al. 2011 that discusses our method for subsequently extracting their spectrum. Combined, our results demonstrate the two principal utilities of using an IFS for high-contrast observations. These techniques are relevant to forth-coming instruments with similar designs, including the Gemini Planet Imager \citep{macintosh_GPI_06}, VLT SPHERE \citep{beuzit_06}, and Subaru HiCIAO \citep{mcelwain_10aas}, which will also use an IFS for chromatic differential imaging and companion characterization.

\label{sec:summary} Central to the issue of generating contrast levels sufficient to directly image extrasolar planets is subtraction of (unavoidable) quasi-static speckles that arise from instrument optical aberrations. Along with the ability to characterize companions, an integral field spectrograph naturally provides the capability for suppressing this dominant noise source. The recently commissioned Project 1640 instrument at Palomar is the first to use an IFS for high-contrast imaging in combination with a coronagraph and adaptive optics. We have written a custom program, the Project 1640 speckle suppression pipeline (PSSP), to improve instrument raw sensitivity via post-processing. The program utilizes a version of chromatic differential imaging, or spectral deconvolution, to perform point-spread-function subtraction, as was first proposed by Sparks \& Ford 2002. To this, we also use the multitude of images provided by the IFS to construct PSF references via the locally optimized combination of images (LOCI) algorithm \citep{lafreniere_loci}. We have quantified the degree of spatial correlation between speckle patterns in adjacent P1640 hyper-cube images as a function of color and time. We find that spectral image diversity provides PSF references that match science images at a level comparable to, if not better than, observations of a nearby calibration star, depending on over-head time and system flexure. Analysis related to image scaling factors, throughput, and signal-to-noise ratio well-match theoretical expectations and results from other speckle diversity techniques that also use LOCI, once relevant parameters are optimized. Further, we are able to perform pseudo-real-time data processing while at the observatory using the Bluedot super-computing cluster at NExScI, We have quantified the improvement afforded by the PSSP with three different data sets. Injecting faint companions, recovering known companions, and calculating contrast levels before and after performing speckle suppression each indicate that on-sky sensitivities consistently improve by at least one order of magnitude at angular separations within the AO control region following post-processing. At larger separations, the intensity of residual starlight approaches the limit set by photon noise. Our sensitivity currently reaches $5\sigma$ contrast levels of $\approx2.1\times10^{-5}$ at $1\arcsec$ in the H-band with 20 minutes of integration time. This result is comparable to other experiments at Palomar and is currently limited by calibration of non-common-path errors \citep{burruss_10,crepp_10}. Preliminary tests using an interferometric wave-front calibration unit show promise to reduce these errors by an order of magnitude when operating in tandem with the PALM-3000 ``extreme" AO system, suggesting that contrast levels of order $\approx10^{-7}$ inside of an arcsecond are feasible with integration times of 1-2 hours (Shao et al. 2011, in prep., \cite{bouchez_09}). IFS data provides the ability to play movies by sequentially stepping through each image in the color dimension. We have performed experiments carefully comparing our ability to detect faint companions using these color-movies and the PSSP. We find that movies perform remarkably well, and that the PSSP generally provides an additional factor of 2-3 in effective contrast in comparison. This gain is most noticeable in close proximity to the star, just exterior to the coronagraphic spot, and is largely a result of implementing the LOCI algorithm to improve PSF matching. We play color-movies as part of our standard reduction package to identify companions with brightness comparable to raw speckle noise. Stepping through wavelength channels is likewise helpful with differenced images. It is necessary to calibrate the effects of partial subtraction resulting from LOCI to accurately measure the relative intensity of a faint companion \citep{marois_08,thalmann_09,currie_10,bowler_10,janson_10gj758,marois_10spie}. This is nominally accomplished by injecting artificial companions with known brightness into each data cube. Our photometric analysis indicates that $10\%$ uncertainties are typical for angular separations within the AO control region and that larger uncertainties, by factors of several, may occur very close to the star where only marginally sufficient wavelength diversity is achieved. High-contrast astrometry is likewise challenging. The inclusion of astrometric grid spots aids with locating the position of the occulted star, but aggressive differential imaging techniques modify the shape and flux distribution of the companion PSF, degrading precision. Our preliminary analysis using artificial companions indicates that systematic errors of size $\approx0.5$ spatial pixels are common, even when using companion masks in reference frames. Carefully chosen reference frames can help to mitigate these effects, but 1 mas astrometry using near-Nyquist sampled data ($\sim$2-4 spatial pixels per diffraction width in this case) may prove to be prohibitive unless more advanced algorithms are developed. This paper demonstrates one of the two important benefits provided by an IFS for high-contrast imaging: the automatic generation of useful PSF references for chromatic differential imaging and the detection of faint companions. A subsequent paper by Pueyo et al. 2011 will discuss the challenges and possible solutions for accurately extracting their spectra.

1012.4016

Project 1640 is a high-contrast imaging instrument recently commissioned at the Palomar observatory. A combination of a coronagraph with an integral-field spectrograph (IFS), Project 1640 is designed to detect and characterize extrasolar planets, brown dwarfs, and circumstellar material orbiting nearby stars. In this paper, we present our data processing techniques for improving upon instrument raw sensitivity via the removal of quasi-static speckles. Our approach utilizes the chromatic image diversity provided by the IFS in combination with the locally optimized combination of images algorithm to suppress the intensity of residual contaminating light in close angular proximity to target stars. We describe the Project 1640 speckle suppression pipeline and demonstrate its ability to detect companions with brightness comparable to and below that of initial speckle intensities using on-sky commissioning data. Our preliminary results indicate that suppression factors of at least one order of magnitude are consistently possible, reaching 5σ contrast levels of 2.1 × 10<SUP>-5</SUP> at 1'' in the H band in 20 minutes of on-source integration time when non-common-path errors are reasonably well calibrated. These results suggest that near-infrared contrast levels of order ≈10<SUP>-7</SUP> at subarcsecond separations will soon be possible for Project 1640 and similarly designed instruments that receive a diffraction-limited beam corrected by adaptive optics systems employing deformable mirrors with high actuator density.

2952678

2011MNRAS.413..908C

1012

1012.2336.txt

Accretion disk outflows play a key role in the physics of quasars and their environs. Quasar outflows may be an integral part of the accretion process and the growth of supermassive black holes (SMBHs), by allowing the accreting material to release angular momentum. These outflows may also provide enough kinetic energy feedback to have an effect on star formation in the quasar host galaxies, to aid in `unveiling' dust-enshrouded quasars, and to help distribute metal-rich gas to the intergalactic medium (e.g., \citealt{diMatteo}, \citealt{moll07}). The location and three-dimensional structure of quasar outflows are poorly understood. Sophisticated models have been developed that envision these outflows as arising from a rotating accretion disk, with acceleration to high speeds by radiative and/or magneto-centrifugal forces (\citealt{murray}, \citealt{proga04}, \citealt{proga07}, \citealt{everett}). Improved observational constraints are needed to test these models and allow estimation of mass loss rates, kinetic energy yields, and the role of quasar outflows in feedback to the surrounding environment. In this paper, we focus on the most prominent signatures of accretion disk outflows that appear in quasar spectra, the broad absorption lines (BALs). BALs are defined to have velocity widths $>$2000 \kms\ at absorption depths $>$10\% below the continuum \citep{balnicity}, and they appear in the spectra of $\sim$10-15\% of quasars (\citealt{reichard03a}, \citealt{trump06}). Since only a fraction of quasar spectra display these features, the presence of BALs could represent a phase in the evolution of a quasar and/or particular orientations where the outflow lies between us and the quasar emission sources. Other studies have found support for the latter scenario, given the similarity of various properties between BALQSOs and non-BALQSOs (e.g., \citealt{balnicity}, \citealt{shen08}). Studying the variability in these absorption features can provide important insight into the structure and dynamics of the outflows. For example, we can use information about short-term variability to place constraints on the distance of the absorbing material from the central SMBH. The more quickly the absorption is varying, the closer the absorber is to the central source, based on nominally shorter crossing times for moving clouds (\citealt{fred08} and \S5 below) or the higher densities required for shorter recombination times \citep{hamann97}. Long-term variability measurements provide information on the homogeneity and stability of the outflow. A lack of variability on long time-scales implies a smooth flow with a persistent structure. General variability results provide details on the size, kinematics, and internal makeup of sub-structures within the flows. Variability studies can also address the evolution of these outflows as the absorption lines are seen to come and go, or one type of outflow feature evolves into another (e.g., \citealt{fred08}). A small number of previous studies have investigated variability in BALs, emphasizing \civ\ $\lambda$1549 because of its prominence and ease of measurement. \citet{bar} carried out a study of 23 BALQSOs, covering time-scales of $\Delta$t $\la$1 yr\footnote{Throughout this paper, all time intervals are measured in years in the rest frame of the quasar.}. \citet{L07} used Sloan Digital Sky Survey (SDSS) spectra of 29 quasars to study \civ\ BAL variability on similar time-scales. \citet{gibson08} studied variability on longer time-scales (3-6 years) by combining data for 13 BALQSOs from the Large Bright Quasar Survey (LBQS) and SDSS. \citet{gibson10} reports on variability on multi-month to multi-year time-scales, using 3-4 epochs of data for 9 BALQSOs. They also compare variability in \siiv\ absorption to variability in \civ\ absorption in their sample. In the present study, we go beyond existing work by carrying out a monitoring program of BALQSOs from the sample of \citet{bar}. This strategy provides BAL variability data covering longer time-scales as well as a wider range in time-scales within individual sources. We include archival spectra from the SDSS, when available, to augment the temporal sampling. This paper reports our results for \civ\ \lam1549 variability with measurements selected to enable comparison between short-term (0.35 to 0.75 yr) and long-term (3.8 to 7.7 yr) behavior. We identify trends in the data with velocity and the depth of the absorption. We avoid using equivalent width (EW) measurements, which can minimize a change that occurs in a portion of a much wider trough. In subsequent papers, we will discuss variability properties in the entire data set, including more epochs on individual quasars, sampling in a much shorter time domain, and comparisons between the \civ\ and \siiv\ $\lambda$1400 variabilities to place constraints on the outflow ionizations and column densities. Section 2 below discusses the observations and the quasar sample. Section 3 describes the steps we followed to identify the \civ\ BALs and where they varied. Section 4 describes our results, and Section 5 discusses the implications of these results with comparisons to previous work. \section[]{Data and Quasar Sample} Between 1988 and 1992, \citet{bar} and his collaborators (e.g., \citealt{barlow92}) obtained spectra for 28 BALQSOs with the Lick Observatory 3-m Shane Telescope. They obtained multi-epoch spectra for 23 of these quasars for a study of short-term BAL variability. They selected these objects from already known BALQSOs, with redshifts of 1.2 to 2.9. When selecting objects for their monitoring program, they gave preference to quasars known to have either absorption line or photometric variability. However, in most cases, this information was not available, and they selected more quasars at higher redshifts due to higher detector sensitivity at longer wavelengths. While this sample of BALQSOs is not randomly selected, it does cover a wide range of BAL strengths (see Table 1 and Figure 1). \citet{bar} used the Lick Observatory Kast spectrograph to obtain spectra with settings for high resolution, $R\equiv\lambda/\Delta\lambda\approx 1300$ (230\kms), and moderate resolution, $R\approx 600$ (530 \kms). Most of the objects were observed at both settings. For our analysis of these data discussed below, we used the high-resolution observations for most sources. In some cases, only a moderate-resolution observation is available or choosing the high-resolution observation would compromise wavelength coverage. Since BALs have a width of at least 2000 \kms, a resolution of 530 \kms\ is sufficient to detect changes in their profiles. Starting in January 2007, we have been using the MDM Observatory 2.4-m Hiltner telescope to reobserve the BALQSOs in the sample of \citet{bar}, with the goal of monitoring BAL variability over a wide range of time-scales, from $<$1 month to nearly 8 years. We also supplemented our data set with spectra from the Sloan Digital Sky Survey (SDSS). So far, we have over 120 spectra for 24 BALQSOs, and we continue to collect more data. We note that two of these objects are not strictly BALQSOs because they have a balnicity index of zero (see \S3.2 below). Nonetheless, we include them in our sample because they do have broad absorption features at velocities that we include in our variability analysis (\S3.3). We used the MDM CCDS spectrograph with a 350 groove per mm grating in first order and a $1^{\prime\prime}$ slit to yield a spectral resolution of $R \approx 1200$ (250 km s$^{-1}$), with wavelength coverage of $\sim 1600$ \AA . ÊThe spectrograph was rotated between exposures to maintain approximate alignment of the slit with the parallactic angle, to minimize wavelength-dependent slit losses. The wavelength range for each observation was determined based on the redshift of the target quasar such that the coverage was blue enough to include the Ly$\alpha$ emission and red enough to include the \civ\ emission feature. The coadded exposure times were typically 1-2 hours per source. We reduced the data using standard reduction techniques with the IRAF\footnote{The Image Reduction and Analysis Facility (IRAF) is distributed by the National Optical Astronomy Observatories (NOAO), which is operated by the Association of Universities for Research in Astronomy (AURA), Inc., under cooperative agreement with the National Science Foundation.} software. The data are flux calibrated on a relative scale to provide accurate spectral shapes. Absolute flux calibrations were generally not obtained due to weather or time constraints. The Sloan Digital Sky Survey is an imaging and spectroscopic survey of the sky at optical wavelengths. The data were acquired by a dedicated 2.5-m telescope at Apache Point Observatory in New Mexico. The spectra cover the observer-frame optical and near-infrared, from 3800 to 9200 \AA, and the resolution is $R\approx 2000$ (150 \kms). For spectra, typically three exposures were taken for 15 minutes each, and more exposures were taken in poor conditions to achieve a target signal-to-noise ratio. The multiple observations were then co-added \citep{sloan}. The SDSS includes observations of 11 of the quasars in the Lick sample. However, 3 of the SDSS observations are unusable because the redshifts of those quasars are too low for the \civ\ BALs to appear in the SDSS spectra. Therefore, we have usable SDSS spectra for 8 of the objects in our sample, and they were taken between 2000 and 2006.

We have analyzed short-term and long-term \civ\ $\lambda$1549 BAL variability in a sample of 24 quasars, and we looked for trends with the outflow velocity and absorption strength. We found that 39\% (7/18) of the quasars varied in the short-term, intervals of 0.35 to 0.75 yr, whereas 65\% (15/23) varied in the long-term, intervals of 3.8 to 7.7 yr. Overall, 67\% (16/24) of the quasars varied. We find a trend for variability to occur more often at higher velocities (Figure 3) and in shallower absorption troughs (or shallower portions of absorption troughs; Figures 4 and 5). When looking at the fractional change in strength of the varying absorption features, there is no apparent significant correlation with velocity (Figure 6), but there is a trend for a larger fractional change in absorption strength in shallower features (Figure 7). We do include two objects with BI=0, but the variability in the broad absorption in these two quasars is consistent with the trends described above. \citet{gibson08} studied \civ\ BAL variability on time-scales of $\Delta$t$\sim$3-6 yrs in 13 quasars with two epochs of data and found that 92\% (12/13) varied. This is larger than our overall percentage of 67\%. This is surprising because the selection criteria adopted by \citet{bar} for the sample used here might be biased toward more variable sources (\S2). Evidently, this bias is small or nonexistent. The difference between our result and \citet{gibson08} might be due to the small numbers of objects studied and/or the more conservative approach we took for identifying variable absorption. \citet{L07}, which is a study of BAL variability on time-scales of $<$1 yr, find an inverse correlation of fractional change in equivalent width (EW) with average EW. While EW is a different calculation than our measure of absorption strength, $\langle A\rangle$, encapsulating the entire BAL profile in one number instead of considering separately the different profile regions (\S3.3), this result is consistent with our findings that fractional change in $\langle A\rangle$ correlates with $\langle A\rangle$. \citet{L07} also find no overall trend between fractional change in EW and outflow velocity, which is again consistent with our results. While not plotted here, we do not find a significant correlation between $|\Delta A|$ and velocity or absorption strength, which is consistent with the analysis of \citet{gibson08}. \citet{gibson08} also calculate the change in absorption strength, with values concentrated in the range 0.15-0.25. We find an average value of $|\Delta A|$ for the long-term data of 0.22 $\pm$ 0.10. Both \citet{gibson08} and \citet{L07} comment on the lack of evidence for acceleration in the BAL features they study. We also find no evidence for acceleration in our sample. The BAL variations only involve the absorption trough growing deeper or becoming shallower. We note, however, that limits on the acceleration are difficult to quantify for BALs because, unlike narrower absorption lines, they do not usually have sharp features to provide an accurate velocity reference. BALs also exhibit profile variability, which can shift the line centroid without necessarily corresponding to a real shift in the velocity of material in the outflow. In this work, we compare two different time intervals, and we do not find a significant difference between the short-term and long-term data in the trends described in \S4. However, the incidence of variability and the typical change in strength is greater in the long-term than in the short-term. This indicates that BALs generally experience a gradual change in strength over multi-year time scales, as opposed to many rapid changes on time-scales of less than a year. This will be investigated more in a later paper where we will include all the epochs from our data set. Variability has been investigated in other types of absorbers that are potentially related to BALs. \citet{narayanan} and \citet{wise04} looked for variability ($\Delta$t$\sim$0.3-6 yrs) in low-velocity narrow absorption lines (NALs) and found that 25\% (2/8) and 21\% (4/19), respectively, of the NALs they studied varied. {\citet{paola10} examined variability in mini-BALs, which are absorption lines with widths larger than those of NALs and smaller than those of BALs. They observed \civ\ mini-BALs in 26 quasars on time-scales of $\sim$1 to 3 years and found variability in 57\% of the sources. This number is similar to our result for BALs, however this comparison is complicated by a differing number of epochs and different time baselines. \citet{paola10} found no correlation between incidence of mini-BAL variability and absorption strength or outflow velocity. However, this comparison might be skewed by the fact that mini-BALs tend to be weaker than BALs. For example, none of the mini-BALs studied by \citet{paola10} have depths $A>0.6$, while some of the BAL features in our sample have $A\sim1$. There are also differences in the velocity distributions. The mini-BAL distribution in \citet{paola10} peaks at $\sim$$-$20,000 \kms, whereas the distribution of BAL absorption in our data peaks at $-$15,000 to $-$12,000 \kms (Figure 3). More work is needed, e.g., with larger samples, to control for these differences in the absorption characteristics and compare BAL and mini-BAL variabilities on equal footing. Comparisons like that could provide valuable constraints on the physical similarities and relationship between BAL and mini-BAL outflows. Two possibilities for the cause of BAL variability are the movement of gas across our line of sight and changes in ionization. We found a trend for lines at higher velocity to vary more often, which suggests movements of clouds since faster moving gas might vary more. Moreover, this is a global trend, dependent on velocity in an absolute sense, rather than relative velocity within a given trough. However, it is possible that the root cause of the variability is associated with the strength of the lines, and weaker lines happen to appear at higher velocities (see Figure 2 and \citealt{korista93}). In either case, faster moving material tends to have lower optical depths or cover less of the continuum source. Measured values of $\langle A\rangle$ thus provide either a direct measure of optical depth, or covering fraction in the case of saturated absorption. With the current results, we cannot yet disentangle the trends in velocity and $\langle A\rangle$ with certainty. Evidence from BAL variability studies nonetheless supports an interpretation of varying covering factor. In addition to greater variability at higher velocities, we found changes over small portions of the BAL troughs, rather than entire BALs varying, which can be understood in terms of movements of sub-structures in the flows. Changes in the ionizing flux should globally change the ionization in the flow, and we would expect these changes to be apparent over large portions of the BAL profiles and not in just some small velocity range. \citet{fred08} looked at the variation in the \civ\ and \siiv\ BALs in one quasar and found that the BALs were locked at roughly a constant \siiv/\civ\ strength ratio. They attribute this result to the movements of clouds that are optically thick in both lines. If the level of ionization is constant between observations or if $\tau\gg$ 1, then the changes in absorption strength in our sample are directly indicative of changes in the covering fraction. We will examine the \siiv/\civ\ line ratio in our BAL data in a subsequent paper. On the other hand, studies of NALs support the possibility of changes in ionization causing the line variability. \citet{misawa07} and \citet{hamann10a} observed quasars where multiple NALs varied in concert, which suggests global changes in the ionization of the outflowing gas. If a change in ionization is causing the variability in the BAL outflows, then the optical depth ($\tau$) of the BALs is changing. Features with lower optical depth are much more sensitive to changes in ionization than those with higher optical depth. The outflowing gas is presumably photoionized by the flux from the continuum source, and changes in the continuum flux could cause a change in the ionization of the outflowing gas. However, studies to date have not found a strong correlation between continuum variability and BAL variability (\citealt{L07}, \citealt{gibson08}, \citealt{bar}, \citealt{barlow92}), casting doubt on ionization changes causing BAL variability. Finally, we use the bolometric luminosities of the quasars in our sample to derive characteristic time-scales of the flows for comparison to the observed time-scales. The bolometric luminosities for the quasars in our sample range from $\sim2\times 10^{46}$ to $3\times 10^{47}$ ergs s$^{-1}$, and the average is $\sim7\times 10^{46}$ ergs s$^{-1}$. These values are based on the observed flux at 1450 \AA\ (rest-frame) in absolute flux-calibrated spectra from \citet{bar}, a cosmology with $H_o = 71$ \kms\ Mpc, $\Omega_M = 0.3$, $\Omega_{\Lambda}=0.7$, and a standard bolometric correction factor $L\approx 4.4\lambda L_{\lambda}(1450 {\rm \AA})$. Based on the average value of $L \sim 7\times10^{46}$ ergs s$^{-1}$, a characteristic diameter for the continuum region at 1550 \AA\ is $D_{1550}\sim 0.006$ pc and for the \civ\ BEL region is $D_{\rm CIV}\sim 0.3$ pc, assuming $L = 1/3L_{edd}$ and $M_{BH} = 1.4\times10^{9}M_{\sun}$ (\citealt{peterson04}, \citealt{bentz07}, \citealt{gaskell08}, \citealt{hamann10b}). If the outflowing gas is located just beyond the \civ\ BEL region and it has a transverse speed equal to the Keplerian disk rotation speed, then the gas cloud would cross the entire continuum source in $\sim$1 yr. The typical time-scale in our short-term data is $\sim$0.5 yr and the typical change in absorption strength is at least 10\%, which nominally requires the moving clouds to cross at least 10\% of the continuum source. This implies transverse speeds $>$1000 \kms\ and thus radial distances that are conservatively $<$6 pc if the transverse speeds do not exceed the Keplerian speed. A typical time-scale in the long-term data is $\sim$6 yr, which is enough time to cross the continuum source $\sim$7 times (if the gas is located just beyond the \civ\ BEL region). This suggests a fairly homogeneous flow for at least the cases where there is no long-term variability. To place better constraints on the outflow properties and understand the underlying cause(s) of the line variabilities, our following papers will make more complete use of our entire data sample. We will examine the ratio of \siiv\ to \civ\ absorption to gain better insight into the cause(s) of the variability, and we will utilize more complete time-sampling instead of just looking at two specific ranges of time intervals. We have also obtained more data on BAL variability at the shortest time-scales, $\Delta$t $\sim$ 1 week to 1 month, in order to determine whether there is a minimum time-scale over which variability occurs. This will provide a better constraint on the location of the outflowing gas.

1012.2336

Broad absorption lines (BALs) in quasar spectra identify high-velocity outflows that likely exist in all quasars and could play a major role in feedback to galaxy evolution. The variability of BALs can help us understand the structure, evolution and basic physical properties of the outflows. Here we report on our first results from an ongoing BAL monitoring campaign of a sample of 24 luminous quasars at redshifts 1.2 &lt; z &lt; 2.9, focusing on C IVλ1549 BAL variability in two different time intervals: 4-9 months (short term) and 3.8-7.7 yr (long term) in the quasar rest frame. We find that 39 per cent (7/18) of the quasars varied in the short-term data, whereas 65 per cent (15/23) varied in the long-term data, with a larger typical change in strength in the long-term data. The variability occurs typically in only portions of the BAL troughs. The components at higher outflow velocities are more likely to vary than those at lower velocities, and weaker BALs are more likely to vary than stronger BALs. The fractional change in BAL strength correlates inversely with the strength of the BAL feature, but does not correlate with the outflow velocity. Both the short-term and long-term data indicate the same trends. The observed behaviour is most readily understood as a result of the movement of clouds across the continuum source. If the crossing speeds do not exceed the local Keplerian velocity, then the observed short-term variations imply that the absorbers are &lt;6 pc from the central quasar.

12131646

2010arXiv1012.2861G

1012

1012.2861.txt

With discovery of inflation as solution of the horizon and flatness problems in cosmology \cite{Gu} it became widely accepted that, apart from gravity, some other homogeneous and isotropic fields have to be present at cosmological scale which mimic variable cosmological constant. Traditionally this role is attributed to a single scalar field, the inflaton, or several scalar fields \cite{rews}. Modern theories provide various candidates for inflaton varying from Higgs field of the standard model to more hypothetical moduli fields originating from compactified supergravity/string theory. Still, physical nature of the inflaton is far from being uniquely understood, and the choice of the inflaton potential remains mostly phenomenological. Similarly, popular current models of dark energy \cite{DE} involve scalar fields with rather exotic properties (quintessence, $\kappa$-essence, phantom) whose physical nature is far from clear. Moreover, no massless elementary scalar field was observed experimentally so far. Meanwhile, vector fields certainly do exist (and are masseless before the spontaneous symmetry breaking) being basic ingredients of the Standard Model and its generalization. Therefore an idea to use vector fields instead or together with scalar ones to model inflation and dark energy seems to be appropriate. In fact, the model of inflation driven by vector field was suggested long ago \cite{Ford}, but remained unnoticed by cosmologists until recently when it was revived in the context of the dark energy problem \cite{Arm}. Formation of YM condensates in superdense matter was discussed long ago by Linde \cite{Linde:1979pr}. There are two major reasons why vector fields were not welcome in cosmology, apart from their relative complexity. Spatially homogeneous configuration of a single (Abelian) vector field evidently can not be isotropic. Therefore, in order to fit the Friedmann-Robertson-Walker (FRW) cosmology one has to introduce (at least) a triplet of vector fields ensuring isotopy of the total energy-momentum tensor. Another problem is conformal invariance of the standard classical YM lagrangian, which leads within the FRW cosmology to the equation of state $w=P/\ep=1/3$, identical with that of the photon gas. Therefore, the solution for the coupled EYM system will be the hot Universe driven by the cold classical matter field \cite{Galtsov:1991un}. Meanwhile, for an accelerated expansion one needs the equation of state $-1\le w\le -1/3$, so the conformal invariance must be broken. However, the first problem is avoided in the non-Abelian case: the SU(2) triplet of YM fields admits an essentially non-Abelian configuration (with non-zero commutator of the matrix-valued potentials) whose stress-tensor exhibits three-dimensional homogeneity and isotropy. The second problem (conformal symmetry breaking) can be overcome passing to effective actions which account for quantum corrections either in the context of gauge theories or string theory. Various attempts to use YM vector fields in constructing dark energy models thus were undertaken during recent years \cite{YMDE}.

1012.2861

We discuss cosmological models involving homogeneous and isotropic Yang-Mills (YM) fields. Such models were proposed recently as an alternative to scalar models of cosmic acceleration. There exists a unique SU(2) YM configuration (generalizable to larger gauge groups) whose energy-momentum tensor is homogeneous and isotropic in space. It is parameterized by a single scalar field with a quatric potential. In the case of the closed universe the coupled YM -- doublet Higgs system admits homogeneous and isotropic configurations too. While pure Einstein-Yang-Mills (EYM) cosmology with the standard conformally invariant YM action gives rise to the hot universe, Einstein-Yang-Mills-Higgs (EYMH) cosmology has a variety of regimes which include inflationary stages, bounces, and exhibits global cycling behavior reminiscent of the Multiverse developed in time. We also discuss other mechanisms of conformal symmetry breaking such as string-inspired Born-Infeld (BI) modification of the YM action or field-theoretical quantum corrections.

901399

2011MNRAS.413.1643P

1012

1012.4586.txt

The supernova remnant (SNR) SN~1006 is one of the most interesting objects for studies of Galactic cosmic rays. It is quite symmetrical with a rather simple bilateral morphology in radio \citep[e.g.][]{pet-SN1006mf}, nonthermal X-rays \citep[e.g][]{SN1006Marco} and TeV \g-rays \citep{HESS-SN1006-2010}. Its prominent feature is the positional coincidence of the two bright nonthermal limbs in all these bands, including TeV \g-rays as demonstrated by recent results of \citet{HESS-SN1006-2010}. Current investigations of SNRs with TeV \g-ray emission demonstrate an ambiguity in the explanation of the nature of TeV \g-rays. Namely, the broad-band (radio-to-\g-rays) spectrum of these SNRs can be fitted by assuming the TeV radiation either as leptonic or as hadronic in origin \citep[e.g. RX J1713.7-3946:][]{RX1713aha2006,RX1713Ber-Volk-06}. The question of the origin of TeV \g-rays is closely related to the problem of the presence and the role of non-linear effects of cosmic rays acceleration by the forward shock. One of the key parameter distinguishing between these two possibilities is the strength (and thus the nature) of the post-shock magnetic field. The classical picture considers only the compression of the typical interstellar magnetic field (ISMF) $B\rs{o}\sim 3\un{\mu G}$ to downstream values of the order of tens $\mu$G. Models including non-linear acceleration (NLA) predict that the ISMF is first amplified upstream due to the back reaction of accelerated protons to $B\rs{o}\sim 30\un{\mu}$G and then compressed above hundred $\mu$G. In the former case the inverse-Compton (IC) emission of electrons would be responsible for most of the TeV \g-rays, in the latter case the proton-origin TeV \g-ray radiation is expected to be dominant. The spectrum of SN~1006 may be explained in these two scenarios. One limiting possibility (we call it `extreme NLA model'), namely the case of ISMF amplified and compressed to $B\rs{s}\approx 150\un{\mu G}$ is considered in details by \citet{SN1006Ber-Ksen-Volk-09}. The model successfully fits the broadband nonthermal spectrum from SN~1006 and the sharpest radial profile of the X-ray brightness. TeV \g-rays are shown to be produced in both the inverse-Compton mechanism and the pion-decay one, the latter is dominant. Here, we present a new method to compare models and observations. In particular, we investigate the origin of the patterns of nonthermal images in radio, X-rays and \g-rays. At present time, this can be done only by using the classic MHD and particle acceleration theories. Therefore, the questions behind the present paper are: may a classical model explain the radio-to-TeV-\g-ray observations of SN~1006 and can one put observational constraints on some properties of the particle kinetics and/or on the MF? In this work, we introduce a ``classic'' model describing SN~1006 and compare the spatial distribution of surface brightness derived from the model with those from observations in different wavelength bands. The comparison will allow us to put some constraints on the parameters of the model, thus deriving some hints on the physical mechanisms governing the cosmic rays acceleration in SN~1006. In the following, the section order is determined by the order of parameters determination: the azimuthal and radial profiles in the radio band are analysed in Sect.~\ref{sn1006cp:radio_section}; the variation of the break frequency in Sect.~\ref{EmaxTheta0}; the broadband spectrum of SN~1006 is calculated in Sect.~\ref{SN1006:vol-sp} to check the consistency of our model and to determine the average MF; the X-ray and \g-ray brightness are investigated in Sect.~\ref{sn1006cp:xray_section}. Finally, we draw our conclusions in Sect.~\ref{sn1006cp:conclusion_section}. %-------------------------------------------------------------- \begin{figure*} \centering \includegraphics[width=17truecm]{regions.eps} \caption{NE and SW limbs of SN~1006 in radio at $\lambda\sim20\un{cm}$ (top panels) and X-rays with energy 2-4.5 keV (bottom panels) \citep{pet-SN1006mf,SN1006Marco}. The maximum value of brightness is 100 times the minimum one, in the radio and X-ray images. Radio image is smoothed with Gaussian with $\sigma=0.4'$ to lower fluctuations. Color straight lines mark the regions used for extraction of the radial profiles of brightness; length of regions shown is from $0.8R$ to $1.1R$. Green lines represent X-ray contours, linearly spaced.} \label{SN1006cp:fig-a} \end{figure*} %-------------------------------------------------------------- %%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%

\label{sn1006cp:conclusion_section} The magnetic field strength in SN~1006 is one of the key parameter in the model. Being related to $E\rs{max\|}$ with Eq.~(\ref{sn1006cp:Emaxpar}) and to the parameter $\epsilon\rs{f\|}$ (which regulates efficiency of the radiative losses of relativistic electrons) with Eq.~(\ref{sn1006cp:etaf}), it influences almost everything in nonthermal spectra and images. \begin{table*} \centering \caption{Summary of the observables used for parameter determination and cross-check$^\mathrm{a}$} \begin{tabular}{lll} \hline Observable & Parameter & Value \\ \hline radio azumuthal profile$^\mathrm{b}$ & aspect angle & $\phi\rs{o}=70^\mathrm{o}\pm4.2^\mathrm{o}$\\ &injection type & isotropic \\ &orientation of ISMF and {SNR morphology} & SE-NW, {barrel-like}\\ radio radial profile & $b$ in $K\rs{s}\propto V^{-b}$ & $-1\lsim b\lsim 0$\\ {local broad-band fits of spectra$^\mathrm{c}$}& local index $s\rs{loc}$ over shock & $s\rs{loc}=2.0$ over most of SNR rim\\ $\nu\rs{break}$ azimuthal profile$^\mathrm{c}$ & model of $E\rs{max}$ & time-limited$^\mathrm{d}$\\ & ratio of the mean free path to Larmour radius & $\eta=1.5$\\ &electron maximum energy at parallel shock & $E\rs{max\|}=7.0(B\rs{o}/12\un{\mu G})^{-1/2}\un{TeV}$\\ &electron maximum energy at perpendicular shock & $E\rs{max\bot}=3.25 E\rs{max\|}$\\ radio and hard X-ray spectrum & MF strength {and index $s\rs{tot}$ for the whole SNR}& ($B\rs{o}=25\un{\mu G}$ and $s\rs{tot}=2.0$) or \\ & & ($B\rs{o}=12\un{\mu G}$ and $s\rs{tot}=2.1$)\\ radio and TeV \g-ray spectrum & \g-ray emission model, MF strength {and index $s\rs{tot}$} & IC with $B\rs{o}=12\un{\mu G}$ and $s\rs{tot}=2.1$\\ X-ray radial profiles & post-shock MF strength in the limbs & $B\rs{s\bot}\simeq 50\un{\mu G}$\\ X-ray azimuthal profile & MF strength & OK \\ & ISMF orientation and aspect angle & OK \\ & model of $E\rs{max}$ & OK \\ \g-ray limbs location & MF strength, aspect angle & OK \\ & \g-ray emission model & OK \\ \hline \multicolumn{3}{l}{$^\mathrm{a}$ the model assumes uniform ISMF/ISM and $\gamma=5/3$}\\ \multicolumn{3}{l}{$^\mathrm{b}$ \citet{pet-SN1006mf,Reynoso2010}}\\ \multicolumn{3}{l}{$^\mathrm{c}$ \citet{SN1006Marco}}\\ \multicolumn{3}{l}{$^\mathrm{d}$ see also \citet{Katsuda2010}}\\ \end{tabular} \label{sn1006cp:tab:methods} \end{table*} We consider a `classic' model of SN~1006, i.e. model which is based on classic MHD and acceleration theories. Since they are better developed compared to NLA approach, they allow us to put observational constraints on the (test-particle) kinetics and MF, to compare the azimuthal variations of the electron maximum energy and the surface brightness in radio, hard X-rays and TeV \g-rays. At the present time, such comparison may not be done in the frame of the NLA theory. We demonstrate that the `classic' model is in agreement with most of the observational data. We try to fix free parameters of the model step-by-step, looking for observations which is mostly sensitive to some of them (Table~\ref{sn1006cp:tab:methods}). In addition to the commonly used broad-band spectrum, the properties of the nonthermal (radio, X-ray and TeV \g-ray) images of SNR {as well as spatially resolved spectral fits} are considered. In particular, the morphology and azimuthal profiles of the radio brightness may determine the orientation of ISMF. Namely, the radio data may be fitted by the model with uniform ISMF which is oriented perpendicular to the Galactic plane with an angle $70^\mathrm{o}$ to the line of sight \citep{pet-SN1006mf,Reynoso2010}. If so, the injection efficiency should be independent of obliquity. The radial distribution of the radio brightness depends now only on the way the injection efficiency varies with time ($K\propto V^{-b}$). The observations however may not definitively fix $b$. It is somewhere between $-1$ and $0$ but accuracy of the data allow also for a bit wider range. Spatially resolved X-ray analysis of regions around the forward shock demonstrate that distribution of $\nu\rs{break}$ may be explained by the time-limited model of $E\rs{max}$; this is in agreement with recent results of \citet{Katsuda2010}. The maximum energy of electrons at the parallel shock is found $E\rs{max\|}=7(B\rs{o}/12\un{\mu G})^{-1/2}\un{TeV}$. It is 3.25 times higher at regions where shock is perpendicular. We obtain expressions for the radio, X-ray and \g-ray spectra from the whole SNR in a form which clearly show which parameter of the model is responsible for the amplitude of the spectrum and which one for its shape. The modification factor of the synchrotron X-ray spectrum -- which shows the deviation of the spectrum from the power law -- may well be explained by the classical model with ISMF strength $B\rs{o}=25\un{\mu G}$ if $s=2.0$ or with $B\rs{o}=12\un{\mu G}$ if $s=2.1$. At the same time, the TeV \g-ray modification factor prefers only the pair $B\rs{o}=12\un{\mu G}$, $s=2.1$; TeV emission is then completely due to IC process. In case $B\rs{o}=25\un{\mu G}$, an additional component in the TeV \g-ray spectrum is needed, from pion decays, as it is in the model of \citet{SN1006Ber-Ksen-Volk-09}. The proton injection in such scenario should increase with obliquity in order to fit the observed azimuthal profiles of TeV \g-ray brightness (ISMF is parallel to the limbs). Could the electron and proton injections have so different dependences on obliquity in the same SNR: isotropic for electrons and quasi-perpendicular for protons? The extreme NLA approach \citep{SN1006Ber-Ksen-Volk-09} predicts $B\rs{o}=30\un{\mu G}$ immediately before the forward shock and $B=150\un{\mu G}$ everywhere inside the SNR. A number of the radial profiles of X-ray brightness obtained from XMM image agree with our model if an ambient MF is $B\rs{o}=12\un{\mu G}$. Around the quasi-perpendicular shock, where the profiles are extracted from, our model predicts the post-shock MF with strength $B\rs{s\bot}\simeq 50\un{\mu G}$. However, in the classic model of SN~1006, this is the value immediately post-shock; after then it rapidly decreases downstream. Therefore, an effective (emissivity weighted average) MF within SN~1006 is estimated to be $32\un{\mu G}$ that agrees well with estimates of \cite{Volk-Ber-2008-MFest} and \citet{HESS-SN1006-2010}. MF in the sharpest Chandra profile is fitted in our model with $B\rs{s}=95\un{\mu G}$; it reflects the local conditions, only within this filament. We found that the broad-band spectrum from the whole SN~1006 is better represented with the electron spectral index $s\rs{tot}=2.1$ while local radio-to-X-ray spectra over the SNR shock prefers $s\rs{loc}=2.0$ \citep{SN1006Marco}. {It is interesting, that similar difference in spectral index is found in the theoretical study of \citet{Schure-et-al-2010}: the spectrum near the shock is flatter than the overall spectrum. The authors attributed this difference to the time evolution of $E\rs{max}$ which was lower at previous times. In contrast, the time-limited model for the maximum energy \citep{Reyn-98} which fits the azimuthal variation of $\nu\rs{break}$ in SN~1006 (Sect.~\ref{EmaxTheta0}) as well as absence of the time variation of the synchrotron flux \citep{Katsuda2010}, suggest what $E\rs{max}$ varies quite slowly in this SNR, at least in the regions close to the shock. The issue of different spectral slopes has to be considered in the future. As to the purpose of the present study, } difference between $s\rs{tot}$ and $s\rs{loc}$ is negligible for azimuthal and radial profiles of radio, X-ray and IC \g-ray brightness. Azimuthal profiles of the X-ray and \g-ray brightness in our model behave in the same way as in the observations. `Classic' model has also few difficulties. \citet{Rotetal04} developed a simple geometrical criterion to distinguish between barrel-like and polar-cap morphology in SN~1006. They have shown analytically that the ratio between the central and the rim brightness should be larger than some value in BarMF case (projected ``barrel'' has to provide enough brightness in the internal regions). In XMM map, this ratio is smaller. The luck of brightness from equatorial belt in the central part is an argument against BarMF morphology. Our model, which strongly prefer BarMF, does not agree with the criterion of \citet{Rotetal04}. Nevertheless, the polar-cap scenario, which is in agreement with this criterion (and is adopted by the NLA model), is unable to explain the observed azimuthal profiles of the break frequency $\nu\rs{break}$and the radio brightness, under assumptions that ISMF/ISM are uniform and the amplified/compressed MF increases with obliquity. Another minor point of the classic model of SN~1006 is the rather large ambient MF, $B\rs{o}=12\un{\mu G}$, which is difficult to expect without MF amplification at the high location of SN~1006 above the Galactic plane. Our model deals with ideal gas with the adiabatic index $\gamma=5/3$ and cannot explain the small distance between the forward shock and the contact discontinuity \citep{chr08,SN1006Marco}. Instead, if acceleration is so efficient that relativistic particles affect hydrodynamics then the adiabatic index may be smaller than ours. The small distance observed may naturally be explained by such, more compressible, plasma with the index like $\gamma=1.1$ \citep{orland-et-al-2010}. It is worth noting that in such case of efficient acceleration, there is no need for MF amplification: the only shock compression to factor $\sigma=21$ (as it is for $\gamma=1.1$) may result in quite large downstream MF even in case of $B\rs{o}$ of few $\un{\mu G}$. The two models, classical and extream NLA, are compared in Table~\ref{sn1006cp:tab:ModelsOfSN1006}. It is evident that none of them explaine the whole set of the SN~1006 properties. A new model of SN~1006 has to include either combination of the two extremes or inclusion of the ISMF/ISM nonuniformity. All results presented here are obtained under assumption that SN~1006 evolve in the uniform ISMF and uniform ISM. It is shown that the scenario of classic MHD/acceleration plus uniform ISMF/ISM strongly prefers the barell-like morphology of SN~1006. However, we also see that nonuniform ISMF/ISM could be an essential element in the model of SN~1006. In particular, slanted lobes, the inversion of the brightness ratio between NE and SW limbs from radio to X-ray band and the higher break frequency in NE limb may only be explained by presence of gradient of ISMF and/or ISM. We expect that the effect of the nonuniform ISMF might dominate the role of some nonlinear effects arising from efficient acceleration of cosmic rays by the forward shock in SN~1006. We like to address this issue in the future. \begin{table*} \centering \caption{Comparison of SN~1006 models} \begin{tabular}{l|l|l} \hline Properties of SN~1006&Classic model$^\mathrm{a}$&Extream NLA model$^\mathrm{b}$\\ &(with uniform ISMF) &(with uniform ISMF)\\ \hline two-limbs in radio and hard X-ray image & YES &YES? \\ &\hspace{0.2cm}SN~1006 is barrel &\hspace{0.2cm}SN~1006 has polar caps\\ &\hspace{0.2cm}ISMF direction: SE-NW &\hspace{0.2cm}ISMF direction: SW-NE\\ location of TeV \g-ray limbs &YES &YES?\\ \citet{Rotetal04} criterion &NO &YES\\ radio spectrum &YES as power low &YES with concave shape\\ hard X-ray spectrum &YES with $\left\langle B\right\rangle\rs{ev}=32\un{\mu G}$ &YES with $\left\langle B\right\rangle\rs{v}\approx 150\un{\mu G}$\\ TeV \g-ray spectrum &YES &YES \\ &\hspace{0.2cm}IC with $\left\langle B\right\rangle\rs{ev}\approx 32\un{\mu G}$ &\hspace{0.2cm}IC with $\left\langle B\right\rangle\rs{v}\approx 150\un{\mu G}$ \\ & &\hspace{0.2cm} and hadronic component\\ radio radial profile &YES &NO? (uniform $B$ inside SNR)\\ {sharpest X-ray radial profile} &YES with $B\rs{s}\approx95\un{\mu G}$ &YES with $B\rs{s}\approx 150\un{\mu G}$\\ radio azimuthal profile&YES &?\\ hard X-ray azimuthal profile &YES &?\\ TeV \g-ray azimuthal and radial profiles &YES &? \\ $\nu\rs{break}$ azimuthal profiles &YES &NO? \\ pre-shock MF strength $B\rs{o}=12\un{\mu G}$& NO &YES\\ &\hspace{0.2cm}(if ISMF around SN~1006 &\hspace{0.2cm}as result of amplification\\ &\hspace{0.2cm} is typical $\sim3\un{\mu G})$&\hspace{0.2cm}(if any)\\ (eventual) concave shape of the radio spectrum$^\mathrm{c}$&NO&YES\\ very close forward shock and contact discontinuity&NO&YES (with $\gamma=1.1$)\\ radio `overbrightness' in the SNR interior &NO &NO? \\ slanted lobes &NO &NO\\ ratio of radio ${\cal R}\rs{r}<1$ and X-ray brightness ${\cal R}\rs{x}>1$ &NO &NO\\ $\nu\rs{break,NE}/\nu\rs{break,SW}>1$&NO &NO\\ \hline {$^\mathrm{a}$} present paper\\ {$^\mathrm{b}$} \citet{SN1006Ber-Ksen-Volk-09}\\ {$^\mathrm{c}$} \citet{allen2008} \end{tabular} \label{sn1006cp:tab:ModelsOfSN1006} \end{table*}

1012.4586

Experimental spectra and images of the supernova remnant SN 1006 have been reported for radio, X-ray and TeV gamma-ray bands. Several comparisons between models and observations have been discussed in the literature, showing that the broad-band spectrum from the whole remnant as well as a sharpest radial profile of the X-ray brightness can be both fitted by adopting a model of SN 1006 which strongly depends on the non-linear effects of the accelerated cosmic rays; these models predict post-shock magnetic field (MF) strengths of the order of 150 ?. Here, we present a new way to compare models and observations, in order to put constraints on the physical parameters and mechanisms governing the remnant. In particular, we show that a simple model based on the classic magnetohydrodynamic (MHD) and cosmic rays acceleration theories (hereafter the 'classic' model) allows us to investigate the spatially distributed characteristics of SN 1006 and to put observational constraints on the kinetics and MF. Our method includes modelling and comparison of the azimuthal and radial profiles of the surface brightness in radio, hard X-rays and TeV γ-rays as well as the azimuthal variations of the electron maximum energy. In addition, this simple model also provides good fits to the radio-to-gamma-ray spectrum of SN 1006. We find that our best-fitting model predicts an effective MF strength inside SN 1006 of ?, in good agreement with the 'leptonic' model suggested by the HESS Collaboration. Finally, some difficulties in both the classic and the non-linear models are discussed. Some evidence about non-uniformity of MF around SN 1006 is noted.

12213025

2011mast.conf..133G

1012

1012.2360_arXiv.txt

Magnetic fields are unexpected in hot, massive stars due to the lack of convection in their outer envelopes. However, a small number of massive B stars host strong, organized magnetic fields, such as the chemically peculiar He strong stars like the archetypical star $\sigma$~Ori~E and the recently discovered B2V star HR~7355 (Oksala et al. 2010; Rivinius et al. 2010). These stars are rapidly rotating and host strong magnetic fields that are coupled to a co-rotating magnetosphere (Townsend, Owocki, Groote 2005).

1012.2360

We discuss a recent detection of a strong, organized magnetic field in the bright, broad-line B2V star, HR 5907, using the ESPaDOnS spectropolarimeter on the CFHT as part of the Magnetism in Massive Stars (MiMeS) survey. We find the rotational period of 0.50833 days, making it the fastest-rotating, non-degenerate magnetic star ever detected. Like the previous rapid-rotation record holder HR 7355 (also discovered by MiMeS: Oksala et al., 2010; Rivinius et al., 2010), this star shows emission line variability that is diagnostic of a structured magnetosphere.

12163087

2011A&A...527A...5C

1012

1012.3199_arXiv.txt

Classical novae are cataclysmic stellar events. Their thermonuclear origin, theorized by Schatzmann (1949, 1951) and Cameron (1959) --- see also Gurevitch \& Lebedinsky (1957) and references therein --- has been established through multiwavelength observations and numerical simulations pioneered by Sparks (1969), who performed the first 1-D, hydrodynamic nova simulation. These efforts helped to establish a basic picture, usually referred to as the {\it thermonuclear runaway} model (TNR), which has been successful in reproducing the gross observational properties of novae, namely the peak luminosities achieved, the abundance pattern, and the overall duration of the event; see Starrfield et al. (2008), Jos\'e \& Shore (2008), Jos\'e \& Hernanz (2007) for recent reviews. Many details of the dynamics of nova explosions remain to be explored. In particular, there are many observed cases of nonspherical ejecta, inferred from line profiles during the early stages of the outburst and from imaging of the resolved ejecta, including multiple shells, emission knots, and chemical inhomogeneities. Although the broad phenomenology of the outburst can be captured by 1-D calculations, it is increasingly clear that the full description requires a multidimensional hydrodynamical simulation of such outbursts. To match the energetics, peak luminosities, and the abundance pattern, models of these explosions require mixing of the material accreted from the low-mass stellar companion with the outer layers of the underlying white dwarf. In fact, because of the moderate temperatures achieved during the TNR, a very limited production of elements beyond those from the CNO-cycle is expected (Starrfield et al. 1998, 2009; Jos\'e \& Hernanz 1998; Kovetz \& Prialnik 1997; Yaron et al. 2005), and the specific chemical abundances derived from observations (with a suite of elements ranging from H to Ca) cannot be explained by thermonuclear processing of solar-like material. Mixing at the core-envelope interface represents a likely mechanism. The details of the mixing episodes by which the envelope is enriched in metals have challenged theoreticians for nearly 40 years. Several mechanisms have been proposed, including diffusion-induced mixing (Prialnik \& Kovetz 1984; Kovetz \& Prialnik 1985; Iben et al. 1991, 1992; Fujimoto \& Iben 1992), shear mixing at the disk-envelope interface (Durisen 1977; Kippenhahn \& Thomas 1978; MacDonald 1983; Livio \& Truran 1987; Kutter \& Sparks 1987; Sparks \& Kutter 1987), convective overshoot-induced flame propagation (Woosley 1986), and mixing by gravity wave breaking on the white dwarf surface (Rosner et al. 2001; Alexakis et al. 2004). The multidimensional nature of mixing has been addressed by Glasner \& Livne (1995) and Glasner et al. (1997, 2005, 2007) with 2-D simulations of CO-novae performed with VULCAN, an arbitrarily Lagrangian Eulerian (ALE) hydrocode capable of handling both explicit and implicit steps. They report an effective mixing triggered by Kelvin-Helmholtz instabilities that produced metallicity enhancements to levels in agreement with observations. Similar studies (using the same initial model as Glasner et al. 1997) were conducted by Kercek et al. (1998, 1999) in 2-D and 3-D, respectively. Their results, computed with the Eulerian code PROMETHEUS, displayed mild TNRs with lower peak temperatures and velocities than Glasner et al. (1997) and insufficient mixing. While Glasner et al. (1997) argue that substantial mixing can naturally occur close to peak temperature, when the envelope becomes fully convective and drives a powerful TNR, Kercek et al. (1998) conclude instead that mixing must take place much earlier: if it occurs around peak temperature, it leads to mild explosions or to events that do not resemble a nova. The differences between these studies have been carefully analyzed by Glasner et al. (2005), who conclude that the early stages of the explosion, before TNR ignition when the evolution is quasi-static, are extremely sensitive to the adopted outer boundary conditions. They show that Lagrangian simulations, in which the envelope is allowed to expand and mass is conserved, lead to consistent explosions. In contrast, in Eulerian schemes with a ``free outflow'' outer boundary condition, the choice adopted in Kercek et al. (1998), the outburst can be artificially quenched. The scenario was revisited by Casanova et al. (2010), who show that simulations with an Eulerian scheme --- the FLASH code --- and a proper choice of the outer boundary conditions can produce deep-mixing of the solar-like accreted envelopes with core material. The puzzling results reported by Kercek et al. (1998) stress the need for a systematic evaluation of the effect that different choices of model parameters (e.g. the intensity and location of the initial temperature perturbation, resolution, or size of the computational domain) may have on the results. To this end, we performed a series of 9 numerical simulations in 2-D aimed at testing the influence of these parameters on the level of metal enhancement of the envelope. Here we report the results of these simulations. Our paper is organized as follows. In Sect. 2 we explain our input physics and initial conditions. Then Sect. 3 is devoted to studying the mixing at the core-envelope interface for our fiducial model. In Sect. 4 the effects of the size of the initial perturbation are analyzed, while in Sect. 5 we discuss the effects of the size of the computational domain. In Sect. 6 we quantify the influence of the grid resolution. Finally, in Sect. 7 we discuss the significance of our results and draw our conclusions.

In this paper we have reported results for a series of nine 2-D numerical simulations that test the influence of the initial perturbation (duration, strength, location, and size), the resolution of the grid, and the size of the computational domain on the results. We have shown that mixing at the core-envelope interface proceeds almost independently of the specific choice of such initial parameters, above threshold values. The study confirms that the metallicity enhancement inferred from observations of the ejecta of classical novae can be explained by Kelvin-Helmholtz instabilities, powered by an effective {\it mesoscopic} shearing resulting from the initial buoyancy. Fresh core material is efficiently transported from the outermost layers of the white dwarf core and mixed with the approximately solar composition material of the accreted envelope. As soon as $^{12}$C and $^{16}$O are dredged up, convection sets in and small convective cells appear, accompanied by an increased nuclear energy generation rate. The size of these convective cells increases in time. Eventually, these cells merge into large convective eddies with a size comparable to the envelope height. The range of mean mass-averaged envelope metallicities obtained in our simulations at the time when the convective front hits the outer boundary, $0.21 - 0.29$, matches the values obtained for classical novae hosting CO white dwarfs. It is, however, worth noting that the convective pattern is actually produced by the adopted geometry (e.g., 2-D), forcing the fluid motion to behave very differently than 3-D convection (Shore 2007; Meakin \& Arnett 2007). Nevertheless, the levels of metallicity enhancement found in our 2-D simulations will likely remain unaffected by the limitations imposed by the 2-D geometry (D. Arnett, private communication). Fully 3-D simulations aimed at testing this hypothesis are currently underway.

1012.3199

Context. Classical novae are explosive phenomena that take place in stellar binary systems. They are powered by mass transfer from a low-mass, main sequence star onto a white dwarf. The material piles up under degenerate conditions and a thermonuclear runaway ensues. The energy released by the suite of nuclear processes operating at the envelope heats the material up to peak temperatures of ~(1-4) × 10<SUP>8</SUP> K. During these events, about 10<SUP>-4</SUP>-10<SUP>-5</SUP>M<SUB>⊙</SUB>, enriched in CNO and other intermediate-mass elements, are ejected into the interstellar medium. To account for the gross observational properties of classical novae (in particular, a metallicity enhancement in the ejecta above solar values), numerical models assume mixing between the (solar-like) material transferred from the companion and the outermost layers (CO- or ONe-rich) of the underlying white dwarf. <BR /> Aims: The nature of the mixing mechanism that operates at the core-envelope interface has puzzled stellar modelers for about 40 years. Here we investigate the role of Kelvin-Helmholtz instabilities as a natural mechanism for self-enrichment of the accreted envelope with core material. <BR /> Methods: The feasibility of this mechanism is studied by means of the multidimensional code FLASH. Here, we present a series of 9 numerical simulations perfomed in two dimensions aimed at testing the possible influence of the initial perturbation (duration, strength, location, and size), the resolution adopted, or the size of the computational domain on the results. <BR /> Results: We show that results do not depend substantially on the specific choice of these parameters, demonstrating that Kelvin-Helmholtz instabilities can naturally lead to self-enrichment of the accreted envelope with core material, at levels that agree with observations. <P />Movie is only available in electronic form at <A href="http://www.aanda.org">http://www.aanda.org</A>

12205430

2011Icar..214..179B

1012

1012.3220_arXiv.txt

\label{sect:intro} Collisional evolution of many-body astrophysical systems in which the relative velocity between colliding objects is large compared to the escape speed and coagulation is unimportant is dominated by fragmentation. Examples of such systems include the asteroid belt \citep{DD, Bottke}, the Kuiper Belt \citep{DavisFarinella, PanSari}, and debris disks around young stars \citep{KW,KB}. As the large bodies are slowly ground down, a collisional cascade is launched, in which mass flows unidirectionally to smaller objects until at some scale it is flushed out of the system by some removal process (e.g. by Poynting-Robertson drag, radiation pressure, or gas drag). In real astrophysical systems there is normally a large dynamic range between the scale $m_{inj}$ at which mass is injected into the cascade and the scale $m_{rm}$ at which mass is removed from the system. The mass $m_{inj}$ can be defined as the mass for which the collisional destruction timescale is comparable to the age of the system. In highly evolved systems, such as the Kuiper Belt, this scale corresponds to the characteristic mass of the largest bodies \citep{PanSari, Fraser}. The collision time for the smallest bodies is usually much shorter than for the largest ones, so a steady-state can be set up for $m_{rm}\ll m\ll m_{inj}$. In such a steady-state, the mass distribution of particles evolves on the collision timescale of bodies at the injection mass scale and can be considered static on shorter timescales. \citet{Dohnanyi} was the first to construct a model of a fragmentation cascade that aimed at explaining the mass distribution of objects in the asteroid belt. He assumed that the internal strength of colliding objects is independent of mass with the implication that $m_B(m_t)$, which we define as the mass of the smallest projectile capable of dispersing one half the mass of a target of mass $m_t$ to infinity, has its mass linearly proportional to $m_t$. Dohnanyi also assumed that $m_0(m_t,m_p)$, the mass of the largest fragment produced in a collision between a target of mass $m_t$ and a projectile of mass $m_p$, scales linearly with $m_t$ and is independent of $m_p$. Under these assumptions, he showed that a fragmentation cascade allows a steady-state power law solution \ba n(m)\propto m^{-\alpha}, ~~~\alpha = -11/6 \label{eq:PL} \ea where $n(m)dm$ is the number of objects in the size distribution with mass between $m$ and $m + dm$. Later, \citet{Tanaka} generalized Dohnanyi's result by assuming a self-similar model of fragmentation, which again entails $m_B(m_t)\propto m_t$, but now $m_0(m_t,m_p) \propto m_t q(m_p/m_t)$, where $q$ is an arbitrary function. They confirmed the power law form of the mass spectrum and showed that the value of $\alpha$ is determined by the mass dependence of the collision rate and that $\alpha$ reduces to $11/6$ if the collision rate is proportional to $m_t^{2/3}$ (geometrical cross-section with mass-independent relative velocities). In reality, however, fragmentation does not have to be self-similar because the minimum energy necessary for disruption is not always linearly proportional to the target mass. For instance, if an object's internal strength is dominated by gravity, then $Q_D^\star$, the energy per unit mass required to shatter an object and disperse half of its mass to infinity, scales as $m_t Q_D^\star \propto m_t^{5/3}$, which is the case for objects larger than $\sim 1$ km \citep{BenzAsphaug, Holsapple, BAR}. \citet{OBG} considered a model in which $Q_D^\star$ scales as a power law in $m_t$, but took $m_0$ linearly proportional to $m_t$. Under these assumptions, \citet{OBG} found that a steady-state power law solution for $n(m)$ still exists, but that $\alpha$ differs from $11/6$ unless $Q_D^\star$ is constant, even when the collision rate scales as $m_t^{2/3}$. The steady-state power law solutions of \citet{Dohnanyi}, \citet{Tanaka}, and \citet{OBG} have since been confirmed numerous times by simulations, which have also shed light on non-power law effects present in astrophysical fragmentation cascades. These effects typically manifest themselves as waves superimposed on top of the steady-state power law solution and are caused either by non-collisional mass sinks, e.g. due to the removal of small particles ($\lesssim 1$ $\mu$m for $L=L_\sun$) by radiation pressure \citep{Thebault, CampoBagatin, DD}; a change in the power law index of $Q_D^\star(m_t)$ induced by a transition from the strength-dominated to the gravity-dominated regime \citep{OBG, OBG1}; or a transition from a primordial to a collisionally-evolved size distribution \citep{Fraser, PanSari, KB2004}. In this work, we describe a new source of non-power law behavior in fragmentation cascades. We consider a model similar to the one used by \citet{OBG}, but with a more general form for $m_0$. More specifically, previous researchers \citep{PetitFarinella, OBG1, Williams, KobayashiTanaka, Brunini} have typically assumed that $m_0 = m_t p(E_{coll}/m_tQ_D^\star)$, where $E_{coll}$ is the kinetic energy of the colliding particles in the center of mass frame, and $p$ is a function that varies from author to author. We instead consider the more general dependence $m_0 = m_t^\mu p(E_{coll}/m_tQ_D^\star)$, which is motivated in \S \ref{setup}. We find that unless $\mu=1$, there is no steady-state power law solution. Instead, $n(m)$ is described by a power law with the same power law index as found by \citet{OBG}, but multiplied by a slowly varying function of mass (i.e. $n(m) \propto m^{-\alpha} \varphi(m), ~~~|d \ln \varphi/d \ln m| \ll 1$). The non-power law effects caused by the slowly varying function show up as a smooth deviation from power-law behavior, which is quite different from the wave-like features described earlier. This deviation is significant when extrapolating over many orders of magnitude in mass, as demonstrated in \S\ref{subsect:applications}. In the course of our investigation, we have also discovered that it is possible for waves to appear and persist in a collisional cascade, even if the particle size distribution does not contain an upper or lower mass cutoff, and the strength is described by a pure power law with no breaks. This is a completely independent type of non-power law behavior from the one caused by $m_0$ not proportional to $m_t$. In astrophysical systems, such waves may be triggered by stochastic collision events between large planetesimals \citep{KB2005, WyattDent, Wyatt}. However, the main focus of our paper is on the non-power law behavior that results when $m_0$ is not proportional to $m_t$, and we defer a detailed discussion of these waves for the future. The paper is organized as follows. In \S\ref{setup}, we introduce the general equations describing fragmentation and discuss specific assumptions relevant to our model. In \S \ref{pls} we demonstrate that pure power law solutions for a fragmentation cascade are indeed possible if $m_0 \propto m_t p(E_{coll}/m_tQ_D^\star)$, consistent with \citet{KobayashiTanaka}. We show in \S \ref{nonpl} that if the assumption $m_0 \propto m_t$ is not satisfied, then solutions are given by the product of a power law with the same index as in the $m_0 \propto m_t$ case, and a slowly varying function of mass. We then find analytic solutions for this slowly varying function for monodisperse (all fragments having the same mass) and power law fragment mass distributions. We find that in the monodisperse case, the solutions for the steady-state distribution are not unique and can support waves. We confirm our analytical results numerically in \S \ref{numerical} and discuss their validity and possible applications in \S \ref{discussion}.

\label{discussion} \subsection{Validity of $|d\ln \varphi/d\ln m| \ll 1$} \label{subsect:validity} In deriving our analytical results, we have assumed that $|d\ln \varphi/d\ln m| \ll 1$ (Eq. \ref{eq:weak_var}). Results from the previous section demonstrate that the qualitative behavior of $\varphi(m)$ is insensitive to the specific form of $m_B$. Thus we can get a sense of when this assumption is valid by using the form of $d\ln\varphi/d\ln m$ for the monodisperse case with $m_B = m$: \ba \frac{d\ln\varphi(m)}{d\ln m}= -\frac{1}{2\ln(m/m_0^\star)}, \label{eq:verif} \ea In order to have $|d \ln \varphi/d \ln m| \ll 1$, we must have $|\ln(m/m_0^\star)| \gg 1$, where $m_0^\star$ was defined in \S\ref{sec:mbm_pow}. In practice, we find that even for $|\ln(m/m_0^\star)| \sim 4$ our analytical solutions give an accurate description of the non-power law behavior. For instance, in Fig. \ref{powlawfig}b it is clear that for the monodisperse fragmentation law, the exact solution for $\varphi(m)$ (Eq. (\ref{phimbm})) works very well, even though we have $|\ln(m/m_0^\star)| = 4.6$ at $m = 10^{-18}$. \subsection{Comparison with existing studies} \label{subsect:compare} We illustrate how our work fits into existing studies of fragmentation cascades with a parameter space plot in $\mu-\beta$ coordinates. Figure \ref{parameterspace} shows the domains of applicability in the $\mu - \beta$ plane for the Dohnanyi and OBG solutions in relation to our analytic solutions for the monodisperse case. Each of our solutions is labeled by its corresponding formula number, and it is evident that our investigation covers the remainder of the $\mu-\beta$ plane. \begin{figure}[!h] \centering \includegraphics[width=.49\textwidth]{parameterspace.eps} \caption{Parameter space plot in the $\mu-\beta$ plane. The case considered by \citet{Dohnanyi,Tanaka} is at the point $\mu = 1, \beta=1$, and the case considered by \citet{OBG} lies on the line $\mu = 1$ (solid line). Our solutions cover the rest of phase space and are labeled with references to corresponding equations in this work. Thus, solutions (\ref{eq:varphi_sol_2_norm}) and (\ref{eq:varphi_sol_1_norm}) (monodisperse, $m_B(m) = Bm$) lie on rays $\mu < 1, \beta=1$ and $\mu > 1, \beta=1$, correspondingly. Solutions (\ref{eq:varphi_sol_4_md}) and (\ref{eq:varphi_sol_3_md}) (monodisperse, $m_B(m) = Bm^\beta$) are valid in the half planes $\beta > 1$ (white) and $\beta < 1$ (gray) respectively. For these solutions $\varphi =$const along the line $\mu = 1$ in agreement with \citet{OBG}.} \label{parameterspace} \end{figure} We next discuss why previous authors have not seen non-power law behavior. The reason is that they have all assumed $m_{0,B} \propto m_t$, and we have shown in \S\ref{subsect:failure} that non-power law behavior only results when $m_{0,B}$ is not proportional to $m_t$. In some studies, the assumption $m_{0,B} \propto m_t$ was explicit such as in \citet{Dohnanyi} and \citet{OBG} who both assumed $m_0 = Cm_t$, and in \citet{PetitFarinella}, \citet{OBG1}, and \citet{Brunini} who all assumed\footnote{Note that $m_{0,B} = m_t/2$ does not follow from $m_{rem}(m_B(m_t),m_t) = m_t/2$, because the remnant does not belong to the distribution of ejecta (\S\ref{setup}).} $m_{0,B} = m_t/2$. In other cases, such as \citet{Tanaka} and \citet{KobayashiTanaka} the scaling $m_{0,B} \propto m_t$ was implicit in assumptions about the form of $f_{ej}$ (\S\ref{subsect:Dohnanyi},\S\ref{subsect:OBG}). We now return to our argument from \S\ref{subsect:m0_mt} that $m_{0,B}$ should have the form (\ref{eq:m0B}). This conclusion was based on an extension of the experimental results of \citet{Fujiwara} beyond the strength-dominated regime. We mention that a number of authors \citep{PetitFarinella, OBG1, Brunini} have considered a different extension of those results, and instead of our Eq. (\ref{eq:Fujiw}), these authors used \ba \frac{m_0(m_t,m_p)}{m_t} &\propto& \left(\frac{E_{coll}(m_t,m_p)/2}{Q_S(m_t) m_t}\right)^{-\gamma}, \label{Ereleq} \ea Here, $Q_S$ is the energy per unit mass required to shatter an object, but not necessarily to disperse its fragments to infinity \citep{OBG1}. If we assume for simplicity that $Q_D^\star \propto Q_S$, then from Eq. (\ref{Ereleq}) and Eq. (\ref{eq:Ecrit}), we have $m_{0,B} \propto m_t$, for any functional form of $m_B$. The reason it is possible to derive two different forms for $m_0(m_t,m_p)$ from the results of \citet{Fujiwara} (Eq. (\ref{eq:Fujiw}) and Eq. (\ref{Ereleq})) is because their experiments were performed over a small range of target masses using a constant projectile mass. The question of how to properly extend their results over a larger mass range can best be settled by more experiments and simulations \citep{Leinhardt,BenzAsphaug,BAR}, which can decisively answer how $m_{0,B}$ varies with mass. However, we point out that unless $m_{0,B}$ is exactly proportional to $m_t$, non-power law behavior will result. As we show in the next section, these deviations from power law behavior can be observationally significant when extrapolating over many orders of magnitude in mass. \subsection{Applications} \label{subsect:applications} Our results clearly demonstrate that one should be somewhat cautious when adopting a pure power law approximation to describe the properties of fragmentation cascades. Even though the non-power law corrections computed in this work scale very weakly with object mass (as the square root of the logarithm of the mass (Eq. (\ref{phimbm}))), one has to keep in mind that in astrophysical systems collisional cascades span $\sim 30$ orders of magnitude in mass. Thus, even a weak deviation from a power law can become important, such as when inferring the disk mass (dominated by the largest bodies) from its infrared luminosity (dominated by the smallest bodies) \citep{Wyatt}. Just for illustration, let us consider a population of $R_{inj}=10$ km objects which get ground down to $R_{rm}=$1 $\mu$m size particles by collisions. Infrared observations give us some idea of the mass in small particles, thus fixing the normalization of the mass spectrum at its low-mass end, and we want to infer from these data the total mass in large bodies feeding this collisional cascade. Connecting the mass contained at the low and high mass ends of the spectrum by a simple power law leads to an error caused by the neglect of the non-power law effects. We can estimate this error $\delta$ by using Eq. (\ref{phimbm}) and taking the ratio of $\varphi$ at the high and low mass ends. Assuming some values of $\mu$ and $C$ in Eq. (\ref{eq:m0B}) that are ``averaged'' over the whole cascade (in practice these parameters will change several times between $R_{rm}$ and $R_{inj}$ because of variations in the internal properties of objects), we have \ba \delta\approx \sqrt{\frac{\ln (R_{rm}/R_0^\star)} {\ln (R_{inj}/R_0^\star)}}=\sqrt{1+\frac{\ln (R_{inj}/R_{rm})} {\ln (R_0^\star/R_{inj})}}, \ea where $R_0^\star$ is the radius of the object with mass $m_0^\star$ defined in \S\ref{subsect:mB=m}. Assuming for illustration that on ``average'' $\mu\approx 1.1$ and that at the high mass end the largest fragments produced in collisions have mass equal to $0.3$ of the target mass ($m_{0,B}(m_{inj})=0.3m_{inj}$) one finds $\ln(R_0^\star/R_{inj})=\ln(m_{0,B}(m_{inj})/m_{inj})/3(\mu-1) \approx 4$ and $\delta\approx 2-3$. Thus, in this particular exercise the neglect of non-power law effects leads to an {\it underestimate} of the mass in large bodies by a factor of several. This also implies that the total disk mass, which is dominated by the mass in large bodies, is underestimated. An underestimate of the disk mass also leads to an underestimate of the disk lifetime, $M_{disk}/\dot{M}_{disk}$. This occurs because if the disk is in steady-state, then $F(m)$ is independent of $m$, which means it is possible to infer $\dot{M}_{disk}$ from infrared observations alone, without extrapolation to large masses \citep{Wyatt}. Supposing that we had correctly inferred $\dot{M}_{disk}$ from observations, but had failed to apply the non-power law correction, and hence underestimated the disk mass, then we would also have underestimated the disk lifetime. Based on the above discussion, breaking the assumption $m_{0,B} \propto m_t$ affects the calculation of disk properties from observations. Conversely, observations of disks can be used to constrain the model parameters (i.e. $\mu$ and $C$ if $m_{0,B}$ is given by Eq. (\ref{eq:m0B})), if e.g. the inferred disk mass is found to be unreasonable for some parameter range. However, as mentioned in \S\ref{subsect:compare}, direct application of our theoretical results to the observed mass spectrum of objects is complicated by the multitude of additional factors playing an important role in real astrophysical systems. Nevertheless, modern calculations \citep{Brunini, OBG1} of the collisional evolution in the asteroid belt and of debris disks \citep{Thebault, Krivov} aim for a precision of tens of percent or less over a broad range of masses. At this level of accuracy, the non-power law effects considered in this work would play a significant role and should be taken into account.

1012.3220

Collisions resulting in fragmentation are important in shaping the mass spectrum of minor bodies in the asteroid belt, the Kuiper Belt, and debris disks. Models of fragmentation cascades typically find that in steady-state, the solution for the particle mass distribution is a power law in the mass. However, previous studies have typically assumed that the mass of the largest fragment produced in a collision with just enough energy to shatter the target and disperse half its mass to infinity is directly proportional to the target mass. We show that if this assumption is not satisfied, then the power law solution for the steady-state particle mass distribution is modified by a multiplicative factor, which is a slowly varying function of the mass. We derive analytic solutions for this correction factor and confirm our results numerically. We find that this correction factor proves important when extrapolating over many orders of magnitude in mass, such as when inferring the number of large objects in a system based on infrared observations. In the course of our work, we have also discovered an unrelated type of non-power law behavior: waves can persist in the mass distribution of objects even in the absence of upper or lower cutoffs to the mass distribution or breaks in the strength law.

12163609

2011A&A...534A..45S

1012

1012.1225_arXiv.txt

As technical possibilities increase, more advanced and complicated instruments are being designed for observational studies of the Sun. New instrumentation is generally designed to offer better spatial resolution, better spectral coverage, and full polarization information. Higher spatial resolution will help resolve, e.g., small-scale magnetic fields in the quiet Sun. In sunspots, the magnetic fields and flows in penumbra, lightbridges, and umbral dots make interesting targets. Better spectral coverage and polarization information are needed to enable meaningful interpretation of the observed features. Adaptive optics (AO) are currently available at the major ground-based solar observatories and have greatly improved the performance of telescopes. Additionally, postprocessing of high-cadence imaging data to correct the blurring effect of seeing has made a huge impact on image quality. Image restoration methods such as MOMFBD \citep[Multi-Object Multi-Frame Blind Deconvolution,][]{vannoort:2005,lofdahl:2002} and Speckle interferometry \citep[see][and references therein]{vonderluhe:1993} have become essential for ground-based, high spatial-resolution solar imaging. Because these methods allow combining all the images obtained within a dataset, while retaining high image quality, it removes the need for frame selection and results in more reliable and higher signal-to-noise data. In this context, Fabry-P\'{e}rot interferometer (FPI) based instruments have become quite popular and are in operation at many solar observatories, e.g.\ CRISP \citep{scharmer:2008} at the Swedish 1-m Solar Telescope (SST), the Telecentric etalon solar spectrometer \citep[TESOS,][]{kantischer:1998}, and G\"ottingen Fabry-P\'erot spectropolarimeter \citep{Puschmann:2006} at the Vacuum Tower Telescope and the Interferometric bidirectional spectrometer \citep[IBIS,][]{cavallini:2006} at the Dunn Solar Telescope. Similar instruments are likely to be developed for the future 4-m class solar telescopes ATST and EST. The main advantages of FPIs over Lyot filters are that they have a high transmission, allow for rapid wavelength tuning, and allow for a dual-beam polarimetric setup with a polarizing beam splitter close to the final focal plane. A disadvantage is that even very well-made FPIs will show variations in the width and central wavelength of the transmission peak over the field of view \citep[FOV, e.g.][]{ibis:2008}. When the FPIs are set up in a telecentric mount, as with CRISP, these variations cause intensity variations on relatively small spatial scales. In collimated mount, such small-scale features disappear, but the central wavelength shifts systematically away from the center of the image. For further comments on FPIs in telecentric vs.\ collimated mount see \citet{scharmer_crisp:2006}. We plan to discuss a theoretical approach to the processing of full-Stokes line scans from FPI-based instruments in a future paper, but here we present the practical application of that approach to real CRISP data of a sunspot observed in the 6302.5 \AA\ \ion{Fe}{i} line, including all the necessary data reduction steps, such as correction for telescope polarization, flatfielding, image restoration, and demodulation. In Sect.\ \ref{sect:pol_acq} we discuss the general setup of the data acquisition procedures, and in Sect.~\ref{sect:instr_calib} the challenges and possibilities for calibration procedures. The proposed calibration procedure is presented in Sect.~\ref{sect:cal_scheme}, results of the procedure applied to CRISP data are presented in Sect.~\ref{sect:results}, and the conclusions in Sect.~\ref{sect:conclusions}.

\label{sect:conclusions} We have presented a method for reducing full Stokes data obtained with CRISP at the SST, an instrument with a fast-tuning Fabry-P\'erot filter in a telecentric mount. Although some calibration issues are inherent to such an instrument used together with image restoration, most problems can be overcome by using the proper calibration procedures. The impact on the flatfields of wavelength shifts over the FOV, owing to cavity errors, can be corrected for. This is important for preventing changes in the line shape caused by errors in the gain tables. For instruments with the FPI in a collimated mount, the cavity errors occur in pupil space, and are therefore not a problem in this respect. These instruments do show radial wavelength shifts over the FOV that have to be accounted for \citep[][]{janssen:2006}. Without a proper calibration scheme, the polarimetric imprint of the flatfielding can cause strong saddle-shaped offsets, fringes, and sharp lines from deviating pixel rows/columns that are enhanced by the deconvolution in the image restoration. By using the proposed calibration scheme these artifacts can almost be completely removed. Any remaining weak artifacts are most likely due to discrepancies in the SST telescope model, which is only accurate to a few tenths of a percent, and nonzero intrinsic polarization. The severity of the artifacts depends on the (variation in the) telescope polarization and the variations in the modulation matrix elements over the FOV. For telescopes that have low and/or constant intrinsic polarization and instrumentation whose modulation matrix elements are nearly constant over the FOV, such flatfielding problems will be much less severe. Velocities, as determined from inversions of the data, can be corrected for cavity errors using a convolved version of the cavity map. However, the PSFs will vary from wavelength to wavelength resulting in a non equidistant wavelength sampling. Both the wavelength sampling and the variations in the width of the FPI transmission peak will vary from pixel to pixel, but can in principle be taken into account during the inversions. An intrinsic limitation is the use of non monochromatic input images for the image restoration, which results in deconvolution artifacts: mixing of information from different wavelengths through the image restoration procedure. Modern developments in FPI instrumentation, such as rapid (kHz) modulation combined with charge shuffling (which may allow for demodulation before image restoration), could reduce calibration problems. However, to reach the polarimetric and spectroscopic accuracy that the future large solar observatories EST and ATST are aiming for with the next generation of Fabry-P\'erot systems, advanced calibration techniques and optimized optical designs to deal with the issues discussed here will be required.

1012.1225

Context. The combination of image restoration and a Fabry-Pérot interferometer (FPI) based instrument in solar observations results in specific calibration issues. FPIs generally show variations over the field-of-view, while in the image restoration process, the 1-to-1 relation between pixel space and image space is lost, thus complicating any correcting for such variations. <BR /> Aims: We develop a data reduction method that takes these issues into account and minimizes the resulting errors. <BR /> Methods: By accounting for the time variations in the telescope's Mueller matrix and using separate calibration data optimized for the wavefront sensing in the MOMFBD image restoration process and for the final deconvolution of the data, we have removed most of the calibration artifacts from the resulting data. <BR /> Results: Using this method to reduce full Stokes data from CRISP at the SST, we find that it drastically reduces the instrumental and image restoration artifacts resulting from cavity errors, reflectivity variations, and the polarization dependence of flatfields. The results allow for useful scientific interpretation. Inversions of restored data from the δ sunspot AR11029 using the Nicole inversion code, reveal strong (~10 km s<SUP>-1</SUP>) downflows near the disk center side of the umbra. <BR /> Conclusions: The use of image restoration in combination with an FPI-based instrument leads to complications in the calibrations and intrinsic limitations to the accuracy that can be achieved. We find that for CRISP, the resulting errors can be kept mostly below the polarimetric accuracy of ~10<SUP>-3</SUP>. Similar instruments aiming for higher polarimetric and high spectroscopic accuracy, will, however, need to take these problems into account.

21032487

2011NCimC..34c..27D

1012

1012.4400_arXiv.txt

The ARGO-YBJ experiment is in stable data taking since November 2007 with a duty cycle $>$85\% and with excellent performance. The large field of view ($\sim$2 sr) and the high duty cycle allow a continuous monitoring of the sky in the declination band from -10$^{\circ}$ to 70$^{\circ}$. Several interesting results are available in cosmic ray (CR) physics as well as in gamma ray astronomy with the first 3 years of operation. We observed TeV emission from 3 sources with a significance greater than 5 standard deviations (s.d.): Crab Nebula, Mrk421 and MGRO J1908+06. In particular we observed the Crab Nebula with a significance of 14 s.d. in $\sim$800 days. A detailed long-term monitoring of the Mrk421 flaring activity has been carried out in 2008-2009. A clear correlation with X-ray data was found. The relation between spectral index and flux has been studied in one single flare (June 2008) \cite{aielli10} and averaging the data over 2 years. Working in scaler mode \cite{aielli08} ARGO-YBJ has performed a search for emission from GRBs in coincidence with 93 events observed by satellites (16 with known redshift), setting upper limits on the fluence between 10$^{-5}$ and 10$^{-3}$ erg cm$^{-2}$ in the 1 - 100 GeV energy range \cite{aielli09a}. A medium-scale CR anisotropy has been observed with a significance greater than 10 s.d. at proton median energy of about 2 TeV. Two excesses (corresponding to a flux increase of $\sim$0.1\%) are observed by ARGO-YBJ around the positions $\alpha\sim$ 120$^{\circ}$, $\delta\sim$ 40$^{\circ}$ and $\alpha\sim$ 60$^{\circ}$, $\delta\sim$ -5$^{\circ}$ \cite{vernetto09}, in agreement with a similar detection reported by the Milagro Collaboration \cite{mil08}. The origin of this medium-scale anisotropy is puzzling. In fact, these regions have been interpreted as excesses of hadronic CRs, but TeV CRs are expected to be randomized by the magnetic fields. Understanding these anisotropies should be a high priority as they are probably due to a nearby source of CRs, as suggested by some authors (see, e.g., \cite{markov}). With 2 years of data we carried out a 2D measurement of the CR large-scale anisotropy to investigate detailed structural information beyond the simple Right Ascension profiles. The p-air cross section has been measured in the range 1 - 100 TeV and the corresponding p-p cross section inferred \cite{aielli09b}. The light-component (p+He) spectrum of primary CRs has been measured in the range 5 - 250 TeV. The preliminary results are in good agreement with the CREAM balloon data. For the first time direct and ground-based measurements overlap for a wide energy range thus making possible the cross-calibration of the different experimental techniques. A measurement of the $\overline{p}/p$ ratio at few-TeV energies has been performed setting two upper limits at the 90\% confidence level: 5\% at 1.4 TeV and 6\% at 5 TeV \cite{cris10}. In the few-TeV range these results are the lowest available, useful to constrain models for antiproton production in antimatter domains. The main results after about 3 years of stable operation are summarized in \cite{vulcano10}. In this paper gamma-ray astronomy results are summarized.

1012.4400

The ARGO-YBJ experiment is a multipurpose detector exploiting the full coverage approach at very high altitude. The apparatus, in stable data taking since November 2007 with an energy threshold of a few hundreds of GeV and a duty-cycle of about 90 %, is located at the YangBaJing Cosmic Ray Laboratory (Tibet, P.R. China, 4300 m a.s.l., 606 g/cm2). A number of interesting results are available in Cosmic Ray Physics and in Gamma Ray Astronomy after the first 3 years of stable data taking. In this paper Gamma-Ray Astronomy results are summarized.

12207297

2011JCAP...02..015B

1012

1012.2652.txt

Cosmological surveys in general and the cosmic microwave background (CMB) in particular are naturally constructed on our celestial sphere. Because of the statistical isotropy of such observations, cosmological statistical properties, such as the angular power spectra or the bispectra, are better captured in reciprocal space, that is in harmonic space. In general however, most of the physical mechanisms at play take place at small scale and are therefore not expected to affect the whole sky properties. For instance, the physics of the CMB is, to a large extent, determined by sub-Hubble interactions corrsponding to sub-degree scale on our observed sky. Decomposition in spherical harmonics, while introducing a lot of complication of the calculations, does not carry much physical insight into these mechanisms so rather blurs the physics at play. In this respect, a flat-sky approximation, in which the sky is approximated by a 2-dimensional plane tangential to the celestial sphere, hence allowing the use of simple Cartesian Fourier transforms, drastically simplifies CMB computations. Such an approximation is intuitively expected to be accurate at small scales. So far this approximation is mostly based on an heuristic correspondence between the two sets of harmonic basis (spherical and Euclidean) which can be summarized for a scalar valued observable by~\cite{1997PhRvD..55.1830Z,2000PhRvD..62d3007H} \begin{equation}\label{corres} \Theta(\vn) = \sum a^{\Theta}_{\ell m}Y_\ell^m(\hat \vn)\to \int \dd^2 {\bf l}\Theta({\bf l})e^{\ii {\bf l}.{\bf \theta}}\,. \end{equation} In the context of CMB computation, the relations between the flat-sky and the full-sky expansions have been obtained at leading order in Ref.~\cite{2000PhRvD..62d3007H}. However, its validity for the angular power spectrum and the bispectrum is not yet understood in the general case and the expression and order of magnitude of the next to leading order terms are still to be computed. The goal of this article is to provide such a systematic construction. In particular, we will show that there exists a two-parameter family of flat-sky approximations for which well-controlled expansions can be built. That allows us to discuss in details their accuracy by performing the computation up to next-to-leading order. In the new route we propose here, we derive the flat-sky expansion directly on the 2-point angular correlation function, instead of relying on the correspondence~(\ref{corres}). One of the technical key step is to expand the eigenfunctions of the spherical Laplacian onto the eigenfunctions of the cylindrical Laplacian in order to relate the (true) spherical coordinates on the sky expansion to the cylindrical coordinates of the flat-sky expansion. This approach proves very powerful since it enables to obtain the full series of corrective terms to the flat-sky expansion. %Our construction will provide %efficient and precise computational tools for CMB, and in particular for non-Gaussianity. Once the method has been developed, it can be generalized to the polarisation and also to the computation of the bispectrum. In this latter case, depending on the way one chooses to describe the bispectrum, the exact form of the corrective terms has not been obtained but we can still provide an approximation whose validity can be checked numerically. %FB Before we enter the details of our investigations, and as the literature can be very confusing regarding the flat-sky approximations, let us present the different levels of approximations we are going to use. The reason there exist at all a flat-sky approximation is that the physical processes at play have a finite angular range. In case of the CMB, most of physical processes take place at sub-horizon scales and within the last scattering surface (LSS) (to the exception of the late integrated Sachs-Wolfe effect) and therefore within 1 degree scale on the sky. Let us denote $\ell_{0}$ the scale, in harmonic space associated with this angular scale. While using the flat sky approximation, the physical processes will be computed in a slightly deformed geometry (changing a conical region into a cylindrical one) introducing a priori an error of the order of $1/\ell_{0}$ (actually in $1/\ell_{0}^2$ depending on the type of source terms as it will be discussed in details below). Another part of the approximation is related to the projection effects which determine the link between physical quantities and observables. It introduces another layer of approximation of purely geometrical origin. For that part the errors behave a priori as $1/\ell$ if $\ell$ is the scale of observation in harmonic space. The resulting integrals do not lead to factorizable properties as it is the case for exact computations, while a factorization property can be recovered taking advantage of two possible limiting situations. First, for most of the small scale physical processes, one can use the fact that the radial extension of the source is much smaller than its distance from the observer. It is then possible to perform a {\em thin-shell approximation} effectively assuming that all sources are at the same distance from the observer. Another limit case corresponds to the situation in which the source terms are slowly varying and spread over a wide range of distances, as e.g. for galaxy distribution or weak-lensing field. In this case the sources support appears very elongated and it is then possible to use the {\em Limber approximation}~\cite{1953ApJ...117..134L,1980lssu.book.....P,2008PhRvD..78l3506L} which takes advantage of the fact that contributing wave modes in the radial direction should be much smaller that the modes in the transverse direction (but as such the Limber approximation can be used in conjunction of the flat-sky approximation or not). These different layers of approximations proved to be useful to compute efficiently the effects of secondaries such as lensing, but also of great help for computing the effects of non-linearities at the LSS contributing to the bispectrum, either analytically~\cite{2009JCAP...08..029B,Pitrou:2008ak} or numerically~\cite{2010JCAP...07..003P} (see Ref.~\cite{2009JCAP...05..014N} to compare to the full-sky expressions) as well as for the angular power spectrum; see e.g. Ref.~\cite{2006PhR...429....1L} for a review and for the relation between the flat-sky and the full-sky expansions in both real and harmonic space. We shall thus detail the expressions and corrective terms of the flat-sky approximation in these two approximations. In particular, we recover the result by Ref.~\cite{2008PhRvD..78l3506L} with a different method in the case of the Limber approximation. This is a consistency check of our new method.\\ %The article is thus organized as follows. Sections~\ref{SecSpectrumScalaire} %to~\ref{SecSpectrumGeneral} deal with the power spectrum of the CMB power spectrum, %Section~\ref{SecPolar} with the polarisation and Section~\ref{SecBispectre} with %the bispectrum. %FB : important changes below First, we consider the computation of the angular power spectrum in Section~\ref{SecSpectrumScalaire} starting with an example of such a construction in order to show explicitly how to construct next to leading order terms whose correction is found to be of the order of $1/\ell^2$. We then show that this construction is not unique and present the construction of a whole family of approximations whose relationship can be explicitly uncovered. In Section~\ref{SecApproximations} we present further computation approximations, e.g. the Limber (\S~\ref{limbersec}) and thin-shell (\S~\ref{thinsec}) approximations. %Section~\ref{SecFlatSkyDirection} describes the way to relate the different %flat-sky approximations and discusses the dependence in the choice of the flat-sky direction. %We emphasize (and compute) the fact that the flat-sky approximation breaks the %isotropy of the sky so that the power spectrum receives non-diagonal contributions, %the amplitude of which depend on the approximation. While in Sections~\ref{SecSpectrumScalaire} and ~\ref{SecApproximations} we have assumed, for clarity but also because it changes the result only at next-to-leading order, that the transfer function was scalar, in Section~\ref{SecSpectrumGeneral} we provide the general case of the flat-sky approximation up to next-to-leading order corrections in $1/\ell$ including all physical effects. Eventually Section~\ref{SecPolar} considers the case of higher spin quantities to describe the CMB polarization. We explore the case of the bispectrum construction in Section~\ref{SecBispectre}. One issue we encountered here is that different equivalent parameterizations can be used to describe bispectra (amplitudes of bispectra depend on both the scale and shape of the triangle formed by three $\ell$ modes that can be described in different manners). We thus present an alternative description of the bispectrum for which the flat-sky approximation can be done in a controlled way. Although we did not do the calculation explicitly, next-to-leading order terms can be then obtained. This is not the case for the reduced bispectrum for which we could nonetheless propose a general flat-sky approximation. Similarly to the case of spectra practical computations can then be done in the thin-shell approximation or the Limber approximation. %FB : Comments to be put in the conclusion ? %Our study thus provides several flat-sky approximations, %which may seems surprinzing at first. But as %usual with approximation, the %formula to use shall depend on the context. For %instance concerning the CMB, as long as the thin-shell approximation is a good %approximation, %one should use Eq.~(\ref{ClJP2}); for large scale structures, as long as the %transfer function is purely scalar as is the case for galaxy catalogs or weak-lensing, one should use %Eq.~(\ref{ClJP}) which requires neither the Limber nor the thin-shell %approximation, and if the Limber approximation applies, one should use %Eq.~(\ref{ClLimber2}).

This article provides a systematic construction of the flat-sky approximation of the angular power spectrum both for the temperature and polarization, in particular it shows that the expansion can be performed to any order. Additionally, we showed that this construction is not unique and that there exists a two-parameter family of flat-sky expansions, depending on the arbitrary choice of the flat-sky directions with respect to the azimuthal angle. As long as the sources are scalar, the first correction term scales as $1/\ell^2$ for a proper choice of flat-sky constraints ($k_\perp r = \ell+1/2$ or $k_\perp r = \sqrt{\ell(\ell+1)}$). In more realistic cases involving direction-dependent sources, such as the Doppler term or the anisotropic stress, the corrections terms were shown to scale as $1/\ell$ whatever the flat-sky constraint and are given in Eq.~(\ref{ClJP2}). Two particular extensions of the flat-sky expansion are particularly useful: the thin-shell and the Limber approximations, depending on the spatial extension of the sources. We checked that the corrective terms obtained in the Limber approximation are consistent with existing literature~\cite{2008PhRvD..78l3506L}, and we recovered the expression derived at leading order in the thin-shell approximation~\cite{Bond1996} where all flat-sky expressions are identical. We discussed here the validity of those approximations. For practical purposes, the ``best'' formula to use depends on the context: \begin{enumerate} \item for the CMB, as long as the thin-shell approximation is a good approximation, that is on the LSS, one should use Eq.~(\ref{ClJP2}) since one cannot neglect the Doppler term. It includes corrections of order $1/\ell$, and its validity is limited by the errors introduced while taking the thin-shell limit which are at most of order $\Delta r_\lss/r_\lss<0.01$. They appear in practice to be less on small scales and on large scales they can be corrected with Eq.~(\ref{CorrectThinshell}). \item for large scale structures such as galaxy catalogs or weak-lensing, for which the Limber approximation is a good approximation, one should use Eq.~(\ref{ClLimber2}). It includes corrections of order $1/\ell^2$ and its validity is limited to $1/\ell^3$~\footnote{Note that this differs from the conclusion of Ref.~\cite{2008PhRvD..78l3506L} where it is claimed without proof that the next order is only at $\ell^{-4}$.}, as long as there are no larger errors introduced while taking the Limber approximation. Note also that for the late integrated effects and the effect of reionization on the CMB, we should also use a Limber approximation. \end{enumerate} Generalization of this construction scheme to the bispectra was found to be more cumbersome. It led us to introduce an alternative description of the bispectra for which the flat-sky approximation is well controlled. This corresponds to the coefficients $\xi_{\ell_{1}\ell_{2}M}$ defined in Eq.~(\ref{xienfonctiondeblll}) whose relation with the usual reduced bispectrum form $b_{\ell_{1}\ell_{2}\ell_{3}}$ can be found in Eq.~(\ref{xienfonctiondeblll}) and inverted in Eq.~(\ref{xienfonctiondeblllInvert}). For this quantity we were able to propose a well controlled flat sky approximation. It actually leads to a specific form of flat-sky approximation for $b_{\ell_{1}\ell_{2}\ell_{3}}$, as described in Eq.~(\ref{bl1l2fs}) in the sense given by Eqs. (\ref{FormuleFlatsky}-\ref{k3perpExression}), the next-to-leading order terms of which remain however obscure (we encounter here exactly the same difficulty as when one tries to use the correspondance of Eq.~(\ref{corres})). In this case also one can further simplify the numerical integrations by using the thin-shell or Limber approximations. The validity of the bispectrum flat-sky leading order expansion was tested numerically in simple cases of separable primordial bispectrum such as the local and flat primordial bispectra. It was found to be very accurate in those cases, below the 1\% level. Such expressions are obviously of great interest for non-separable shapes of primordial non-Gaussianity since no fast full-sky method is known yet. % %This study provides an efficient tool for CMB computation, particularly for %non-Gaussianity. The level of approximation is well-controlled %and can be brought below the 1\% level.

1012.2652

This article constructs flat-sky approximations in a controlled way in the context of the cosmic microwave background observations for the computation of both spectra and bispectra. For angular spectra, it is explicitly shown that there exists a whole family of flat-sky approximations of similar accuracy for which the expression and amplitude of next to leading order terms can be explicitly computed. It is noted that in this context two limiting cases can be encountered for which the expressions can be further simplified. They correspond to cases where either the sources are localized in a narrow region (thin-shell approximation) or are slowly varying over a large distance (which leads to the so-called Limber approximation). Applying this to the calculation of the spectra it is shown that, as long as the late integrated Sachs-Wolfe contribution is neglected, the flat-sky approximation at leading order is accurate at 1% level for any multipole. Generalization of this construction scheme to the bispectra led to the introduction of an alternative description of the bispectra for which the flat-sky approximation is well controlled. This is not the case for the usual description of the bispectrum in terms of reduced bispectrum for which a flat-sky approximation is proposed but the next-to-leading order terms of which remain obscure.

12164745

2011AJ....141...94G

1012

1012.3470_arXiv.txt

\label{sec:introduction} Clusters of galaxies are ideal tracers of the largest density fluctuations in the Universe, and their abundance (and its evolution with cosmic time) may be used to place constraints on cosmology (e.g., \citealt{ecf96}). They also provide ideal laboratories for studying galaxy evolution. Originally used in this way since they contain a large number of galaxies all in the same location, it has since become clear that the properties of their member galaxies are markedly different from galaxies in the general field (e.g., \citealt{dressler80}, \citealt{1999ApJ...527...54B}, \citealt{Ellingson:2001zo}), implying that mechanisms which truncate star formation and transform galaxy morphology operate on cluster scales (e.g., \citealt{Treu:2003ty} and references therein). Constructing large, well-defined samples of galaxy clusters has a long and varied history. The first systematic searches involved visual identification of overdensities of optical galaxies on photographic plates \citep{abell,aco}. In the 1970s, with the advent of X-ray telescopes above the Earth's atmosphere, selection of clusters from their extended X-ray emission found favour \citep{mitchel76,serlemitsos77}. Recently, a combination of large format CCD detectors and objective algorithms to search efficiently for signatures of galaxy clusters has led to a revival in the use of optical selection in cluster surveys \citep{pdcs,kep,gal,gy00,2004MNRAS.348..551G}. A variety of techniques have been suggested to exploit the expected luminosity and/or color distribution of galaxies in clusters. The big advantage of these surveys compared with the older visual searches is that the detection method could be automated and characterised, meaning that the survey selection function could be quantified. Arguably the most efficient method is that of \citet{gy00} which uses the fact that the cores of galaxy clusters are dominated by galaxies with old stellar populations, forming a tight red-sequence in color magnitude space \citep{visv, bow92}. A number of other realisations of red-sequence based cluster finding algorithms exist (e.g., \citealt{Koester:2007ek}) differing in some details but all relying on accurate colors from imaging in two or more filters. The observed color of this sequence provides an accurate distance estimate. The application of this method led to the construction of the first Red-sequence Cluster Survey (RCS-1, \citealt{Gladders:2005oi}), a 72 square degree imaging survey in two bands ($R_C$ and $z^\prime$) designed to locate galaxy clusters from $0.2\lsim z \lsim1.1$ using the technique of \citet{gy00}. Not only is optical selection of galaxy clusters undergoing a revival, but astronomy in general is entering an era of `survey science' where an unprecedented number of wide-field optical (and NIR) surveys are currently underway or planned, such as LSST \citep{LSST-Science-Collaborations:2009uq}; Pan-STARRS\footnote{see http://pan-starrs.ifa.hawaii.edu/public/}; UKIDSS \citep{Lawrence:2007rp}; and DES\footnote{see https://www.darkenergysurvey.org/}. In addition, many of these wide-field optical surveys are specifically targeted at areas surveyed for clusters using other methods, such as the Blanco Cosmology Survey \citep{High:2010qy} of the South Pole Telescope (SPT, \citealt{Carlstrom:2009ys}) Sunyaev Zel'dovich effect (SZ)-selected cluster survey. For surveys using other methods (such as SZ selection), the optical data are critical for the verification of the cluster candidates found and for the determination of photometric redshifts. Furthermore, surveying the same areas with multiple techniques allows important comparisons of the different selection effects and the resulting properties of the clusters found (e.g., \citealt{don01,2004MNRAS.348..551G,Rasmussen:2006qy}). In this paper we describe the second Red-sequence Cluster Survey (RCS-2), the largest survey of this new generation for which imaging has already been completed. This builds on the methodology of RCS-1. The RCS collaboration has invested a large amount of work in attempting to characterize the selection function and the properties of clusters selected with this technique. Many of these results are directly applicable to RCS-2 (such as mass--richness calibrations) and so it is useful to summarize some of the RCS work to date. The efficiency of the selection method employed by the red-sequence surveys is that it can locate and estimate the redshifts of clusters using only one color (two filter) data, given the appropriate choice of filters. It is impractical to obtain mass estimates from follow-up observations of the $\sim$30 000 clusters which will be found in RCS-2, so the survey data themselves must be used to produce a proxy for cluster mass. Significant, representative samples of clusters from RCS-1 have been followed up using a variety of mass estimators such as dynamical \citep{gilbank:07a,felipe07}, X-ray \citep{hicks07}, strong and weak-lensing (from an ACS snapshot programme, PI:Loh; ACS SNe Cosmology project PI:Perlmutter) and SZ observations. In this way, the relationship between our mass proxy (optical richness from the survey data) and cluster mass can be understood. One of the primary goals of RCS-1 was to place constraints on cosmological parameters ($\Omega_M$, $\sigma_8$, \citealt{Gladders:2007us}) via the growth of the cluster mass function. This demonstrated for the first time the feasibility of such an approach using an optically-selected cluster sample. This approach used the measured relation between mass and richness, but also showed that meaningful constraints could be obtained using a self-calibration technique \citep{Majumdar:2004nj} to estimate the form of this relation from the survey data themselves. These authors demonstrate that the best constraints are obtained when accurate mass estimates are available for a subsample of clusters within the survey. It is worth emphasising that even if there is significant scatter in the relation between mass and the proxy (as we have found for optical richness), it is only important that the size of the scatter be well understood. With an order-of-magnitude larger survey than RCS-1, it becomes feasible to also constrain the equation of state of dark energy, $w$, \citep{Majumdar:2004nj} and this is in part the motivation for RCS-2. RCS-1 also produced a significant sample of strongly gravitationally lensed arcs around the clusters found. The number and redshift distribution of these lensing clusters were used to argue about the physical properties of the clusters responsible for their lensing cross-section and the relevance of such systems to constraining cosmology \citep{Gladders:2003xo}. The identification of such high surface brightness, strongly-lensed galaxies is another primary science driver for RCS-2. The massive clusters can be used as gravitational telescopes for studying high-redshift galaxies \citep[e.g.,][]{Pettini:2000fk,Wuyts:2010ul} which would otherwise be too faint to observe in detail. The giant arcs can also be used as probes of the properties of the cluster lenses themselves. With a statistical sample of galaxy clusters, such as in RCS-1, it is possible to study the properties of their member galaxies (e.g., their luminosity functions and blue fractions) by stacking subsamples built from the survey data themselves \citep{Gilbank:2007rq,Loh:2007be}. With the order-of-magnitude larger RCS-2, it becomes feasible to measure much weaker trends and push measurements of cluster galaxies down to much lower overdensities. The addition of photometric redshifts (e.g., \citealt{Hsieh:2005fq}) will allow these techniques to be extended to the field environment. Such galaxy evolution studies will be explored in future work with RCS-2. The outline for this paper is as follows. In \S2 we give an overview of the survey design and observations; \S\S3 \& 4 deal with CCD pre-processing, reduction, object detection and photometry; \S5 describes the photometric calibration via accurate fits to the star colors in our survey fields; \S6 outlines the procedure for and accuracy of the astrometric calibration. \S\S7 \& 8 describe the incorporation of additional data into our primary RCS-2 catalogs: $i$-band data which covers a large ($\sim$70\%) subsample of the primary $g$, $r$, $z$ survey area; and public imaging data from the CFHTLS-Wide survey, respectively. \S9 describes the final cleaning of the photometric catalogs: stitching into contiguous patches, removing duplicate data between overlapping pointings, and masking of artefacts. \S10 summarises and describes ongoing and future work for the survey.

We have presented the methodology for the reduction and precise calibration of the second Red-sequence Cluster Survey (RCS-2). The primary purpose of this survey is to build the largest sample of optically-selected galaxy clusters out to lookback times around half the age of the Universe. Results from cluster-finding (which will be presented in a forthcoming paper) show that the red-sequence redshifts are better than in RCS-1, due to the better uniformity and accuracy of the photometry in RCS-2, with typical $|z_{RCS}-z_{spec}| \ll 0.02$. A large number of projects aimed at characterizing clusters found with our technique have been ongoing for several years. An accurate understanding of the mass--richness relation will allow the cluster survey to place constraints on cosmological parameters. Future papers will present not only the cosmology results, but results on galaxy evolution, both within the clusters and the field. In addition to the large number of cluster science works possible, many other studies including those concerned with the properties of the Milky Way, such as the stellar populations and searches for dwarf galaxies, are possible. The large and accurate dataset provided by RCS-2 make it an important wide-field imaging survey of the new generation.

1012.3470

The second Red-sequence Cluster Survey (RCS-2) is a ~1000 deg<SUP>2</SUP>, multi-color imaging survey using the square-degree imager, MegaCam, on the Canada-France-Hawaii Telescope. It is designed to detect clusters of galaxies over the redshift range 0.1 &lt;~ z &lt;~ 1. The primary aim is to build a statistically complete, large (~10<SUP>4</SUP>) sample of clusters, covering a sufficiently long redshift baseline to be able to place constraints on cosmological parameters via the evolution of the cluster mass function. Other main science goals include building a large sample of high surface brightness, strongly gravitationally lensed arcs associated with these clusters, and an unprecedented sample of several tens of thousands of galaxy clusters and groups, spanning a large range of halo mass, with which to study the properties and evolution of their member galaxies. This paper describes the design of the survey and the methodology for acquiring, reducing, and calibrating the data for the production of high-precision photometric catalogs. We describe the method for calibrating our griz imaging data using the colors of the stellar locus and overlapping Two Micron All Sky Survey photometry. This yields an absolute accuracy of &lt;0.03 mag on any color and ≈0.05 mag in the r-band magnitude, verified with respect to the Sloan Digital Sky Survey (SDSS). Our astrometric calibration is accurate to Lt0farcs3 from comparison with SDSS positions. RCS-2 reaches average 5σ point-source limiting magnitudes of griz = [24.4, 24.3, 23.7, 22.8], approximately 1-2 mag deeper than the SDSS. Due to the queue-scheduled nature of the observations, the data are highly uniform and taken in excellent seeing, mostly FWHM &lt;~ 0farcs7 in the r band. In addition to the main science goals just described, these data form the basis for a number of other planned and ongoing projects (including the WiggleZ survey), making RCS-2 an important next-generation imaging survey.

12133278

2010arXiv1012.5851Z

1012

1012.5851_arXiv.txt

In models of stellar structure, situations are found where the heavier products of nuclear burning provide stability to a zone which otherwise would be unstable to convective overturning. Such a zone, or part of it, would become convective if something managed to mix its composition (R.J.\ Tayler, 1953). The question whether such a zone should be treated as if it were mixed or not has become known as the {\em semiconvection} problem. Answers to this question differ substantially. In practice, recipes are used containing a free parameter that allows the degree of mixing to be varied. Calculations in which such a parameter is adjusted to match observations are then called `with semiconvection'. Commonly used prescriptions are those of Langer (1985) and Maeder (1997). The presence of a semiconvective zone has only a minor effect on the thermal structure of the star. The assumed amount of mixing of composition is critical, however, because the evolution of the star is sensitive to the precise distribution of products of nuclear burning with depth in the star. The main goal of a theory for semiconvection is thus a good determination of the rate of mixing. From the perspective of the stellar evolutionist, the theory would ideally provide a formula for the rates of mixing and energy transport (the effective diffusivities), as functions of local thermodynamic state and composition, and their gradients. In Spruit (1992, hereafter S92) such formulas were derived, adapting the known physics of \textit{double-diffusive convection} to the case of a stellar interior. In the following, numerical simulations are used to measure mixing rates and their dependence on astrophysical conditions. They are compared with S92, and used as a basis for updated fitting formulas for the mixing rates in semiconvective zones of stars.

1012.5851

A grid of numerical simulations of double-diffusive convection is presented for astrophysical conditions. As in laboratory and geophysical cases convection takes place in a layered form. A translation between the astrophysical fluid mechanics and incompressible (Boussinesq) approximation is given, valid for thin layers. Its validity is checked by comparison of the results of fully compressible and Boussinesq simulations of semiconvection. A fitting formula is given for the superadiabatic gradient as a function of this parameter. The superadiabaticity depends on the thickness $d$ of the double diffusive layers, for which no good theory is available, but the effective He-diffusion coefficient is nearly independent of $d$. For a fiducial main sequence model (15 $M_\odot$) the inferred mixing time scale is of the order $10^{10}$ yr.

12207181

2011JCAP...10..007B

1012

1012.2476_arXiv.txt

It is commonly believed that Ultra-high Energy (UHE) neutrinos should arrive at the Earth coming from very distant sources like Active Galactic Nuclei (AGN) and Gamma Ray Bursts (GRB) following straight trajectories. Neutrinos interact only through weak interactions. They are not charged and if they have a nonzero magnetic moment it must be very small. Hence, there would be no interaction preventing them to travel cosmological distances~\cite{Halzen:2002pg}. The efforts to improve sensitivity to the UHE neutrino flux may test the existence of the Berezinsky-Zatsepin (BZ) neutrinos~\cite{Berezinsky}. Those neutrinos are generated through the same process that predicts the Greisen-Zatsepin-Kuzmin (GZK) cutoff~\cite{Greisen:1966jv} of Cosmic Ray (CR) flux, which has been observed by HiRes~\cite{hires} and the Auger Observatory~\cite{Abraham:2008ru}. The most energetic CRs, most probably composed by protons and nuclei, may have an origin at relatively close sources, of the order of 100~Mpc for protons and even less for nuclei~\cite{Hooper:2006tn}. In fact, UHE CRs are most probably arriving from nearby AGNs~\cite{Cronin:2007zz}. But we observe that despite all the efforts to detect high energy extragalactic neutrinos, i.e., neutrinos with energies higher than $10^{3}$~TeV, more and more restricted limits on their flux have been reported by several experiments due to the non observation of these neutrinos~\cite{Ackermann:2005sb,Desai:2007ra, Abraham:2007rj,BlanchBigas:2009zz,Abbasi:2011qc}. In addition, based on astronomical and cosmological observations, it is difficult to deny the existence of Dark Matter (DM) and Dark Energy (DE). Cosmological observations of clusters of galaxies indicate that the density fraction of DM is $\Omega_{\rm DM} \approx 0.227 \pm 0.014$~\cite{wmap}. The two most popular solution to the DM problem are {\it i)} a slight modification on the Newtonian dynamics \cite{milgron} and {\it ii)} relic particles~\cite{Bertone:2004pz}. A DM relic particle candidate must be non-relativistic, must be stable on the cosmological time scale, and must interact weakly with other particles. Some of the possible candidates are: WIMPS~\cite{Jungman:1995df}, axions~\cite{Peccei:1977ur,Peccei:1977hh,Mikheev:2008zz,Andrianov:2009kj,dfsz}, MeV-scalar fields~\cite{Boehm:2004,Boehm:2003hm}, technicolor candidates ~\cite{Foadi:2008qv,techndm1,techndm2}, and ultra-light scalar fields~\cite{Matos:2008ag,Arbey:2001qi,Amendola:2005ad,sflocal,Lesgourgues:2002hk,Hu:2000ke}. An ultra-light scalar field is motivated as DM because it can alleviate some of the problems that arise at galactic scale in the standard paradigm of cold DM, namely, the origin of cuspy halos~\cite{Moore:1999nt} and the overproduction of substructure~\cite{Lin:2000qq}. Neutrino interaction with DM, $\nu$-DM for short, could have strong implications at cosmological scales. Interactions of neutrinos with light scalar fields have been studied with some interesting implications noticed, such as a reduction of the relic neutrino density, leading to a neutrinoless universe~\cite{Beacom:2004yd}, or a modification on the CMB spectra~\cite{Mangano:2006mp,Serra:2009uu}, or even a connection between the smallness of neutrino mass and a MeV-mass scalar field DM~\cite{Boehm:2006mi}. Furthermore, $\nu$-DM interaction might affect the flux of UHE neutrinos. In particular, such interaction may suppress the neutrino flux resulting in a kind of GZK cutoff for neutrinos. Many DM candidates were analyzed in this context: heavy neutrinos as dark matter~\cite{Weiler:1983xx,Weiler:1992fm,Roulet:1992pz}, lightest supersymmetric particles (LSP) discussed in Ref.~\cite{Weiler:1992fm} and updated in Ref.~\cite{Barenboim:2006dj}, and MeV-mass scalar field~\cite{Boehm:2004,Mangano:2006mp,Boehm:2006mi}. In all these cases the suppression is small, not interfering in the propagation of UHE neutrinos. Nevertheless, we show here that a coupling between relic ultra-light scalar fields and neutrinos may imply a suppression of the UHE neutrino flux, in which case there may be a confusion between the flux limit at the source and a reduction of the UHE neutrino flux during propagation. Previous analysis have put constraints on the $\nu$-DM interaction couplings for those models by using, for instance, SN1987A neutrino data or possible imprints on the angular power spectra of CMB anisotropies \cite{Mangano:2006mp}. Nevertheless, those limits do not apply to our case, since they were obtained by assuming a mass of the scalar field $m_\phi > 10$ MeV while we will explore the possibility that ultra-light scalar field with $m_\phi \ll 1$ eV can couple to neutrinos. Such a small mass for the scalar particle gives a cross section with different behavior compared the one previously reported in other works~\cite{Boehm:2004,Boehm:2003hm}, allowing for new phenomena like a flux suppression for reasonable values of the coupling constants. Consequences of this type of ultra-light scalar fields have been studied in other astrophysical contexts like in the equilibrium of degenerate stars~\cite{Grifols:2005kv}. In this work we use the $\nu$-DM coupling as described by the Feynman rule \cite{Boehm:2003hm} shown in figure~\ref{feyn-rule} where $\,P_R\,$ denote the chiral projector, $\phi$ stands for the spin-0 dark matter field and $g_{\nu\phi}$ denote the strength of its Yukawa coupling. \begin{figure} \center{ \begin{picture}(80,80)(0,50) \ArrowLine(0,70)(60,90) \DashArrowLine(0,120)(60,90){5} \Text(5,65)[]{F} \ArrowLine(60,90)(100,90) \Text(105,85)[]{f} \Text(5,110)[]{$\phi$} \Text(180,90)[]{$g_{\nu\phi} P_R$} \end{picture} } \caption{Feynman rule for the interaction of a neutrino with an ultra-light scalar field.} \label{feyn-rule} \end{figure} In the next Section we show how the neutrino flux may vary due to a nonstandard interaction of neutrinos with DM. We also introduce the space of parameters, namely, cross section and DM mass, that lead to an important effect on the neutrino propagation through cosmological distances. In Section~\ref{crosssec} we show the elastic scattering cross section for self and non-self-conjugate scalar fields while Section~\ref{couplimits} is devoted to study the suppression to the neutrino flux due to the proposed $\nu$-DM interaction. Finally in Section~\ref{conclusion} we discuss our results and conclusions.

There are several experiments, like IceCube and the Pierre Auger Observatory, expecting to detect extragalactic neutrinos. But neutrinos with energies above $10^{15}$~eV, coming from extragalactic sources, have not been detected yet. We study the possibility that UHE neutrinos could be absorbed while traveling from their sources to the Earth. In particular we illustrate this idea by considering an ultra-light particle as a component of the Dark Matter in the Universe. We consider a mechanism for the neutrino interaction based on a scalar field dark matter model and we show that in this case the propagation of extragalactic neutrinos from sources $100$~Mpc or farther from the Earth may be affected. This would give negative results on neutrino telescopes or UHE neutrino detectors. On the other hand, despite neutrinos from a nearby supernova could interact with the DM halo around the collapsing star, the scale involved would not be sufficient for an absorption like the one proposed here. Although nonstandard interactions of neutrinos with Dark Matter particles had been considered in the literature before, no important effect on neutrino propagation had been predicted. In most of the literature, light scalar field DM had been considered to be relativistic and the coupling to neutrinos was constrained due to measurable effects on the CMB spectra. In this work we have considered non-relativistic ultra-light scalar fields, proposed in the literature, that besides their gravitational effects, may not have other measurable astrophysical consequences. To our knowledge, this is the first example of a possible suppression of the extragalactic neutrino flux due to propagation effects. Therefore, care must be taken when using the limits obtained by such experiments, since those limits can be due to two factors: source limit and/or absorption due to UHE neutrino-ultra light scalar field Dark Matter interaction during neutrino propagation from the source to the Earth. On the other hand, a positive signal of UHE neutrinos could be useful to put restrictions on models that contains a light scalar field DM candidate. Similar arguments could be applied for particles other than neutrinos. For instance, in~\cite{Andrianov:2009kj} it was analyzed the suppression of charged particles due to the interaction with a pseudoscalar and it was shown that the axion can play the role of a shield for high energy cosmic rays.

1012.2476

We study the possible suppression of the extragalactic neutrino flux due to a nonstandard interaction during its propagation. In particular, we study neutrino interaction with an ultra-light scalar field dark matter. It is shown that the extragalactic neutrino flux may be suppressed by such an interaction, leading to a new mechanism to reduce the ultra-high energy neutrino flux. We study both the cases of non-self-conjugate as well as self-conjugate dark matter. In the first case, the suppression is independent of the neutrino and dark matter masses. We conclude that care must be taken when explaining limits on the neutrino flux through source acceleration mechanisms only, since there could be other mechanisms for the reduction of the neutrino flux.

12164580

2011AJ....141..131H

1012

1012.1580_arXiv.txt

An excess of faint blue stars in Quadrant 1 (Q1) of the Galaxy compared to complemetary fields in Quadrant 4 (Q4) was initially recognized by \citet{lar96}. \citet{par03} subsequently extended the survey to a much larger contiguous region covering several hundred square degrees. They confirmed the original findings, mapped the stellar excess in Q1 from {\it l} $\sim 20 - 55\arcdeg$ and {\it b} $\sim 20 - 45 \arcdeg$, and argued for a comparable asymmetry in Q1 below the plane. The stars showing the excess were probable Thick Disk stars, 1 - 2 kpc from the Sun. \citet{par04} also reported an associated kinematic signature. Velocities and metallicities of stars in 12 fields in Q1 and Q4 showed that the Thick Disk stars in Q1 have a much slower effective rotation rate $\omega$, compared to the corresponding Q4 stars, with a significant lag of 80 to 90 km s$^{-1}$ in the direction of Galactic rotation, greater than the expected lag of 30 to 50 km s$^{-1}$ \citep{Reid98,CandB,Carollo2010} for the canonical Thick Disk population. Interpretation of the asymmetry, now referred to as the Hercules Thick Disk Cloud \citep{lar08}, is not straightforward. It is tempting to assume that the asymmetry is the fossil remnant of a merger, however the star counts were also consistent with a triaxial Thick Disk or Inner Halo, with its axis in Q1, as well as with a gravitational interaction with the stellar bar in the disk \citep{1992ApJ...384...81W,2000MNRAS.317L..45H}. The latter is especially intriguing given the corresponding asymmetry in the kinematics. A triaxial Thick Disk could also yield different effective rotation rates because of noncircular streaming motions along its major axis. The star count excess appears to terminate near {\it l} $\sim$ 55$\arcdeg$ \citep{par03}. To search for evidence of triaxiality, in Paper I \citep{lar2010a} we extended the star counts to fainter magnitudes, corresponding to greater distances, from {\it l} of 50$\arcdeg$ to 75$\arcdeg$. The fields at 55$\arcdeg$ to 75$\arcdeg$, show no significant excess, including the faintest magnitude interval, and therefore, do not support a triaxial interpretation of the asymmetry. The outstanding question relevant to the origin of the Hercules Cloud is its spatial extent along the line of sight and its associated kinematics. In Papers I and II \citep{lar2010b} in this series we describe our faint CCD survey to map the spatial extent of the over-density. The photometric survey covers 47.5 square degrees in 63 fields in Q1 and Q4 above and below the Galactic plane. Except for fields with {\it b} $\sim$ 30 -- 40$\arcdeg$ in Q1, most of these regions are not covered by the Sloan Digital Sky Survey (SDSS, \citet{York}). \citet{lar2010b} find that the over-density or star count asymmetry in Q1 extends to approximately 5 kpc along our line of sight, and that the regions showing the excess interestingly are above the near side of the density contours for the bar in the Disk \citep{1992ApJ...384...81W}. We have also extended our corresponding spectroscopic survey to fainter magnitudes and greater spatial coverage. We have obtained additional medium-resolution spectra and now have radial velocities and metallicity estimates for more than 4000 stars in 31 fields. In the next section we describe the observations and data reduction. In \S {3} we discuss the kinematics and confirm the asymmetry between the stars in Q1 and Q4, and in \S {4} we use the metallicity information from the spectra to separate the stars into the different populations. We determine the rotational lag for the different populations in \S {5}, and in the concluding section we summarize the kinematic and spatial asymmetries in the Hercules Cloud and its interaction with the Galactic bar in the Disk.

1012.1580

In the first two papers of this series, Larsen et al. describe our faint CCD survey in the inner Galaxy and map the overdensity of thick disk stars in Quadrant 1 (Q1) to 5 kpc or more along the line of sight. The regions showing the strongest excess are above the density contours of the bar in the Galactic disk. In this third paper on the asymmetric thick disk, we report on radial velocities and derived metallicity parameters for over 4000 stars in Q1, above and below the plane, and in Quadrant 4 (Q4) above the plane. We confirm the corresponding kinematic asymmetry first reported by Parker et al., extended to greater distances and with more spatial coverage. The thick disk stars in Q1 have a rotational lag of 60-70 km s<SUP>-1</SUP> relative to circular rotation, and the metal-weak thick disk stars have an even greater lag of 100 km s<SUP>-1</SUP>. Both lag their corresponding populations in Q4 by ≈30 km s<SUP>-1</SUP>. Interestingly, the disk stars in Q1 also appear to participate in the rotational lag by about 30 km s<SUP>-1</SUP>. The enhanced rotational lag for the thick disk in Q1 extends to 4 kpc or more from the Sun. At 3-4 kpc, our sight lines extend above the density contours on the near side of the bar, and as our lines of sight pass directly over the bar the rotational lag appears to decrease. This is consistent with a "gravitational wake" induced by the rotating bar in the disk which would trap and pile up stars behind it. We conclude that a dynamical interaction with the stellar bar is the most probable explanation for the observed kinematic and spatial asymmetries. <P />Based on observations obtained at the MMT Observatory, a joint facility of the Smithsonian Institution and the University of Arizona, and at the Cerro Tololo Inter-American Observatory (NOAO) operated by the Association of Universities for Research in Astronomy (AURA).

428063

2013SSRv..176..101F

1012

1012.0586_arXiv.txt

\label{intro} The newest frontier in space exploration is close to home, where the outflowing magnetized solar wind plasma meets and mixes with the interstellar cloud at the heliosphere's boundaries. The evolving heliosphere is formed by the relative ram (dynamic) pressures of the solar wind and partially ionized low density, $\sim 0.3$ \cc, interstellar cloud flowing around (interstellar ions and magnetic field) and through (interstellar neutrals) the heliosphere at 26.3 \kms. The interstellar neutrals dominate the mass-density in the heliosphere beyond 10--15 AU, and are able to penetrate deep into the inner heliosphere. Gravitational focusing of heavy elements in the flow forms the helium focusing cone that the Earth traverses early every December. The Interstellar Boundary Explorer (IBEX) mission \citep{McComasetal:2009ssr} has recently mapped the energetic neutral atoms (ENAs) that are formed by charge-exchange between heliosphere plasmas and interstellar neutrals in the heliosheath regions. These maps are now reshaping our understanding of the heliospheric interaction with the interstellar medium. IBEX has also made the first $in~situ$ detection of the flow of interstellar oxygen through the heliosphere \citep{Moebius:2009sci}. The presence of interstellar neutrals in the heliosphere was established 40 years ago. OGO 5 mapped the \HI\ Lyman alpha sky background and showed a diffuse component attributed to interstellar neutral hydrogen in the inner heliosphere \citep{ThomasKrassa:1971,BertauxBlamont:1971}. A $Copernicus$ spectrum of this weak Ly$-\alpha$ emission firmly established it as interstellar \citep{AdamsFrisch:1977}. Measurements of the fluorescence of solar 584 A emission from interstellar \HeI\ in the heliosphere revealed the helium focusing cone \citep{WellerMeier:1974}. The discovery of helium pickup ions \citep{Moebiusetal:1985} showed that the ionization of interstellar neutrals inside of the heliosphere produces energetic ions that trace the neutrality of the interstellar cloud around the Sun. Inside of the heliosphere, interstellar neutrals are ionized through charge-exchange with the solar wind, photoionization, and for those neutrals surviving to 1 AU electron-impacts \citep{Rucinskietal:1996}. Pickup ions are a crucial diagnostic of the interaction between the heliosphere and the interstellar medium. ENAs formed from charge-exchange between interstellar neutrals and solar wind ions were recognized as an important remote diagnostic of the distant heliosphere boundary regions \citep{Gruntman:1993,HsiehGruntman:1993}. ENAs with energies 55--80 keV, originating in the heliosheath, were discovered by CELIAS/HSTOF on SOHO \citep{Hilchenbachetal:1998ena}. IBEX maps of ENAs produced in the heliosphere boundaries, and the unexpected discovery of the IBEX Ribbon, requires a new paradigm for understanding the interaction between the heliosphere and interstellar medium \citep{McComas:2009sci,Funsten:2009sci,Fuselier:2009sci,Schwadron:2009sci}. Interstellar neutral helium is the best marker for the upwind direction of the ISM flowing through the heliosphere, usually denoted the heliosphere nose direction. \citet{Moebiusetal:2004} combined several data sets to obtain an upwind direction toward \elon,\elat$ \sim 255^\circ, ~ 5^\circ$ (or in galactic coordinates, $ \ell \sim 4^\circ, ~ b \sim 15^\circ$). Photoionization models of the interstellar cloud surrounding the heliosphere, that are constrained by the observed properties of interstellar material inside and surrounding the heliosphere, show that the circumheliospheric interstellar cloud is low density, partially ionized, and warm, \nHI$\sim 0.20$ \cc, \nel$\sim 0.07$ \cc, \npro$\sim 0.06$ \cc, and $T \sim 6300$ K \citep[Model 26 in ][]{SlavinFrisch:2008}. If thermal and magnetic pressures are similar, the interstellar magnetic field (ISMF) strength is $\sim 2.7$ \microG. \begin{figure}[t!] \begin{center} \includegraphics[width=0.950\textwidth]{fig1.eps} \end{center} \caption{ IBEX measurements of the first all-sky maps of ENA fluxes over central energies of 0.2 keV (Lo) through 4.3 keV (Hi). The maps are Molleweide projections in ecliptic coordinates (J2000), and are centered near the longitude of the heliosphere nose. Also shown are the galactic plane (red line in upper figure) and the current locations of the Voyager 1 (at \elon,\elat=$255^\circ, ~ 34^\circ$) and Voyager 2 ($289^\circ, -29^\circ$) satellites. The prominent arc of bright emission is the IBEX Ribbon. The Ribbon is not ordered by either ecliptic or galactic coordinates, but instead appears to form where the sightlines are perpendicular to the direction of the interstellar magnetic field draping over the heliosphere. The magnified region shows fine scale structure in the Ribbon, which has been identified by smoothing each $0.5^\circ$ angle in spin phase along a measurement swath by the amount needed to reach 10\% counting statistics. The ENA energies are plotted in the spacecraft frame. Pixels in the small area near $\lambda \sim 120^\circ - 150^\circ$ are contaminated by the magnetosphere, and are filled in by extrapolation from adjacent regions. \citep[The figure is from ][]{McComasetal:2010var}.} \label{fig:1} \end{figure}

1012.0586

The Interstellar Boundary Explorer (IBEX) mission is exploring the frontiers of the heliosphere where energetic neutral atoms (ENAs) are formed from charge exchange between interstellar neutral hydrogen atoms and solar wind ions and pickup ions. The geography of this frontier is dominated by an unexpected nearly complete arc of ENA emission, now known as the IBEX `Ribbon'. While there is no consensus agreement on the Ribbon formation mechanism, it seems certain this feature is seen for sightlines that are perpendicular to the interstellar magnetic field as it drapes over the heliosphere. At the lowest energies, IBEX also measures the flow of interstellar H, He, and O atoms through the inner heliosphere. The asymmetric oxygen profile suggests that a secondary flow of oxygen is present, such as would be expected if some fraction of oxygen is lost through charge exchange in the heliosheath regions. The detailed spectra characterized by the ENAs provide time-tagged samples of the energy distributions of the underlying ion distributions, and provide a wealth of information about the outer heliosphere regions, and beyond.

12133598

2010arXiv1012.5706H

1012

1012.5706_arXiv.txt

Since discovery of the sunspots by Galileo, genesis of their 22 year cyclic activity in general and, their formation and decay during their evolutionary stages in particular still remain a mystery. The study of sunspots' origin, formation and decay is important owing to the observed fact that variation of sunspot occurrence activity is related with the solar irradiance that in turn affects the earth's environment and the climate (Prabhakaran Nayar {\em et. al.} 2002; Hiremath and Mandi 2004; Soon 2005; Badruddin, Singh and Singh 2006; Perry 2007; Feymann 2007; Tiwari and Ramesh 2007 and references there in; Scafetta and West 2008 Komitov 2009, Hiremath 2009b and references there in). Present general consensus is that the sunspots originate below the solar surface due to an unknown dynamo mechanism. Due to very high conductivity of the solar plasma and assuming that raising flux tube does not acquires extra flux from the ambient medium, sunspots isorotate with the internal plasma and due to buoyancy raise toward the surface along the path of rotational isocontours. This implies that sunspots are very good tracers of the internal dynamics and structure of the solar interior. Hence if the sunspots that have first and second days appearance on the surface, and if one computes their initial rotation rates, then one can infer rotation rate of the internal solar plasma where the sunspots' foot points are anchored. Recent studies (Javaraiah and Gokhale 1997; Javaraiah 2001; Hiremath 2002; Sivaraman {\em et. al.} 2003) show that variation of initial rotation rates obtained from the daily motion of sunspot groups with respect to their life spans on the surface is almost similar to the radial variation of the internal rotation profile of the solar plasma. \begin{figure}[h] \begin{center} \includegraphics[width=20.5pc,height=18pc]{1874-1976_allarea_l0-l40_rotin.ps} \caption{The dashed and the dotted curves are the variation of initial rotation rates of the sunspot groups with respect to their life spans (Hiremath 2002). The continuous curve is the radial variation of the internal rotation as inferred from the helioseismologyi (Antia, Basu and Chitre 1998). } \end{center} \label{sunspot_area_age} \end{figure} From Hiremath's (2002) paper, results are reproduced in Fig 1 that illustrates a comparison between the variation of initial rotation rates of the sunspot groups for different life spans and radial variation of internal rotation profile as inferred (Antia, Basu and Chitre 1998) from the helioseismology. Note the striking similarity between these two profiles. In order to reach closer to the reality of the physics of convection zone and dynamics of the flux tubes, in the same study, the rate of change of initial rotation rates of the sunspot groups (that represent the acceleration or deceleration of the flux tubes in the ambient plasma) are compared (the Fig 5(b) of Hiremath (2002) with the radial profile of gradient of rotation that is computed from the radial variation of rotation of the plasma inferred from the helioseismology). Again we get a striking similarity between these profiles. To conclude from that study, for different life spans, initial sunspot dynamics over the surface represents the internal dynamics in different layers of the convection zone. For example initial anchoring of a flux tube whose life span is 10 days is near base of the convection zone and initial anchoring of a flux tube whose life span is 5 days is in the middle of the convective envelope. Observations show that there are three important stages in the sunspot's evolutionary history : (i) a well developed sunspot (that consists of umbra and penumbra) is formed due to coalescing of the emerging flux regions, (ii) once stabilized sunspot is formed, its area increases and reach the maximum value and, (iii) decay of the sunspot from it's maximum area to minimum area and ultimately disintegrating into smaller active regions and diffusion of the flux on the surface. As for the first and last stages, there are many studies that explain the formation and decay parts of the sunspot's evolutionary history. The first stage is supposed to be due to convective collapse, a kind of instability that has been invoked to explain the kilo gauss fields on the surface (Parker 1978; Spruit 1979; Hasan 1985). Once flux element is formed, different adjacent flux elements coalesce and sunspot is formed. Owing to their strong magnetic field structure, sunspots inhibit the ambient convection resulting in reduction of temperature and density. Ultimately lower density of the flux tube results in raising (due to buoyancy) from the convection zone to the surface. Contrary to this conventional view, Parker (1992) has proposed that sunspots are basically formed due to coalescence of magnetic elements by the vortices. According to him, flux tubes are surrounded by vortex flows that attract other vortices leading to coalescence of different flux elements. On the other hand Meyer {\em et. al.} (1974), have different view on the formation of the flux tubes. According to them a strong converging flow is necessary to form the sunspots. That means sunspots might be formed at the boundary of the convective cells, with an outflow at the surface and an inflow in the deeper layers. Where as Hiremath (2009b; 2010), by updating Alfven's (1943) seminal idea of sunspot formation, came to the conclusion that sunspots are formed due to superposition of Alfven wave perturbations of the underlying steady part of large scale toroidal magnetic field structure and travel along isorotational contours in order to reach at the proper activity belt on the surface. There are many studies on the decaying phase of the sunspot. Cowling (1946) was the first person to investigate the decay part of the sunspot area. Bumba (1963) obtained a linear decay law for the recurrent spot groups and exponential decay law for the non-recurrent spot groups. Where as some of the previous studies (Petrovay and Moreno-Insertis 1997; Petrovay and van Driel-Gesztelyi 1997) indicate the quadratic decay ({\em i.e.}, sunspot area as quadratic function of time) and other studies (Solanki 2003 and references there in) indicate the linear decay law. Moreno-Insertis and Vazquez (1988) and Martinez Pillet, Moreno-Insertis and Vazquez (1993) conclude that the present sunspot data do not allow any distinction between either linear or quadratic decay law. To add to these decay laws, log-normal distribution (Martinez Pillet, Moreno-Insertis and Vazquez 1993) also fit the decay of umbrae. There are following theoretical studies to understand the sunspot decay. First theoretical study in supporting the results of linear decay laws is by Zwan and Gokhale (1972). Such a linear decay law suggests that flux loss takes place everywhere within the spot irrespective of their different sizes. Zwan and Gokhale (1972) assumed a current sheet around the sunspot and turbulent diffusion inside the tube. In this case Ohmic diffusion dictates the decay of the current sheet and hence as spot decays to smaller area, thickness of the current sheet reduces. In fact such current sheets around the sunspots have been observed by Solanki, Rueedi and Livingston (1992). In contrast, Simon and Leighton (1964) and Schmidt (1968) propose that the sunspots are decayed by the erosion of the sunspot boundary which implies that $dA/dt$ is proportional to $A^{1/2}$, where A is area of spot. Supporting the erosion model, Petrovay and Moreno-Insertis (1997) proposed that turbulent diffusivity depends strongly on the field strength. Their model predicts the quadratic decay and spontaneous current sheet around the sunspot. Though there are many studies on the first and last phases of the sunspot evolution, the second stage of a sunspot, viz., physics of a growth phase, during it's life time is not understood. Moreover, it is not clear whether all the three phases in a sunspot's life time remain same or different over the whole solar cycle. That means: is there any year to year variations in the area gradients (rate of change of area $dA/dt$ with respect to time, where A(t) is time dependent area of the sunspots and $t$ is time variable) of the sunspots during it's increasing (second phase) and decaying (last phase)? Is there any connection between the evolutionary history of the sunspots and underlying deeper dynamics or this phenomenon is simply due to surface convection. Some of these important issues are addressed in this study. As for year to year changes in gradient of sunspots' area, for the year 1955-1965 and for different life spans, Hiremath (2005, with summer student Mr. Subba Rao), computed both growth $dA_{1}/dt$ and the decay $dA_{2}/dt$ rates of the sunspots and came to the following conclusions. For the same life span, to reach their maximum areas sunspots take different times as cycle progresses. That is, in the beginning of the cycle, area-age curves are nearly gaussian and as cycle progresses area-age curve follow the simple linear decay law. Further they conclude: (i) during the beginning of the solar cycle, sunspots' {\em rate of growth} and {\em rate of decay} are larger compared at end of the solar cycle, (ii) in the beginning of the solar cycle, in order to reach their maximum areas, sunspots increase their area at a rate of $\sim$ 100 mh (millionths hemisphere)/day where as at the end of solar cycle sunspots increase their area at the rate of $\sim$ 50 mh/day and, (iii) sunspots decay faster ( $\sim$ 75 mh/day) in the beginning of the solar cycle compared to the end of the solar cycle ($\sim$ 25 mh/day). The last conclusion is similar to the conclusion from the recent study (Hathaway and Choudhary 2008) on decay of area of the sunspot groups. Active regions are centers of solar activity ranging from flares to CMEs. They are believed to be locations where magnetic flux bundles erupt from deep interior in the convection zone to emerge at solar surface in the form of sunspots due to magnetic buoyancy. Further the complexity of sunspot groups plays an important role in determining the active regions (Zirin 1988). The difference in the total energy output between a solar minimum and maximum, indirectly associated with the sunspots, is about 0.1\%. Even this small energy changes in the sun's output over 11-year solar cycle can intensify wind and rainfall activities and therefore has a major impact on global weather and climatic patterns in the Earth's climatic parameters. Therefore it is useful to investigate how the sunspot groups themselves eventually grow and decay. The growth and decay of sunspot groups also play an important role in the day to day irradiance variations (Wilson 1981). If decay of sunspots were purely by ohmic dissipation, sunspots would have lifetimes of about 300 years by considering their size and photospheric conductivity (Cowling 1946). However, the sunspots have shorter life span of $\sim$ weeks for non-recurrent spot groups and $\sim$ months for recurrent spot groups. How to reconcile these observed phenomena, viz., three phases of formation, growth and decay of the sunspots. In the recent study (Hiremath 2009b; Hiremath 2010), it is proposed that sunspots are formed by the superposition of many Alfven wave perturbations of the embedded toroidal magnetic field structure. Once sunspots are formed, due to buoyancy, at different depths in the convective envelope raise along isorotational contours and reach the surface at different latitudes. One can notice from Fig 1 that the internal rotational profile (continuous curve), as inferred from helioseismology (Antia, Basu and Chitre 1998), has two rotational gradients, {\em viz.,} a positive rotational gradient from base of the convective envelope to $0.935 R_{\odot}$ and a negative rotational gradient from $0.935 R_{\odot}$ to $1.0 R_{\odot}$. From the magnetic induction equation, it is proposed in this study that growth and decay of either sunspots' area or magnetic flux is due to interplay of both convective source term (that in turn depends mainly upon fluctuations in the positive rotational gradient and convection) and sink term (that in turn depends upon fluctuations in negative rotational gradient, magnetic eddy diffusivity and radiation effects near the surface). That means sunspots that are formed in the region of positive rotational gradient, while raising toward the surface, accumulate magnetic flux from the ambient magnetic turbulent medium and reduction of magnetic flux in the region of negative rotational gradient. The net magnetic flux of the sunspot that is formed in the region of positive rotational gradient in the convective envelope while raising it's anchoring feet and reaching toward $0.935 R_{\odot}$, should increase and, magnetic flux should decrease as flux tubes' anchoring feet lifts from $0.935 R_{\odot}$ to $1.0 R_{\odot}$. On the other hand, the sunspots that are formed in the region of negative rotational gradient while raising toward the surface mainly experience decay phase only. These reasonable ideas will be clear in the following section. In order to understand and test these ideas on growth and decay phases of the sunspots, magnetic induction equation is solved by considering the source and the sink terms separately. In section 2, formulation of the equations are presented. With reasonable approximations, analytical solution of magnetic induction equation for the growth of area is presented in section 3 and solution for decay part is presented in section 4. In section 5, both the analytical solutions are fitted with observed sunspots' growth and decay phases of the sunspots and conclusions from these results are presented. \section {Formulation of the equations} It is assumed that, in the convective envelope, fluid is incompressible. We also assume that the magnetic eddy diffusivity $\eta$ with value represented by the appropriate average. Magnetic field ${\bf B}$ and the velocity field ${\bf V}$ vectors are expressed as \begin{equation} {\bf B} = {P}{\bf \hat{\hbox{I}}}_\vartheta + {T}{\bf \hat{\hbox{I}}}_\varphi \ , \end{equation} \begin{equation} {\bf V} = {U}{\bf \hat{\hbox{I}}}_\vartheta + {r \Omega sin\theta}{\bf \hat{\hbox{I}}}_\varphi \ , \end{equation} where ${\bf \hat{\hbox{I}}}_\vartheta$ and ${\bf \hat{\hbox{I}}}_\varphi$ are the unit vectors along heliographic latitude $\vartheta$ and longitude $\varphi$ of the sun; $P$, $T$, $U$ and $\Omega$ are scalar functions. $P$ and $T$ are scalar functions that represent poloidal and toroidal parts of the the magnetic field structures and $U$ and $\Omega$ are scalar functions that represent poloidal (meridional) and toroidal (angular velocity) parts of the velocity field structures. Equation of continuity is \begin{equation} {{\partial \rho}\over{\partial t}} + \rho \nabla . {\bf V} = 0 . \end{equation} As the life spans ($\sim$ weeks to months) of sunspots are very much larger than the time scales ($\sim$ minutes) of ambient density perturbations in the convective envelope, we have ${{\partial \rho}\over{\partial t}}=0$ and the resulting equation is \begin{equation} {\Large\bf {{ \rho \nabla . {\bf V} = 0}}} \end{equation} where $\rho$ is the density. Similarly as magnetic diffusivity is assumed to be constant, magnetic induction equation is \begin{equation} {{\partial {\bf B}}\over{\partial t}} = curl ({\bf V}\,\times {\bf B}) + \eta \nabla^{2} {\bf B} . \end{equation} This magnetic induction equation determines growth and decay of the sunspot. The first term on right hand side (RHS) is the source term that enhances the magnetic flux of the sunspot and second term on RHS is the sink term that attempts to destroy the generated magnetic flux. As magnetic induction equation in turn depends upon velocity and diffusivity $\eta$, these two source and sink terms are important and dictate the growth and decay of the sunspots. We solve the induction equation by considering the source and sink terms separately for the following reasons. In case of region of positive rotational gradient from base of convective envelope to $0.935R_{\odot}$, rate of increase of magnetic flux that mainly depends upon fluctuations of increase in rotational gradient is dominant compared to magnetic diffusivity. As for region of negative rotational gradient from $0.935R_{\odot}$ to $1.0R_{\odot}$, fluctuations in decreasing rotational gradient, increasing magnetic diffusivity ( as magnetic diffusivity $\eta \, \sim T^{-3/2}$, where $T$ is ambient temperature) and dominant radiational effects near the surface remove and destroy the magnetic flux.

In order to test results of the physical ideas on the growth and decay of the sunspots that are presented in the previous sections, data of time evolution of corrected areas of non-recurrent sunspot groups from Greenwich Photoheliographic Results (GPR) are used. For the four latitude zones of $0-10$, $10-20$, $20-30$ and $30-40$ degrees two spot groups that lie between $\pm$ 70 degree from the central meridian and life spans in the range of 8-10 days are considered. In Figures 2-9, time evolution of growth of area of the non-recurrent sunspot groups are presented. It is assumed that sunspot area grows linearly, quadratically and exponentially and relevant laws are fitted with the observed growth of area of the sunspot groups. As measured uncertainties in the areas of sunspot groups are not available in GPR, it is assumed that growth and decay of area curves follow the Poisson distribution and hence uncertainty in each of measured area $A(t)$ (where $t$ day of observation) is taken as $[A(t)]^{1/2}$. By knowing area $A(t)$ values and their uncertainties, all the three laws are fitted to the observed sunspots' area growth curves and are over plotted on top of the each plot. In all the Figures 2-9, the plots in the top are for the linear and quadratic fits and the the plot at the bottom is fit for the exponential growth law. Similarly, in Figures 10-17, observed decay of area of the sunspot groups for all the four latitude zones are presented. In addition to three (viz., linear, quadratic, exponential) decay laws, a law of log-normal distribution is also considered for fitting the observed decay curves. In all the Figures 10-17, first and second plots in the top row are for linear and quadratic decay fits respectively. In the second row of Figures 10-17, log-normal and exponential decay fits are presented. As for growth of the sunspots, it is interesting to note that among all the Figures 2-9, exponential fit is best one. This is also clearly evident from the $ \chi^2 $ values presented in Table 1. In the Table 1, first column represents latitude of occurrence of the sunspot, second column represents lifespan and, columns 3-5 represent $ \chi^2 $ values for linear, quadratic and exponential fits. It is to be noted that low value of $ \chi^2 $ means, observed and expected curves are almost similar. In Table 2, constants $C_{1}$ and $C_{2}$ of exponential growth and decay parts of the area curve are presented. First column represents latitude of occurrence of the spot group, second and third columns represent the constants $C_{1}$ and $C_{2}$ that are determined from the exponential growth and, fourth and fifth columns represent the constants that are determined from fits of exponential decay of the sunspot area curve respectively. Another important property, according to theoretical expectations presented in section 3 regarding growth of the sunspot, as is evident from Table 2 (see the third column) that the exponent of the exponential fit for the high heliographic latitude is high compared to the exponential fits for the low heliographic latitudes. That means the spots that formed at the high latitudes grow fast (with exponential growth) and the spots that are formed near the low heliographic latitudes grow slowly. \begin{table} \caption{Values of constants obtained from growth and decay of the exponential fits.} \label{congd} \centering \begin{tabular}{p{2.3cm}p{2.3cm}p{2.3cm}p{2.3cm}p{2.3cm}} \hline & Growth & & Decay& \\ \hline Latitude& $C_{1}$& $C_{2}$& $C_{1}$& $C_{2}$\\ \hline 0 - 10$^\circ$ & 39.65$\pm$15.93& 0.26$\pm$0.66& 395.44$\pm$6.17 &0.93$\pm$0.56\\ 0 - 10$^\circ$ & 4.57$\pm$3.22&0.59$\pm$0.41 & 208.51$\pm$4.71 & 0.98$\pm$0.28\\ 10 - 20$^\circ$ &12.31$\pm$4.18 & 0.72$\pm$0.58 & 138.38$\pm$1.08& 0.43$\pm$0.35\\ 10 - 20$^\circ$ & 60.95$\pm$7.10& 0.54$\pm$0.93& 403.434$\pm$6.17 &0.55$\pm$0.48\\ 20 - 30$^\circ$ & 36.23$\pm$5.31& 0.63$\pm$0.54& 5.02$\pm$1.45& 0.14$\pm$0.27\\ 20 - 30$^\circ$ & 12.18$\pm$3.86 & 0.65$\pm$0.34 & 212.73$\pm$6.23& 0.27$\pm$0.53\\ 30 - 40$^\circ$ &29.08$\pm$5.53 & 0.14$\pm$0.62& 92.76$\pm$4.02& 0.38$\pm$0.26\\ 30 - 40$^\circ$ &30.27$\pm$5.10 & 0.68$\pm$0.61& 391.51$\pm$6.11& 0.46$\pm$0.50\\ \hline \end{tabular} \end{table} \begin{table}% \caption{$ \chi^2 $ fit for the laws of linear, quadratic, log-normal and exponential decay of the sunspot.} \label{Dsunch2} \centering \begin{tabular}{p{2cm}p{2.cm}p{2.cm}p{2.cm}p{2.cm}p{2cm}} \hline Latitude & Life span & Linear & Quadratic & Log-normal& Exponential\\ &(Days)& & & & \\ \hline 0 - 10$^\circ$& 8 & 19.82& 7.77 & 1.17& 0.29\\ 0 - 10$^\circ$& 9 & 9.12 & 3.29 & 0.05& 0.03\\ 10 - 20$^\circ$& 10 & 53.27 & 6.37& 0.23& 0.08\\ 10 - 20$^\circ$& 8& 101.64& 62.45 & 0.26& 0.13\\ 20 - 30$^\circ$& 10& 16.60& 17.28& 0.13& 0.08\\ 20 - 30$^\circ$& 10 & 24.56 & 3.29& 0.08& 0.05\\ 30 - 40$^\circ$& 8 & 17.38& 17.28& 2.44& 0.91\\ 30 - 40$^\circ$& 8 & 36.84& 31.51& 0.36& 0.12\\ \hline \end{tabular} \end{table} As for decay of the sunspots, even though log-normal fit appears to be a very good fit among all the Figures 10-17, from criterion of goodness of fit of $\chi^{2}$ value, exponential decay fit is best one. In fact this result is also clearly evident from the values of $ \chi^2 $ presented in Table 3. Similarly we have another important property from these results. According to theoretical expectations presented in section 4 regarding decay of the sunspots, exponent of the exponential fit (see the $5^{th}$ column of Table 2) for the high heliographic latitude is very low compared to exponent of the exponential fits for the low heliographic latitudes. That means the spots that are formed at the high latitudes decay slowly compared to the spots that are formed near the low heliographic latitudes. Even with approximations by neglecting fluctuations in poloidal (meridional) and toroidal (angular) components of velocity fields, theoretical solutions of growth (equation 16) and decay (equation 28) parts of sunspot's area evolutionary phases match with the observed area evolutionary phases. In order to understand a unique single solution for understanding growth and area decay curve, one should solve consistently full set of MHD equations (as the neglected fluctuations ${\partial \Omega^{'} \over \partial t}$ in turn depend upon fluctuations in the momentum equations). From the observed characteristics of growth and decay of the sunspots at different latitudes on the surface and from the theoretical ideas presented in this study, one can safely conclude that sunspots are formed due to constructive interference of toroidal Alfven wave perturbations and, after attaining a critical strength in the convective envelope, due to buoyancy, sunspots raise along isorotational contours and reach the respective latitudes. It is understood from this study that growth and decay phases of the sunspots not only depend upon the surface physical characteristics , as this problem (especially decay part) was treated by the earlier studies, but also evolutionary history of internal dynamics and magnetic field structure of the sunspots while they raise toward the surface. As the sunspot is a three dimensional structure whose evolutionary history not only depends upon its internal structure but also on the ambient dynamic properties of the solar convective envelope that ultimately yields a combined solution of growth and decay of the sunspot.

1012.5706

From the previous study (Hiremath 2009b; Hiremath 2010), on the genesis of solar cycle and activity phenomena, it is understood that sunspots are formed at different depths by superposition of Alfven wave perturbations of a strong toroidal field structure in the convective envelope and after attaining a critical strength, due to buoyancy, raise toward the surface along the rotational isocontours that have positive (0.7-0.935 $R_{\odot}$) and negative (0.935-1.0 $R_{\odot}$) rotational gradients. Owing to physical conditions in these two rotational gradients, from the equation of magnetic induction, sunspot's area growth and decay problem is solved separately. It is found that rate of growth of sunspot's area during its evolution at different depths is function of steady and fluctuating parts of Lorentzian force of the ambient medium, fluctuations in meridional flow velocity, radial variation of rotational gradient and $cot(\vartheta)$ (where $\vartheta$ is co-latitude). While rate of decay of sunspot's area at different depths during its evolution mainly depends upon magnetic diffusivity, rotational gradient and $sin^{2}(\vartheta)$. Gist of this study is that growth and decay of area of the sunspot mainly depends upon whether sunspot is originated in the region of either positive or negative rotational gradient. For different latitudes and life spans of the sunspots on the surface during their evolutionary history, both the analytically derived theoretical area growth and decay curves match reasonably well with the observed area growth and decay curves.

12131944

2010arXiv1012.0912R

1012

1012.0912_arXiv.txt

In this paper I review the current understanding of the mergers of double neutron stars (DNS) and neutron star black hole systems (NSBH). I will collectively refer to them as ``compact binary'' mergers, and thus exclude systems with white dwarf components. Only binary systems that merge under the influence of gravitational wave emission will be discussed. Dynamical collisions as they occur in dense regions such as the cores of globular clusters will not be discussed here. Two excellent reviews have recently appeared \cite{faber09,duez10a} which put their emphasis on numerical relativity, I want to round up the picture by focusing on the (nuclear) astrophysics aspects of compact binary mergers.\\ To date 10 binary systems are known where at least the mass function and the periastron advance are consistent with both stars being neutron stars \cite{lorimer08}. Five of these systems have small enough orbital separations so that the constant leakage of orbital angular momentum due to gravitational wave emission will cause a coalescence within a Hubble time. Incompletely understood physical mechanisms, poorly known parameters and hard to quantify selection effects make it difficult to estimate the rates at which such events occur. Rates derived from the observed systems roughly agree with those from population synthesis models, about 40 to 700 Myr$^{-1}$ in a Milky Way equivalent galaxy \cite{kalogera04a,belczynski07}, but with an uncertainty of about an order of magnitude in each direction. Even less secure is the rate for NSBH systems, to date none has been observed and population synthesis studies have predicted values from an order of magnitude more \cite{bethe98} to about two orders of magnitude less \cite{belczynski07} than the DNS merger rate.\\ \noindent It is difficult to overrate the importance of this type of binary system: \bi \vspace*{-0.3cm} \i The first discovered system, PSR 1913+16, has delivered --via the measured decay of the binary orbit-- the first unambiguous evidence for the existence of gravitational waves, in excellent agreement with the predictions of General Relativity \cite{taylor89}. \vspace*{-0.3cm} \i The measurement of at least two Post-Keplerian parameters allows the measurements of {\em individual} neutron star masses. For example, the masses of PSR 1913+16 are $m_1= 1.4414$ \Msun and $m_2= 1.3867$ \Msun $\pm \; 0.0002$ \Msun \cite{weisberg05}. Accurately known masses provide stringent tests for the hadronic physics inside a neutron star \cite{lattimer04}. This regime of high densities, but low temperatures is hardly accessible to any laboratory experiment, but it can be probed via accurately known neutron star masses. \vspace*{-0.3cm} \i The large "compactness" of neutron stars, $\zeta \equiv G M_{\rm ns}/R_{\rm ns} c^2\approx 0.2$ (for comparison: the compactness of the Sun is $\sim 10^{-6}$), and their high orbital velocities, $v \sim 10^{-3} c$ ($c$ being the speed of light) make DNS excellent laboratories for strong gravity. They allow for accurate tests that have the potential to distinguish General Relativity from other theories of gravity \cite{kramer09}. \vspace*{-0.3cm} \i The last stages of the inspiral of a DNS system are a prime candidate for the first {\em direct} detection by the ground-based gravitational wave detectors \cite{sigg08,arcenese08,grote08} that have now finished their first complete science runs. Population synthesis models \cite{belczynski07} have predicted detection rates for the Advanced LIGO project near 10 for DNS and around one event per year for NSBH systems. The detection efficiency could be further enhanced by the simultaneous detection of accompanying signals in other channels. \vspace*{-0.3cm} \i The astrophysical production site of the r-process is still an unsolved problem. For many years supernovae, in particular the neutrino-driven winds from a new-born neutron star, were considered very promising sites to forge r-process material, but recent studies find it difficult to reproduce the observed abundance patterns with parameters that are considered plausible for core-collapse supernovae \cite{arnould07,roberts10}. The main competitor are the neutron-rich ejecta that seem unavoidable in a compact binary merger \cite{lattimer76,lattimer77,rosswog99,ruffert01,oechslin07a,metzger08}. \vspace*{-0.3cm} \i Since the very beginning, compact binary mergers have been considered a prime candidate for the central engine of Gamma-ray bursts (GRBs) \cite{blinnikov84,paczynski86,goodman86,eichler89,paczynski91,narayan92} and they have survived being confronted with a wealth of observational results. Although the case is far from being settled, they still are the major candidate for the central engine of short GRBs \cite{piran05,nakar07,lee07,gehrels09}. \ei

The past decade has seen a tremendous progress in our ability to reliably model the mergers of compact binary systems. On the one side, fully relativistic 3D simulations have become possible and are now performed regularly. On the other side, many physical processes have been explored from (exotic) high-density nuclear physics, over both neutrino cooling and heating to nuclear reactions and magnetic fields.\\ Technically, the ``hydrodynamics plus gravity'' part of the models is likely to see a couple of substantial improvements. Eulerian approaches are today substantially hampered in their ability to accurately follow lower density material, to a large extent due to heavy computational burden, so that it is difficult to afford large computational volumes. Although some refinement techniques are already used \cite{baiotti08,anderson08a}, the field will most likely profit a lot from the implementation of fully adaptive mesh refinement schemes. Closely related is the treatment of ``numerical vacuum'' and its impact on low-density regions. The latter include also the accretion disks which are thought be crucial in the transformation of gravitational binding energy into observable radiation. As outlined above, this technical progress will have major implications for both nucleosynthesis and GRB questions. Another natural improvement would be the exploration of higher order methods, though there is a tradeoff between the order of the method and the affordable resolution.\\ Lagrangian, and in astrophysics this usually means Smoothed Particle Hydrodynamics, methods, are currently somewhat lagging behind in their treatment of dynamical space-time evolution with CFA currently being the most advanced gravity approach \cite{faber06a,oechslin02,bauswein10}. The method has recently made much progress with the most advanced sets of relativistic equations following directly from ideal fluid Lagrangians \cite{monaghan01,rosswog09b,rosswog10b,rosswog10a} and hybrid approaches that couple SPH with grid-based space-time solvers seem promising for an interesting class of problems. Moreover, adaptive Lagrangian-Eulerian (ALE) approaches may be worth the development effort.\\ Also on the non-gravity side remains much to do. Problems where the hard to reach numerical resolution is the major stumbling block need apart from massive parallelism further algorithmic developments and suitable approximation methods. Today, still many simulations are run with polytropic equations of state with fixed adiabatic exponents which are insufficient approximations for the required broad spectrum of physical conditions in a merger. With resolution becoming better, the simulated times becoming longer and low-density phenomena such as winds now becoming feasible, also the presently available physical/nuclear equations of state reach their limits, mainly at low densities and temperatures. Closely related is the question of nucleosynthesis, another crucial facet in the multi-messenger picture of compact binary mergers. Despite its importance it is still in its infancy and many pressing questions are essentially unexplored. For some related technical problems, say the implementation of nuclear reactions into existing hydrodynamics codes, the development effort should be moderate. For others, say for self-consistent 3D calculations of neutrino-driven winds, a serious amount of work is needed to achieve computationally affordable and reasonably accurate results.\\ In anticipation of the first {\em direct} gravitational wave detection, one can overall be optimistic that the field will keep its high current momentum, so that hopefully in time for the first gravitational wave detections reliable multi-physics models will be in place.

1012.0912

This paper reviews the current understanding of double neutron star and neutron star black hole binaries. It addresses mainly (nuclear) astrophysics aspects of compact binary mergers and thus complements recent reviews that have emphasized the numerical relativity viewpoint. In particular, the paper discusses different channels to release neutron-rich matter into the host galaxy, connections between compact binary mergers and short Gamma-ray bursts and accompanying electromagnetic signals.

12167612

2011ApJ...728...28W

1012

1012.0592_arXiv.txt

Long-term and often short-term variability in both flux and spectral shape, across the electromagnetic spectrum, is a fundamental characteristic of the spectra of active galactic nuclei (AGN). While variability is ubiquitous in AGN, our understanding of long-term spectral changes in local AGNs is the result of studies of a few dozens of sources focusing primarily on variations in broad optical emission lines in conjunction with reverberation mapping studies (e.g.,~\citealt{2004ApJ...613..682P}) for local AGNs ($z < 0.1$). It is important to extend these studies in temporal space (e.g.,~acquiring a larger sample of observations over longer periods of time), the wavebands covered (e.g.,~multi-wavelength studies), and the number/type of sources sampled, in order to understand the physics behind the AGN emission -- the accretion processes, nature of outflows, and associated feedback. In an effort to extend the variability studies of individual AGN, we present an analysis of the available UV and X-ray observations of Mrk 817, a nearby Seyfert 1.5 galaxy with a redshift of $z = 0.031158$ \citep{1999PASP..111..438F}. Its optical spectrum exhibits complex emission in the narrow-line region (NLR) and broad-line region (BLR), with outflowing NLR [\ion{O}{3}] $\lambda\lambda 4959, 5007$\AA~emission measured at a systemic velocity of $-450$\,km\,s$^{-1}$ \citep{2006MNRAS.371.1610I}. In the UV, observations with both the {\it Hubble Space Telescope} (HST) and the {\it Far Ultraviolet Spectroscopic Explorer} (FUSE) also revealed outflowing gas. HST Goddard High-Resolution Spectrograph (GHRS) observations from 1997 of Ly$\alpha$ indicated the presence of an intrinsic AGN outflow \citep{2000ApJS..130..121P}. Later observations with FUSE, in 2000--2001, show weak absorption in \ion{O}{6} with a radial systemic velocity component of $-$4100\,km\,s$^{-1}$, the largest offset measured in a Seyfert \citep{2007AJ....134.1061D}. Analysis of multiple observations from FUSE shows that the weak absorption is variable in that the equivalent width of the absorber increases with decreasing continuum flux \citep{2008AJ....136.1201D}. Additionally, a second and third absorption component ($v \approx -3000$\,\km) appeared in the final FUSE observation, potentially due to a change in ionization. Owing to its UV brightness and sight line through a complex of Galactic high velocity clouds (HVC Complex C, \citealt{2001AJ....122.3280G,2003ApJ...585..336C}), Mrk 817 was a selected target for the guaranteed time observations with the HST Cosmic Origins Spectrograph (COS). Two high signal-to-noise ratio observations were taken with COS, separated by $\sim 5$ months. Additionally, UV observations of this source taken with the {\it International Ultraviolet Explorer} (IUE) and HST GHRS allow a search for variability over a period of nearly 30 years. Details of the UV observations, are presented in Section~\ref{sect-uv}. We include an analysis of the intrinsic AGN properties, including the UV continuum and emission-line properties from each of these observations. In Section~\ref{sect-xray}, we present an analysis of X-ray spectra of Mrk 817, along with simultaneous optical/UV photometry, from XMM-Newton and Swift. Although the source is bright in the X-ray band, it has only recently been targeted for observation in the X-ray with XMM-Newton. Finally, we present our conclusions on the spectral analysis and UV and X-ray variability of Mrk 817 in Section~\ref{sect-conclusion}. We assume a standard cosmology with $H_0 = 75$\km\,Mpc$^{-1}$, throughout.

In this paper, we present an analysis of several UV and X-ray spectra of the Seyfert 1.5 Mrk 817. The UV spectral analysis includes new high-resolution spectra from the Cosmic Origins Spectrograph, as well as archived spectra from IUE and HST GHRS. These observations span nearly 30 years and allow us to search for spectral variability on decade time scales. The X-ray spectral analysis includes a spectrum from XMM-Newton, taken within two weeks of the second COS observation, as well as five lower signal-to-noise ratio spectra with Swift XRT. These observations allow us to probe X-ray spectral variability on the 2.5\,year time scale covered. In addition, luminosities from ROSAT allow us to determine variability in the soft X-ray band over a 20\,year period. Finally, an analysis of the OM/UVOT observations, which were simultaneous with the X-ray observations, allows us to determine variability in the UV/X-ray SED. In this section, we summarize and discuss the results of our analysis. \subsection{Spectral Shape and Fundamental Properties} The HST COS observations of Mrk 817 exhibit the broad UV emission lines (Ly$\alpha$, \ion{N}{5}, \ion{Si}{4} + \ion{O}{4}], and \ion{C}{4}) characteristic of an optical broad-line Seyfert. Our analysis of the ERO and GTO observations, separated by $\sim 5$\,months, finds similar power-law slopes ($F_{\lambda} \propto \lambda^{-1.30}$) and FWHMs for the emission lines (e.g., we find a broad component of $\sim 3000$\km\,and a narrower component of $\sim 1000$\km\,for the \ion{C}{4} emission). The \ion{C}{4} emission is blue-shifted by values similar to other AGNs (blue-shifted by 50--340\,\km), but there is no correlation between the shifts in the \ion{C}{4}/\ion{Si}{4}/\ion{N}{5}/Ly$\alpha$ emission components. Further, these velocity shifts change between observations. No strong intrinsic absorption lines (i.e., Ly$\alpha$, \ion{N}{5}, \ion{C}{4}) are detected in the COS observations. However, four high-velocity cloud absorbers and four intrinsic Ly$\alpha$ absorption lines are detected in the GHRS and COS spectra (as well as four absorbers detected at low-significance, see \S~\ref{sect-absorbers}). Corresponding features are not detected in the more ionized lines. The X-ray spectra of Mrk 817 are well-characterized by a blackbody + power-law model. The temperature of the blackbody component, used to characterize the soft excess, of $kT \approx 0.1$\,keV, is typical of Seyfert 1s (e.g., \citealt{2005AA...444...79M,2009ApJ...690.1322W}) and quasars \citep{2004MNRAS.349L...7G}. However, the power-law slope is steeper than the average photon index of similar AGN. We find $\Gamma \ga 2.0$ in all but the lowest flux observations; while typical values for local Seyfert 1s detected in the hard X-rays with Swift's BAT (of which, Mrk 817 is included in the 22-month Swift BAT catalog) are $\Gamma = 1.78 \pm 0.24$ \citep{2009ApJ...690.1322W}. Other than the steep slope, there are no unusual features in the spectrum. Additionally, we find no evidence for either neutral or ionized absorption. Using the flux and continuum properties derived from our spectral analysis, we determine fundamental properties of Mrk 817 to determine how it compares to local Seyferts. The X-ray 2--10\,keV luminosity, which we compute using the 2--10\,keV flux of $1.45 \times 10^{-11}$\,\flux~from the XMM-Newton spectrum, can be used to estimate the bolometric luminosity of the AGN. Using a bolometric correction of 35, typical of unabsorbed AGN \citep{2005AJ....129..578B} and within the range of parameters determined from our SED fits to the simultaneous optical/UV/X-ray data from Swift and XMM-Newton, we estimate $L_{bol} = 9.58 \times 10^{44}$\,erg\,s$^{-1}$. This value is consistent with local ($\langle{z}\rangle = 0.03$) broad-line Seyferts selected in the 14--195\,keV band with the Swift BAT detector, where the average bolometric luminosity is $7.4 \times 10^{44}$\,erg\,s$^{-1}$, with a range of values from $\sim 9 \times 10^{43}$ -- $6 \times 10^{45}$\,erg\,s$^{-1}$ \citep{2010ApJ...710..503W}. The black hole mass of Mrk 817, $\log {M}/{M}_{\sun} = 7.69 \pm 0.07$ as derived from reverberation mapping \citep{2009ApJ...697..160B}, is also typical of local AGN \citep{2002ApJ...579..530W} and the BAT AGN sample in particular ($\langle \log {M}/{M}_{\sun}\rangle = 7.87 \pm 0.66$; \citealt{2010ApJ...710..503W}). However, the estimated accretion rate ($L_{bol}/L_{Edd}$, where $L_{Edd} = ({M}/{M}_{\sun}) \times 1.3 \times 10^{38}$\,erg\,s$^{-1}$) of 0.14 is higher than the average accretion rate of 0.034 (with a range of values from $\sim$ 0.001--0.5) determined for broad-line Seyferts in the BAT sample \citep{2010ApJ...710..503W}. Thus, Mrk 817 is in many ways (mass, luminosity, UV spectral shape) a typical broad-line Seyfert. Since the accretion rate is believed to determine the hard X-ray spectral slope in AGNs, with high spectral slopes correlating with high accretion rates \citep{2006ApJ...646L..29S}, it is not surprising that Mrk 817 has both a slightly higher photon index and accretion rate than the average local Seyfert 1. \subsection{Intrinsic Absorption} Four intrinsic absorption features were detected embedded in the broad Ly$\alpha$ emission of Mrk 817. These features include a $-4250$\km, $-4100$\km, $-3550$\km, and $-2600$\km\,absorber. The strength of these features vary between the 1997 GHRS and 2009 COS observations, with the ratio of EW between GHRS/COS being $5.7 \pm 1.3$, $2.2 \pm 1.1$, $3.9 \pm 3.0$, and $\sim 0.3$, respectively. The width of the absorption features are consistent between the GHRS and COS observations. No intrinsic \ion{C}{4}~$\lambda\lambda$1548.20,1550.77 and \ion{N}{5}~$\lambda\lambda$1238.82,1242.80 absorption features are detected at the same velocities. However, given the weakness of the intrinsic Ly$\alpha$ absorbers in the COS spectra (EW $\sim 35$\,m\AA), this absence is expected. For similar ionic column densities, the ratio of optical depth between two lines should be proportional to the ratio of their $f\lambda$ values (where $f$ is the oscillator strength and $\lambda$ is the wavelength of the transition). Under these conditions, $\tau$(Ly$\alpha$)/$\tau$(\ion{C}{4} $\lambda$1548)=1.7 and $\tau$(Ly$\alpha$)/$\tau$(\ion{N}{5} $\lambda$1238)=2.6. For a photo-ionized plasma with a fixed total hydrogen column density, the column density of each ion is a strong function of the ionization parameter $U$. However, as shown in figure 2 of \citet{2007ApJ...658..829A}, for solar abundances, the column densities of \ion{C}{4} and \ion{N}{5} are always smaller than that of \ion{H}{1}, irrespective of the value of $U$. Combined with the stronger expected optical depth of the Ly$\alpha$ line, a weak absorption feature of the latter should not show related features in \ion{C}{4} and \ion{N}{5}. The same figure shows that under similar conditions \ion{O}{6} absorption is expected for $\log{U}\sim-0.5$, where for solar abundances the \ion{O}{6} column density is 8 times larger than that of \ion{H}{1}, while $\tau$(Ly$\alpha$)/$\tau$(\ion{O}{6} $\lambda$1032)=3.7 for the same column densities. Under these conditions, we expect $\tau$(\ion{O}{6} $\lambda$1032) to be twice that of Ly$\alpha$. This may explain why \ion{O}{6} and Ly$\beta$ features are detected in the same FUSE data (For the same column densities, $\tau$(Ly$\beta$)/$\tau$(\ion{O}{6} $\lambda$1032)=0.6.). The photo-ionization curves also show that $\tau$(\ion{O}{6} $\lambda$1032)/$\tau$(Ly$\alpha$) drops below 1 for $\log{U}>0.4$ and $\log{U}<-1$. Similarly, with such a weak Ly$\alpha$ absorption feature, we do not expect to detect the \ion{O}{7} and \ion{O}{8} absorption edges in the near-simultaneous X-ray data. The strongest GHRS Ly$\alpha$ absorber at $-4250$\,\km, which is $\sim 5$ times weaker in the COS spectra, is likely associated with the strongest \ion{O}{6} absorber seen in the four FUSE observations from 2000--2001 \citep{2008AJ....136.1201D}. As we found for the GHRS/COS observations, the \ion{O}{6} absorber, whose velocity changes slightly from $-4198$\km~to $-4144$\km~throughout the observations, also varies in EW (from 0.60\,\AA~to 0.33\,\AA). Both the strength and velocity of the \ion{O}{6} absorber decreased over time -- however, this decrease was not correlated with changes in the Far-UV luminosity of Mrk 817. Additional weak \ion{O}{6} absorption lines were also found to appear throughout the set of four FUSE observations, with no obvious correlation to the UV luminosity. In a similar way, a comparison of the COS observations with the GHRS observation shows no correlation between UV luminosity (in the continuum or emission lines) and the equivalent width of the absorber. % Mrk 817 is not alone in exhibiting variable absorbers. Variability in UV absorption lines has been detected in several Seyfert 1s, including NGC 4151 (e.g.,~\citealt{1981MNRAS.196..857P,1985MNRAS.215....1B,2006ApJS..167..161K}), NGC 3516 (e.g.,~\citealt{1983ApJ...267..515U}), NGC 5548 (e.g.,~\citealt{1993ApJ...416..536S,2009ApJ...698..281C}), and NGC 3783 (e.g.,~\citealt{1996ApJ...465..733M,2005ApJ...631..741G}). The cause of variability of intrinsic absorbers is attributed to a change in ionization state or shielding of the photo-ionizing continuum of an outflowing wind or bulk motions of an absorbing cloud/clouds. A potential cause of the change in the absorbers in Mrk 817 is from bulk motions of clouds along our line-of-sight. As pointed out in \citet{2008AJ....136.1201D}, weak absorption -- as in the \ion{O}{6} absorbers in Mrk 817 -- is seen in spectra where this is the case. Further, the fact that there is little change in the UV luminosity between the GHRS and COS spectra, while the strength of the Ly$\alpha$ absorber decreases significantly, also favors this model. This would suggest that ionization is not causing a change in the absorber. Based on the assumption of an absorbing cloud moving radially out of the line of sight in the 12.5 years between the GHRS and COS observations, this implies that the cloud moved a distance about $10^{17}$\,cm ($\ll 1$\,pc). Based on the X-ray measured upper limits on the \ion{O}{7} and \ion{O}{8} absorption edges, we can estimate an upper limit on the column density of the absorbing cloud to compare with UV-derived measurements. Using the cross-sections ($\sigma$) for \ion{O}{7} and \ion{O}{8} from \citet{1996ApJ...465..487V} and the relation that $\tau = \sigma {\rm N}_{\rm O, ion}$, we find upper limits on the column density as N$_{\rm O VII} \la 4.0 \times 10^{17}$\,cm$^{-2}$ and N$_{\rm O VIII} \la 4.7 \times 10^{17}$\,cm$^{-2}$. The typical range of column densities in \ion{C}{4} is well below the limit we find for \ion{O}{7} and \ion{O}{8}, with N$_{\rm C IV} = (0.1$--$14) \times 10^{14}$\,cm$^{-2}$ \citep{1999ApJ...516..750C}. This X-ray derived limit is also much higher than the estimated \ion{H}{1} column from the Ly$\alpha$ absorption line ($\sim 4.3 \times 10^{13}$\,cm$^{-2}$ in the GHRS and $\sim 7.1 \times 10^{12}$\,cm$^{-2}$ in COS, assuming a fully covering absorber). Therefore, it is possible that an outflow is present, but it has a column density too low to be detected in the CCD spectra analyzed. \subsection{Variability} With available observations of Mrk 817 from multiple observatories spanning decades, we were able to probe variability in flux and spectral shape across the UV and X-ray bands. In the UV, from analysis of the IUE, GHRS, and COS spectra, we do not find drastic changes in either the line or continuum flux of Mrk 817. As with other AGN, the continuum flux is correlated with the emission line flux of \ion{C}{4} in accordance with the Baldwin effect \citep{1977ApJ...214..679B,1994ApJ...436..678O}. Measured changes in UV flux are limited to a factor of $\la 2.3$ between the minimum and maximum measured values over the 30\,year time-span probed. Additionally, in support of our spectral analysis, we find that the UV photometry from XMM-Newton OM and Swift UVOT observations are at a similar level as the UV continuum measured from the COS spectra. Variability in the X-ray spectra is much more pronounced than in the UV. In particular, we find that the X-ray spectrum of Mrk 817 varies in both luminosity and spectral shape. The photon index measured from a power-law component varies from $\Gamma \approx 1.5$--$2.1$ between the five Swift XRT and the XMM-Newton spectra, which span 2.5\,years. The variability in the spectral slope is correlated strongly with the observed X-ray luminosity, where the steepest slopes are observed at the highest X-ray luminosities. Two ROSAT flux measurements from the early 1990s have allowed us to search for long-term variability in the soft X-ray band luminosities from the ROSAT, Swift, and XMM-Newton observations. We find that the X-ray luminosity of Mrk 817 changes markedly. Shorter term variability is observed in the Swift XRT observations, with a factor of four change in variability over $\sim 2$\,years. On the long term, we find that the 1990 ROSAT luminosity is $\sim 40$ times smaller than the luminosity measured from the 2009 XMM-Newton observation. The factor of 40 variability seen over 20 years in the X-ray observations contrasts with the factor of 2.3 seen in the UV spectra. Thus, Mrk 817 is clearly more variable in the X-ray than UV band. To search for a possible relationship between the UV and X-ray luminosities, we analyzed simultaneous optical/UV/X-ray data from XMM-Newton and Swift. We found that the optical/UV fluxes were at similar levels in the four observations with OM/UVOT photometry. The optical-X-ray SED of Mrk 817 changes between observations, as indicated by $\alpha_{OX}$. However, these changes are dominated by the change in X-ray luminosity/spectral shape. Therefore, it appears that the UV is not correlated with the X-ray and that even in our simultaneous UV/X-ray data, we find more variability in the X-ray than the UV. Unfortunately, however, there is no UVOT data for the two Swift observations with the lowest X-ray luminosities, making it impossible to determine whether the UV luminosity is low when the X-ray luminosity is at a much lower level than sampled with the available simultaneous data. Finally, our analysis of the UV and X-ray observations suggests that more data are needed before we can rule out the possibility that the Ly$\alpha$ absorbers are the result of an ionized outflow. While we did not find an X-ray \ion{O}{7} or \ion{O}{8} absorption edge strongly detected in our data, it is possible that a high signal-to-noise grating observation could reveal ionized X-ray absorption/emission at a lower column density than we could probe with the CCD spectra. Additionally, while the UV luminosity does not change significantly, it is possible that the X-ray spectrum during the 1997 GHRS observation was significantly different than in 2009, when the COS spectrum was taken. Further, our results of both a lack of a correlation between the UV and X-ray luminosities and a nearly constant UV luminosity over 30\,years while the X-ray luminosity changes by an order of magnitude, coupled with the possibility that changes in ionization of an outflow are the result of changes in the X-ray emission, suggest that the factor of 5 change in the Ly$\alpha$ absorption is potentially explained by a change in the X-ray spectrum. Therefore, future multi-wavelength observations, with COS and X-ray grating data from Chandra/XMM-Newton, of Mrk 817 are vital to unveil both the true nature of the absorbers and any potential connection between the UV and X-ray emission. %

1012.0592

We present an investigation of the ultraviolet and X-ray spectra of the Seyfert 1.5 galaxy Markarian 817. The ultraviolet analysis includes two recent observations taken with the Cosmic Origins Spectrograph (COS) in 2009 August and December, as well as archival spectra from the International Ultraviolet Explorer and the Hubble Space Telescope. Twelve Lyα absorption features are detected in the 1997 Goddard High Resolution Spectrograph (GHRS) and 2009 COS spectra—of these, four are associated with high-velocity clouds in the interstellar medium, four are at low significance, and the remaining four are intrinsic features, which vary between the GHRS and COS observations. The strongest intrinsic absorber in the 1997 spectrum has a systemic velocity of ~-4250 km s<SUP>-1</SUP>. The corresponding feature in the COS data is five times weaker than the GHRS absorber. The three additional weak (equivalent width from 13 to 54 mÅ) intrinsic Lyα absorbers are at systemic velocities of -4100 km s<SUP>-1</SUP>, -3550 km s<SUP>-1</SUP>, and -2600 km s<SUP>-1</SUP>. However, intrinsic absorption troughs from highly ionized C IV and N V are not detected in the COS observations. No ionized absorption signatures are detected in the ~14 ks XMM-Newton EPIC spectra. The factor of five change in the intrinsic Lyα absorber is most likely due to bulk motions in the absorber, since there is no drastic change in the UV luminosity of the source from the GHRS to the COS observations. In a study of the variability of Mrk 817, we find that the X-ray luminosity varies by a factor of ~40 over 20 years, while the UV continuum/emission lines vary by at most a factor of ~2.3 over 30 years. The variability of the X-ray luminosity is strongly correlated with the X-ray power-law index, but no correlation is found with the simultaneous optical/UV photometry.

2775278

2011A&A...527A...8S

1012

1012.5181_arXiv.txt

The transiting extrasolar planet (TEP) \object{WASP-7}\,b was discovered by the WASP consortium \citep[][hereafter H09]{Hellier+09apj} through the detection of transits in front of its F5\,V parent star. It is a challenging target for acquiring high-precision transit photometry, due to the brightness of the parent star ($V = 9.5$), the paucity of good nearby comparison stars, the transit duration (3.8\,hr), and the orbital period ($4.95$\,d) which is both comparatively long and close to an integer number of days. The characterisation of WASP-7 therefore relied upon the photometry obtained by the WASP-South telescope \citep{Pollacco+06pasp}. The relatively large scatter in the discovery data meant that the transit shape was poorly delineated. Because of this, the analysis by H09 included an additional constraint in the form of a main sequence mass--radius relation for the host star \citep[e.g.][]{Anderson+10apj}. \reff{The radius, surface gravity and density of the planet resulting from their analysis are $R_{\rm b} = \er{0.91}{0.046}{0.040}$\Rjup, $g_{\rm b} = \er{26.4}{4.4}{4.0}$\mss\ and $\rho_{\rm b} = \er{1.26}{0.25}{0.21}$\pjup, respectively. These values placed WASP-7\,b in an outlier position in the mass--radius diagram of TEPs, having one of the largest densities within the main planet population (masses $\la$2\Mjup). This was interpreted by H09 as evidence that WASP-7\,b has a massive heavy-element core.} In this work we present the first follow-up photometric observations obtained for WASP-7. The high precision of our observations (0.68\,mmag scatter) allows us to obtain the physical properties of the transiting system without needing to impose any constraints on the parameters of the parent star. We find a substantially larger radius, \reff{and therefore a lower density and surface gravity}. We also greatly improve the orbital ephemeris for the system, so transit midpoints in the 2011 observing season can be predicted to within 45\,s instead of 27\,min.

\begin{table} \begin{center} \caption{\label{tab:final} Final physical properties of WASP-7 compared to those found by H09. The second errorbars represent systematic errors.} \begin{tabular}{l r@{\,$\pm$\,}c@{\,$\pm$\,}l r@{\,$\pm$\,}l} \hline \hline \ & \mcc{This work (final)} & \mc{H09} \\ \hline $M_{\rm A}$ (\Msun) & 1.276 & 0.061 & 0.022 & \erc{1.28}{0.09}{0.19} \\ $R_{\rm A}$ (\Rsun) & 1.432 & 0.092 & 0.008 & \erc{1.236}{0.059}{0.046} \\ $\log g_{\rm A}$ (cgs) & 4.232 & 0.047 & 0.003 & \erc{4.363}{0.010}{0.047} \\ $\rho_{\rm A}$ (\psun) & \mcc{$0.434 \pm 0.074$} & \mc{ } \\[2pt] $M_{\rm b}$ (\Mjup) & 0.96 & 0.13 & 0.01 & \erc{0.96}{0.12}{0.18} \\ $R_{\rm b}$ (\Rjup) & 1.330 & 0.093 & 0.008 & \erc{0.915}{0.046}{0.040} \\ $g_{\rm b}$ (\ms) & \mcc{$13.4 \pm 2.6$} & \erc{26.4}{4.4}{4.0} \\ $\rho_{\rm b}$ (\pjup) & \mcc{$0.41 \pm 0.10$} & \erc{1.26}{0.25}{0.21} \\[2pt] \Teq\ (K) & \mcc{$1487 \pm 48$} & \erc{1379}{35}{23} \\ \safronov\ & 0.070 & 0.011 & 0.000 & \mc{ } \\ $a$ (AU) & 0.0617 & 0.0010 & 0.0004 &\erc{0.0618}{0.0014}{0.0033}\\ Age (Gyr) &\ermcc{2.4}{0.8}{1.0}{0.3}{0.4}& \mc{ } \\ \hline \end{tabular} \end{center} \end{table} \begin{figure} \includegraphics[width=\columnwidth,angle=0]{M2R2.eps} \caption{\label{fig:m2m3} Plot of the masses and radii of the known TEPs. The blue symbols denote values from the homogeneous analysis of \citet{Me10mn} and the red symbols results for the other known TEPs. WASP-7 is shown in black with an open circle (H09) and a filled circle (this work). Grey dotted lines show where density is 1.0, 0.5 and 0.25 \pjup.} \end{figure} \begin{figure} \includegraphics[width=\columnwidth,angle=0]{Porb-g2.eps} \caption{\label{fig:pg2} Plot of the orbital periods and surface gravities of the known TEPs. The symbols are the same as for Fig.\,\ref{fig:m2m3}.} \end{figure} We have observed a transit of the WASP-7 system using defocussed photometry techniques, resulting in a light curve with an rms scatter of only 0.68\,mmag. These data were modelled with the {\sc jktebop} code, and the results used to determine the physical properties of the planet and its host star. The higher quality of our photometry has allowed us to measure the system properties without constraining the host star to follow a main-sequence mass-radius relation. Compared to H09, we find both a larger radius for the host star and a larger ratio of the radii of the planet and star. This results in a sizeable increase in the planetary radius, from $\er{0.915}{0.046}{0.040}$\Rjup\ to $1.330 \pm 0.093$\Rjup, which in turn means a lower surface gravity and mean density. We have collected literature measurements of the physical properties of the 96 known TEPs as of 2010/11/18. For 30 of these objects we used the results of the homogeneous analysis performed by \citet{Me10mn}. In Fig.\,\ref{fig:m2m3} we plot their masses and radii, plus the values from WASP-7 found in this work (filled circle) and from H09 (open circle). The H09 results for WASP-7\,b indicated that it was one of the densest known TEPs below $\sim$2\Mjup, along with the more recently-discovered planets CoRoT-13\,b \citep{Cabrera+10aa} and HAT-P-15\,b \citep{Kovacs+10apj}. The outlier status has now been lifted: our results place WASP-7\,b in a well-populated part of the mass--radius diagram and demonstrate that high-quality data is required to obtain reliable measurements of the properties of TEPs. Fig.\,\ref{fig:pg2} shows the surface gravities of the known TEPs as a function of their orbital periods. The existing planet population shows an inverse correlation between period and surface gravity \citep{Me++07mn}, at least for the dominant population with periods $\la$10\,d and masses $\la$3\Mjup\ \reff{(see also \citealt{Fressin++09aa})}. The revised properties of WASP-7\,b move it from outlier status to within the sprawl of parameter space occupied by the general planet population. The theoretical models of irradiated gas giant planets \citet{Bodenheimer++03apj}, \citet{Fortney++07apj} and \citet{Baraffe++08aa} predict radii of no more than 1.16\Rjup\ for a 1.0\Mjup\ planet with a range of chemical compositions and core masses, and without an arbitrary additional heating source. Our upward revision of the radius of WASP-7\,b means that it no longer matches the predictions of these models at the 2$\sigma$ level. \reff{These conclusions could be strengthened by the provision of more radial velocity measurements as well as the acquisition of further photometric observations.}

1012.5181

We present the first high-precision photometry of the transiting extrasolar planetary system WASP-7, obtained using telescope defocussing techniques and reaching a scatter of 0.68 mmag per point. We find that the transit depth is greater and that the host star is more evolved than previously thought. The planet has a significantly larger radius (1.330 ± 0.093 R<SUB>Jup</SUB> versus ; R<SUB>Jup</SUB>) and much lower density (0.41 ± 0.10 ρ<SUB>Jup</SUB> versus ; ρ<SUB>Jup</SUB>) and surface gravity (13.4 ± 2.6 m s<SUP>-2</SUP> versus ; m s<SUP>-2</SUP>) than previous measurements showed. Based on the revised properties it is no longer an outlier in planetary mass-radius and period-gravity diagrams. We also obtain a more precise transit ephemeris for the WASP-7 system. <P />Based on data collected by MiNDSTEp with the Danish 1.54 m telescope at the ESO La Silla Observatory.Lightcurves data is only available in electronic form at the CDS via anonymous ftp to cdsarc.u-strasbg.fr (130.79.128.5) or via <A href="http://cdsarc.u-strasbg.fr/viz-bin/qcat?J/A+A/527/A8">http://cdsarc.u-strasbg.fr/viz-bin/qcat?J/A+A/527/A8</A>

3831654

2011A&A...528A..45H

1012

1012.5654_arXiv.txt

The study of periods of activity of Soft Gamma-ray Repeaters \citep[SGRs, for a recent review see ][]{2008A&ARv..15..225M} showing recurrent bursts with sub-second duration and much more extreme events known as giant flares (on rare occasions), may have an important input to our understanding of neutron stars (NSs). In particular, the detection and analysis of the presence of quasiperiodic oscillations (QPOs), up to a few kHz, has triggered a number of theoretical studies for prediction and direct comparison with various equations of state for superdense matter. These oscillations have been interpreted initially as torsional oscillations of the crust \citep{SA2007,SKS2007}. While later the Alfv\'en oscillations of the fluid core have been taken into account \citep{2007MNRAS.377..159L,2008MNRAS.385L...5S,CBK2009,DSF2009,2010arXiv1006.0348V,2010arXiv1007.0856G,CK2010Pre}. These studies suggested how the observations can constrain both the mass, the radius, the thickness of the crust and the strength of magnetic field of NSs. In particular, the timing analysis of the decaying tail of the unprecedent giant flare of SGR 1806$-$20 on 2004 Dec 27, allowed to detect QPO frequencies approximately at 18, 26, 30, 92, 150, 625, and 1840 Hz \citep{I2005,WS2006a,SW2006,WS2006b} in different time intervals, different rotation phase and different amplitudes of oscillations, by means of computation and analysis of the averaged power spectrum. Here, we present the results of a Bayesian approach of timing analysis, of the giant flare data set of SGR\, 1806$-$20 registered on 2004 Dec 27 by \emph{RXTE PCA}, for detecting a periodic signal of unknown shape and period developed by \cite{GL1992} \citep[for its sensitivity and advantages see further and also][]{GL1996}.

\begin{enumerate} \item We found new {\bf oscillationsof transient nature} applying Bayesian timing analysis method of the decaying tail of the giant flare of SGR\, 1806$-$20 2004 Dec 27, observed by \emph{RXTE PCA}, not yet reported in the literature. \item Some of those {\bf oscillations} frequencies ($\mathrm{f}_{QPOs} \sim 17,22,37,56,112 \mathrm{Hz}$) are predicted by the theoretical study of torsional Alfv\'en oscillations of magnetars \citep[see, Table 2, by][]{CBK2009}, suggesting $APR_{14}$ \citep{1998PhRvC..58.1804A} EoS\footnote{Neutron star models based on the models for the nucleon-nucleon interaction with the inclusion of a parameterized three-body force and relativistic boost corrections, estimating maximum mass and stiffness \citep[see,also][]{1999ApJ...525L..45H}.} of SGR\, 1806$-$20. \item These preliminary results are very promising and we plan to extend our high frequency oscillations research (both the theoretical predictions as well as the observations) to both activity periods, as well as to the quiescent state of SGRs and AXPs, as well as to the isolated neutron stars with comparatively smaller magnetic fields. \end{enumerate}

1012.5654

<BR /> Aims: By detecting high-frequency quasi-periodic oscillations (QPOs) and estimating their frequencies during the decaying tail of giant flares from soft gamma-ray repeaters (SGRs) useful constraints for the equation of state (EoS) of superdense matter may be obtained via comparison with theoretical predictions of eigenfrequencies. <BR /> Methods: We used the data collected by the Rossi X-Ray Timing Explorer (RXTE/XTE) Proportional Counter Array (PCA) of a giant flare of <ASTROBJ>SGR 1806-20</ASTROBJ> on 2004 Dec. 27 and applied an existing Bayesian periodicity detection method to search for oscillations of a transient nature. <BR /> Results: In addition to the already detected frequencies, we found a few new frequencies (f<SUB>QPOs</SUB> ~ 16.9,21.4,36.4,59.0,116.3 Hz) of predicted oscillations based on the APR<SUB>14</SUB> EoS for <ASTROBJ>SGR 1806-20</ASTROBJ>.

12213442

2011MNRAS.415.3119G

1012

1012.4652_arXiv.txt

\label{sec:intro} The standard unified scheme for active galactic nuclei (AGN) attempts to explain the observational appearance of different AGN types mainly as an inclination effect \citep{antonucci1993,urry1995}. The continuum emission being produced by the accretion disc of the central supermassive black hole (SMBH) irradiates the surrounding broad line region (BLR) and gives rise to broad optical/UV line emission. Obscuring equatorial dust at larger distances plays a crucial role in the unified model as at higher inclinations it covers the BLR and thereby disentangles type-1 AGN, which spectroscopically show broad line emission, from type-2 objects, which do not. The accretion process onto the SMBH partly leads to re-ejection of matter and supposedly causes ionised winds in the polar direction. The presence of these winds has brought strong support to the unified model as it allowed to indirectly detect hidden BLRs in obscured type-2 objects by scattering of BLR light around the torus \citep{antonucci1985}. Systematic studies in optical spectropolarimetry lead to a classification of Seyfert galaxies according to the direction of the optical polarisation angle \citep[see e.g.][and references therein]{antonucci1984,smith2002,smith2004} that can be explained in the framework of the unified scheme when including equatorial scattering inside the torus funnel. The simplest approach of the unified model assumes that the rotation axis of the accretion disc is aligned to the axis of the torus and of the outflows. The alignment of these components is justified by assuming a symmetric mass transfer from the inner torus towards the outer accretion disc as well as by the symmetric collimation effect that the torus should have on the polar outflows. But recent work by \citet{raban2009} surprisingly suggests that the outflows of the well-studied Seyfert-2 galaxy NGC~1068 are inclined by $18\degr$ with respect to the apparent axis of the torus \citep[see Fig.~9 in][]{raban2009}. If this result is confirmed it has important consequences for our understanding of the accretion and ejection mechanisms close to active SMBHs. In this work, we explore to which extend the misalignment of the outflows with respect to the torus in NGC~1068 can be probed by X-ray polarimetry observations. With the advent of a new generation of X-ray polarimeters based on photo-electric effects, X-ray imaging polarimetry becomes feasible across the 2--35~keV energy range \citep{soffitta2010}. The X-ray emission from AGN is much less sensitive to stellar contributions from the host galaxy than optical/UV radiation is, and therefore it offers a better probe of the reprocessing geometry in the innermost regions of NGC~1068. One could expect that the radiative coupling of centrally emitted X-rays to the torus and the outflows is somewhat similar to the observed behaviour in the optical/UV as in both wavebands electron scattering plays an important role. But in the hard X-ray band electron scattering becomes wavelength-dependent, which should have an impact on the polarisation. In the following, we build up a sophisticated X-ray reprocessing model of NGC~1068. We start out with a simple irradiated slab representing the accretion disc and an elevated primary source. Then, we add dust obscuration in a torus geometry testing for different optical depths. Next, we include inclined polar outflows and conduct calculations for different optical depths of these ionisation cones. In the soft X-ray band with photon energies $E < 10$~keV absorption processes in the torus are very important and at high inclination angles the primary source and reflection from the accretion disc should only be seen indirectly by scattering in the outflows. The scattering in the ionisation cones should then imprint a polarisation position angle that is perpendicular to the axis of the outflow. At higher energies, when scattering off the torus becomes more prominent, the position angle should change and be determined by the torus geometry. This switch of the polarisation angle depends on the relative Stokes fluxes coming from the ionisation cones and from the torus. Therefore, it must vary with the optical depths in both components as well as with the inclination angle. Finally, we also examine the impact of additional equatorial scattering in the farther away components of the accretion flow that are located between the inner surfaces of the torus and the outer accretion disk. Such an additional scattering component produces polarised reprocessing by itself but also changes the irradiation geometry of the inner surfaces of the torus. We test the resulting spectra and polarisation properties for different optical depths of the obscuring dust and of the scattering electrons in the ionisation cones. Aside from the optical depths, our model comprises other interesting parameters such as the half-opening angle of the torus or of the outflows. Their impact is going to be examined elsewhere. This study is particularly designed to predict the expected X-ray polarisation signal if the geometry inferred by \citet{raban2009} is correct. We therefore fix all dimensions and angles according to their findings. The radiative transfer code applied here is presented in Sect.~\ref{sec:tools} and the details of the model geometry are described in Sect.~\ref{sec:geom}. The results for different combinations of the individual reprocessing regions and then for the case of NGC~1068 follow in Sect.~\ref{sec:results}. We close with a discussion in Sect.~\ref{sec:discuss} where we also shed light on the prospects for future X-ray polarimetry observations of NGC~1068.

\label{sec:discuss} \subsection{Directly measuring the orientation of the inner outflows} \label{sec:discuss_orient} We have conducted accurate modelling of the expected X-ray polarisation induced by complex reprocessing in the active nucleus of NGC~1068. This work is motivated by the apparent misalignment of the ionisation cones with respect to the torus axis as discussed by \citet{raban2009}. The orientation of the ionisation cones was inferred from slit spectroscopy of the narrow line region (NLR) using the {\it HST STIS} in combination with kinetic modelling \citep{das2006}. But \citet{raban2009} also point out that this orientation is still a matter of debate. From the observation of infrared emission lines, Poncelet, Sol \& Perrin (2008) deduce a different position angle of the NLR that is more in agreement with the {\it HST} imaging results of [OIII] emission reported by \citet{evans1991}. But then the UV-imaging probably suffers from significant foreground absorption. In addition to that, we want to point out that all these measurements are taken on large spatial scales reaching out to 100~parsec from the central engine. It is not straightforward to derive from these observations what the actual geometry of the ionisation cones across a few parsec is. One could imagine that the outflows become deflected or twisted on intermediate scales and thus their orientation would vary with distance. As \citet{raban2009} conclude, further investigation is needed to really constrain the geometry of the innermost outflows. In this context, X-ray polarimetry is going to give unambiguous constraints. The scattered X-ray emission must come from very close to the central engine, as is indicated from X-ray spectroscopy results revealing the presence of high-ionisation emission lines of iron \citep{marshall1993,matt2004}. Furthermore, any foreground absorption effects seen in the UV should much less interfere in the X-ray range so that X-ray polarimetry enables a direct view on the innermost parts of the outflow. Our modelling results have shown that the polarisation position angle in the soft X-ray range is clearly dominated by the polar scattering, which enables a precise measurement of the orientation of the scattering material. \subsection{Determining the misalignment with the torus} Another aim of this investigation is to verify if the rotation $\Delta \psi$ of the X-ray polarisation angle between soft and hard X-ray energies can be used to measure the misalignment between the polar scattering regions and the obscuring torus. We have found that coherent modelling of the different scattering components leads to values of $\Delta \psi$ that at heavily obscured viewing directions allows one to distinguish the torus from the ionisation cones. The relation between $\Delta \psi$ and the actual misalignment also depends on $i$. The highest values of $\Delta \psi$, which are most favourable for future X-ray polarimetry observations, are found at moderate viewing angles. At strictly edge-on viewing directions, $\Delta \psi$ becomes very small, especially for higher optical depths of the ionisation cones. But, as we discuss in the following, our modelling approach is rather conservative, and in reality we should expect more favourable conditions for the measurement of $\Delta \psi$. The results we present in this work give lower limits on $\Delta \psi$ for NGC~1068 assuming the least favourable conditions for a rotation of the polarisation angle between low and high photon energies. \subsubsection{Optical depth and geometry of the outflow} \label{sec:discuss_cones} The net Stokes flux and therefore the resulting polarisation angle is determined by the competition between the disc/torus system and the polar scattering cones. Especially at higher inclination the polar scattering has a strong impact as the scattering-induced polarisation rises towards orthogonal scattering angles. The efficiency of polar scattering depends on $\tau_{\rm cone}$ and is low for optically thin scattering cones. It rises with increasing $\tau_{\rm cone}$ and then goes through a maximum until multiple scattering effects become important and diminish the resulting polarisation. Note also, that the Stokes flux from the double-cone results from integrating photons over its entire opening angle. This leads to combining scattered photons having a large range of polarisation orientations and also diminishes the net polarised flux. Overall, the case of $\tau_{\rm cone} = 0.3$ that we investigate produces a strong Stokes flux but at most inclinations it remains comparable to the polarised flux coming from the disc/torus system and therefore the rotation $\Delta \psi$ is significant. Having in mind the unified scheme of AGN, it is important that $\tau_{\rm cone}$ remains low enough as the ionised winds are seen in transmission in Seyfert-1 galaxies, where they produce warm absorption in the soft X-ray band. Observations of the warm absorber rule out that the medium has a higher column density than $10^{23} {\rm cm}^{-2}$ for most objects \citep[see][and references therein]{turner2009}. Only a few exceptions exceed this threshold so that our choice of $\tau_{\rm cone} = 0.3$, which corresponds to a column density around $4.5 \times 10^{23} {\rm cm}^{-2}$, is a good estimate for the upper limit. Note also that for NGC~1068 the column density of the polar scattering medium has been constrained from the observed emission measure distribution to $4 \times 10^{21} {\rm cm}^{-2} < \log{N_{\rm H,Cone}} < 7 \times 10^{22} {\rm cm}^{-2}$ \citep{kinkhabwala2002}. Furthermore, we assume in our modelling that the matter is uniformly distributed inside the double-cone, whereas the approach by \citet{das2006} is based on outflows with a shell-structure, such as predicted for magnetically driven outflows \citep{konigl1994}. For all these reasons, the Stokes flux coming from the ionisation cones of NGC~1068 should be lower than we assume in our modelling, which enables a larger $\Delta \psi$. The misalignment of the ionisation cones with respect to the torus axis as derived by \citet{raban2009} refers to the projection of the double-cone onto the plane of the sky. But if the outflows are rotated in azimuth with respect to the $yz$-plane, it is possible that the real misalignment with respect to the torus axis is larger than $18\degr$. Note that a slight inclination of the double-cone towards the observer is suggested by the kinematic modelling of \citet{das2006}. Then, the average scattering angle for the polar electron scattering is no longer close to $90\degr$ but will shift towards forward and backward scattering. This reduces the polarised flux coming from the outflows. At the same time, a stronger inclination of the cones also requires a larger half-opening of the torus, otherwise the outflows would be blocked. The scattering distribution of the torus is thus geometrically flatter in this case, which should increase its Stokes flux. Both effects work in favour of higher values for $\Delta \psi$ than found from the models presented here. \subsubsection{Clumpiness of the torus material} \label{sec:discuss_torus} We are also pessimistic in our assumptions about the torus scattering properties. In our modelling we assume a maximum optical depth by setting $N_{\rm H,tor} = 10^{27} {\rm cm}^{-2}$. This agrees with previous spectroscopic studies showing that the obscuring material along the line-of-sight toward NGC~1068 is entirely opaque. But for most viewing angles we are still more than one order above the lower threshold of the observed hydrogen column density, which was given by \citet{matt1997} and \citet{bianchi2001} based on {\it ASCA} and {\it BeppoSAX} data. It is thus likely that the torus is less optically thick than we assumed in our modelling. A lower $N_{\rm H,tor}$ increases the polarised flux of the torus as we show in Fig.~\ref{fig:torPF}. Between the cases of lower and higher $N_{\rm H,tor}$ there is an increase in Stokes flux at 20~keV by a factor of 2.5 at $i \approx 76\degr$ and of more than a factor of 10 at $i \approx 87\degr$. \begin{figure} \centering \includegraphics[width=8.2cm]{disk_torus1e25_1e27_no_cones_PF.eps} \caption{Modelled polarised flux spectra for a system of an irradiated accretion disc and an equatorial torus with $N_{\rm H,tor} = 10^{27} {\rm cm}^{-2}$ (solid line) and $N_{\rm H,tor} = 10^{25} {\rm cm}^{-2}$ (dashed line) for two different obscured viewing angles.} \label{fig:torPF} \end{figure} The exact geometry of the obscuring dust region remains a matter of debate, although infrared interferometry observations start putting robust constraints on the geometrical size of the dusty region \citep{jaffe2004,ramos2009,hoenig2010a}. The latest generation of infrared radiative transfer models allows for a satisfying description of the data by assuming that the dust region is clumpy \citep{nenkova2008a,nenkova2008b,schartmann2008,hoenig2010b}. Such a clumpy composition of the torus could also address criticism that the unified model could not properly explain how the observed spatial extension of the dust region can be maintained against self-gravity. The clumpy structure would argue for a dynamical model in which the torus is not in hydrostatic equilibrium \citep{krolik1988}. In the results presented here we do not include the effect of clumpiness on the resulting polarisation. But if the torus medium is clumpy and at the same time dynamical, there is a better chance for fluctuations in the column density along a given line-of-sight. Observing NGC~1068 with a lower opacity $N_{\rm H,tor}$ is thus more likely if the medium is clumpy instead of uniform. As shown above, a lower opacity enhances the Stokes flux from the torus and thus the rotation $\Delta \psi$. \subsubsection{The primary source and equatorial scattering} \label{sec:discuss_primary} We assume that the polarisation coming from the accretion disc is based on reprocessing across the whole disk surface. The extended disk has a large radius compared to the height of the X-ray source above. Then, the outer disc regions produce polarisation with $\psi = 90\degr$. This aligned polarisation is being diminished by the Stokes flux coming from the disk centre and having $\psi = 0\degr$. As discussed in Sect.~\ref{subsec:irrad_disc}, such a setup assumes that the primary source is located far above the black hole so that the central hole of the disk does not have any significant impact. When relaxing this assumption and considering sources that are much closer to the black hole, the reprocessing geometry changes \citep[see e.g.][]{martocchia1996,miniutti2004,schnittman2010,dovciak2011} and the missing part of the disc inside the marginally stable orbit becomes significant for the Stokes flux. The net polarisation with $\psi = 90\degr$ of the disc is thus stronger for such a case, which works in favour of a larger value for $\Delta \psi$. Assuming a different irradiation geometry, for instance a central but more extended hot corona, would not change this result much as is implied by the results of \citet{matt1993}. The main reason for this is that the disc remains very large with respect to the size of the corona. We have shown that generally equatorial scattering leads to a lower $\Delta \psi$, in particular at high inclinations of the observer. But note that we assume the equatorial scattering being entirely due to electrons. This is an extreme example as beyond the BLR also moderately ionised and neutral medium should exist. Less ionised matter has a different scattering efficiency and should absorb more strongly in the soft X-ray band and thereby maintain higher values of $\Delta \psi$. The two exemplary cases shown in Sect.~\ref{sec:equatscat} rather define lower limits of the expected $\Delta \psi$ when equatorial scattering is important. \subsection{Encouraging prospects for future X-ray polarimetry of NGC~1068} The observational window of X-ray polarimetry will soon be (re-)opened by the NASA {\it Gravity and Extreme Magnetism Small Explorer (GEMS)} that is currently prepared for launch in 2014 \citep{swank2009,jahoda2010}. The satellite will be equipped with a soft X-ray polarimeter reaching up to 10~keV. Confirmation of the polarisation position angle induced by scattering in the misaligned outflows of NGC~1068 will thus principally be possible with {\it GEMS}. For the practical measurement, the collecting area of the mirrors might still be a limiting factor. On the other hand side, {\it GEMS} is entirely devoted to X-ray polarimetry and allows for long exposure times. The technology is also ready for mid-size observatories that include imaging polarimetry at high angular resolution and over the whole range of 2~keV to 35~keV. Such X-ray polarimeters could efficiently observe near-by AGN. We use the response matrices of the polarimeters designed for the formerly proposed {\it New Hard X-ray Mission} \citep{tagliaferri2010,soffitta2010}. It turns out that a 300~ks observation of NGC~1068 would be sufficient to constrain the polarisation angle with a statistical uncertainty of $5\degr$ if the polarisation degree is 10 per cent across 6~keV to 35~keV. For higher values of $P$, the precision on the measurement rises; a polarisation of 20 per cent constrains the angle with an uncertainty of $2.5\degr$. At lower photon energies, the performance of the X-ray polarimeter is expected to be even better. The soft X-ray polarisation obtained by our modelling is at least comparable to and often significantly higher than 20 per cent as shown in Table~\ref{tab:dpsi}. Therefore, a future mid-size X-ray observatory with broad band polarimetry capabilities could directly measure the orientation of the innermost ionisation cones as described in Sect.~\ref{sec:discuss_orient}. If we assume a viewing angle of $i = 70\degr$ towards NGC~1068, as found by \citet{hoenig2007} from applying a clumpy torus radiative transfer model to infrared data, then the expected values for $\Delta \psi$ given in Table~\ref{tab:dpsi} are also within the measurable limits. Should the inclination angle be higher as might be suggested by water maser observations \citep{gallimore1996,greenhill1996} then the possibility to measure $\Delta \psi$ depends on the model details. We consider the values given in Table~\ref{tab:dpsi} as lower limits that are based on a conservative modelling approach. Note that if one decides to include both, a soft and a hard X-ray polarimeter (up to about 25~keV) in the large mission design of the {\it International X-ray Observatory} \citep{barcons2011}, then measurements of the broad band polarisation still further improve. The method described here might then be applied systematically to a larger sample of objects. In summary, we therefore think that future X-ray polarimetry over a broad energy range is a promising tool to disentangle the details of the inner AGN geometry.

1012.4652

We model the expected X-ray polarization induced by complex reprocessing in the active nucleus of the Seyfert 2 galaxy NGC 1068. Recent analysis of infrared interferometry observations suggests that the ionized outflows ejected by the central engine are not aligned with the symmetry axis of the obscuring torus. This conclusion was obtained by extrapolating the apparent orientation of the narrow-line region to the inner parts of the ionization cones. We show that future measurements of the soft X-ray polarization vector unambiguously determine the orientation of the ionization cones. Furthermore, X-ray polarimetry across a broad photon energy range may independently verify the misalignment between the ionization cones and the axis of the torus. To model the expected polarization percentage and position angle, we apply the radiative transfer code STOKES. Reprocessing of the primary X-ray radiation takes place in the accretion disc, the surrounding equatorial torus and the inclined, ionized outflows. We also examine additional equatorial scattering occurring in between the accretion disc and the inner surfaces of the torus. Radiative coupling between the different reprocessing components is computed coherently. The resulting polarization properties depend on the optical depth of the reprocessing regions and on the viewing angle of the observer. We show that even under unfavourable conditions the misalignment of the outflows with respect to the torus axis can be determined from a rotation of the polarization position angle between softer and harder X-rays. We argue that the misalignment of the outflows with respect to the torus axis in NGC 1068 may be constrained by a future X-ray mission if equipped with a broad-band polarimeter.

12164533

2011AJ....141..105P

1012

1012.4187_arXiv.txt

With the completion of the Sloan Digital Sky Survey \citep[SDSS;][]{York00}, the era of high quality, homogeneous CCD imaging over large fractions of the sky has arrived. The final data release from the SDSS, \citep[DR7;][]{Abazajian09} contains over 11,000 deg$^{2}$ of imaging data, with 357 million unique objects being identified in a $\sim 60$ Terabyte database. \begin{table*} \begin{center} \caption{Descriptions of previous catalogs of quasars. $^{a}$O/I, Optical or Infrared; P/S, Photometric or spectroscopic survey. $^{b}$SDSS Faint Quasar Survey. $^{c}$The FIRST-2MASS Red Quasar Survey. $^{d}$The SDSS DR7 Quasar catalog provides NIR photometry from 2MASS for \hbox{53,584} objects, \hbox{29,551} or which are detected in the {\it K}-band. $^{e}$The 2MASS Second Incremental Data Release. } \setlength{\tabcolsep}{4pt} \begin{tabular}{lrrccl} \hline \hline Survey & Area (deg$^{2}$) & N$_{\rm Q}$ & Magnitude Range & O/I \& P/S$^{a}$ & Reference \\ \hline COMBO-17 & 0.8 & \hbox{ 192} & $R<24$ & O/P & \citet{Wolf03} \\ SFQS$^{b}$ & 4 & \hbox{ 414} & $g<22.5$ & O/S & \citet{Jiang06} \\ SDSS-ULAS DR1 & 189 & \hbox{ 2 873} & $K<18.2$ & I/S & \citet{Chiu07} \\ 2SLAQ QSO & 190 & \hbox{ 8 764} & $18.00 < g < 21.85$ & O/S & \citet{Croom09a} \\ 2QZ & 700 & \hbox{ 23 338} & $18.25<b_{\rm J}<20.85$ & O/S & \citet{Croom04} \\ SDSS-ULAS DR3 & 1 200 & \hbox{ 74 351} & $K<18.4$ & I/P/S & this work \\ FIRST-2MASS RQS$^{c}$ & 2 716 & \hbox{ 57} & $K \leq 14.3, (R-K)>4, (J-K)>1.7$ & I/S & \citet{Glikman07} \\ SDSS DR6pQ & 8 342 & \hbox{1 015 082} & $i<21.3$ & O/P & \citet{Richards09} \\ SDSS DR7Q$^{d}$ & 9 380 & \hbox{ 29 551} & $K \lesssim 17.0$ & I/S & Schneider et al. 2010\\ SDSS DR7Q & 9 380 & \hbox{ 105 807} & $i<20.2$ & O/S & Schneider et al. 2010 \\ 2MASS 2IDR$^{e}$ & $\sim30$ 000 & \hbox{2 277} & $K \leq 15$ & I/S & \citet{Barkhouse01} \\ \hline \hline \label{tab:previous_surveys} \end{tabular} \end{center} \end{table*} The identification of quasars was a major focus of the SDSS project; utilizing the broad 5-filter photometry \citep{Richards02} led to efficient selection \citep{Vanden_Berk05, Richards06} of low, $z<2.2$, and high, $z\gtrsim 3.5$, redshift quasars, identified via their spectroscopic signatures \citep[][ and references therein]{Schneider10}. However, quasars can also be efficiently identified by their SDSS imaging properties alone \citep{Richards01}, due to their point-source appearance, but non-stellar location in color-color space. With the most recent catalog of \citet[][hereafter R09]{Richards09}, based on imaging from the SDSS, the number of {\it photometrically identified} quasars now stands at over 1 million. Both the SDSS spectroscopic and photometric quasar catalogs have been used to investigate global quasar properties such as the luminosity function \citep[QLF; ][]{Fan01, Richards06, Croom09b} and clustering \citep{Myers06, Myers07a, Shen07, Shen09, Ross09}. The emergence of large surveys has not been confined to the optical regime. In the near-infrared (NIR; $\lambda \approx 1-5\mu$m) quasar catalogs \citep[e.g.][]{Barkhouse01, Cutri02, Francis04, Ofek07, Kouzuma10a, Kouzuma10b} have been available since the completion of the 2 Micron All Sky Survey \citep[2MASS; ][]{Skrutskie06}. Table~\ref{tab:previous_surveys} summarizes how previous quasar surveys, both in the optical and NIR, compare by area, number of objects and magnitude range. Quasar observations in the NIR are particularly important for individual objects that are seen only in the reddest, or indeed potentially none, of the optical filters \citep[e.g. ][]{Fan06, Venemans07, Willott09}. Observations of the general quasar population in the {\it K}-band are key since this links the rest-frame ultraviolet (UV)/optical to the mid-infrared (MIR; $\lambda \approx 5-30\mu$m); the former being where there is the peak in the radiative output for Type I, non-obscured quasars, \citep[e.g. ][]{SS73, Sanders89, Kishimoto08, RIchards09b}, and the latter where reprocessed light heats intrinsic dust to $\sim 30-300$ K \citep[e.g. ][]{Pier93, Efstathiou95, Lacy04}. Observations in the observed {\it K}-band also measures the rest-frame {\it i} and {\it g}-bands at redshifts $z\sim1.9$ and $z\sim3.7$, respectively. However, it is only the bright, $g\lesssim16$ magnitude quasars that are detected by the relatively shallow limits of the 2MASS survey, and the majority of known quasars are fainter than this in the NIR bands. The UKIRT Infrared Deep Sky Survey \citep[UKIDSS; ][]{Lawrence07}, a seven-year sky survey which began in 2005 May, has five different survey components. The UKIDSS ``Large Area Survey'' (ULAS) aims to reach $\sim 4$ magnitudes deeper than 2MASS over an area of up to 4,000 deg$^{2}$, directly overlapping the optical imaging footprint of the SDSS. The primary goal of this paper is to create a catalog of over 130,000 quasars with optical, {\it ugriz} \citep{Fukugita96}, and NIR, {\it YJHK} photometry, with an areal coverage of \hbox{1 200}~deg$^{2}$. By matching the catalogs of R09 in the optical regime to that of the ULAS in the NIR we produce, by a factor of at least two, the largest catalog of quasars detected in the {\it K}-band. The major motivation for the catalog will be its utilization in an future study where we measure the $K$-band quasar luminosity function (Peth, Ross et al. in prep.). By concentrating on the $\approx 200$~deg$^{2}$ area from the SDSS known as ``Stripe 82'', our {\it K}-band quasar luminosity function will show the evolution from redshift of zero to two. The analysis by \citet{Trammell07} is an example of the synergy produced by the matching and production of a multi-wavelength catalog. These authors match $\sim$6000 SDSS quasars to UV data provided by the {\it Galaxy Evolution Explorer} \citep[{\it GALEX}; ][]{Martin05} satellite and find that over 80$\%$ of the optically detected quasars have near-UV detections. The quasars are well separated from stars in UV-optical color-color space. The large sample size allows for the construction of SEDs in bins of redshift and luminosity, which shows the median SED becoming bluer at UV wavelengths for quasars with lower continuum luminosity. \citet{Ball07} also perform catalog matching using SDSS quasar data and {\it GALEX} UV photometry, with the goal of understanding quasar photometric redshift properties. The studies by \citet{Warren00}, \citet{Croom01KX}, \citet{Sharp02} and more recently, \citet{Maddox08} and \citet{Smail08}, are another motivation why a sample of quasars with NIR photometric properties is desired. These authors show that using a ``KX selection'', where the quasar SED shows an excess in the {\it K}-band compared to a stellar SED, can successfully identify quasar candidate objects that would be normally excluded from the SDSS (optical) quasar selection algorithm - even for dust reddened quasars. This is an important result since selecting complete quasar samples via the KX method opens up the possibility of investigating the ``Quasar Epoch'' over the redshift range of $2.2<z<3.5$, where current, usually optically selected, quasar samples are particularly poorly represented. A similar project to \citet{Maddox08} was \citet{Nakos09}, who also select quasar candidates using the KX-technique, where these authors identify quasars on the basis of their optical ($R$ and $z^{'}$) to NIR ($K_{\rm s}$) photometry and point-like morphology. \citet{Jurek08} also test the KX method and find that it is more effective than the traditional ``UV Excess'' (UVX) selection method at finding red, ($b_{\rm J} - K) \geq 3.5$, quasars. There are two comparable studies and samples to our own work. \citet{Chiu07} and \citet{Souchay09}, the latter recently producing the ``Large Quasar Astrometric Catalog'' (LQAC). Our work differs from these studies, in two key ways; ({\it i}) we have over an order of magnitude more objects in our sample compared to \citet{Chiu07}; ({\it ii}) our catalog uses UKIDSS data, as opposed to 2MASS data \citep{Souchay09}, where the former is much better matched to the SDSS imaging depth. This paper is organized as follows. In Section 2 we present our sample, giving a brief overview of the SDSS and UKIDSS. Section 3 lays out our catalog. In Section 4 we present $N(z)$, color-redshift and color-color relations from our matched catalog. In Section 5 we present analysis of high-redshift quasars and calculations of the {\it i} and {\it K}-band number counts. The Appendix gives further details on magnitude conversions and cross-checks of our study. We assume the currently preferred flat, ``Lambda Cold Dark Matter'' ($\Lambda$CDM) cosmology where $\Omega_{\rm b}$ =0.042, $\Omega_{\rm m}$ = 0.237, $\Omega_{\Lambda}$ = 0.763 \citep{Sanchez06, Spergel07} and quote distances in units of $\hmpc$ to aid in ease of comparisons with previous results in the literature. Where a value of Hubble's Constant is assumed, e.g. for absolute magnitudes, this will be quoted explicitly. All optical magnitudes are based and quoted on the AB zero-point system \citep{Oke83}, while all NIR magnitudes are based in the {\it Vega} system, with conversions from AB to Vega given in the Appendix.

\begin{itemize} \item{The positional standard deviation of the SDSS Quasar to ULAS matches is $\delta_{\rm R.A.} = 0.1370''$ and $\delta_{\rm Decl.}= 0.1314''$. We find an absolute systematic astrometric offset between the SDSS Quasar catalog and the ULAS, of $|{\rm R.A._{offset}}| = 0.025''$, and $|{\rm Decl._{offset}}| = 0.040''$; we suggest the nature of this offset to be due to the matching of catalogs, rather than image level data.} \item{Our matched catalog has a surface density of $\sim$108 deg$^{-2}$, for objects detected in any of the four NIR bands, and $\approx53$ deg$^{-2}$ for objects $K\leq18.27$. This compares to the $\approx122$ deg$^{-2}$ found from the R09 DR6pQ catalog and $85-150$ deg$^{-2}$ down to $K\leq 20$ from \citet{Smail08}.} \item{Tests using our matched catalog, along with data from the UKIDSS DXS, implies that our limiting magnitude is $i\approx20.6$ and that the reddest, $(r-K) \geq 5.0$ objects turn out to be either associated with, or contaminated by, a foreground extended source.} \item{We plot redshift histograms for the Stripe 82 subsample, $K$-matches sample and total matched sample. The photometric $N(z\sim2.3)$ spike seen for the ``total sky'' sample appears to be dramatically reduced (although does not disappear completely) for the {\it K}-band matches. The photometric and spectroscopic histograms now seem to be in reasonable agreement, though there are possible deficiencies of photometric objects at $z\sim0.7$ and in particular, $z\sim1.9$.} \item{Color-redshift diagrams, for the optical and NIR, show the close agreement between our matched catalog and the ``average'' models of \citet{Hewett06}, at redshift $z\lesssim2.0$. At higher redshifts, the models generally appear to be generally bluer than the mean observed quasar colors. We {\it tentatively} suggest this is the same affect as has recently been reported in \citet[][]{Worseck10}, namely, that the original SDSS color-selection that was used to select quasars (and ultimately our own matched catalog) may be systematically biased towards missing blue $(u-g)\leq2.0$ quasars at $z\sim2.2-3.5$, due to the selection preferentially selecting intervening \ion{H}{1} Lyman limit systems.} \item{Since stellar contamination is of great interest for ongoing quasar surveys such as the SDSS-III:BOSS and AUS, we plot a test set of stellar data with our matched catalog data in $gJK$ and $giK$ color space. We confirm findings from previous studies, in particular \citet{Maddox08}, that $(a)$ the stellar locus traces out a clear band in $gJK$ color-space, related to the different stellar spectral type and $(b$ that in general, quasars, especially those with $2.2 < z < 3.2$ mostly lie in a distinct region of $gJK$ color space than stars.} \item{We plot the $iYJ$ colors of our matched catalog for high, $z>4.6$, redshift spectroscopically confirmed quasars, again comparing to model tracks, though no obvious trends are seen. Finally, just using the ULAS DR3 and the very high, $z>5.7$, redshift objects reported in \citet{Fan06}, \citet{Jiang08}, \citet{Jiang09} and \citet{Mortlock09}, we find 6 (5) out of 13 (12) quasars have NIR ($K$) band detections.} \end{itemize} It is worthwhile to mention that many of our tests have also been performed very recently in \citet{Wu10}, and although we have not performed any direct comparisons between the results found herein and their study, we see very good general agreement, e.g. for color-$z$ relations, between the two investigations. The major motivation for the construction of the matched catalog presented here, was its utilization in a future study where we measure the $K$-band quasar luminosity function. As quasars are measured further in the NIR, the flux due to host galaxies is no longer negligible but will rather constitute a sizeable percentage of the total bolometric flux from a quasar. Our future work (Peth, Ross et al. in prep.), will address this, and other issues in order to construct and measure the observed $K$-band quasar luminosity function. Advancement can come from technological breakthroughs as well as new theoretical insights. Future telescopes and surveys, e.g. VISTA-VHS and the SASIR, will not only cover more of the sky but should also be able to observe to greater depths. Future observations will warrant an updated analysis of quasar properties in multiple bands, and will be both necessary and essential to further understand the formation and evolution of quasars.

1012.4187

We present a catalog of over 130,000 quasar candidates with near-infrared (NIR) photometric properties, with an areal coverage of approximately 1200 deg<SUP>2</SUP>. This is achieved by matching the Sloan Digital Sky Survey (SDSS) in the optical ugriz bands to the UKIRT Infrared Digital Sky Survey (UKIDSS) Large Area Survey (LAS) in the NIR YJHK bands. We match the ≈1 million SDSS DR6 Photometric Quasar catalog to Data Release 3 of the UKIDSS LAS (ULAS) and produce a catalog with 130,827 objects with detections in one or more NIR bands, of which 74,351 objects have optical and K-band detections and 42,133 objects have the full nine-band photometry. The majority (~85%) of the SDSS objects were not matched simply because these were not covered by the ULAS. The positional standard deviation of the SDSS Quasar to ULAS matches is δ<SUB>R.A.</SUB> = 0farcs1370 and δ<SUB>decl.</SUB> = 0farcs1314. We find an absolute systematic astrometric offset between the SDSS Quasar catalog and the UKIDSS LAS, of |R.A.<SUB>offset</SUB>| = 0farcs025 and |decl.<SUB>offset</SUB>| = 0farcs040; we suggest the nature of this offset to be due to the matching of catalog, rather than image, level data. Our matched catalog has a surface density of ≈53 deg<SUP>-2</SUP> for K &lt;= 18.27 objects; tests using our matched catalog, along with data from the UKIDSS Deep Extragalactic Survey, imply that our limiting magnitude is i ≈ 20.6. Color-redshift diagrams, for the optical and NIR, show a close agreement between our matched catalog and recent quasar color models at redshift z &lt;~ 2.0, while at higher redshifts, the models generally appear to be bluer than the mean observed quasar colors. The gJK and giK color spaces are used to examine methods of differentiating between stars and (mid-redshift) quasars, the key to currently ongoing quasar surveys. Finally, we report on the NIR photometric properties of high, z &gt; 4.6, and very high, z &gt; 5.7, redshift previously discovered quasars.

12147222

2010nuco.confE..82A

1012

1012.3072_arXiv.txt

Half of the elements heavier than iron are produced by rapid neutron captures in a yet unknown astrophysical scenario. Galactic chemical evolution models favor core-collapse supernovae, since they occur early and frequently enough to account for the abundances observed in old halo stars and in the solar system \cite{Ishimaru.etal:2004,Qian.Wasserburg:2007}. Although the necessary conditions to produce heavy elements ($A>130$) are identified \cite{Meyer92} (high entropies, low electron fractions, and short expansion timescales), these are not found in the most recent long-time supernova simulations \cite{arcones.janka.scheck:2007,Fischer.etal:2010}. When a supernova explodes, matter surrounding the proto-neutron star is heated by neutrinos and expands very fast reaching sometimes even supersonic velocity \cite{Thompson.Burrows.Meyer:2001}. This neutrino-driven wind moves through the early supernova ejecta and eventually collides with it. The interaction of the wind with the slow-moving ejecta results in a wind termination shock or reverse shock where kinetic energy is transformed into internal energy. Therefore, the expansion velocity drops and the temperature (and thus the entropy) increases after the reverse shock. The matter near the proto-neutron star consists mainly of neutrons and protons due to the high temperatures in this region. When a mass element expands, its temperature decreases and neutrons and protons recombine to form alpha particles. At lower temperatures some of the alpha particles can form $^{12}$C either by the triple alpha reaction or by the sequence $\alpha(\alpha n,\gamma){}^9\mathrm{Be}(\alpha,n)^{12}$C. The carbon nuclei will capture additional alpha particles (alpha-process) until iron group or even heavier nuclei are produced~\cite{Woosley.Hoffman:1992}. The amount of these seed nuclei depends on the entropy and the expansion timescale of the ejecta. Once the formation of $^{12}$C nuclei freezes out the remaining neutrons can be captured by the newly formed seed nuclei and the r-process starts.

\label{sec:conclusions} We have explored the impact of the long-time dynamical evolution and of nuclear masses on the r-process abundances. We have found that the relevance of the different nuclear physics inputs depends on the long-time dynamical evolution \cite{Arcones.Martinez-Pinedo:2010}. If an $(n,\gamma)$-$(\gamma,n)$ equilibrium is reached (hot r-process), nuclear masses have a big influence on the final abundances. While for a cold r-process there is a competition between neutron capture and beta decay and these two process become relevant. This rises the importance of future experiments to measure nuclear masses that will provide a direct input for network calculations and constraints for the theoretical mass models. In both cases, as matter decays to stability, neutron captures become key to understand the final abundances and beta-delayed neutron emission becomes important not only for the redistribution of matter, but also for the supply of neutrons. The neutron captures during the decay to stability are required to explain the rare earth peak. More experimental effort is necessary to test the validity of the current theoretical cross sections and more sensitivity studies of the impact of the neutron capture rates on the final abundances will give rise to new insights. \providecommand{\newblock}{}

1012.3072

We present nucleosynthesis studies based on hydrodynamical simulations of core-collapse supernovae and their subsequent neutrino-driven winds. Although the conditions found in these simulations are not suitable for the rapid neutron capture (r-process) to produce elements heavier than A$\sim$130, this can be solved by artificially increasing the wind entropy. In this way one can mimic the general behavior of an ejecta where the r-process occurs. We study the impact of the long-time dynamical evolution and of the nuclear physics input on the final abundances and show that different nuclear mass models lead to significant variations in the abundances. These differences can be linked to the behavior of nuclear masses far from stability. In addition, we have analyzed in detail the effect of neutron capture and beta-delayed neutron emission when matter decays back to stability. In all our studied cases, freeze out effects are larger than previously estimated and produce substantial changes in the post freeze out abundances.

2798607

2011A&A...526A.155T

1012

1012.3591_arXiv.txt

\begin{figure} \centering \includegraphics[width=73mm]{figure1.eps} \caption{Upper panel~(a): SDSS imaging of M\,31 in $g$ filter, corrected for background variations. Middle panel~(b): the same as upper panel, with intrinsic extinction effects removed (see Fig.~\ref{fig:tau}b for extinction map). Lower panel~(c): the same as middle panel, showing also the masked regions (white circles).} \label{fig:m31_pix} \end{figure} \begin{figure*} \centering \includegraphics[width=175mm]{figure2.eps} \caption{Elliptically averaged observed surface brightness profiles from the SDSS photometry, converted to the $U\!BV\!RI$ system (black solid lines) together with uncertainties (grey regions). Conversion to the $U\!BV\!RI$ system is done according to \citet{Blanton:07}. For comparison, some earlier measurements are also shown (points with error bars; for references, see Sect.~\ref{results}). The earlier measurements in $U\!BV\!R$ are from elliptically averaged data, while $I$ and outermost parts of $V$ measurements by \citet{Irwin:05} have been made along the minor axis only. Left panel: major axis; right panel: minor axis. For clarity, the $BV\!RI$ profiles have been shifted by the indicated values. } \label{fig:obs_ubvri} \end{figure*} The primary information about the structure of galaxies and their stellar populations comes from the distribution of surface brightnesses and colour indices. Thus, it is only natural that the photometry of the largest nearby galaxy M\,31 has been studied thoroughly for almost a century \citep[for references, see][]{van-den-Bergh:91,Tenjes:94,Tempel:10}. More recently, satellite observations in the far-infrared have become available and attempts have been made to consider the impact of dust extinction on the observable properties of M\,31 \citep{Xu:96,Haas:98,Montalto:09,Tempel:10}. While the general luminosity distributions of different studies are in good agreement, the distribution of colours is not so well settled. Small systematic deviations between different datasets lead to considerable uncertainties of colour indices -- too big for studying in detail the stellar populations or the star formation history of the galaxy with the help of chemical evolution models. Recently, within the Sloan Digital Sky Survey \citep[SDSS; ][]{York:00}, a contiguous strip covering the entirety of M\,31 through $ugriz$ filters has been observed. The combination of five filters on a single, well-calibrated telescope and roughly uniform atmospheric conditions during the observations provide a uniquely homogeneous dataset for a thorough analysis of the extensive galaxy. In this paper, we use the SDSS observations to study the detailed luminosity and colour distribution of M\,31. We use also far-infrared imaging by the Spitzer Space Telescope and Infrared Astronomical Satellite to correct the derived photometry from dust extinction. Throughout the paper, we assume the distance of M\,31 to be 785\,\mbox{kpc} \citep{McConnachie:05}, corresponding to the scale \hbox{1$\arcmin$ = 228\,\mbox{pc}}. The inclination angle of the galaxy has been taken 77\fdg5 \citep{Walterbos:88,deVaucouleurs:91}. According to \citet{Walterbos:87}, \citet{Ferguson:02}, and our analysis of the SDSS and Spitzer images, we set 38\fdg1 as the major axis position angle. All luminosities and colour indices have been corrected from extinction in the Milky~Way according to \citet{Schlegel:98} and are given in AB-magnitudes for the $ugriz$ filters and in Vega magnitudes for the $U\!BV\!RI$ filters, as usual. For the Sloan filters, the extinction is derived from the \citet{Schlegel:98} estimates and the Galactic extinction law by linear interpolation. The absolute solar luminosity for each filter was taken from \citet{Blanton:07}. The applied solar luminosities and the Galactic extinction for each filter are presented in Table~\ref{table:filters}. \begin{table} \caption{The applied solar luminosities (in AB~magnitudes for $ugriz$ filters and in Vega magnitudes for $U\!BV\!RI$ filters) and Galactic extinctions~($A$) for each filter.} \label{table:filters} \centering \begin{tabular} {llllll} \hline\hline & $u$ & $g$ & $r$ & $i$ & $z$ \\ \hline M$_{\sun}$ & 6.38 & 5.12 & 4.64 & 4.53 & 4.51 \\ $A$ & 0.320 & 0.250 & 0.175 & 0.135 & 0.095 \\ \hline & $U$ & $B$ & $V$ & $R$ & $I$ \\ \hline M$_{\sun}$ & 5.55 & 5.45 & 4.78 & 4.41 & 4.07 \\ $A$ & 0.337 & 0.268 & 0.206 & 0.166 & 0.120 \\ \hline \end{tabular} \end{table}

1012.3591

<BR /> Aims: The objective of this work is to obtain an extinction-corrected distribution of optical surface brightness and colour indices of the large nearby galaxy M 31 using homogeneous observational data and a model for intrinsic extinction. <BR /> Methods: We process the Sloan Digital Sky Survey (SDSS) images in ugriz passbands and construct corresponding mosaic images, taking special care of subtracting the varying sky background. We apply the galactic model developed in Tempel et al. (2010, A&amp;A, 509, A91) and far-infrared imaging to correct the photometry for intrinsic dust effects. <BR /> Results: We obtain observed and dust-corrected distributions of the surface brightness of M 31 and a map of line-of-sight extinctions inside the galaxy. Our extinction model suggests that either M 31 is intrinsically non-symmetric along the minor axis or the dust properties differ from those of the Milky Way. Assuming the latter case, we present the surface brightness distributions and integral photometry for the Sloan filters as well as the standard UBVRI system. We find the following intrinsic integral colour indices for M 31: (U-B)_0 = 0.35; (B-V)_0 = 0.86; (V-R)_0 = 0.63; (R-I)_0 = 0.53; the total intrinsic absorption-corrected luminosities of M 31 in the B and the V filters are 4.10 and 3.24 mag, respectively.

12224875

2011PhRvD..83f4027S

1012

1012.1307_arXiv.txt

A fundamental limitation of the cosmological models based on general relativity (GR) is the occurrence of singularities. Perhaps one of the simplest examples is the case of an expanding homogeneous and isotropic universe filled with a matter satisfying strong energy condition such as dust or radiation. Independent of the intrinsic geometry of the universe, be it closed, flat or open, the past evolution of such a universe from arbitrary initial conditions leads to an initial singularity: the big bang, where the classical dynamical equations break down and the physics stops. Another example is the case of inflationary universe in which even though the evolution is almost deSitter, the classical spacetime is past incomplete \cite{bgv}. In recent years, various new singularities have been found in classical cosmology \cite{bigrip,sudden,sudden1,future}. Unlike the big bang and big crunch singularities, these singularities do not occur when scale factor vanishes. These occur either at finite values of the scale factor or when it diverges. Recall that for the matter satisfying weak energy condition, the latter is not possible as the spacetime curvature goes to zero when scale factor goes to infinity. However, if matter violates weak energy condition, for example in the case of phantom fields, then spacetime curvature will diverge as scale factor becomes infinite. For homogeneous and isotropic models with matter equation of state in the form of a perfect fluid these exotic singularities come in four types. Big rip (type I) where the energy density and pressure diverge along with a divergence in the scale factor, sudden singularity (type II) occurring at a finite value of the scale factor and energy density with a divergence in pressure, big freeze (type III) where energy density and pressure diverge at a finite value of the scale factor and big brake (type IV) singularity where scale factor, energy density and pressure are finite but there is a divergence in the time derivative of the pressure or rate of change of energy density. Lack of successful resolution of these singularites renders classical cosmological models incomplete. It is generally believed that existence of these singularities is a result of assuming the validity of GR even in the regime of large spacetime curvature where the effects due to quantum gravity may become important and lead to significant departures from the classical theory. It is thus hoped that incorporation of quantum gravitational effects may result in a possible resolution of these singularities. Loop quantum gravity (LQG) is one of the candidate theories of quantum gravity which attempts to address this issue. It is a non-perturbative and background independent quantization, with a key prediction that the continuum differential geometry of the classical theory is replaced by a discrete quantum geometry in the quantum theory. Perhaps one of the best illustrations of the novel effects of quantum geometry is captured in loop quantum cosmology (LQC) which is a quantization of homogeneous spacetimes based on LQG \cite{ashtekar:lqc_review,bojowald:livingreview,singh:review1}. A key prediction of LQC is that the big bang singularity is replaced by a big bounce, which is a direct consequence of the underlying quantum geometry % \cite{aps,aps1,aps2,apsv}. These results which were first obtained for homogeneous and isotropic models (for all values of spatial curvature) with a massless scalar field have been extended to inflationary potential % \cite{aps3}, anisotropic spacetimes % \cite{cv:bianchi1,aw:bianchi1} and also certain inhomogeneous situations % \cite{gowdy}. Further, using an exactly solvable model it has been shown that the expectation values of energy density have a universal upper bound for a dense subspace in the physical Hilbert space \cite{slqc}. There are strong constraints on the change in relative fluctuations of quantum observables across the bounce % \cite{recall}. Recently, much stronger constraints on the change in dispersions have been obtained by Kaminski and Pawlowski \cite{polish_recall}. These results show that a universe like ours i.e. macroscopic at late times bounces from a a similar universe at very early times (in the contracting branch) and the universe recalls almost of all its state through the bounce. Interestingly, the loop quantum dynamics admits an effective description on a continuum spacetime which can be obtained using coherent state techniques \cite{vt,st}. An important feature of this analysis is that one can obtain an effective Hamiltonian from which one can obtain modified Friedmann and Raychaudhuri equations as the Hamilton's equations. The modified set of dynamical equations inherit quantum geometric effects via higher order non-perturbative corrections which vanish at small spacetime curvatures. It is important to note that various numerical simulations have shown that effective equations capture the underlying quantum evolution very accurately for universes which become macroscopic at late times. These thus prove to be useful tools to understand the physics in LQC, such as whether the underlying theory has well defined ultra-violet and infra-red limits. It turns out that even though there exist various quantization ambiguities, there is a unique quantization leading to a consistent unambiguous physical description \cite{cs1,cs2} (the improved dynamics \cite{aps2,apsv}: which is being considered here). Using effective equations, we can ask various questions regarding the generality of singularity resolution in LQC. For example, one can ask whether spacetime curvature is always bounded in LQC? Here we should note that a universal bound on energy density (as in LQC), does not imply that the spacetime curvature is also bounded. This is easy to understand for the classical cosmology, where the Ricci scalar, which provides us a complete information about the spacetime curvature in the homogeneous and isotropic spacetime, depends both on the energy density and pressure. Though for most matter-energy configurations, the behavior of equation of state is such that an upper bound in energy density is sufficient to control the divergence in pressure and hence the spacetime curvature, it is not difficult to come up with counter examples with a more general equation of state \cite{portsmouth,generic}. Hence, an upper bound in energy density is not sufficient to prevent a divergence in the spacetime curvature. A pertinent question is whether this divergence signals the end of spacetime in LQC. In order to answer this question, we recall that even in GR we encounter events where spacetime curvature blows up but there is no associated physical singularity. This can happen if the tidal forces are not strong enough to cause a complete destruction of in-falling objects in to the singularity and geodesics can be extended beyond such events. It turns out that the events where spacetime curvature diverges in flat isotropic LQC are weak singularities and geodesics can be extended beyond them. In flat isotropic LQC, the divergence of spacetime curvature occurs only for sudden singularities which are caused by a divergence in pressure at a finite scale factor and energy density. It is straight forward to show that the expansion parameter in this case is bounded and the spacetime is geodesically complete in the flat isotropic LQC \cite{generic}. In this paper we take the first step to generalize the above result by including intrinsic curvature in the spacetime. This is done by considering the effective dynamics of loop quantized spatially closed and open models in the Robertson-Walker geometry. In the classical Friedmann dynamics, intrinsic curvature term enters in form of $1/a^2$ term in the dynamical equations. Thus it is quite straightforward to understand the expected modifications from the results in the spatially flat model. In LQC, the quantization of intrinsic curvature brings non-trivial modifications to the effective description and makes the resulting form of effective dynamical equations less straight forward to analyze. Though one expects that at small intrinsic curvatures one recovers the results of flat isotropic LQC, physics may be bring up surprises when intrinsic curvature is large. As we will show, this is indeed what happens in the case of $k=\pm 1$ models in LQC. Our analysis will be based on considering a sufficiently general phenomenological model for the equation of state which was proposed in Ref.\cite{not}. This model allows a study of all of the exotic singularities by the choice of different parameters and was earlier used for investigation of resolution of strong curvature singularities in the flat isotropic LQC \cite{generic}. As we will see, the effective dynamical equations of spatially curved models approximate those of the flat model in LQC at large volumes because the contribution from intrinsic curvature becomes negligible in this limit. Thus for future singularities, at large volumes, the resulting physics is similar for models with or without spatial curvature in LQC. However for certain values of parameters the spatially closed model in LQC permits two separate physical branches, a short lived baby universe at small volume and a parent universe which evolves to a macroscopic size. This branch is absent in the classical theory and has a pure quantum geometric origin. To completely capture the new physics from inclusion of intrinsic curvature, it is important to study exotic singularities in the past evolution when they occur at small volumes. Our analysis of past and future exotic singularities shows that all strong curvature singularities are resolved in $k=\pm 1$ isotropic LQC. The scale factors at which big rip and big freeze singularities occur are excluded from the allowed range by loop quantum effects. As in the flat model, the spacetime curvature can diverge in spatially curved LQC, however when ever this happens one has a weak singularity which is known to be harmless. In almost all cases these singularities are ignored by LQC. The only exception to this occurs in spatially closed model where weak singularities occurring in the past evolution may be resolved. This occurs purely because of the non-trivial role of intrinsic curvature effects in LQC. \vspace{5pt} We organize our paper as follows: In Sec.\ref{classical} we revisit the classical equations for spatially curved model in classical cosmology and introduce the phenomenological ansatz of the equation of state. To facilitate the reader to follow the derivation of effective equations in LQC, we derive the Friedmann and Raychaudhuri equations in the Hamiltonian framework. In Sec.\ref{quantum}, we derive the effective equations for spatially curved LQC starting from the effective Hamiltonian % \cite{apsv,kv,szulc}. Using these modified equations, we numerically obtain solutions and discuss the new physics in Sec.\ref{numerics}. Here we show that all exotic strong curvature singularities, irrespective of whether they occur in the past or the future, are resolved in spatially curved LQC. We summarize our results in Sec.\ref{summary}.

\label{summary} A fundamental question in quantum gravity is whether spacelike singularities of the classical theory are resolved. Since not all such singularities signal end of the spacetime, it is important to understand the role of quantum gravitational effects in resolution of strong singularities (those beyond geodesics can not be extended) and weak singularities (those beyond which geodesics can be extended). These issues were addressed in the loop quantization of cosmological models which are spatially flat and it was found the non-perturbative loop quantum effects resolve all strong singularities and ignore weak singularities \cite{generic}. The aim of the present analysis was to investigate these issues for spatially curved models using phenomenological model of equation of state permitting exotic singularities such as big rip, sudden singularities, big freeze singularity and the big brake singularity. In order to capture the role of intrinsic curvature, we considered exotic singularities both in the future and the past evolution. To our knowledge, even in the classical theory exotic singularities had not been studies earlier for the spatially curved model. For the singularities occurring in the future, the contribution of the intrinsic curvature is expected to become very small and we expect results to agree with the spatially flat case. This turns out to be true. However, more interesting are cases where the exotic singularities occur in past. Here one would expect effects due to quantization of intrinsic curvature to play a non-trivial role. In fact we encounter some surprising results. Though strong singularities are always resolved in LQC, it turns out that for the closed model weak singularities occurring in the past evolution may also be resolved. Thus LQC, does not always ignores weak singularities. This is an intriguing result which deserves further investigation. Another peculiar feature of the curved models is the appearance of a small branch for type II and type IV singularities for certain values of the parameters. This ``baby-universe" is bounded, and is devoid of any singularities. It will be interesting to analyze these additional branches which we find for both spatially open and closed models in more details, in particular by taking in to account inverse scale factor effects which may play some role when scale factor is below the Planck length. Our results extend earlier results on generic resolution of strong curvature singularities in spatially flat model in LQC to the spatially curved models. They also bring some important lessons, the primary one being that quantization of intrinsic curvature may throw some novel unexpected results and we need to gain more insights on when quantum gravity may ignore or resolve a weak curvature singularity. Another lesson is that as for the spatially flat model, spacetime curvature invariants may diverge for the spatially curved models and yet there may be no physical singularity. These results strengthen the case for a generic resolution of strong singularities in LQG. To achieve this goal, the next step will be to include anisotropies and then inhomogeneities. The latter will require us to go beyond the minisuperspace approximation considered here. This brings up additional challenges such as the complete classification of the strong and weak singularities in inhomogeneous situations. Two promising avenues where such an analysis can be undertaken would be the Gowdy models \cite{gowdy} and in the spinfoam paradigm \cite{spinfoams}. We hope that these studies will also provide insights on the deeper relation of these frameworks with LQC. \vskip0.5cm

1012.1307

We investigate the occurrence of various exotic spacelike singularities in the past and the future evolution of k=±1 Friedmann-Robertson-Walker model and loop quantum cosmology using a sufficiently general phenomenological model for the equation of state. We highlight the nontrivial role played by the intrinsic curvature for these singularities and the new physics which emerges at the Planck scale. We show that quantum gravity effects generically resolve all strong curvature singularities including big rip and big freeze singularities. The weak singularities, which include sudden and big brake singularities, are ignored by quantum gravity when spatial curvature is negative, as was previously found for the spatially flat model. Interestingly, for the spatially closed model there exist cases where weak singularities may be resolved when they occur in the past evolution. The spatially closed model exhibits another novel feature. For a particular class of equation of state, this model also exhibits an additional physical branch in loop quantum cosmology, a baby universe separated from the parent branch. Our analysis generalizes previous results obtained on the resolution of strong curvature singularities in flat models to isotropic spacetimes with nonzero spatial curvature.

12207443

2011JCAP...05..003M

1012

1012.4008_arXiv.txt

The non-commutativity of time evolution and spatial averaging in general relativity implies that the average expansion of a universe with structures does not evolve in time like the uniform Hubble expansion in a homogeneous Friedmann model \cite{Ell84,EllisStoeger,Ellis:2005uz}. This effect can be quantified by a backreaction term in generalized Friedmann equations and understood as a consequence of the nonlinearity of gravity \cite{Buchert:1999er}. Although the cosmological backreaction is conceptually well-understood, the complexity of the structure formation at the nonlinear level means its magnitude in the real universe is difficult to evaluate and is hence widely debated \cite{Ellis:2008zza}: some studies have found that the backreaction is small \cite{Ishibashi:2005sj,Paranjape:2008jc,Clifton:2009jw,Clifton:2010fr,Alonso:2010zv,Green:2010qy}, while other studies suggest it can have a significant effect on the cosmological dynamics and observations \cite{Rasanen:2006kp,Wiltshire:2007jk,Mattsson:2007tj,Rasanen:2008it,Wiegand:2010uh}, even to the extent of accounting for the observed cosmic acceleration entirely without additional effects \cite{Rasanen:2006kp,Rasanen:2008it,Wiegand:2010uh}. As many of the current standard values for the cosmological parameters -- not just the cosmological constant -- rely on the hypothesis of negligible backreaction, its evaluation has become a central issue in cosmology today \cite{Sarkar:2007cx,Blanchard:2010gv}. A widely used scheme to estimate the backreaction is to average the scalar parts of the Einstein equation on spacelike hypersurfaces, defined by constant proper time of the freely-falling dust particles. In this so-called Buchert approach, the backreaction term is given by the (positive) variance of the expansion rate minus the (positive) average shear \cite{Buchert:1999er}. In studies that have estimated the backreaction to be significant \cite{Rasanen:2006kp,Wiltshire:2007jk,Rasanen:2008it,Wiegand:2010uh}, the shear on the boundaries between regions characterized by different expansion rates has been neglected or, equivalently, the boundary regions or matching conditions have been ignored. On the other hand, in perturbative studies that do not neglect the shear the backreaction has been found to be small \cite{Brown:2008ra,Clarkson:2009hr,Umeh:2010pr}; however, see \cite{Rasanen:2010wz} for a discussion on the possible shortcomings of the perturbation theory in modelling the structure formation. Considering that on physical grounds, shear is expected to occur on the boundaries between regions of different expansion rates, evaluating the effect of the shear on the backreaction is evidently one of the key issues in the problem. This work is a continuation to our previous work \cite{Mattsson:2010vq}, hereafter paper I, in which the effect of shear on the cosmological backreaction was studied in the context of matching voids (with $\Omega_{{\rm v}} = 0$ and $H t = 1$) and walls (with $\Omega_{{\rm w}} = 1$ and $H t = 2/3$) together using the exact inhomogeneous Lema\^itre-Tolman-Bondi or LTB solution. Whereas neglecting the exact matching can lead to significant backreaction, the main conclusion of paper I was that the shear arising from the exact matching suppresses the backreaction by the squared ratio of the void size to the horizon size, $(r_0/t_0)^2$, thus making it small for voids of the observed size $r_0/t_0 \lesssim 10^{-2} $ \cite{Hoyle:2003hc}. Here we generalize the study of paper I by considering: \begin{enumerate} \item The backreaction as a function of the void and wall density parameters in the range $0 \leq \Omega_{{\rm v}} \leq \Omega_{{\rm w}} \leq 1$, thus relaxing the priors $\Omega_{{\rm v}}=0$ and $\Omega_{{\rm w}}=1$ of paper I. \item The backreaction for voids of arbitrary size $r_0$, thus relaxing the condition $r_0 \ll t_0$ assumed in paper I. \item A network of voids with different densities $\Omega_{{\rm v}}$ and radii $r_0$, thus giving rise to relative variance of the expansion rate between the different void-wall pairs. \end{enumerate} Note that since the backreaction of the quasi-spherical Szekeres model reduces to the LTB model, the results of this work apply to the quasi-spherical Szekeres solution as well \cite{Bolejko:2008zv,Bolejko:2010wc}. Although we focus here on the effect of the structure formation on the average dynamics of the universe, the ultimate goal of the research is to evaluate the total effect of the structure formation on the cosmological observations. In addition to the dynamical backreaction considered in this work, known effects of the structure formation include modifications on the propagation of light not directly determined by the volume-averaged expansion \cite{Mattsson:2007tj,GarciaBellido:2008nz,GarciaBellido:2008gd,GarciaBellido:2008yq,Rasanen:2008be,Enqvist:2009hn,Clifton:2009jw,Kainulainen:2009dw,Blomqvist:2009ps,Rasanen:2009uw,Krasinski:2010rc,Biswas:2010xm,Clarkson:2010ej,Yoo:2010hi,Bolejko:2010nh,Davis:2010jq,Nadathur:2010zm} and effects due to our non-average location in the universe \cite{Wiltshire:2007jk,Wiltshire:2008sg,Wiltshire:2009db,Smale:2010vr}. These effects may be important but we do not try to evaluate them in this work. The paper is organized as follows. The necessary background of the Buchert averaging method and the LTB solution are introduced in Sects.\ \ref{averaging} and \ref{ELTEEBEE}, respectively. Sect.\ \ref{model} provides the definition of the void-wall LTB model, which we apply to study the effect of shear on the backreaction in Sect.\ \ref{backreactioninLTB}: The general analytic expression of the backreaction for the void-wall LTB model is derived in Sect.\ \ref{generalQ} and its expansion as a power series and a simple but accurate fitting formula are considered in Sect.\ \ref{powerexpansionofQ}. The result for the backreaction is then applied to a network of different voids in Sect.\ \ref{voidnetwork} and to large voids in Sect.\ \ref{largevoids}. In Sect.\ \ref{kompensoimatonna}, we discuss uncompensated voids. Finally, the results are summarized in Sect.\ \ref{konkluusiot} and the conclusions are given in Sect.\ \ref{konkluusiot2}.

\label{konkluusiot2} By considering exact solutions of the Einstein equation consisting of one or more LTB solutions, we have pinpointed the issue of small versus large cosmological backreaction to the question of matching conditions: while the variance of the expansion rate alone can induce significant backreaction, the shear arising from matching together the regions with different expansion rates seems to bring down the backreaction by at least five orders of magnitude for voids of the observed size. The crucial question is whether the suppression of the backreaction due to the shear is a general property of all realistic cosmological solutions of general relativity or just a special property of the matching in the considered particular solutions. This issue has to be addressed with solutions more sophisticated than the LTB-based models employed here.

1012.4008

We study the effect of shear on the cosmological backreaction in the context of matching voids and walls together using the exact inhomogeneous Lemaȋtre-Tolman-Bondi solution. Generalizing <A href="10.1088/1475-7516/2010/10/021">JCAP 1010 (2010) 021</A>, we allow the size of the voids to be arbitrary and the densities of the voids and walls to vary in the range 0 &lt;= Ω<SUB>v</SUB> &lt;= Ω<SUB>w</SUB> &lt;= 1. We derive the exact analytic result for the backreaction and consider its series expansion in powers of the ratio of the void size to the horizon size, r<SUB>0</SUB>/t<SUB>0</SUB>. In addition, we deduce a very simple fitting formula for the backreaction with error less than 1% for voids up to sizes r<SUB>0</SUB>gtrsimt<SUB>0</SUB>. We also construct an exact solution for a network of voids with different sizes and densities, leading to a non-zero relative variance of the expansion rate between the voids. While the leading order term of the backreaction for a single void-wall pair is of order (r<SUB>0</SUB>/t<SUB>0</SUB>)<SUP>2</SUP>, the relative variance between the different voids in the network is found to be of order (r<SUB>0</SUB>/t<SUB>0</SUB>)<SUP>4</SUP> and thus very small for voids of the observed size. Furthermore, we show that even for very large voids, the backreaction is suppressed by an order of magnitude relative to the estimate obtained by treating the walls and voids as disjoint Friedmann solutions. Whether the suppression of the backreaction due to the shear is just a consequence of the restrictions of the used exact models, or a generic feature, has to be addressed with more sophisticated solutions.

12137816

2010idm..confE..64D

1012

1012.0184_arXiv.txt

The exciting possibility of detecting dark matter particle candidates with IceCube is based on the assumption that, if they constitute the dark matter in the halo, they can be gravitationally trapped in the deep gravitational wells of heavy objects, like the Sun or the Galactic Center~\cite{DMSun}. Subsequent pair-wise annihilation into Standard Model particles could lead to a detectable neutrino flux. This is a clear signal for a neutrino telescope: it is directional and has a different energy spectrum than the atmospheric neutrino background flux. Popular dark matter candidates are stable relic particles that arise in supersymmetric extensions of the Standard Model or in theories with extra spacial dimensions. In some flavours of the minimal supersymmetric extension of the Standard Model, the MSSM, a viable dark matter candidate is the lightest neutralino, the lightest particle in the super-partner quadruplet of the gauge bosons and neutral Higgs. Neutralinos are stable, interact only weakly and gravitationally and, as relics of the Big Bang, can form a dark matter halo in the galaxy. The current lower limit on the neutralino mass, $m_{\chi}\gtrsim 46$ GeV, comes from accelerator searches~\cite{DELPHI:03a}, while an upper limit of a few hundred TeV can be set based on unitarity constrains on the mass of any thermally produced relic~\cite{Griest:90a}. We will not discuss further other common supersymmetric scenarios where the gravitino is the lightest supersymmetric particle, since they do not provide a signal in neutrino telescopes, but we will consider another thermal relic arising in the scenario of universal extra dimensions, the lightest Kaluza-Klein particle (LKP)~\cite{Hooper:07a}. We have considered the LKP in models with one additional space dimension, associated with the first excitation of the hypercharge gauge boson. The mass of the LKP is inversely proportional to the 'size' of the extra dimension and can lie in the range few hundred GeV to about a TeV. The model thus defined has only two parameters; the LKP mass and the relative mass difference, $\Delta_q$, between the LKP and the first Kaluza-Klein quark excitation. This parameter controls the strength of possible co annihilations and influences the predicted relic density of LKPs. A third kind of candidates we have considered are Simpzillas, superheavy dark matter in the form of strongly-interacting relic particles in the mass range 10$^4$~GeV -- 10$^{18}$~GeV. Strongly-interacting in this context simply means non-weakly (as opposed to the usual assumption for WIMPs) and it should not be understood as a QCD-like interaction. Unlike neutralinos or LKPs, Simpzillas are produced non-thermally at the end of inflation~\cite{Chung:98a}, and the unitarity constraint on their mass can therefore be avoided. Masses up to the unification scale can be generated without violating any fundamental law.

IceCube has an active program of searches for dark matter, both from candidates accumulated in the Sun as well as in the Galactic Halo or center. We have tested the data from the 22-string and 40-string configurations of IceCube for an excess neutrino flux from these objects and interpret the results in terms of several dark matter candidates. With the 40-string detector we have been able to search the Galactic Center for the first time . The size of the detector allows to use a fraction of the instrumented volume as veto region, which enables the identification of starting tracks and an efficient reduction of the atmospheric muon background. This technique will be used in its full potential with the complete 86-string detector in the future. The low-energy extension DeepCore which has been already deployed in the center of the IceCube array will allow to significantly lower the energy threshold of IceCube and extend the dark matter searches in a competitive way to the interesting region of candidate masses below 100~GeV.

1012.0184

The construction of the IceCube neutrino observatory is practically terminated. At the time of this writing, and with 79 strings taking data out of the 86 foreseen, we are one deployment season away from completion. The detector, however, has been taking data since 2006 in different partial configurations. We have evaluated these data for evidence of dark matter annihilations in the Sun, in the Galactic Center and in the Galactic Halo, searching for an excess neutrino flux over the expected atmospheric neutrino background. This contribution reviews the results of dark matter searches for WIMPs, Kaluza-Klein modes and superheavy candidates (Simpzillas), using the 22- and 40-string configurations of IceCube. The results are presented in the form of muon flux limits, constrains on the candidates' spin-dependent cross-section with protons, and constrains in the self-annihilation cross section. These results are presented in the context of direct searches and searches in space

12138619

2010int..workE..25M

1012

1012.3378_arXiv.txt

Cygnus X-3 is a unique microquasar composed of a compact object and a Wolf-Rayet star \cite{vk1}. It has a strong 4.8 hour orbital modulation and is known to produce relativistic jets \cite{pgk,ma}. It also exhibits relatively strong radio emission most of the time ($\rm \sim 100~mJy$) \cite{we}. {\it AGILE} \cite{tm} and {\it Fermi} \cite{fl} have recently shown that Cygnus X-3 is a transient gamma-ray source ($\rm >~100~MeV$). Flaring gamma-ray emission appears to be connected with major radio flares and their associated X-ray states that precede the major radio flares \cite{tm, mkh}. Studies \cite{we,mm} in the radio and hard X-ray (HXR) have found four states: quiescent, minor flaring, quenched and major flaring. Table 1 shows how these states are correlated with the simultaneous X-ray observations. Throughout all states of activity the HXR and soft X-ray (SXR) show an anti-correlation or spectral pivoting around 10 keV \cite{mm2}. Additional studies have shown that the X-ray emission can be broken down into additional states. A study of the X-ray and radio behavior arrive at six states (quiescent, minor flaring, suppressed, post-flare, major flaring, and quenched) \cite{szm}. A recent study which incorporated X-ray spectral hardness (shape of the spectrum) in addition to the X-ray and radio arrive at three major states, Quiescent, Flaring, and Hypersoft, with several substates \cite{khm}. The Hypersoft is a new state which has important ties to the quenched state, gamma-ray emission, and jet production.

Through multi-wavelength campaigns of Cygnus X-3 important insights are being gained. Cygnus X-3 has been shown to be one of the few XRBs to produce gamma-ray emission. This emission is coupled to state behavior and major radio flares. The infrared is providing an important bridge between the radio and the X-ray. These studies are allowing us to probe some of the most extreme environments in nature: high mass flow in a strong gravitational field.

1012.3378

Cygnus X-3 is a unique microquasar which shows X-ray state changes, strong radio emission, and relativistic jets. It is also an unusual X-ray binary with the mass-donating companion being a high mass star Wolf-Rayet but the orbital modulation (as inferred from X-ray emission) is only 4.8 hours, a value more common in low-mass systems. It has recently been shown by AGILE and Fermi that Cygnus X-3, is a transient gamma-ray source (&gt;100 MeV). To understand the environment, nature, and behavior of Cygnus X-3 multi-wavelength observations are necessary. In this proceedings we present the results achieved so far from multi-wavelength campaigns.

12207061

2011JCAP...08..030H

1012

1012.4558_arXiv.txt

Introduction} In the history of the early universe, various types of phase transitions have presumably occurred as a consequence of the spontaneous breaking of symmetries in fundamental physics. During this process, some nontrivial vacuum structures, called topological defects, would be formed, and they can affect the evolution of the universe. For example, the formation of cosmic strings has rich phenomenological consequences for cosmology such as the structure formations~\cite{1994csot.book.....V}. On the other hand, domain walls, which arise when a discrete symmetry become spontaneously broken, are considered as cosmologically undesirable objects~\cite{1974JETP...40....1Z}. These topological defects arise in theories beyond the standard model of particle physics. A well-known example is the theory of axions~\cite{2010RvMP...82..557K}. Axion~\cite{1978PhRvL..40..223W, 1978PhRvL..40..279W} is the pseudo-Nambu-Goldstone boson which arise as a consequence of the Peccei-Quinn (PQ) mechanisms~\cite{1977PhRvL..38.1440P, 1977PhRvD..16.1791P} which is introduced to solve the CP violation problem in quantum chromodynamics (QCD). In this model, one can make the CP violating phase dynamically into zero value by introducing a new symmetry called U(1)$_{\mathrm{PQ}}$. In the early universe, this symmetry has been spontaneously broken when the temperature of the universe has fallen below some energy scale (the PQ scale). This is called the PQ phase transition. At this time the networks of global strings called axionic strings are formed. Furthermore, at the time of the QCD phase transition, the axion acquires a mass, which causes the spontaneous breaking of discrete $Z_{N_{\mathrm{DW}}}$ subgroup of U(1)$_{\mathrm{PQ}}$, leading to the formation of domain walls attached to the strings. The integer parameter $N_{\mathrm{DW}}$ describes the number of degenerate vacua and the number of walls attached to the string. The value of this number depends on the models. For example, in the Dine-Fischler-Srednicki-Zhitnitsky (DFSZ) models~\cite{1981PhLB..104..199D, Zhitnitsky1980}, $N_{\mathrm{DW}}=2N_g$, where $N_g$ is the number of quark generations (this can be reduced to $N_{\mathrm{DW}}=N_g$, depending on the structure of the Higgs sector~\cite{1990PhRvD..41.3848G}). On the other hand, in the Kim-Shifman-Vainshtein-Zakharov (KSVZ) models~\cite{1979PhRvL..43..103K, 1980NuPhB.166..493S}, one can get $N_{\mathrm{DW}}=1$. There are two possibilities for the fate of such string-wall networks. One is the case with $N_{\mathrm{DW}}=1$. In this case, walls bounded by strings quickly slice themselves into small peaces as discussed in \cite{1982PhRvL..48.1867V, 1984PhRvD..30.2036V}. Another possibility is the case with $N_{\mathrm{DW}}>1$. In this case the networks survive for a long time, which is cosmologically disastrous~\cite{1982PhRvL..48.1156S} since they overclose the universe and distort the cosmic microwave background (CMB) observed today~\cite{1974JETP...40....1Z}. Therefore, axion models with $N_{\mathrm{DW}}>1$ seem to conflict with the standard cosmology. This is called the axionic domain wall problem. There are several ways to avoid this problem. The simplest way is just to consider the model with $N_{\mathrm{DW}}=1$. In this model, the walls decay immediately after the QCD phase transition, producing a radiation of barely relativistic axions \cite{1999PhRvD..59b3505C}. This can be realized in the KSVZ models. Also, it is possible to take $N_{\mathrm{DW}}=1$ in the variant axion models~\cite{1986PhLB..172..435P, 1986PhLB..173..189K}, and its collider implications are recently investigated in~\cite{2010JHEP...06..059C}. A more intricate solution is to embed the discrete subgroup $Z_{N_{\mathrm{DW}}}$ of U(1)$_{\mathrm{PQ}}$ in the center of another continuous group (so called the Lazarides-Shafi mechanism~\cite{1982PhLB..115...21L,1987PhR...150....1K}). In this model, the degenerate vacua are connected to each other by another symmetry transformation. However, in this kind of model, one have to choose the symmetry group, Higgs representations and U(1)$_{\mathrm{PQ}}$ charge judiciously, which seems to be unlikely to occur~\cite{1982PhLB..116..227B}. The third possibility is to suppose that inflation has occurred after the PQ phase transition to dilute the number density of strings as well as domain walls. In this case, however, a constraint which comes from the bound on isocurvature perturbations in CMB observation might be severe~\cite{2008PhRvD..78h3507H}. Alternatively, one can eliminate domain walls by introducing a tiny $Z_{N_{\mathrm{DW}}}$ breaking term (i.e. a bias) which lifts the vacuum degeneracy \cite{1982PhRvL..48.1156S, 1999PhRvD..59b3505C, 2008LNP...741...19S}. In this case, domain walls collapse due to the pressure force acting between different vacua~\cite{1981PhRvD..23..852V}. In this paper we mainly consider the last possibility, biased domain walls, as a solution to the domain wall problem. It is natural to expect such a explicit $Z_{N_{\mathrm{DW}}}$ breaking terms in the context of the quantum theory of gravity such as string theory: The quantum gravitational effects induce higher-dimensional operators suppressed by powers of the Planck mass, which may alter the PQ solution to the CP violating problem~\cite{1992PhLB..282..137K, 1992PhLB..282..132H, 1992PhRvD..46..539B, 1997PhRvD..55.5826D}. It is argued that allowed region in the parameter space for this solution is narrow, though it is not ruled out~\cite{2008LNP...741...19S}. The cosmological scenario in which domain walls decay at early time is also interesting from the observational point of view. It has been pointed out that the emission of gravitational waves is likely to occur due to the decay of domain walls~\cite{1999PhRvD..59b3505C}. Indeed, the recent numerical study implies that long-lived unstable domain walls can radiate gravitational waves with the amplitude enough to observe~\cite{2010JCAP...05..032H}. If it is true, it might be possible to probe axion models by gravitational wave experiments in the next decades. We would like to address these issues more quantitatively, based on the field theory lattice simulations. The numerical simulations of axionic domain walls have been performed by several authors. The simulation of decaying domain wall with $N_{\mathrm{DW}}=1$ was examined in~\cite{1999PhRvD..59b3505C, 1993PhLB..318...53N, 1994PhRvD..50.4821N}, neglecting the expansion of the universe. Also, the evolution of string-wall networks with $N_{\mathrm{DW}}\ge1$ models in the expanding background was investigated in~\cite{1990ApJ...357..293R}. In addition to these simulations, we investigate the decay process of the networks due to the explicit $Z_{N_{\mathrm{DW}}}$ breaking term, give the numerical confirmation of the decay of the networks with $N_{\mathrm{DW}}>1$, estimate their lifetime, and constraint the parameter space in which the CP violation problem and the domain wall problem are solved simultaneously. This paper is organized as follows. In section 2, we describe the model which we consider, and briefly review the property of string-wall networks. In section 3, we present the results of the numerical simulations and estimate the decay time of the networks. Based on these results, we give observational constraints for the bias parameter and the PQ scale in section 4. In section 5, we shortly discuss about a possibility to observe gravitational waves from domain wall networks in the future experiments. Then we conclude in section 6. {\bf Notation}: we work in the spatially flat Friedmann-Robertson-Walker (FRW) background in which the metric is given by \begin{equation} ds^2 = dt^2 - a^2(t)[dx^2+dy^2+dz^2]. \nonumber \end{equation} We denote the cosmic time as $t$ and the conformal time as $\tau$, where $d\tau = dt/a(t)$.

1012.4558

We study the cosmological evolution of domain walls bounded by strings which arise naturally in axion models. If we introduce a bias in the potential, walls become metastable and finally disappear. We perform two dimensional lattice simulations of domain wall networks and estimate the decay rate of domain walls. By using the numerical results, we give a constraint for the bias parameter and the Peccei-Quinn scale. We also discuss the possibility to probe axion models by direct detection of gravitational waves produced by domain walls.

12230922

2011PThPS.190...62E

1012

1012.1711_arXiv.txt

In the curvaton mechanism \cite{curvaton}, primordial perturbations originate from quantum fluctuations of a light scalar field which gives a negligible contribution to the total energy density during inflation. This field is called the curvaton $\sigma$. Inflation is driven by another scalar, the inflaton $\phi$, whose potential energy dominates the universe. After the end of inflation, the inflaton decays into radiation. If the inflationary scale is low enough, $H_{*}\ll 10^{-5} \sqrt{\epsilon_{*}}$, the density fluctuations of the radiation component are much below the observed amplitude $\delta\rho/\rho\sim 10^{-5}$ and the fluid is for practical purposes homogeneous. While the dominant radiation energy scales away as $\rho_{r}\sim a^{-4}$, the curvaton contribution to the total energy density may increase and the initially negligible curvaton perturbations get imprinted into the metric. The standard adiabatic hot big bang era is recovered when the curvaton eventually decays and thermalizes with the existing radiation. The mechanism can be seen as a conversion of initial isocurvature perturbations into adiabatic ones and, depending on the parameters of the model, is capable of generating all of the observed primordial perturbation. The scenario sketched above represents the simplest possible realization of the curvaton mechanism, and a wide range of different variations of the idea have been studied in the literature. For example, the inflaton perturbations need not be negligible \cite{mixed}, there could be several curvatons \cite{many_curvatons}, the curvaton decay can result into residual isocurvature perturbations \cite{LUW,isocurvature} and inflation could be driven by some other mechanism than slowly rolling scalars \cite{alt_inflation}. It is however well known that the predictions of the curvaton model are quite sensitive to the form of the curvaton potential \cite{Dimopoulos:2003ss,kesn,kett,Enqvist:2008be,Huang:2008zj,Kawasaki:2008mc,Chingangbam:2009xi,us1,us2,Chambers:2009ki, us3}. In particular, even small deviations from the extensively studied quadratic potential can have a significant effect, at least when considering non-Gaussian effects \cite{kesn, kett}. One can also encounter strong scale-dependence of the non-gaussianity parameters \cite{byrnesetal}. When the initial curvaton field value lies far in the non-quadratic part of the potential, the non-linear nature of the evolution equation will in general result in a very rich structure of phenomena in the parameter space, as has been discussed in detail in \cite{us1,us2}. The simplest non-quadratic curvaton potential is given by \begin{equation} V = \frac{1}{2}m^2\sigma^2 + \lambda{\sigma^{n+4}\over M^n} \ , \label{curvatonpot} \end{equation} where $n$ is an even integer to keep the potential bounded from below, and the interaction term is suppressed by a cut-off scale $M$. For non-renormalizable operators $n > 0$, we set the cut-off scale to be the Planck scale $M = M_{\rm P}\equiv 1$, and the coupling to unity, $\lambda = 1$. For the renormalizable quartic case $n = 0$, the coupling $\lambda$ can be treated as a free parameter. The potential (\ref{curvatonpot}) is reasonably well motivated by generic theoretical arguments. Indeed, the curvaton should have interactions of some kind as it eventually must decay and produce Standard Model fields. The curvaton needs to be weakly interacting to keep the field light during inflation. This however only implies that the effective curvaton potential should be sufficiently flat in the vicinity of the field expectation value during inflation but does not a priori require the interaction terms in (\ref{curvatonpot}) to be negligible. Moreover, as typically the inflationary energy scale is relatively high, the field can be displaced far from the origin and therefore feels the presence of the higher order terms in the potential. The interactions could arise either as pure curvaton self-interactions involving the curvaton field $\sigma$ alone, or more generically as effective terms due to curvaton couplings to other (heavy) degrees of freedom that have been integrated out. An example of a possible physical setup which could lead to (\ref{curvatonpot}) is given by flat directions of supersymmetric models that have been suggested as curvaton candidates \cite{curvaton_flat}. These would lead to a potential of the form (\ref{curvatonpot}) with typically a relatively large power for the non-renormalizable operator. The amplitude and non-gaussianity of the perturbation depend on the curvaton decay time. Here we assume for simplicity a perturbative curvaton decay characterized by some effective decay width $\Gamma$ (for non-perturbative decay, see \cite{curvatondecres, BasteroGil:2003tj}). When the interaction term dominates in (\ref{curvatonpot}), the curvaton oscillations start in a non-quadratic potential and the curvaton energy density always decreases faster than for a quadratic case. For non-renormalizable interactions, the decrease is even faster than the red-shifting of the background radiation and the curvaton contribution to the total energy density is decreasing at the beginning of oscillations. Consequently, the amplification of the curvaton component is less efficient than for a quadratic model. For the same values for $m$ and $\Gamma$, the curvaton typically ends up being more subdominant at the time of its decay than in the quadratic case. Despite the subdominance, the curvaton scenario can yield the correct amplitude of primordial perturbations as the relative curvaton perturbation $\delta\sigma_{*}/\sigma_{*}$ produced during inflation can be much larger than $10^{-5}$. For a quadratic model, it is well known that the curvaton should make up at least few per cents of the total energy density at the time of its decay in order not to generate too large non-gaussianities \cite{LUW,curvaton_ng}. This bound does not directly apply to the non-quadratic model (\ref{curvatonpot}) since the dynamics is much more complicated. Although the subdominant curvaton scenario implies relatively large perturbation $\delta\sigma_{*}/\sigma_{*}$, the higher order terms in the perturbative expansion of curvature perturbation can be accidentally suppressed \cite{kesn,kett, us2}.

\label{jaarittelua} It may appear surprising that even very small deviations from the quadratic form of the curvaton potential can affect the curvature perturbation in a significant way. However, one should bear in mind that the small curvature perturbation is really the difference of two large numbers. The number of e-folds generated during curvaton oscillations is typically $N\sim {\cal O}(10)$, whereas the difference that gives rise to the non-gaussianity is $\Delta N\lesssim 10^{-8}$. Since self-interactions imply non-linearities in the evolution of the curvaton field and in the number of e-folds $N$, one can understand that even small changes can have profound effects in the difference $\Delta N$. In particular, as discussed here, the non-gaussianities turn out to be quite different as compared with the simplest quadratic model. There the magnitude of $\fnl$ in the limit $\rdec \ll 1$ is determined by the curvaton energy density at the time of its decay, $\fnl \sim 1/\rdec$. However, with self-interactions the prediction for $\fnl$ can significantly deviate from this simple estimate. Even if $\rdec \ll 1$, there exists regions in the parameter space with $|\fnl| < \mathcal{O}(1)$. This is because the value of $\fnl$ oscillates and changes its sign. Nevertheless, $\gnl$ can then be very large and one has a rather non-trivial non-Gaussian statistics characterized by a large trispectrum and a vanishing bispectrum. Such a situation, discussed already in \cite{kett}, appears to be rather generic in self-interacting curvaton models, and is possible for a wide, albeit restricted, range of model parameters. Large non-gaussianities can be generated even if the curvaton dominates the energy density at the time of its decay. In general, in the presence of self-interactions the relative signs of $\fnl$ and $\gnl$ and the functional relation between them are typically modified from the quadratic case. Thus the non-linearity parameters taken together, in possible conjunction of other cosmological observables such as tensor perturbations, may offer the best prospects for constraining the physical properties of the curvaton. A TeV mass curvaton is a rather special case. An important constraint, valid also for higher mass curvatons, is that it has to decay before the CDM freeze-out. This, together with observational constraints, fixes the range of the initial conditions for the curvaton field which turn out to be such that the quadratic term in the curvaton potential cannot dominate over possible higher-order terms for the whole dynamical range. One finds\cite{us3} that the only viable curvaton potential that satisfies all the constraints is $V=m^2\sigma^2/2+\sigma^8/M^4$. Moreover, the curvaton decay rate should be in the range $\Gamma=10^{-15}- 10^{-17}$ GeV. Note that in the case where the curvaton energy density is subdominant at the time of decay, the curvaton does not necessarily have to decay before baryogenesis, which can be a process that takes place among the inflaton decay products. However, the decay should be able to produce thermal CDM particles so that the CDM perturbation is adiabatic. Note also that what really matters is the equation of state, not the time of decay. Thus if the curvaton decays too early, the perturbations might still generated if the decay products have the equation of state of matter. An example of this could be the MSSM flat direction fragmenting into Q-balls, which would then slowly decay.

1012.1711

The evolution of the curvature perturbation is highly non-trivialfor curvaton models with self-interactions and is very sensitive to the parameter values. The final perturbation depends also on the curvaton decay rate Γ. As a consequence, non-gaussianities can be greatly different from the purely quadratic case, even if the deviation is very small. Here we consider a class of polynomial curvaton potentials and discuss the dynamical behavior of the curvature perturbation. We point out that, for example, it is possible that the non-gaussianity parameter f<SUB>NL</SUB> ≃ 0 while g<SUB>NL</SUB> is non-zero. In the case of a curvaton with mass m ∼ O(1) TeV we show that one cannot ignore non-quadratic terms in the potential, and that only a self-interaction of the type V<SUB>int</SUB> = σ^8/M^4 is consistent with various theoretical and observational constraints. Moreover, the curvaton decay rate should then be in the range Γ = 10<SUP>-15</SUP> - 10<SUP>-17</SUP> GeV.

2084882

2010arXiv1012.5733D

1012

1012.5733_arXiv.txt

\label{introduction} This document describes the scientific requirements for the the VAO SED builder and analysis tool. It is divided in three sections describing the 'SED building' (section \ref{sedbuilder}, 'SED analysis' (section \ref{sedanalysis}) and 'SED visualization' (section \ref{sedvisualizationandediting}) capabilities of the tool respectively. Each section is split in multiple subsections, each one describing a distinct requirement of the tool. Each distinct requirement discussed in this document is associated to a label indicating the general section of the document (SED builder - SED analysis - SED visualization) where the requirement can be found, and a unique index (for example, {\bf SED.an.3.1} for the first sub-requirement of the third requirement of the analysis section). These labels are used to provide a quick reference to the different parts of the requirements and provide a handle to the hierarchical structure of the document. The hierarchy of requirements is also shown in the tree-graph in figure \ref{plot:sed_requirements_tree}, which is associated to the break-down scheme adopted throughout this document. The labels (in boldface in the document) are also used in tables \ref{table:comparison}, \ref{table:prioritizationSEDbui}, \ref{table:prioritizationSEDan} and \ref{table:prioritizationSEDvis}.

1012.5733

This document describes the scientific requirements for the SED builder and analysis tool that will be designed and built by the US Virtual Astronomical Observatory (VAO). VAO is the VO effort based in the US, whose primary emphasis is to provide new scientific research capabilities to the astronomy community. The near-term goal of the VAO is to put useful and efficient tools in the hands of research astronomers as soon as possible. The VAO has identified eight major research initiatives that are of high priority, including the science obtained with the creation and analysis of the spectral energy distributions of astronomical sources. This document contains the high-level scientific requirements that will drive the design and implementation of the VAO SED tool.

12148513

2010PhDT.......196R

1012

1012.0266_arXiv.txt

\label{ch:introduction} \setlength{\epigraphwidth}{9cm} \epigraph{ { And if the fixed Stars are the Centers of other like systems, these, being form'd by the like wise counsel, must be all subject to the dominion of One, [...]. }}{\tiny Isaac Newton, General Scholium, translated by Motte, 1729} \noindent In 1713 Isaac Newton wrote this sentence in an essay attached to the third edition of his famous Principia Mathematica. In other words, he expected to see planets around other stars. Almost three centuries later, astronomers discovered the first planet beyond our own solar system, a so-called exo-planet. The number of known planets has increased rapidly ever since. To date, 461 extra-solar planets have been discovered \citep{exoplanet}. At least 10\% of all nearby solar type stars host planets \citep{Cumming2008}. With this tremendous observational success, it is now the theoreticians' turn to explain the discovered systems. One important aspect is to find out if, and if so why, these systems formed differently compared to the solar system. In this chapter, first the discoveries of exo-planets in recent years are presented. Then, suggested formation scenarios of planets, planetary systems and their evolution are reviewed. This thesis discusses stochastic phenomena in a range of astrophysical systems. The analytic model presented in chapter \ref{ch:randwalk} is the key in understanding the effects of stochastic forces and turbulence in those systems. It forms the basis of the physical understanding in many celestial systems in which stochastic forces are present. A detailed formation scenario of the planetary system HD45364 is presented in chapter \ref{ch:threetwo}. In chapter \ref{ch:randwalk} another formation scenario is presented, this time for the planetary system HD128311. Both systems are resonant systems and their dynamical states provide important constraints on their formation history. HD45364 formed most likely in a massive disc and had a phase of rapid convergent migration. On the other hand, the current observed orbital parameters of the system HD128311 are consistent with the formation in a strongly turbulent disc. In chapter \ref{ch:moonlet}, these formalism developed in chapter \ref{ch:randwalk} is applied to Saturn's rings and a moonlet. Saturn's rings also exert stochastic forces. The rings, together with embedded moons, resemble a small scale version of the proto-planetary disc. Finally, the issue of numerical convergence in simulations of planetesimal formation is discussed in section \ref{ch:planetesimals}. Planetesimals are likely to form in turbulent proto-stellar discs via gravitational instability. It turns out that the system can be simulated consistently only if the relevant small scale processes are included. The dynamical evolution is then very similar to Saturn's rings, except that the final clump is gravitationally bound. In chapter \ref{ch:summary} we summarise the results. The main numerical codes that have been used in chapters \ref{ch:threetwo}, \ref{ch:randwalk}, \ref{ch:moonlet} and \ref{ch:planetesimals} are described in the appendices \ref{app:dpmhd3d} and \ref{app:gravtree}.

The planets in the multi-planetary system HD45364 are most likely in a 3:2 mean motion resonance. This poses interesting questions on its formation history. Assuming that the planets form far apart and migrate with a moderate migration rate, as predicted by standard planet formation and migration theories, the most likely outcome is a 2:1 mean motion resonance, contrary to the observation of a 3:2 MMR. In this chapter, we investigated a possible way around this problem by letting the outer planet undergo a rapid inward type~III migration. We presented an analytical estimate of the required migration rate and performed both N-body and hydrodynamical simulations. We find that it is indeed possible to form a 3:2 MMR and avoid the 2:1 resonance, thus resembling the observed planetary system using reasonable disc parameters. Hydrodynamical simulations suggest that the system is more likely to sustain the resonance for high aspect ratios, as the migration of the inner planet is slowed down, thus avoiding divergent migration. Finally, we used the orbital configuration found in the hydrodynamical formation scenario to calculate a radial velocity curve. This curve was then compared to observations and the resulting fit has an identical $\chi^2$ value to the previously reported \textit{best fit}. Our solution is stable for at least a million years. It is in a dynamically different state, both planets having lower eccentricities and a different libration pattern. This is the first time that planet migration theory can predict a precise orbital configuration of a multiplanetary system. This might also be the first direct evidence for type III migration if this scenario turns out to be true. The system HD45364 remains an interesting object for observers, as the differences between the two orbital solutions can be measured in radial velocity within a couple of years. \chapter{The survival of mean motion resonances in a turbulent disc} \label{ch:randwalk} \epigraph{ Every body continues in its state of rest, or of uniform motion in a right line, unless it is compelled to change that state by forces impressed upon it. }{\tiny Isaac Newton, Principia Mathematica, 1687} \noindent As shown for the system HD45364 in chapter \ref{ch:threetwo}, resonant configurations can be established as a result of dissipative forces acting on the planets which lead to convergent migration. The previous chapter assumes a laminar disc. However, in general the disc is thought to be turbulent because, as discussed in more detail in chapter \ref{ch:introduction}, the molecular viscosity is not large enough to drive the observed outward angular momentum flux. The most promising mechanism to drive turbulence is the magneto-rotational instability, MRI \citep{BalbusHawley91}, although other possibilities are still being discussed \citep[see e.g.][]{LesurPapaloizou2009}. In this chapter, we study the effects of density perturbations in the disc resulting from any kind of turbulence on planetary systems. We do not solve the full 3D magneto hydrodynamics equations, but rather assume a parametrisation for the forces acting on planets. This has two main advantages. First, we can understand the relevant dynamical processes much more easily. Second, there is a large diversity of simulations of the MRI in the literature which exhibit different diffusion coefficients, varying by several orders of magnitude. By parametrising the forces, we avoid these difficulties altogether. Using this approach, we study the response of both a single planet and planetary systems in a 2:1 mean motion commensurability to stochastic forcing. The forcing could for example result from MRI turbulence. We first develop an analytic model from first principles. We can isolate two distinct libration modes for the resonant angles which react differently to stochastic forcing. Systems are quickly destabilised if the magnitude of the stochastic forcing is large. The growth of libration amplitudes is parametrised as a function of the diffusion coefficient and other relevant physical parameters. We also perform numerical N-body simulations with additional stochastic forcing terms to represent the effects of disc turbulence. These are in excellent agreement with the analytic model. Stochastic forcing due to disc turbulence may have played a role in shaping the configurations of observed systems in mean motion resonance. It naturally provides a mechanism for accounting for the HD128311 system, for which the fast mode librates and the slow mode is apparently near the borderline between libration and circulation.

1012.0266

The increasing number of extra-solar planets opens a new opportunity for studies of the formation of planetary systems. Resonant systems are of particular interest because their dynamical configuration provides constraints on the unobservable formation and migration phase. In this thesis, formation scenarios for the planetary systems HD128311 and HD45364 are presented. N-body simulations of two planets and two dimensional hydrodynamical simulations of proto-planetary discs are used to model the migration phase and the capture into resonance. The results indicate that the proto-planetary disc has a larger surface density than previously thought. Proto-planets are exposed to stochastic forces, generated by density fluctuations in the disc. A generic model of both a single planet, and two planets in resonance, being stochastically forced is presented. The system GJ876, for example, is stable for reasonable strengths of the stochastic forces, but systems with lighter planets can get disrupted. Even if they are not disrupted completely, stochastic forces create characteristic, observable libration patterns. Turbulence plays also an important role during the early phases of the planet formation process. Due to the large separation of scales, the gravitational collapse in the core accretion model is very hard to model numerically. A scaled method is presented, that allows for the correct treatment of self-gravity in a marginally collisional system by taking into account the relevant small scale processes. Interestingly, this system is dynamically very similar to Saturn's rings. The stochastic migration of small bodies in Saturn's rings is also studied. Analytic predictions of the interactions of a moonlet/propeller with ring particles are compared to collisional N-body simulations with up to a million particles. The random walk is fast enough to be directly observable by the Cassini spacecraft.

1214960

2011A&A...530A...6S

1012

1012.1867_arXiv.txt

Cataclysmic variables (CVs) are close binary systems consisting of a late-type star transferring material onto a white dwarf (WD) companion via Roche-lobe overflow, and typically have orbital periods of the order of hours. Mass transfer is driven by the loss of orbital angular momentum, which in turn is thought to be caused by magnetic braking of the secondary star, except at the shortest periods ($\sim$ 2 hours or less), where gravitational radiation dominates. This mass transfer leads to the formation of an accretion disc subject to various instabilities, which can then produce outburst events such as dwarf novae (\citealt{warner}) Magnetic CVs (mCVs) are a subset of the catalogued CVs ($\approx 10\% - 20\%$, \citealt{downes}; \citealt{RKcat}), and generally fall into two categories: polars (or AM Her types after the prototype system) and intermediate polars (IPs or DQ Her types). The WDs in polars have sufficiently strong magnetic fields ($10^7 - 10^9$ Gauss) to prevent the formation of an accretion disc and to lock the secondary late-type star, synchronising the whole system (for a review of polars, see \citealt{cropper}). The strong magnetic field in these systems is confirmed by strong linear and circular polarisation, together with measurements of cyclotron humps (\citealt{warner}). For IPs, the lack of strong optical polarisation implies a weaker magnetic field, which is not powerful enough to synchronise the secondary star (for a review of IPs, see \citealt{patterson}). In these systems, material leaving the $L_{1}$ point only forms an accretion disc up to the point where the magnetic pressure exceeds the ram pressure of the accreting gas. From this point onwards the accretion dynamics are governed by the magnetic field lines, which channel the material onto the WD magnetic poles. The nature of these systems is confirmed by the detection of coherent X-ray and/or optical modulations associated with the spin period of the WD. It is however not uncommon for the dominant optical frequency to differ from the X-ray frequency. When this occurs the X-ray signal is usually associated to white dwarf spin frequency ($\omega$), and the optical frequency is most often observed to be $\omega - \Omega$, where $\Omega$ is the orbital frequency. These optical modulations are thought to arise from X-ray reprocessing within the surrounding medium, where the X-ray emission from the WD pole illuminates an unknown structure fixed in the reference frame rotating with the binary (\citealt{warner81}). Other ``orbital side bands'' (e.g. $\omega + \Omega$) can also be produced, and have been detected in some IPs as well (e.g. \citealt{warner86}). Hard X-ray surveys such as the {\it INTEGRAL}/IBIS survey (\citealt{cat4}) have proven remarkably efficient in detecting mCVs, and in particular intermediate polars with spin-to-orbital ratios below 0.1. (\citealt{scaringi,barlow}). The low and persistent hard X-ray flux of IPs above 15 keV was also expected because previous X-ray spectroscopy at lower energies revealed a hard X-ray excess (\citealt{lamb,chanmugam}). Here we report on the optical photometry and time-resolved spectroscopy of one candidate IP, 1RXS J165443.5$-$191620 (hereafter RXJ1654). This \textit{Rosat} detection has been also catalogued in the {\it INTEGRAL}/IBIS survey (\citealt{cat4}); but remained unidentified until a tentative classification by \cite{masettiVIII} based on optical (follow-up) spectroscopy. The spectra displayed clear Balmer and HeI emission lines, with HeII~$\lambda$4686\AA/H$\beta$ equivalent width ratios $\ge0.5$, suggesting an intermediate polar classification for this CV. Recently \cite{lutovinov} have obtained optical photometry for RXJ1654 that displays clear modulation of $\approx 550$ seconds, consistent with that of the WD spin, which enforces the tentative classification of this object as an IP, but does not confirming its true nature. \cite{pretorius} has already discussed the dangers of inferring IP classifications based on the optical spectra and detection above 15 keV alone. Out of a sample of five hard X-ray emitting systems with tentative IP classification from optical spectroscopy, three showed clear signs of orbital and spin modulations, whilst two did not. Strong corroboration of the suspected IP status can come from long time-span optical photometry showing the expected sideband structure, and time-resolved radial velocity measurements showing periodic variations on the orbital period. Photometry yields incontrovertible orbital periods when eclipses are present, but other modulations can masquerade as orbital periods as well. Time-resolved spectroscopy is thus the most reliable technique for determining orbital periods of CVs. This together with the detection of the $\omega - \Omega$ frequency, and/or coherent X-ray modulations consistent with $\omega$ (\citealt{buckley00}), would then unambiguously confirm the nature of an IP, excluding any other origin to the observed photometric variations. In Section~2 we will present preliminary observations obtained with the IAC80 during the university of Southampton student field trip. The data displayed a candidate WD spin modulation, which then motivated us towards obtaining a better sampled lightcurve spanning a wider range in time for this system to search for orbital periodicities. Section~3 describes the effort and data analysis to achieve this on a variety of telescopes with the global network of astronomers part of the Center for Backyard Astrophysics (CBA). In this respect, we confirm and update the \cite{lutovinov} candidate spin period, and also discuss the detection of the orbital period and the $\omega - \Omega$ signal. Additional time-resolved spectroscopy was also obtained on the Hiltner Telescope at MDM Observatory and this is described in Section~4, together with the data analysis and clear H$\alpha$ radial velocity shifts associated to the orbital period, which undoubtedly confirms the class membership of RXJ1654 as an IP.

Here we have presented optical follow-up observations of the candidate CV source RXJ1654. We updated the nature of this system as another hard X-ray emitting IP, confirming the \cite{lutovinov} detection of a $\approx$550 seconds in the system as the WD spin signal. This finding adds to the small, but fast growing, subset of IPs all displaying persistent hard X-ray flux above 15 keV, spin periods below 30 minutes, and spin-to-orbital ratios below 0.1.

1012.1867

<BR /> Aims: We investigate the physical nature of the X-ray emitting source <ASTROBJ>1RXS J165443.5-191620</ASTROBJ> through optical photometry and time-resolved spectroscopy. <BR /> Methods: Optical photometry is obtained from a variety of telescopes all over the world spanning ≈27 days. Additionally, time-resolved spectroscopy is obtained from the MDM observatory. <BR /> Results: The optical photometry clearly displays modulations consistent with those observed in magnetic cataclysmic variables: a low-frequency signal interpreted as the orbital period, a high-frequency signal interpreted as the white dwarf spin period, and an orbital sideband modulation. Our findings and interpretations are further confirmed through optical, time-resolved spectroscopy that displays Hα radial velocity shifts modulated on the binary orbital period. <BR /> Conclusions: We confirm that <ASTROBJ>1RXS J165443.5-191620</ASTROBJ> is an intermediate polar with a spin period of 546 s and an orbital period of 3.7 h. In particular, <ASTROBJ>1RXS J165443.5-191620</ASTROBJ> is part of a growing subset of intermediate polars that display hard X-ray emission above 15 keV, white dwarf spin periods below 30 min, and spin-to-orbital ratios below 0.1.

2457813

2011MNRAS.412L..45P

1012

1012.4445_arXiv.txt

Very little is known about the magnetic fields of hot, massive OB stars. Due at least in part to the challenges of measurement, direct evidence of magnetic fields in massive stars is remarkably rare - there are perhaps 30 known examples - and few have been studied in detail. Even as a member of the elusive class of magnetic massive stars, the B0.2\,V star $\tau$\,Sco is recognized to be a peculiar and outstanding object. First, $\tau$~Sco is distinguished by its magnetic field, which is unique because it is structurally far more complex than the mostly-dipolar (multipole mode $l=1$) fields usually observed in magnetic OB stars. According to the Zeeman Doppler Imaging performed by \citet{2006MNRAS.370..629D}, the magnetic topology of $\tau$\,Sco presents significant power in multipole modes up to $l=5$ with a mean surface field strength of $\sim300$\,G, and the extrapolated 3D field structure shows an intricate assortment of open field lines and closed magnetic loops \citep[illustrated in Fig. 11 and 12 of][]{2006MNRAS.370..629D}. Secondly, $\tau$\,Sco also stands out from the crowd of early-B stars because of its stellar wind anomalies, as diagnosed through its odd UV spectrum and X-ray emission. The optical and non-resonance UV lines of $\tau$~Sco, that probe the conditions of the photosphere, are typical of an early-B dwarf star. However, the UV resonance lines arising from the stellar wind are not those typical of such a star. Figure \ref{fig|uv} shows a comparison of the UV spectrum of $\tau$\,Sco (second spectrum) with the spectrum of a normal B0 main sequence star (bottom spectrum) and a B0 giant (top spectrum). The resonance wind lines of C\,\textsc{iv} (right panel), Si\,\textsc{iv} and N\,\textsc{v} (left panel) are indicated with dashed lines. The UV wind line morphology of normal, early B stars conforms to a 2-D spectral grid \citep{1995NASRP1363.....W}. Typically, C\,\textsc{iv} strengthens with increasing temperature and luminosity. N\,\textsc{v} is very weak on the main sequence at spectral type B0\,V but strengthens with temperature and luminosity. For stars with fixed C\,\textsc{iv} and N\,\textsc{v} profiles, Si\,\textsc{iv} is strictly luminosity dependent and breaks the degeneracy. $\tau$\,Sco does not fit into this grid. It has strong N\,\textsc{v} indicating that it should be well above the main sequence. However, its C\,\textsc{iv} lines are only slightly stronger than typical and not distinctly wind-like, suggesting a near main sequence luminosity. Finally, its Si\,\textsc{iv} profiles are unique, a bit stronger than typical class V stars, but unlike normal, early giants. As a result, this stars lie outside of the normal classification grid, suggesting a more highly ionized outflow than typical. The hard X-ray emission of $\tau$\,Sco also suggests hot plasma, in excess of 10\,MK \citep{1994ApJ...421..705C}. We suspect that this is connected to the strange UV properties. Like its magnetic field, the stellar wind of $\tau$\,Sco possesses observational characteristics that make it unique among OB stars. Interestingly, the wind lines of $\tau$~Sco have been shown to vary periodically with the star's 41~d rotation period \citep{2006MNRAS.370..629D}, which is interpreted to mean that the wind is structured and modulated by the magnetic field. Clearly the magnetic field exerts an important influence on the wind dynamics. What is not clear is whether the wind-line anomalies described above are a consequence of the unusual complexity of $\tau$ Sco's magnetic field, a general consequence of wind confinement in this class of star, or perhaps even unrelated to the presence of a magnetic field. Because such wind anomalies have never been observed in any other star, magnetic or not, this issue has remained unresolved. The identification and analysis of additional stars with wind properties similar to $\tau$~Sco would therefore represent an important step toward understanding the origin of these peculiarities. In this paper, we present two early B-type stars - HD\,66665 and HD\,63425 - that we identified to be the first $\tau$\,Sco analogues based on their UV spectra, which are strikingly similar to the UV spectrum of $\tau$\,Sco. Figure \ref{fig|uv} shows \text{IUE} archives spectra of these two stars (third to fifth spectra from the top). The discovery of wind anomalies naturally led to an investigation of the magnetic properties of these stars by the Magnetism in Massive Stars (MiMeS) collaboration \citep{2010arXiv1009.3563W}. \begin{figure*} \includegraphics[width=87mm]{fig_uv2.ps} \includegraphics[width=87mm]{fig_uv1.ps} \caption{\label{fig|uv} IUE spectra of $\tau$\,Sco and its analogues HD\,66665 and HD\,63425 (second to fifth spectra). For comparison, typical spectra of a B0 dwarf and a B0 giant are shown (bottom and top respectively). The dashed lines indicate the wind lines N\,\textsc{v}\,$\lambda \lambda 1239, 1243$, Si\,\textsc{iv}\,$\lambda\lambda 1393, 1403$ and C\,\textsc{iv}\,$\lambda \lambda 1548, 1550$. The best-fit synthetic spectra, from which stellar parameters were derived from for HD\,66665 and HD\,63425, are shown in red. } \end{figure*} \begin{table} \caption{\label{tab|obslog} Journal of ESPaDOnS observation listing the UT date, the heliocentric Julian date (2\,400\,000+), the total exposure time and the s/n per 1.8\,km\,s$^{-1}$ velocity bin at 540\,nm. Column 5 gives the probability that the observed signal in Stokes V LSD profiles is due only to noise, and column 6 the global longitudinal field component obtained from the first moment of the Stokes V profiles. The available archival \textit{IUE} spectra are also listed.} \begin{tabular}{c c c c c D{*}{\,\pm\,}{3,3} } \hline Date & HJD & $t_\mathrm{exp}$ & s/n & FAP & \mc{$B_l~^1$} \\ (UT) & & (s) & & & \mc{(G)} \\ \hline \multicolumn{6}{c}{\textbf{HD\,66665}} \\ 2010/02/26 & 55\,253.97 & $4\,400$ & 728 & $<10^{-8}$ & -93*4 \\ 2010/03/02 & 55\,257.98 & $1\,200$ & 236 & $<10^{-8}$ & -3*15 \\ 2010/03/04 & 55\,259.98 & $1\,200$ & 330 & $<10^{-8}$ & 9*10 \\ 2010/03/05 & 55\,260.98 & $1\,200$ & 276 & $2\times10^{-6}$ & 41*13 \\ 1983/12/04 & 45\,436.77 & 900 & \multicolumn{3}{c}{IUE swp\,21680} \\ 1987/03/11 & 46\,866.49 & 420 & \multicolumn{3}{c}{IUE swp\,30484} \\ 1987/03/12 & 46\,866.58 & 1\,080 & \multicolumn{3}{c}{IUE swp\,30486} \\ \hline \multicolumn{6}{c}{\textbf{HD\,63425}} \\ 2010/02/28 & 55\,255.84 & $1\,200$ & 520 & $<10^{-8}$ & 117*8 \\ 2010/03/03 & 55\,258.88 & $1\,200$ & 289 & $<10^{-8}$ & 113*15 \\ 2010/03/07 & 55\,262.76 & $1\,200$ & 435 & $<10^{-8}$ & 132*10 \\ 2010/03/08 & 55\,263.83 & $1\,200$ & 361 & $<10^{-8}$ & 125*12 \\ 1987/03/12 & 46\,866.53 & 600 & \multicolumn{3}{c}{IUE swp\,21679} \\ 1983/12/04 & 45\,672.74 & 720& \multicolumn{3}{c}{IUE swp\,30485} \\ \hline \end{tabular} $^1$ A detectable Stokes V signal can be observed with high-resolution instruments even if the global longitudinal field is null \citep[see][]{2010arXiv1010.2248P}. \end{table} \begin{figure} \includegraphics[width=84mm]{fig_paper.ps} \caption{\label{fig|vis} Fits to the optical He\,\textsc{i-ii} and H$\beta$ lines for HD\,66665 (left) and HD\,63425 (right) } \end{figure} \begin{table} \caption{\label{tab|param} Summary of stellar properties of $\tau$\,Sco, HD\,66665 and HD\,63425. } \begin{tabular}{l c c c} \hline & $\tau$\,Sco$^1$ & HD\,66665 & HD\,63425 \\ \hline Spec. type & B0.2\,V & B0.5\,V & B0.5\,V\\ $T_\mathrm{eff}$ (kK)& $31\pm1$ & $28.5\pm1.0$ & $29.5\pm1.0$ \\ $\log g$ (cgs)& $4.0\pm0.1$ & $3.9\pm0.1$ & $4.0\pm0.1$ \\ $R_\star$ (R$_\odot$) & $5.6\pm0.8$ & $5.5\stackrel{+3.3}{_{-2.7}}$ & $6.8\stackrel{+4.8}{_{-2.0}}$ \\ $\log(L_\star/\mathrm{L}_\odot)$ & $4.47\pm0.13$ & $4.25\stackrel{+0.40}{_{-0.60}}$ & $4.50\stackrel{+0.46}{_{-0.30}}$ \\ $M_\star$ (M$_\odot$) & $11\pm4$ & $9\stackrel{+5}{_{-7}}$ & $17\stackrel{+34}{_{-9}}$ \\ $v\sin i$ (km\,s$^{-1}$) & $<13$ & $\lesssim10$ & $\lesssim15$ \\ $\dot{M}$ ($10^{-9}$\,M$_\odot$\,yr$^{-1}$) & $61\stackrel{+10}{_{-2}}$ & $< 0.45$ & $< 0.75$ \\ $v_\infty$ (km\,s$^{-1}$)& $\sim2000$ & $\sim1400$ & $\sim1700$ \\ \hline \end{tabular} \medskip $^1$ Paramters from \citet{2006A&A...448..351S}, and \citet{2005A&A...441..711M} for $\dot{M}$ and $v_\infty$. \end{table}

We have presented the characteristics of two stars - HD\,66665 and HD\,63425 - which we believe are analogues to the magnetic massive star $\tau$\,Sco for two main reasons: (i) The UV spectra of these stars are similar to the once-unique spectrum of $\tau$\,Sco. (ii) We have shown that these two stars host magnetic fields. We have shown that all three of these stars have similar fundamental properties. However, the mass-loss rates we estimate for HD\,66665 and HD\,63425 are lower than the value adopted for $\tau$\,Sco by \citet[][$6\times10^{-8}$\,M$_\odot$\,yr$^{-1}$]{2006MNRAS.370..629D} in their analysis of that star. However, that mass-loss rate \citep{2005A&A...441..711M} is determined from the optical spectrum of $\tau$\,Sco, and could differ systematically from that determined from UV line profiles -- in particular in the presence of a magnetic field. The derived mass-loss rates for HD\,66665 and HD\,63425 are sufficiently low that they should be considered as stars presenting the weak wind problem \citep{2005A&A...441..735M,2009A&A...498..837M}, since the expected mass-loss rate is about $10^{-8}$\,M$_\odot$\,yr$^{-1}$ -- about 20 time larger than the inferred rates \citep[following the recipe provided by][]{2000A&A...362..295V}. Our modelling of the LSD Stokes V profiles assuming an inclined dipole model results in surface field strengths\footnote{The surface-averaged modulus of a dipolar magnetic field is equal to 0.77 times the dipole polar field strength.} that are comparable to, or perharps slightly larger than, the mean surface field strength of $\tau$\,Sco. The current observations can be acceptably reproduced by the dipole model, although more phase-resolved observations are required in order to assess the potential complexity of their magnetic fields, and verify if the wind anomalies are linked to the field complexity.

1012.4445

The B0.2 V magnetic star τ Sco stands out from the larger population of massive OB stars due to its high X-ray activity, peculiar wind diagnostics and highly complex magnetic field. This Letter presents the discovery of the first two τ Sco analogues - HD 66665 and HD 63425, identified by the striking similarity of their ultraviolet (UV) spectra to that of τ Sco. ESPaDOnS spectropolarimetric observations were secured by the Magnetism in Massive Stars CFHT Large Program, in order to characterize the stellar and magnetic properties of these stars. CMFGEN modelling of optical ESPaDOnS spectra and archived IUE UV spectra showed that these stars have stellar parameters similar to those of τ Sco. A magnetic field of similar surface strength is found on both stars, reinforcing the connection between the presence of a magnetic field and wind peculiarities. However, additional phase-resolved observations will be required in order to assess the potential complexity of the magnetic fields and verify if the wind anomalies are linked to this property. Based on observations obtained at the Canada--France--Hawaii Telescope (CFHT) which is operated by the National Research Council of Canada, the Institut National des Sciences de l'Univers of the Centre National de la Recherche Scientifique of France and the University of Hawaii.

12134266

2010ASPC..430...29A

1012

1012.0500_arXiv.txt

When considering the direct imaging of extrasolar planetary systems, one is faced with two main challenges: the small angular separation and the high contrast between exoplanets and their host stars. If the goal is to characterise the mid-infrared emission of exoplanets in the habitable zone of nearby stars, the typical angular resolution of 50\,mas required to resolve a Sun-Earth system at 20\,pc leads to an impractical aperture size of 40\,m. The only option is then stellar interferometry, which synthesises the resolving power of a larger aperture by using multiple telescopes separated by an appropriate distance called the interferometric baseline $B$. The associated angular resolution equals $\lambda/2B$ (to be compare with $\lambda/D$ for a single aperture of diameter $D$). The need for interferometry to image planetary systems in the infrared was recognised already in the late 70s \citep{Bracewell78}. Its application to the search and characterisation of habitable worlds was proposed 15~years later by \citet{Leger93}. Reaching the appropriate angular resolution to separate the signals of the planet from that of its host star is only half of the solution, and the high dynamic range still needs to be addressed. In this review, we discuss the various techniques that can be implemented with infrared interferometers to reach the required dynamic range to image extrasolar planetary systems of various kinds (from dusty discs and hot giant planets down to Earth-like planets).

1012.0500

In this paper, we review the various ways in which an infrared stellar interferometer can be used to perform direct detection of extrasolar planetary systems. We first review the techniques based on classical stellar interferometry, where (complex) visibilities are measured, and then describe how higher dynamic ranges can be achieved with nulling interferometry. The application of nulling interferometry to the study of exozodiacal discs and extrasolar planets is then discussed and illustrated with a few examples.

12213515

2011MNRAS.414.2354F

1012

1012.2958_arXiv.txt

Evidence is growing for magnetic fields on larger and larger scales in the Universe (see e.g. the reviews \cite{Giovannini:2003yn,Subramanian:2008tt}). In galaxies, the fields have strength of order $\mu$Gauss, ordered on scales $\sim 1-10\,$kpc. Fields of strength $\sim 1-10^{-2}\mu$G on scales $\sim 0.1-1\,$Mpc have been detected in galaxy clusters, and there is evidence of magnetic fields in superclusters. Recently, new evidence has been presented for intergalactic magnetic fields: high energy gamma-rays from distant sources can initiate electromagnetic pair cascades when interacting with the extragalactic photon background; the charged component of the cascades will be deflected by magnetic fields, affecting the images of the sources. Using observations from FERMI, a lower bound of order $10^{-16}\,$G has been claimed for the strength of fields in the filaments and voids of the cosmic web \cite{2010ApJ...722L..39A,Neronov:1900zz,2010arXiv1009.1782D,Essey:2010nd}. The origin of these fields is still unclear (see e.g. \cite{Brandenburg:2004jv,Kulsrud:2007an,Kandus:2010nw}). They could have been generated via astrophysical processes during the nonlinear collapse stage of structure formation. There remain unresolved difficulties in explaining how these astrophysical seed fields lead to fields of the observed strength and coherence scales. Alternatively, the fields could be primordial seed fields -- created in the very early Universe, during inflation, or during subsequent phase transitions. In principle inflation can generate fields on all scales -- but unknown physics must be invoked to achieve non-minimal coupling of the electromagnetic field. The electroweak and QCD transitions can only produce fields on very small scales, up to the Hubble radius at magnetogenesis (and their amplitude is strongly constrained by their gravitational wave production before nucleosynthesis \cite{Caprini:2009pr}). Primordial magnetogenesis also takes place in the cosmic plasma after particle/anti-particle annihilation. This avoids the problem of exotic physics that faces inflationary magnetogenesis -- standard Maxwell theory and standard cosmological perturbations in the cosmic plasma inevitably lead to magnetic fields. It also avoids the small coherence scale problem facing electroweak and QCD fields. However, the problem is the weakness of the fields, since this effect occurs at second and higher order in cosmological perturbations. The key question is how weak is the field and how does it vary with scale? Differing qualitative estimates of the field strength have been given by \cite{Hogan:2000gv,Berezhiani:2003ik,Gopal:2004ut,Siegel:2006px,Kobayashi:2007wd,Maeda:2008dv}. The power spectrum was first numerically computed by \cite{Matarrese2004}, which differs significantly from ours. More recently, \cite{Ichiki:2007hu} presented a power spectrum that is closer to our result. We discuss below the differences between previous results and ours. Our analysis is the first complete general relativistic computation of the power spectrum, taking into account all effects. Our result is shown in Fig. \ref{fig1one}. The power spectrum behaves as \be \sqrt{k^3 P_B} \propto \left\{ \begin{array}{ll} k^4 & k\ll k_{\rm eq} \\ k^{0.5} & k \gg k_{\rm eq}\,. \end{array} \right. \ee On cluster scales the comoving field strength is \be\label{b1mpc} B_{1\,{\rm Mpc}} \sim 3\times 10^{-29}\,{\rm G}. \ee \begin{figure*}[!htp] \includegraphics[width=8.5cm]{MagnplotOneCurve.eps}\quad \includegraphics[width=8.5cm]{BatscaleOneCurve.eps} \caption{{\em Left:} Magnetic field spectrum today. {\em Right:} Comoving magnetic field strength today at a given scale.} \label{fig1one} \end{figure*} Thus the field generated around recombination is too weak to act as a seed for the observed field strength of order $\mu$G. Adiabatic contraction of the magnetic flux lines during nonlinear collapse of structures provides an enhancement of $\sim 10^3$, while the nonlinear dynamo mechanism has an amplification factor $\sim 10^{8}$ (with many remaining uncertainties). Note that hydrodynamical and turbulence effects during nonlinear collapse themselves generate a field of order $10^{-20}\,$G -- which is also too small to account for the observed galactic and cluster fields \cite{Kulsrud:2007an}. The field (\ref{b1mpc}) is also too weak to imprint detectable effects on the CMB. Nevertheless it is a real property of the standard cosmological model, and may have some impact on early structure formation during the `dark ages' if it is the only primordial field. (See e.g. \cite{Sethi:2009dd,Schleicher:2010ph} for the role of magnetic fields in structure formation during the dark ages.) As shown below, the magnetic field is given by \bea \label{Maxwell-0} \left(a^2 B^i\right)'&=&-a^2 \epsilon^{ijk}\partial_{j}\Big[\left(1+\Phi-\Psi \right)E_k \Big],\\ \label{GenerationE-0} E^i &\approx& -\frac{4\rho_\gamma \st}{3 e}\Big(\Delta v_{\ib\gamma}^i +\frac{2}{5}\Theta^i_jv_{\ib}^j\Big), \eea where $\Phi,\Psi$ are first-order metric perturbations, $\Delta v_{\ib\gamma}^i=v_{\ib}^i - v_{\gamma}^i$ is the photon-baryon velocity difference, and $ \Theta^i_j$ is the photon quadrupole moment, from anisotropic stress. This leads to three types of source terms for magnetogenesis: \bea \label{b-sources} \left(a^2 B\right)'&=& S_1\big[ \Delta v_{\ib\gamma}^{(2)}\big]+ S_2\big[\big\{\delta_\gamma^{(1)}+\Phi^{(1)}-\Psi^{(1)} \big\}\Delta v_{\ib\gamma}^{(1)} \big] \nonumber\\ &&~~ + S_3\big[\Theta_\gamma^{(1)} v_{\ib}^{(1)}\big]. \eea The first source term is second-order, while the other two are quadratic in first-order quantities. The contributions of the source terms to the power spectrum are shown in Fig. \ref{fig1} (left). Our paper builds on the physical analysis of nonlinear plasma dynamics presented in \cite{Maartens:1998xg,Matarrese2004,Ichiki:2007hu,Kobayashi:2007wd,Takahashi:2007ds,Maeda:2008dv,Pitrou2008}. The key features of the dynamics are as follows. \begin{itemize}\itemsep=-4pt \item The electric field ensures that the proton-electron relative velocity is always strongly suppressed in comparison with the photon-electron relative velocity -- even at high energies when the Compton interaction is stronger than the Coulomb interaction. \item Vorticity induced in the electron fluid is thus transferred almost entirely to the protons, and the baryon vorticity evolution is determined by the two-fluid dynamics of photons and baryons, which is very close to the equations of CMB dynamics. We use the second-order Boltzmann code of \cite{Pitrou2008}. \item The limit $v_{\rm e} - v_{\rm \gamma}\to 0$ and $v_{\rm p}-v_{\rm e}\to 0$ is not equivalent to setting $v_{\rm p}=v_{\rm e}=v_{\rm \gamma}$ in the momentum exchange equations, and the limit must be taken consistently. \item At first order, cosmological vector perturbations are zero after inflation, in the standard model. Magnetogenesis requires vortical currents, and these can therefore only be generated at second order, via mode-mode coupling of first-order scalar perturbations. This remains true even in the presence of topological defects, which are active sources for vector perturbations: at first order, the vector perturbations induced by the defects cannot break vorticity conservation in the cosmic plasma \cite{Hollenstein:2007kg}. \item On large scales there is some cancellation amongst the source terms in (\ref{b-sources}) (this is evident from Fig. \ref{fig1}). Neglecting any of the effects can thus lead to unreliable results. \item The magnetic field continues to be created after recombination, due to the residual nonzero ionization fraction. If the numerical integration is stopped at recombination, then the comoving field is under-estimated by a factor $\sim 10$ (see Fig. \ref{fig1}). \end{itemize} The plan of the paper is as follows. In the next section we review and clarify the magnetic and electric field generation beyond the tight-coupling limit. In Sec.~\ref{results}, we detail the numerical integration of the differential evolution equations at second order in cosmological perturbations that we perform in order to solve for the magnetic field spectrum. We also provide analytical insight into the time and scale behaviors of the numerical results. We compare our results with previous work in Sec.~\ref{SecComparison}. Details of some calculations are given in the Appendices.

\label{conclusions} We have performed for the first time a full numerical computation of the seed magnetic field generated by nonlinear dynamics, taking into account all general relativistic effects and all source terms. We discussed the range of applicability of the mechanism on cosmological scales and concluded that the generation of the magnetic field is directly related to the Compton drag by photons on baryons. Even in the tight coupling regime, photons exchange vorticity with baryons and the magnetic field is created. Since the electric field that sources the magnetic field does not depend on the fraction of free electrons, the magnetic field is still generated after recombination, given that there is a relic fraction of charged particles, and we find that the largest production takes place in this final stage. Our results are summarized in Fig. \ref{fig1one}. The power spectrum (left plot) behaves as \be \sqrt{k^3 P_B} \propto \left\{ \begin{array}{ll} k^4 & k\ll k_{\rm eq} \\ k^{0.5} & k \gg k_{\rm eq} \end{array} \right. \ee On cluster scales the comoving field strength is (right plot) \be B_{1\,{\rm Mpc}} \sim 3\times 10^{-29}\,{\rm G}. \ee

1012.2958

Non-linear dynamics creates vortical currents when the tight-coupling approximation between photons and baryons breaks down around the time of recombination. This generates a magnetic field at second order in cosmological perturbations, whose power spectrum is fixed by standard physics, without the need for any ad hoc assumptions. We present the fully general relativistic calculation of the magnetic power spectrum, including the effects of metric perturbations, second-order velocity and photon anisotropic stress, thus generalizing and correcting previous results. We also show that significant magnetogenesis continues to occur after recombination. The power spectrum ? decays as k<SUP>4</SUP> on large scales, and grows as k<SUP>0.5</SUP> on small scales, down to the limit of our numerical computations, ∼1 Mpc. On cluster scales, the created field has a strength of ∼3 × 10<SUP>-29</SUP> G.

12203208

2011CQGra..28h5018M

1012

1012.2443_arXiv.txt

1012.2443

Recently, the exact solutions of wormhole geometries supported by a nonminimal curvature-matter coupling were found, where the nonminimal coupling minimizes the violation of the null energy condition of normal matter at the throat. In this brief report, we present a solution where normal matter satisfies the energy conditions at the throat and it is the higher order curvature derivatives of the nonminimal coupling that are responsible for the null energy condition violation, and consequently for supporting the respective wormhole geometries. For simplicity, we consider a linear R nonmiminal curvature-matter coupling and an explicit monotonically increasing function for the energy density. Thus, the solution found is not asymptotically flat, but may in principle be matched to an exterior vacuum solution.

12163991

2011AIPC.1331..254F

1012

1012.1992_arXiv.txt

\noindent The transit technique involves searching for periodic dips in stellar light-curves due to the orbital revolution of a transiting body, blocking a fraction of the stellar light. For a given planetary radius, the transit depth ($\delta$) is proportional to (R$_{pl}$/R$_*)^2$. Therefore, planets orbiting solartype stars have extremely shallow eclipses, blocking $\sim$1\% of the light for a giant planet and $\sim$0.01\% of the light for an Earth-sized planet. Current ground-based wide-field surveys can achieve the necessary photometric accuracy of better than 1\%, only for the brightest stars ($V\sim$9-12 in the case of WASP), so the bulk of the planets discovered by transit surveys around main-sequence stars have radii in the range R$_{pl}\sim$0.8$-$1.8 R$_{jup}$. A strong advantage over main sequence star primaries is offered by white dwarf stars. White dwarfs (WD) are compact degenerate objects, with approximately the same radius as the Earth, and represent the final stage of evolution of main-sequence stars with masses~$<$~8M$_{\odot}$ (i.e. $\sim$97\% of all stars in our Galaxy). Any sub-stellar or gas giant companion in orbit around a white dwarf will completely eclipse it, while bodies as small as the Moon will have relatively large transit depths ($\sim$~3\%), with the only caveat being that it remains unclear as to whether any such systems survive beyond the latter stages of stellar evolution.\\ \noindent Sub-stellar companions to white dwarfs are rare (\citealt{Farihi05} using 2MASS estimated that $<0.5\%$ of WDs have L dwarf companions). At the time of writing only three wide white dwarf + brown dwarf (WD+BD) systems have been spectroscopically confirmed, GD\,165 \citep{Becklin88}, PHL5038 \citep{Steele09}, and LSPM~$1459+0857$\,AB \citep{Day-Jones} and two detached, non-eclipsing, short-period WD+BD systems are currently known, WD$0137-349$ (\citealt{Maxted06}, \citealt{wd0137b}, P$\approx$116mins), and GD1400 (\citealt{Farihi04}, \citealt{Dobbie05}, \citealt{Burleigh10}, P$\approx$9.9h). The latter, is currently the lowest mass ($\sim$50M$_{jup}$) object known to have survived CE evolution. Although infrared surveys such as UKIDSS, VISTA and WISE, and observatories such as Spitzer hope to reveal many more such binaries, they remain difficult to identify either as infra-red excesses or through radial velocity measurements. In addition the detection of a significant number of eclipsing WD+BD binary systems might help uncover the hypothesised population of `old' cataclysmic variables (CVs) in which the companion has been reduced to a sub-stellar mass (e.g. \citealt{Patterson98}; \citealt{Patterson05}; \citealt{Littlefair03}). These systems are undetectable as X-ray sources and difficult to identify in optical and infra-red surveys. \citet{Littlefair06} confirmed the first such system through eclipse measurements, while \citet{Littlefair07} showed that another eclipsing CV, SDSS~J$150722.30+523039.8$, was formed directly from a detached WD+BD binary.\\ \noindent Several theoretical studies discuss post-main sequence evolution of planetary systems and show that planetary survival is not beyond possibility (\citealt{Duncan98}; \citealt{Debes02}; \citealt{Burleigh02}; and \citealt{Villaver07}). Radial velocity observations of red giants indicate that planets in orbits beyond the red giant's envelope can survive stellar evolution to that stage (see \citealt{Frink02}; \citealt{Hatzes05}, \citealt{Sato03}). Moreover, \citet{Silvotti07} reported the detection of a $\sim$3M$_{jup}$ planet orbiting an extreme horizontal branch star, and \citet{Mullally08} found convincing evidence of a 2M$_{jup}$ planet in a 4.5 year orbit around a pulsating WD. Furthermore, \citet{Beuermann10} reported the detection of two planetary companions (M$_c$=6.9M$_{Jup}$ and M$_d$=2.2M$_{Jup}$) in the post common envelope binary NN Ser (ab) via measurements of a light-travel-time effect superposed on the linear ephemeris of the binary; showing that planets do survive stellar evolution.\\ \noindent Short-period rocky companions to white dwarfs may seem less likely. \citet{Villaver07} suggested that planets in orbit within the reach of the AGB envelope will either evaporate or in rare cases, more massive bodies may accrete mass and become close companions to the star. Planets in wide orbits that escape engulfment by the red giant or asymptotic giant will move outwards to conserve angular momentum (as described by Jeans 1924). \citet{Duncan98} found that for WD progenitors experiencing substantial mass loss during the AGB phase, planetary orbits become unstable on timescales of $\leq$10$^8$~year. \citet{Debes02} found that the mass loss is sufficient to destabilise planetary systems of two or more planets and that the most likely result is that one planet would be scattered into an inner orbit (occupied, before the RGB phase, by a `now evaporated' inner planet), while the other would either be boosted into a larger orbit, or ejected from the system altogether.\\ \noindent The above scenario provides a plausible explanation for the recent detection of silicate-rich dust discs around a growing number of white dwarfs at orbital radii up to $\sim$1R$_{\odot}$ (e.g. \citealt{Reach05}; \citealt{Farihi07}, 2008; \citealt{Jura03}). \citet{Jura03} suggests that the formation of dust discs around white dwarfs is most probably due to the tidal disruption of an asteroid or larger body which has strayed too close to the parent star. (\citealt{Jura09}) suggest that the disc around GD362 originated from the tidal destruction of a single massive body such as Callisto or Mars.\\ \noindent The detection of short period sub-stellar and planetary mass companions to white dwarfs, will open an exciting chapter in the study of exoplanet evolution, constraining theoretical models of CE evolution and helping us to understand the ultimate fate of hot Jupiter systems as well as the fate of our own solar system in the post main-sequence phase. Here we present some of the results of our study which investigated the detection limits for transiting sub-stellar and terrestrial companions in close orbits around white dwarfs (for more details see \citealt{Faedi10}).\\ \noindent In $\S$2 we discuss the characteristics of the transit signals, the parameter space investigated and our detection method. In $\S$3 we analysed a sample of 194 WDs in the WASP archive. Finally in $\S$4 we discuss our conclusions.

We have investigated the detection limits for sub-stellar and planetary companions to white dwarfs using in the WASP survey. We found that Mercury-sized bodies at small orbital radii can be detected with good photometric data even in the presence of red noise. For smaller bodies red noise in the light-curves becomes increasingly problematic, while for larger orbital periods, the absence of significant numbers of in-transit points, significantly decreases our detection sensitivity. Application of our modified BLS algorithm to search for companions to WDs in our sample of 194 stars available in the WASP archive, did not reveal any eclipsing or transiting sub-stellar or planetary companions. We have used our results, to place upper limits to the frequency of sub-stellar and planetary companions to WDs. While no useful limits can be placed on the frequency of Mercury-sized or smaller companions, slightly stronger constraints can be placed on the frequency of BDs and gas giants with periods $<0.1-0.2$days, which must certainly be relatively rare ($<10\%$). More stringent constraints would requires significantly larger WD samples. Our key conclusion from simulations and analysis, using WASP data, suggests that photometric precision is of secondary importance compared to a high cadence and continuous coverage. The short duration of eclipses and transits of WDs compared to the WASP observing cadence, appears to be the main factor limiting the transit detection rate in a survey optimised for planetary transits of main sequence stars. Future surveys such as Pan-STARRS and LSST will be capable of detecting tens of thousands of WDs. However, we emphasise the importance of high cadence and long baseline observation when attempting to detect the signature of close, eclipsing and transiting sub-stellar and planetary companions to WDs. Space missions such as {\it COROT}, {\it Kepler} (see \citealt{DiStefano}) and, especially, {\it PLATO} may therefore be better suited to a survey of white dwarfs as they deliver uninterrupted coverage at high cadence and exquisite photometric precision ($\sim 10^{-4}-10^{-5}$) and could at least in principle detect the transits of asteroid-sized bodies across a white dwarf.

1012.1992

We used photometric data from the WASP (Wide-Angle Search for Planets) survey to explore the possibility of detecting eclipses and transit signals of brown dwarfs, gas giants and terrestrial companions in close orbit around white dwarfs. We performed extensive Monte Carlo simulations and we found that for Gaussian random noise WASP is sensitive to companions as small as the Moon orbiting a V~12 white dwarf. For fainter stars WASP is sensitive to increasingly larger bodies. Our sensitivity drops in the presence of co-variant noise structure in the data, nevertheless Earth-size bodies remain readily detectable in relatively low S/N data. We searched for eclipses and transit signals in a sample of 194 white dwarfs in the WASP archive however, no evidence for companions was found. We used our results to place tentative upper limits to the frequency of such systems. While we can only place weak limits on the likely frequency of Earth-sized or smaller companions; brown dwarfs and gas giants (radius~=RJup) with periods &lt;=0.2 days must certainly be rare (&lt;10%). More stringent constraints requires significantly larger white dwarf samples, higher observing cadence and continuous coverage. The short duration of eclipses and transits of white dwarfs compared to the cadence of WASP observations appears to be one of the main factors limiting the detection rate in a survey optimised for planetary transits of main sequence stars.

901402

2011MNRAS.413.1657P

1012

1012.4579_arXiv.txt

Non-thermal images of SNRs are rich sources of information about the properties of interstellar magnetic field (ISMF) behavior and kinetics of charged particles in vicinity of the strong non-relativistic shocks. Despite of their importance, images of SNRs -- in contrast to broad-band spectra -- are not well studied. Synchrotron X-ray brightness profiles were used as diagnostic tools for the estimate of the post-shock magnetic field in some SNRs \citep[e.g.][]{Ber-Volk-2004-mf}. Radio azimuthal profiles were used for determination of some properties of SN~1006 \citep[][hereafter Paper I]{pet-SN1006mf} and X-ray radial profiles were used to detect the shock precursor in SN~1006 and prove particle acceleration \citep{Morlino-etal-2010}. A detailed approach to modeling the {synchrotron} images of adiabatic SNRs in uniform ISMF and uniform interstellar medium (ISM) is developed by \citet{Reyn-98}. \citet{reyn-fulbr-90,Reyn-98,Reyn-04} use modeled synchrotron maps of SNRs to put constraints on properties of accelerated particles. Properties of the inverse-Compton (IC) \g-ray maps are investigated and compared to radio images in \citet[][hereafter Paper II]{thetak}. The influence of nonuniform ISM and/or nonuniform ISMF on the thermal X-ray morphology of adiabatic SNRs are studied in \citet{1999A&A...344..295H}, on the radio maps in \citet{Orletal07} and on the synchrotron X-ray and IC \g-ray images in {\citet{Orletal10}}. All studies of SNR maps assume classic MHD and test-particle theory of acceleration. Though they neglect effects of the back-reaction of the efficiently accelerated particles, they are able to explain general properties of the distribution of the surface brightness in radio, X-rays and \g-rays. This is because the classic theory, in contrast to the non-linear one, is able to deal with oblique shocks, that is vital for synthesis of SNR images. At present, the theory which considers effects of accelerated particles on the shock and on acceleration itself is developed for the initially quasi-parallel shocks only. One may therefore model the only radial profiles of brightness, in the rather narrow region close to the shock (in order to be certain that obliquity does not introduce prominent modifications). Effects of non-linear acceleration on the radial profiles of brightness are considered in \citet{Ell-Cassam2005-profiles,Decours-2005-prof,Ell2008-images,Zirakashv-Aha-2010}. Future studies on SNR morphology should take into account the NLA effects. Nevertheless, the classic approach is able to reveal the general properties of SNR maps determined by MF behavior and particle acceleration. Beside that, it is important to know the properties of the `classic' images because they create the reference base for investigation of the efficiency of NLA effects in the surface brightness distribution of SNRs. In Paper I, we introduced a method to derive an aspect angle of ISMF from the radio brightness of SNR. In Paper II, we synthesized radio and IC \g-ray maps and concluded that {coincidence of the position of the \g-ray and radio limbs is not a common case in theoretical models, because different parameters are dominant in determintion of the radio and \g-ray brightness variations. On the other hand, radio, X-ray and \g-ray observations \citep{pet-SN1006mf,SN1006Marco,HESS-sn1006-2010} show that radio, X-ray and $\gamma$-ray limbs coincide in SN 1006. As discussed in Paper II, such coincidence might be attributed to a combination of obliquity dependences of magnetic field and properties of emitting particles, as well as orientation versus the observer.} In the present paper, we make a step forward extending our model of non-thermal leptonic emission of Sedov SNRs in uniform medium to the X-ray band. We also derive brightness profiles for representative parameters which are suited for the comparison with adiabatic SNRs. Moreover, we derive analytical approximations of the azimuthal and radial profiles of radio, X-ray and \g-ray brightness which are extremely useful to demonstrate their dependence on the acceleration parameters, magnetic field orientation and the viewing geometry; they can also be directly and very easily fitted to SNR images to derive estimations for the best-fit quantities. While the analytical approximation cannot substitute the more accurate numerical simulations, we show that they retain enough accuracy to represent an effective diagnostic tool for the study of non-thermal SNR shells.

\label{xmaps:conclusion_section} The present paper extends analysis of properties of the surface brightness distribution of spherical adiabatic SNRs started in Paper I (radio band) and Paper II (IC \g-rays) to the nonthermal X-rays. It also generalizes the method of approximate analytical description of the azimuthal and radial profiles of brightness introduced in these papers. Synchrotron images of adiabatic SNR in X-rays are synthesized for different assumptions about obliquity variations of the injection efficiency, MF and maximum energy of accelerated electrons. We analyze properties of these images. Different models of electron injection (quasi-parallel, isotropic and quasi-perpendicular) as well as models of the electron maximum energy (time-limited, loss-limited and escape-limited) are considered. The azimuthal variation of the synchrotron X-ray and IC \g-ray brightness is mostly determined by variations of $\varsigma$, $\sigma\rs{B}$ and $E\rs{max}$, of the radio brightness by $\varsigma$ and $\sigma\rs{B}$ only. In general, higher $B$ increases X-ray and decreases IC \g-ray brightness. Really, higher MF is a reason of larger losses of emitting electrons (i.e. decrease of their number) and thus of the smaller brightness due to IC process. In contrast, X-rays are more efficient there because $S\rs{x}\propto B^{3/2}$. The radial profiles of brightness depend on a number of factors. It is quite sensitive to the adiabatic index: $\gamma<5/3$ makes plasma more compressible. Therefore, the brightness profile is thinner due to larger compression factor, larger gradient of density and MF downstream of the shock and larger radiative losses. The role and importance of various factors on the surface brightness in radio, synchrotron X-rays and IC \g-rays are demonstrated by the approximate analytical formulae. They accurately represent numerical simulations close to the shock and are able to account for some non-linear effects of acceleration if necessary. This makes the approximations a powerful tool for quick analysis of the surface brightness distribution due to emission of accelerated electrons around SNR shells. The application of the approximate formulae to the case of SN1006 yields measures of the aspect angle and the post-shock MF in good % agreement with more accurate analysis found in the literature. \begin{figure} \centering \includegraphics[width=7.6truecm]{radial_xray_NE_ap.eps} \caption{Radial profiles of the hard X-ray brightness of SN~1006. Experimental data \citep[from Fig. 4A in][]{Long-et-al-2003} are shown by histogram. Profiles given by approximation (\ref{xmaps:xray_brightness}) for $B=50,100,150\un{\mu G}$ are shown by blue, green and red lines respectively. Other parameters: $\varepsilon=1.2\un{keV}$ and $\nu\rs{break\|}=150\un{eV}$, $\phi\rs{o}=68^\mathrm{o}$, $\varphi=70^\mathrm{o}$, % ${\cal E}\rs{max}=2.9$, $q=0$, $b=0$, $\alpha=1$. } \label{xmaps:fig_xray_ap_sn1006} \end{figure}

1012.4579

Distributions of non-thermal surface brightness of supernova remnants (SNRs) contain important information about the properties of magnetic field and acceleration of charged particles. In the present paper, the synchrotron radio, X-ray, and inverse-Compton (IC) γ-ray maps of adiabatic SNRs in uniform interstellar medium and interstellar magnetic field are modelled and their morphology is analysed, with particular emphasis on comparison of azimuthal and radial variations of brightness in radio, X-rays and γ-rays. Approximate analytical formulae for the azimuthal and radial profiles of the synchrotron radio and X-ray as well as the IC γ-ray brightness are derived. They reveal the main factors which influence the pattern of the surface brightness distribution due to leptonic emission processes in shells of SNRs and can account for some non-linear effects of acceleration if necessary. These approximations provide observers and theorists with a set of simple diagnostic tools for quick analysis of the non-thermal maps of SNRs.

12214409

2011MNRAS.415.2553Z

1012

1012.3445_arXiv.txt

Since the pioneer works of \citet{einasto74}, \citet{bo78} and \citet{dressler}, stating that early-type systems tend to concentrate in high density regions, it has become clear that galaxy properties depend on the local environment. This dependence on environment must hold important information about the history of galaxy formation, so it is important to study the connection between the properties of the galaxies and their location in the Universe. In particular, a galaxy property that has been widely used in the literature is the luminosity, mainly through the analysis of the galaxy luminosity function (LF). This function describes the distribution of luminosities of a given population of galaxies and, in most cases, the shape of that distribution can be fully described by a function with two parameters \citep{schechter76}: the characteristic absolute magnitude $M^{\ast}$ and the faint end slope $\alpha$. Although these statistical measurements themselves do not give physical explanations about galaxy formation and evolution, they provide important constraints on various physical processes involved in such galaxy life stages (e.g. \citealt{benson03}). During the 1990's a lot of efforts have been made in order to compute the LF for galaxies in different environments such as the field, groups and clusters of galaxies (see for instance \citealt{mhg94,lin96,zucca97,lopcruz97,valotto97,ratclif98,mvl98,rauzy98,tren98}). However, the advent of large surveys of galaxies, such as the Sloan Digital Sky Survey \citep{sdss} and the Two degree Field Galaxy Redshift Survey (2dFGRS, \citealt{2df}), allowed for much better determinations of the LF \citep{blanton01,blantonlf,blanton05,norb02,madgwick,trentu02,m02,cz03,eke04}. There is a broad consensus that the LF of galaxies in the field is mainly flat ($\alpha \sim -1$), meanwhile a brighter characteristic magnitude $M^{\ast}$ and a steeper faint end slope $\alpha$ have been found in galaxy systems for absolute magnitudes in the range $M_r\la-16$. Other authors argue that a very steep faint end slope, at $M_r\ga-17$, can be measured in rich clusters and galaxy groups \citep{pop05,gonz05} when using only photometric information. It should be remarked that the methods that do not use spectroscopic redshifts are sensitive to the background computation, so, they are less reliable. \begin{figure} \includegraphics[width=90mm]{f1.eps} \caption{ Distributions of the properties SDSS DR7 groups. Redshift distribution ({\em upper left panel}), virial radii ({\em upper right panel}), line of sight velocity dispersion ({\em lower left panel}) and virial mass ({\em lower right panel}). Vertical dashed lines are the median of the distributions. } \label{fig1} \end{figure} Many works in the last few years have been concentrated on the effect of the environment on the LF. For instance, \citet{croton05,hoyle05,xia06,park07,phleps07} studied the dependence of the LF on the density contrast within spheres of different radii. A brightening of the characteristic magnitude and a steepening of the faint end slope are observed when moving from underdense to overdense regions. \cite{deng07} found that the dependence of galaxy luminosity on a dense environment is much weaker than that on an underdense environment. \cite{choi07} argued that the LF shows significant fluctuations due to large scale structures, while the morphological fraction as a function of luminosity is relatively less sensitive and thus seems to be more universal. The importance of the large scale environment was also established by \cite{yang09} and \cite{tempel09} showing strong environmental dependencies. One of the main benefits of working with large spectroscopic samples of galaxies is that they are very suitable for the construction of large galaxy group catalogues. \citet{zmm06} used the main galaxy sample of the SDSS DR4 to construct a large galaxy group catalogue, and they have deepen the study of galaxy luminosities as a function of environment. Their analysis comprised the variation of the Schechter function parameters, for different galaxy populations, as a function of the galaxy group virial mass. Their results showed clear variations of $M^{\ast}$ and $\alpha$ with group virial mass, and proved that these variations are mainly caused by the red population of galaxies in groups. More recently, \cite{robot10} studied the dependence of the LF of galaxies in 2dFGRS groups on the group virial mass and multiplicity. They also found clear trends for decreasing $\alpha$ when increasing the masses and/or multiplicity for early type galaxies, while a much suppressed relation was observed for late type population. At present, the largest galaxy redshift survey is the Seventh Data Release of the Sloan Digital Sky Survey (hereafter DR7; \citealt{dr7}). This catalogue covers a very wide area on the sky and has high quality photometric and spectroscopic information. From this galaxy catalogue we can extract one of the largest galaxy group catalogued to date. Therefore, the main aim of this work is two fold: firstly, improving the results obtained for the mass dependence galaxy LFs in the SDSS DR4 \citep{zmm06} by using the galaxies in groups in the SDSS DR7 in order to obtain more reliable and detailed results; and secondly, analysing galaxy LFs of different galaxy types using several criteria to classify them as well as study the influence of local and global environment on the LF. All this information is intended to provide a better understanding of galaxies in a wide range of density environments and, consequently, to clarify the scenario of galaxy evolution. The layout of this paper is as follows. In section 2 we describe the galaxy sample and the group identification process. The detailed analysis of the LFs is in section 3. {\bf Finally, in section 4 we discuss possible implications of our results and summarize them in section 5.}

In the last few years, the mass dependence of the LF of galaxies in groups has been analysed by several authors (e.g \citealt{zmm06,robot10}). Our findings for the overall LF are in agreement with the trends found in those previous works, but the use of a larger sample of groups allowed us an important statistical improvement. We have been able to disentangle the dependence of the LF on galaxy types, the location of galaxies within groups and the large scale environments of groups. Our results strongly suggest that the population of red spheroids is the most related to the group environment. Among the different galaxy types, only red spheroids have LF that strongly correlates with group mass. On the other hand, the late types (blue and red), have LFs that are independent of mass. The mass dependence of the LF of red spheroids and their increasing fraction with mass are the responsible for the mass dependence of the overall LF of galaxies in groups. There are several physical mechanisms proposed in the literature to explain the presence of red spheroids in groups, some of them are more effective in higher mass systems, which may explain the increasing fraction of red spheroids with group mass. For instance, in high mass groups, strangulation (e.g \citealt{larson80,balogh00}) and ram pressure (e.g. \citealt{gunn72,abadi}) can transform and quench star formation, which, in turn, affect primarily galaxy colour, while galaxy harassment transforms disks into spheroids (e.g. \citealt{farouki81,moore96}). Regardless the masses of the systems, galaxy major mergers (e.g. \citealt{toomre72,hopkins08}) combined with a mechanism to prevent subsequent star formation (such as AGN feedback, \citealt{bower06,croton06}) can be responsible for the formation of red spheroids in groups of any masses. Our work raises another interesting point. We find a clear indication of luminosity segregation, since galaxies in the inner parts have brighter characteristic magnitudes than their counterparts in the outer regions. This agrees with the well-known fact that the most luminous galaxies exist preferentially in the densest regions of the Universe (e.g \citealt{davis88,hamilton88,loveday95,benoist99,zehavi05,deng07}). Moreover, the luminosity segregation becomes more pronounced for massive systems as a consequence of a strong brightening of $M^{\ast}$ with the group mass for galaxies in the inner parts. These facts resemble the results of \citet{skibba07} in which, by testing the halo model predictions in groups of galaxies, they found that the luminosity of the central objects increases with halo mass, while non central galaxies show almost no mass dependence. Regarding the connection between the large scale group surrounding environment and the LF, high density regions show brighter and steeper values than the observed at low density regions, for both, $M_{\ast}$ and $\alpha$ respectively. A similar result was observed by \citet{tempel10} when analysing galaxy luminosities in the SDSS DR7 at different global density environments. They found that $M_{\ast}$ clearly becomes brighter from voids to superclusters, however, they do not observe any variation in $\alpha$. Furthermore, an interesting result arises in our work when analysing the variation of the LF parameters with group mass, since two distinct behaviours are observed: meanwhile galaxies in groups within high density regions have LF parameters that exhibit almost no changes with group mass, LF of galaxies in groups inhabiting low density regions experience a significant variation with mass. Our results suggest a plausible scenario for galaxy evolution in which both, large scale and local environments play important roles. There is evidence in the literature stating that groups in high density regions formed earlier \citep{harker06}. This gives the galaxy members a longer time to evolve, producing brighter bright galaxies ($M^{\ast}$), mostly through mergers, and also a larger number of faint galaxies ($\alpha$), affected by dynamical friction and with depleted gas reservoirs. For these groups, the effect of large scale environment could be the main responsible for galaxy evolution, with the group mass (local environment) playing a secondary role in the final result. On the other hand, for groups inhabiting low density regions, the group mass plays a more important role in the observed galaxy luminosities. For a fixed mass, it might be inferred that the difference in the formation time between groups in high and low density regions, is the key point to understand the differences in the LF, i.e., we are observing different stages of the same galaxy evolution. However, the formation time can not be entirely blamed for all differences, since groups in high density regions may have accreted a larger amount of material from their surroundings during their evolution, while groups in low density regions have had more limited access to fresh material. Thus, galaxy evolution in groups may follow different paths depending on where the groups inhabit.

1012.3445

We perform an exhaustive analysis of the luminosities of galaxies in groups identified in the Sloan Digital Sky Survey (SDSS) Data Release 7. Our main purpose is to perform a detailed study of the Schechter luminosity function parameters: the characteristic absolute magnitude and the faint-end slope, as a function of group virial mass in order to quantify their dependence on environment. We analyse the trends of the Schechter parameters as a function of group mass for different photometric bands, galaxy populations, galaxy positions within the groups and the group surrounding large-scale density. We find that the characteristic magnitude brightens and the faint-end slope becomes steeper as a function of mass in all SDSS photometric bands, in agreement with previous results. From the analysis of different galaxy populations, we observe that different methods to split galaxy populations, based on the concentration index or the colour-magnitude diagram, produce quite different behaviours in the luminosity trends, mainly for the faint-end slope. These discrepancies and the trends with mass mentioned above are explained when analysing the luminosity function of galaxies classified simultaneously according to their concentrations and colours. We find that only the red spheroids have a luminosity function that strongly depends on group mass. Late-type galaxies, whether blue or red, have luminosity functions that do not depend on group mass. The intrinsic change in the luminosity function of spheroids and the varying number contributions of the different types explain the overall trend of the luminosity function with group mass. On the other hand, dividing the galaxy members in the inner and outer regions of the groups do not introduce a significant difference in the Schechter parameter trends, except for the characteristic absolute magnitude in the high group virial mass range (?) which is an indication of luminosity segregation in massive groups. Finally, we also analyse the possible influence of the large-scale surrounding environment on the luminosity function of group galaxies. We find that galaxies inhabiting groups at low-density regions experience more pronounced variations on the Schechter parameters as a function of groups mass, while galaxies in groups at high-density regions show an almost constant behaviour. We discuss the possible implications of our findings in the galaxy evolution scenario.

12168197

2011ApJ...730L..15Y

1012

1012.1395_arXiv.txt

The Crab pulsar wind nebula is powered by its central pulsar born from a supernova explosion in 1054. It is a very luminous source in almost all wavelengths, from radio to the very high energy (VHE) $\gamma$-ray bands. The broadband non-thermal emission spectrum is well modeled by a synchrotron-inverse Compton (IC) scenario \citep{1996MNRAS.278..525A, 2010A&A...523A...2M} with the transition from the synchrotron to IC component occurring at a few hundred MeV \citep{2010ApJ...708.1254A}. The overall emission seems to be steady, so that the Crab nebula has been adopted as a standard candle in high energy astrophysics to calibrate observations from different instruments. However, detailed images of the Crab nebula in optical and X-ray bands indicated dynamical structures at small scales. Observations by the Hubble Space Telescope revealed wisps and knots in the nebula, and resolved some substructures of the highly variable wisps \citep{1995ApJ...448..240H}. X-ray observations by ROSAT and Chandra uncovered a jet-torus structure of the inner nebula \citep{1995ApJ...448..240H,2000ApJ...536L..81W}, and the equatorial ring is found moving outwards with a speed of $\sim 0.5c$ \citep{2002ApJ...577L..49H}. At higher energies in the $\gamma$-ray band, detailed structures of the nebula generally can not be resolved. However, COMPTEL and EGRET observations indicated that the synchrotron radiation in $1-150$ MeV is variable on a timescale of $\sim1$ yr \citep{1995A&A...299..435M,1996ApJ...457..253D}. In September 2010, AGILE collaboration reported a $\gamma$-ray flare above $100$ MeV from the direction of the Crab nebula, which lasted for about 3 days (ATel \#2855; Tavani et al. 2011). The flux during the flare period is about $2-3$ times higher than the average one. This flare was soon confirmed by the Fermi/LAT collaboration (ATel \#2861), and archive search of the Fermi/LAT data revealed another flare in February 2009 \citep{2010arXiv1011.3855F}. There were many other simultaneous or follow-up measurements for the flare in September 2010 in X-ray, optical, infrared and radio bands. However, no significant flux enhancement was discovered. In the VHE $\gamma$-ray energies, ARGO-YBJ collaboration claimed the detection of a flux enhancement around TeV, with a possibly longer duration (ATel \#2921). However, MAGIC and VERITAS observed the source with a shorter duration during the flare phase and did not reveal any enhancement in flux (ATel \#2967, \#2968). It was proposed that the flares was due to synchrotron emission of ultra-relativistic electrons with energies up to $\sim$PeV \citep{2010arXiv1011.3855F}. Considering the fact that the rest frame synchrotron radiation in a magnetic field dominated acceleration regime can not exceed $\sim 70$ MeV due to the fast synchrotron cooling of high energy electrons \citep{1970RvMP...42..237B, 2010MNRAS.405.1809L}, the observed GeV emission implies Doppler shift of the radiation region or other acceleration mechanisms instead of shock acceleration \citep{2010arXiv1011.3855F}. \cite{2010arXiv1011.1800K} proposed that the $\gamma$-ray variability (flare) originated from the ``inner knot'' of the Crab nebula (``knot 1'' as defined in \citet{1995ApJ...448..240H}), with mildly Doppler-boosted emission. The instability of the termination shock may cause the variability as revealed by magnetohydrodynamic (MHD) simulations. \cite{2010arXiv1011.4176B} suggested that electrons are accelerated in a region behind the shock, and the variability was attributed to changes in the maximum energy of accelerated electrons, electron spectral index or the magnetic field. In this Letter we employ a statistical approach to model the $\gamma$-ray variability of the nebula. Fermi/LAT observations have shown that the low-energy synchrotron component is variable on monthly time scale, while the high-energy IC component seems to be stable \citep{2010arXiv1011.3855F}. Since the synchrotron $\gamma$-rays are produced by the highest energy electrons, these observations indicate that fluctuations at the high-energy end of the electron distribution might be responsible for the variability. It is natural to expect that events that can generate the highest energy electrons are rarer, and therefore would lead to the largest fluctuation. Therefore the variability and flares in the sub-GeV $\gamma$-rays can be simply due to the statistical fluctuation of the highest energy electrons achievable in the electron accelerators. Lower energy electrons do not suffer from significant fluctuations since many more accelerators can contribute to them simultaneously, which gives rise to a ``steady-state'' emission in both the lower energy synchrotron component and the higher energy IC component.

In summary, we propose a statistical picture to explain the observed $\gamma$-ray variability and flares of the Crab nebula, using the fluctuations of the highest end of the electron spectra. The electrons are thought to be accelerated in a series of knots, with a size distribution $P(r_i)\propto r_i^{-\beta}$ and a distribution of the Doppler factor. The maximal energy of the electrons in the co-moving frame is assumed to be proportional to the size of the knots. Thus the rare knots with large sizes and high Doppler boosts may generate electrons up to $\sim$PeV, and hence, be responsible for the observed $\gamma$-rays flares at Fermi/LAT. On the other hand, the low energy electrons are generated by many smaller knots. The average effect smooths the fluctuation, so that the low energy synchrotron component and the IC component do not change significantly. This scenario can naturally explain both the variability in the MeV-GeV band, and the relatively steady emission in lower (optical, X-ray) and higher (TeV) energies. The expected variability of the monthly bin fluxes above 100 MeV are well consistent with that observed by Fermi/LAT. The two large $\gamma$-ray flares can also be naturally accounted for without additional assumptions. In the simulation, we do not consider the detailed structure (e.g. jet, torus) of the nebula. The knots are assumed to have a uniform and isotropic distribution. Considering a more realistic morphology of the nebula would not affect our conclusion noticeably, as long as the distributions of the knot sizes and Doppler factors are similar to those introduced here. In more general terms, these knots may be interpreted as individual electron acceleration events distributed throughout the nebula. We point out that at low energies contributions from large amount of knots are equivalent with previous studies that electrons are accelerated continuously inside the whole nebula. Although the model does not discuss the absolute time scales of flares and $\gamma$-ray fluctuations from the first principle, they may be related to timescales of developing MHD instabilities at different spatial scales.

1012.1395

A statistical scenario is proposed to explain the γ-ray variability and flares of the Crab Nebula, which were observed recently by the Fermi/LAT. In this scenario electrons are accelerated in a series of knots, whose sizes follow a power-law distribution. These knots presumably move outward from the pulsar and have a distribution in the Doppler boost factor. The maximal electron energy is assumed to be proportional to the size of the knot. Fluctuations at the highest energy end of the overall electron distribution will result in variable γ-ray emission via the synchrotron process in the ~100 MeV range. Since highly boosted larger knots are rarer than smaller knots, the model predicts that the variability of the synchrotron emission increases with the photon energy. We realize such a scenario with a Monte Carlo simulation and find that the model can reproduce both the two γ-ray flares over a period of ~1 year and the monthly scale γ-ray flux fluctuations as observed by the Fermi/LAT. The observed γ-ray spectra in both the steady and flaring states are also well reproduced.

12167803

2011ApJ...728..112W

1012

1012.3503_arXiv.txt

} Intense magnetic fields around sunspot regions drive the most energetic examples of solar activity -- large solar flares and coronal mass ejections -- which can produce hazardous space weather conditions close to Earth. Severe space weather storms present a variety of hazards including increased radiation levels for space travellers and crew on polar air flights, the risks of disabling of GPS systems and damage to power grids \citep{2008sswe.rept.....C}, and potential large economic losses due to damage to communications satellites (Odenwald et al.\ 2006). These effects motivate a need to understand and model the source regions of space weather threats at the Sun. Spectro-polarimetric measurements of magnetically sensitive photospheric lines may be used to infer the vector magnetic field at the surface of the Sun \citep{LanLan04}, and the resulting maps of the field are called vector magnetograms. A new generation of instrumentation is set to provide vector magnetogram data over the next solar cycle with unprecedented quality, resolution, and temporal cadence, including the Solar Optical Telescope Spectro-Polarimeter (SOT/SP) on the Hinode spacecraft (Tsuneta et al.\ 2008) and the Helioseismic and Magnetic Imager on the Solar Dynamics Observatory (SDO/HMI) \citep{2006cosp...36.1469S,2007SoPh..240..177B}. In principle these measurements provide boundary values for coronal magnetic field modeling, or `magnetic field reconstruction'. There is a pressing need for this modeling capability, which has many potential applications in solar research. However, studies have revealed basic difficulties preventing the construction of reliable coronal field models from vector magnetograms using present modeling techniques. Specifically, a popular and well-developed approach to magnetic field reconstruction, nonlinear force-free modeling, has been shown to produce inconsistent results when applied to Hinode/SOT vector magnetogram data \citep{2008ApJ...675.1637S,2009ApJ...696.1780D,WheEtAl2010}. Different modeling codes produce different results based on the same data, and the results from individual codes are also internally inconsistent. As described in more detail in section~\ref{sec2}, vector magnetograms allow for two sets of boundary conditions for the force-free model, and the two solutions do not in general agree. It is also difficult to construct the solutions: the boundary conditions on electric current present large currents which prevent the force-free methods from converging strictly. As a result the `solutions' may be inaccurate with respect to the model. The modeling has been shown to be unable to provide a reliable estimate of the free magnetic energy of a solar active region. Further details of the problems are provided in section~\ref{sec2}. One likely cause of the failure of nonlinear force-free modeling is a basic inconsistency between the boundary data and the model: the model includes only magnetic forces, an assumption that may be warranted in the corona, but which is thought not to apply at the photospheric level where the data originate \citep{1995ApJ...439..474M}. Recently a new approach to coronal magnetic field reconstruction was presented to address this problem \citep{2009ApJ...700L..88W}. The `self-consistency' procedure (described in more detail in section~\ref{sec2}) identifies a nonlinear force-free solution with boundary values close to those implied by the vector magnetogram data. The method constructs a single solution to the nonlinear force-free model which strictly converges. The solution does not match exactly the field at the photospheric level, but may describe the coronal magnetic field. In \citet{2009ApJ...700L..88W} the method was demonstrated in application to Hinode/SOT data for NOAA active region 10953. A solution to the model was constructed with energy $E/E_0\approx 1.02$, where $E_0$ is the energy of the reference potential (current-free) magnetic field matching the boundary conditions on the vertical component of the field. The self-consistency procedure is capable of taking into account relative uncertainties in the solar boundary data, so that the boundary values for the solution match the vector magnetogram data more closely at points with small uncertainties. However, uncertainties were omitted from the \citet{2009ApJ...700L..88W} calculations, and this omission accounts in part for the low value of the magnetic free energy of the constructed solution ($2\%$ of the energy of the potential field), as explained in section~\ref{sec24}. For this reason the results in \citet{2009ApJ...700L..88W} were argued to represent a `proof in concept' of the self-consistency procedure, rather than a full demonstration of its modeling capability. In this paper we return to the problem of modeling AR 10953, including uncertainties in the calculation. The detailed presentation of this paper is as follows. Section~\ref{sec2} provides a more detailed account of the background to the modeling presented here, incuding a summary of the nonlinear force-free model (section~\ref{sec21}), a description of the failure of the model in application to solar data (section~\ref{sec22}), and an account of the self-consistency procedure (section~\ref{sec23}) and its initial application to AR 10953 by \citet{2009ApJ...700L..88W} (section~\ref{sec24}). A reader familiar with the background may skim or omit section~\ref{sec2}, but it provides necessary background for a reader new to the topic. Section~\ref{sec3} describes the modeling of AR 10953 taking into account uncertainties, with section~\ref{sec31} presenting the data, and section~\ref{sec32} the results of the modeling. Section~\ref{sec4} discusses the results.

} A solution to the nonlinear force-free model is constructed for NOAA solar active region AR 10953, based on a vector magnetogram derived from Hinode/SOT observations at 22:30~UT on 30 April 2007. The solution applies the `self-consistency' procedure of \citet{2009ApJ...700L..88W}, for the first time taking into account uncertainties in the boundary data. The boundary data are subject to improved techniques for data merging by comparison with earlier studies using the same Hinode/SOT observations \citep{2009ApJ...696.1780D,2009ApJ...700L..88W}. The self-consistency procedure addresses the problem that vector magnetogram data are inconsistent with the force-free model \citep{2009ApJ...696.1780D,WheEtAl2010}. In particular, the solar boundary data provide two possible force-free solutions (the $P$ and $N$ solutions, corresponding to choosing boundary conditions for the vertical electric current density $J_z$ from the positive or the negative polarities of the boundary field respectively), which are generally found to be significantly different. The $P$ and the $N$ soutions are distinctly different for the vector magnetogram data used here. The self-consistency procedure modifies the boundary conditions on $J_z$ during a sequence of cycles in which $P$ and $N$ solutions are constructed using Grad-Rubin iteration \citep{2007SoPh..245..251W}, and arrives at a solution to the force-free model with boundary values close to, but not exactly matching, the observational vector magnetogram data. When uncertainties in the boundary data are taken into account, the boundary values of $J_z$ are preserved more closely at points having smaller uncertainties. Self-consistency modeling has two advantages over conventional force-free modeling. First, force-free methods tend to fail to strictly converge when applied to solar data, and hence do not accurately solve the force-free model. The self-consistency procedure achieves strict convergence. Second, the method identifies a single solution, rather than two solutions. Active region AR 10953 was previously modelled using the self-consistency procedure in \citet{2009ApJ...700L..88W}, but in that case uncertainties were not included, and the results were regarded as a proof of concept of the method, rather than a full demonstration of its capabilities. This paper presents a more comprehensive demonstration and investigation. The self-consistent solution obtained for active region AR 10953 is substantially non-potential, with a coronal magnetic energy $E/E_0\approx 1.08$, where $E_0$ is the energy of the potential field with the same vertical component of the magnetic field in the boundary. This energy is significantly larger than the energy $E/E_0\approx 1.02$ of the solution obtained in \citet{2009ApJ...700L..88W}, which is attributed to the neglect of uncertainties in the earlier paper. The non-potentiality of the region is due to strong currents in the inner parts of the region, in particular associated with the southern part of the the negative polarity leading spot, as shown in Figure~3. These currents were reduced in the earlier solution by the self-consistency averaging procedure [Equations~(\ref{eq:bayes_avg})] due to the omission of uncertainties (or rather, the treatment of all boundary points as having equal uncertainty). With uncertainties included correctly, these currents are preserved because they are in regions where the field, and hence current, is well-determined, as shown in Figure~\ref{fig:VaryGRIT}. The energy of the new self-consistent solution is intermediate between the energies of the $P$ and $N$ solutions constructed from the vector magnetogram data ($E/E_0\approx 1.03$ for the $P$ solution, and $E/E_0\approx 1.17$ for the $N$ solution), and is also in the middle of the range of energies found by \citet{2009ApJ...696.1780D} using force-free codes applied to boundary data derived from the same Hinode/SOT observations but subject to the preprocessing procedure \citep{2006SoPh..233..215W,2008SoPh..247..249W}, which is not used here. Recently \citet{2010ApJ...715.1566C} presented a more detailed account of force-free modeling of AR 10953 with two Grad-Rubin codes applied to the $N$-polarity \citet{2009ApJ...696.1780D} boundary data, and reported energies $E/E_0\approx 1.27$ and $E/E_0\approx 1.31$. The energies obtained by these authors are different from the energy of the initial $N$ solution obtained here for a number of reasons. The Hinode/MDI boundary data used here is prepared differently, as discussed in section~\ref{sec31}, and preprocessing has not been applied. The lack of convergence of the initial $P$ and $N$ solutions obtained here means that a range of energies is possible, depending on the choice of stoppping iteration (in practise we find energies for the $N$ solution in the range $E/E_0\approx 1.16$ to $E/E_0\approx 1.22$). The present method treats the side and top boundaries in a different way to the \citet{2010ApJ...715.1566C} method, and in particular we neglect currents on field lines which cross the side and top boundaries, which tends to reduce the energy. Finally, the stated \citet{2010ApJ...715.1566C} and \citet{2009ApJ...696.1780D} energies refer to a smaller sub-region of the computational domain. Because of these differences we do not attempt a more detailed comparison with the earlier results. A degree of arbitrariness is introduced into the modeling by the lack of strict convergence of the Grad-Rubin iterations used to construct the $P$ and $N$ solutions at early and intermediate self-consistency cycles. The results of the self-consistency procedure may depend on the choice of the number of Grad-Rubin iterations. This effect is investigated by repeating the calculation with $N_{\rm GR}=20$ and $N_{\rm GR}=40$ iterations at each cycle (rather than $N_{\rm GR}=30$). The two new self-consistent solutions are remarkably similar to the first solution, and in particular all three solutions have energy $E/E_0\approx 1.08$. This suggests that self-consistency modeling provides a solution to the problem of reliable estimation of the coronal magnetic free energy of an active region, once uncertainties in the boundary data are incorporated. The inclusion of uncertainties preserves boundary conditions on current at points in the boundary where the currents are most accurately determined. The three self-consistent solutions obtained with different numbers of Grad-Rubin iterations have very similar boundary distributions of $J_z$. The self-consistency procedure (including uncertainties) is a promising candidate for routine reconstruction of coronal magnetic fields. The technique is shown to produce a highly non-potential and accurately force-free coronal field model for NOAA AR 10953 from Hinode/SOT vector magnetogram data. The estimate of the free energy of the magnetic field appears to be robust based on the results obtained with varying numbers of Grad-Rubin iterations per self-consistency cycle. The availability of high-quality vector magnetogram data over the next solar cycle from Hinode/SOT and SDO/HMI will provide ample opportunities for further testing, development, and application of the method, to investigate many questions in the physics of solar activity.

1012.3503

A nonlinear force-free solution is constructed for the coronal magnetic field in NOAA solar active region (AR) 10953 based on a photospheric vector magnetogram derived from Hinode satellite observations on 2007 April 30, taking into account uncertainties in the boundary data and using improved methods for merging multiple-instrument data. The solution demonstrates the "self-consistency" procedure of Wheatland &amp; Régnier, for the first time including uncertainties. The self-consistency procedure addresses the problem that photospheric vector magnetogram data are inconsistent with the force-free model, and in particular that the boundary conditions on vertical electric current density are overspecified and permit the construction of two different nonlinear force-free solutions. The procedure modifies the boundary conditions on current density during a sequence of cycles until the two nonlinear force-free solutions agree. It hence constructs an accurate single solution to the force-free model, with boundary values close, but not matched exactly, to the vector magnetogram data. The inclusion of uncertainties preserves the boundary conditions more closely at points with smaller uncertainties. The self-consistent solution obtained for AR 10953 is significantly non-potential, with magnetic energy E/E <SUB>0</SUB> ≈ 1.08, where E <SUB>0</SUB> is the energy of the reference potential (current-free) magnetic field. The self-consistent solution is shown to be robust against changes in the details of the construction of the two force-free models at each cycle. This suggests that reliable nonlinear force-free modeling of ARs is possible if uncertainties in vector magnetogram boundary data are included.

12205598

2011IJBC...21.2211A

1012

1012.4609_arXiv.txt

The three body problem (TBP), consisting of three bodies with masses $m_0$, $m_1$ and $m_2$ that interact with gravitational forces, has been widely studied in the literature. It is considered either as a simple dynamical system with complex behaviour or as a model for explaining the evolution of planets and other celestial bodies. In this work, we consider the planar case, where all bodies move on the same plane and the indices 0, 1 and 2 refer to the particular body. The simplest, yet not trivial, version of TBP is the {\em circular restricted TBP} (CRTBP), where we assume that one of the bodies is massless (say $m_2$=0) and the two massive bodies ($m_0\neq 0$, $m_1\neq 0$), called {\em primaries}, revolve in circular orbits on the plane. When the primaries revolve in elliptic orbits of eccentricity $e_p$ (=$e_0=e_1)$ and period $T_p$ (=$T_0=T_1$), we have the {\em elliptic restricted TBP} (ERTBP). When the massless body is replaced by a massive one, we get the general planar model of the TBP (GTBP). Especially, when one body has very large mass in comparison with the other two bodies (say $m_0\gg m_1$, $m_0\gg m_2$), we get the so-called ``planetary'' GTBP. Continuation and existence of periodic orbits in this problem has been studied many years ago by Hadjidemetriou (1975) and Bozis and Hadjidemetriou (1976) and recently their results found a fruitful field of applicability in the dynamics of resonant extrasolar systems (e.g. Psychoyos and Hadjidemetriou, 2005; Haghighipour \emph{et al.}, 2003; Ferraz-Mello \emph{et al.}, 2005). In most cases, periodic orbits are associated with resonant planetary motion and their determination in the general TBP can be based on the unperturbed circular problem (i.e. both planets are massless bodies). However, not all solutions can be constructed in this way. In Voyatzis \emph{et al.} (2009), it is shown that the main set of families of periodic orbits in the 2:1 resonance can be constructed by the continuation of periodic orbits from the restricted problem to the general one. In this scheme, the role of families of periodic orbits of the restricted TBP is crucial. The origin of periodic orbits in the {\em circular restricted TBP} can be assigned to periodic orbits of the unperturbed problem. The structure of families of {\em symmetric} periodic orbits is well-known (see e.g. Bruno, 1994; H\'enon, 1997). Previous studies (e.g. Beaug\'e, 1994; Voyatzis \emph{et al.}, 2005)) indicate that {\em asymmetric} periodic orbits exist only in resonances of the form $\frac{T_2}{T_1}=\frac{1}{q}$, $q\in \mathbb{N}$, called {\em asymmetric}, where for $q>1$ the massive body ($m_1\neq 0$) is the inner planet and the massless body ($m_2=0$) is the outer one. In the {\em elliptic restricted TBP}, families of periodic orbits bifurcating from periodic orbits of the circular problem also exist having period $T=\frac{k}{\ell} T_p$, where $k$ and $\ell$ are prime integers (Broucke, 1969). Many periodic orbits for the elliptic model associated to the dynamics of asteroids and Kuiper belt objects have been computed (e.g. Hadjidemetriou, 1999; Voyatzis \emph{et al.}, 2005). All the periodic orbits found were symmetric. However, Voyatzis and Kotoulas (2005) found many bifurcation points along the families of periodic orbits and conjectured the generation of families of asymmetric orbits. In this paper, we compute and show the existence of such families in the elliptic restricted model and examine their continuation to the general model. In the following section, we present our model and some basic notions on the description, existence and continuation of periodic orbits in the particular system. In section 3, we show how families of asymmetric periodic orbits can appear in the elliptic model and in section 4, how such families are continued to the general problem. Finally, in section 5, we conclude our results.

We have studied some new types of bifurcation of periodic orbits in the planar three body problem consisting of a heavy body (star) and two bodies of small or negligible mass (planets). When one planet has zero mass (massless body) but the second one is massive (primary body), we have the simplest models of the circular (CRTBP) or the elliptic (ERTBP) restricted three body problem. Considering these models in an appropriate rotating frame, we can compute resonant periodic orbits and then perform continuation to the general problem where both planets are massive. It is known that periodic orbits in the ERTBP are generated from periodic orbits of the CRTBP and are continued parametrically, by varying the eccentricity of the primary body, $e_p$. These orbits are generally symmetric. In this paper, we have shown the existence of families of asymmetric periodic orbits in the ERTBP. There are two types of bifurcations of such orbits, which we called as type I and II bifurcations. In a type II bifurcation, the asymmetric orbits of the ERTBP are generated from an asymmetric periodic orbit of the CRTBP. It is known, that such periodic orbits in the CRTBP exist only in the resonances of the form $1/q$ and therefore such families of periodic orbits exist for the RTBP only in these resonances. However, in a bifurcation of type I, the asymmetric families are generated from symmetric periodic orbits of the ERTBP. Actually, these bifurcation points are critical orbits with respect to their linear stability and can be met in any resonance. We presented such families for the 1:2 and 1:3 resonance and also for the resonance 3:2, which is not of the form $1/q$. The families of the CRTBP or ERTBP are continued by varying the planetary masses. So, they become families of the general planar three body problem. Their continuation with respect to the planetary masses can follow two different schemes. The continuation scheme I holds for the families related to type I bifurcation in the ERTBP. Namely, these families continue smoothly by varying the mass preserving their topological structure in phase space for adequately small planetary masses. In the continuation scheme II, families of both the CRTBP and ERTBP models involve in order to generate a family in the GTBP. Following the generation and the continuation of the above mentioned families we complete the net of families of periodic orbits in the planar general three body model and we explain their origin. Periodic orbits of the TBP correspond to the so-called "exact resonances" in the planetary dynamics and are important since an exosolar planetary system can be trapped in such orbits after a physical dissipation process.

1012.4609

We consider the planar three body problem of planetary type and we study the generation and continuation of periodic orbits and mainly of asymmetric periodic orbits. Asymmetric orbits exist in the restricted circular three body problem only in particular resonances called "asymmetric resonances". However, numerical studies showed that in the general three body problem asymmetric orbits may exist not only for asymmetric resonances, but for other kinds, too. In this work, we show the existence of asymmetric periodic orbits in the elliptic restricted problem. These orbits are continued and clarify the origin of many asymmetric periodic orbits in the general problem. Also, we illustrate how the families of periodic orbits of the restricted circular problem and those of the elliptic one join smoothly and form families in the general problem, verifying in this way the scenario described firstly by Bozis and Hadjidemetriou (1976).

12165863

2011AN....332..128C

1012

1012.1038_arXiv.txt

\sloppy The photospheric solar abundances of some elements have been studied more often than others; this is mainly due to the importance of the elements to explain nucleosynthesic processes. Moreover, the difficulty of extracting suitable lines with accurately known transition probabilities in the visible range of the solar spectrum can also explain why some elements have been studied less extensively. For the triad Sr-Y-Zr, there exist only a few detailed studies of their solar abundances (an ADS\footnote{The Astrophysics Data System, ads.harvard.edu} search yields 4 papers for Sr, 5 for Y, and 10 for Zr) in spite of their importance in Galactic chemical evolution. Sr-Y-Zr, Ba-La, and Pb are located at three abundance peaks of the s-processes producing the enrichment of these elements in the Galaxy. Knowledge of the present Sr-Y-Zr abundances in stars of different metallicity and age is required to understand the complicated nucleosynthesis of these elements (for details see \citealt{travaglio04}). Zirconium is present in the solar spectrum with lines of \ion{Zr}{i} and \ion{Zr}{ii}. The dominant species is \ion{Zr}{ii}. According to our 3D model, and in agreement with our 1D reference model, about 99\% of zirconium is singly ionised in the solar photosphere. Departure from local thermodynamical equilibrium (LTE) probably affects \ion{Zr}{i} lines, producing a too low Zr abundance under the assumption of LTE. In fact NLTE effects, even on weak unsaturated \ion{Zr}{i} lines, have been found by \citet{brown83} for G and K giants. However, both \citet{biemont81} and \citet{bogdanovich96} analysed in LTE \ion{Zr}{i} lines (34 and 21 lines, respectively) and \ion{Zr}{ii} lines (24 and 15 lines, respectively) in the solar photosphere, and found an excellent agreement between the abundances derived from both ionisation stages. Very recently \citet{velichko10} performed a NLTE analysis of zirconium in the case of the Sun and late type stars. According to their computations the NLTE abundance is larger than the LTE one, by up to 0.03\,dex for \ion{Zr}{ii} and by 0.29\,dex for \ion{Zr}{i}. Owing to the small number of \ion{Zr}{i} and \ion{Zr}{ii} lines in metal poor stars \citep{gratton94}, it is important to analyse both of them in the solar photosphere to derive the solar abundance from both ionisation stages and to assess the agreement between the derived abundances. In spectra of cool stars it is easier to observe \ion{Zr}{i} lines. For example, \citet{goswami10} realised that none of the \ion{Zr}{ii} lines were usable in their analysis of the cool Pop.\,II CH star HD~209621, and the Zr abundance is derived from the only \ion{Zr}{i} line in their spectrum at 613.457\,nm. A similar situation had been encountered by \citet{vanture02} in their study of the Zr/Ti abundance ratio in cool S stars, where only \ion{Zr}{i} lines in the red part of the spectrum can be used.

We analysed the zirconium abundance in the solar photosphere, investigating a selected sample of lines of both \ion{Zr}{i} and \ion{Zr}{ii}. The \ion{Zr}{i} lines are weak and heavily blended so that only four of them are acceptable for abundance work. However, the present analysis of the zirconium abundance relies primarily on 15 lines of \ion{Zr}{ii} that we found suitable for this purpose. We have applied three different fitting strategies to derive the abundance from Zr lines that are blended by lines of other elements. In the case that all components making up the blend are weak, the different methods give consistent abundances. If, however, stronger lines are involved (for an example see Sect.\,\ref{szrii4496}), the methods that ignore saturation effects may severely underestimate the Zr abundance by up to $0.05$\,dex, even though the result of the fitting may appear pleasing to the eye, and the reduced $\chi^2$ may be close to one. We find a good agreement between A(Zr) derived from the \ion{Zr}{i} and the \ion{Zr}{ii} lines, but, due to the scarcity of \ion{Zr}{i} lines, we consider this result as fortuitous. The abundance from disc-centre spectra is systematically higher than the one from disc-integrated observations. A similar result is also found for other heavy elements like Fe, Th, Hf, both with 3D and 1D model atmospheres. This may indicate a problem with the thermal structure of the models, rather than a physical abundance gradient in the solar atmosphere. Further investigations are necessary to find an explanation for this small discrepancy. Our recommended solar zirconium abundance is based on the 3D result for 15 \ion{Zr}{ii} lines, and is A(Zr)$=2.62\pm 0.06$, where the uncertainty is the line-to-line scatter of the selected sample of \ion{Zr}{ii} lines. This value is at the upper end of the solar zirconium abundances found in the literature. Still, within the mutual error bars, this result is in good agreement with the meteoritic zirconium abundance of A(Zr)=$2.57\pm 0.04$ \citep{lodders09}.

1012.1038

Zirconium (Zr), together with strontium and yttrium, is an important element in the understanding of the Galactic nucleosynthesis. In fact, the triad Sr-Y-Zr constitutes the first peak of s-process elements. Despite its general relevance not many studies of the solar abundance of Zr were conducted. We derive the zirconium abundance in the solar photosphere with the same CO<SUP>5</SUP>BOLD hydrodynamical model of the solar atmosphere that we previously used to investigate the abundances of C-N-O. We review the zirconium lines available in the observed solar spectra and select a sample of lines to determine the zirconium abundance, considering lines of neutral and singly ionised zirconium. We apply different line profile fitting strategies for a reliable analysis of Zr lines that are blended by lines of other elements. The abundance obtained from lines of neutral zirconium is very uncertain because these lines are commonly blended and weak in the solar spectrum. However, we believe that some lines of ionised zirconium are reliable abundance indicators. Restricting the set to Zr II lines, from the CO<SUP>5</SUP>BOLD 3D model atmosphere we derive A(Zr) {=2.62± 0.06}, where the quoted error is the RMS line-to-line scatter.

12226230

2011PhRvD..84d3507F

1012

1012.3667_arXiv.txt

Observations of the cosmic microwave background (CMB) radiation have ushered in a new era of precision cosmology. Full-sky temperature maps produced by the {\em Wilkinson Microwave Anisotropy Probe} (WMAP)~\cite{Bennett:2003ba} have confirmed with high precision that the observed temperature fluctuations are consistent with a nearly Gaussian and scale invariant primordial power spectrum, as predicted by inflation. The recently launched {\em Planck} satellite~\cite{Tauber2010} has a resolution three times better than that of WMAP, with an order of magnitude greater sensitivity, and significantly wider frequency coverage (allowing for far more robust foreground removal, and therefore reduced systematics). These high quality data sets allow for the possibility of observing deviations from the standard inflationary paradigm, some of which could have drastic consequences for our understanding of the universe and its origins. Perhaps the largest gap in our description of the early universe lies in an understanding of its initial conditions. One possibility, motivated by the proliferation of vacua in compactifications of string theory (known as the string theory landscape~\cite{Susskind:2003kw}), is that our observable universe is only a tiny piece of a vast multiverse, the majority of which is still inflating. This picture of {\em eternal inflation} (for a review, see, e.g., Ref.~\cite{Aguirre:2007gy}) arises when the rate at which local regions exit an inflating phase is outpaced by the accelerated expansion of the inflating background. Eternal inflation is a fairly generic consequence of any theory containing positive vacuum energy and multiple vacua, highlighting the importance of understanding how this scenario might be confronted with observational tests. The first attempts to embed our cosmology inside an eternally inflating universe led to ``open inflation"\cite{Bucher:1994gb,Gott:1982zf}; see Ref.~\cite{GarciaBellido:1997uh} for a review. In this scenario, a scalar field (or set of scalar fields) has a potential with a high energy metastable minimum that drives the eternally inflating phase. Transitions out of this vacuum proceed via the Coleman-de Luccia (CDL) instanton~\cite{Coleman:1977py,Coleman:1980aw}, resulting in expanding bubbles inside which the scalar field rests on an inflationary plateau. The symmetries of the CDL instanton ensure that there is a very nearly homogeneous and isotropic open universe inside the bubble; inflation, reheating, and standard cosmological evolution follow. In any given bubble, the future light cone of the nucleation event forms the ``Big Bang'' (where the scale factor vanishes) of an open FRW universe. The eternally inflating phase outside our bubble can therefore be thought of as a pre-Big Bang epoch, and one might expect inflation to erase any of the scant observational evidence of our parent vacuum. In single bubble open inflation, various anomalies are induced in the CMB temperature power spectrum (see Ref.~\cite{GarciaBellido:1997uh} and references therein), but unfortunately, the size of these effects decreases with the present energy density in curvature (related to the number of inflationary $e$-folds), rendering them negligible at all but the lowest multipoles where cosmic variance dominates. However, our bubble does not evolve in isolation. There are other nucleation events from the false vacuum, containing a phase that might be identical to ours, or perhaps very different. If one of these secondary nucleation events occurs close enough to our bubble wall, then a collision inevitably results. In fact, since our bubble grows to reach infinite size, there are an infinite number of collisions~\cite{Guth:1981uk,Guth:1982pn,Gott:1984ps,Garriga:2006hw} (a finite subset of which are causally accessible to any one observer). This raises the possibility that if such collisions are both survivable and only small perturbations on top of standard cosmology, they might leave observable signatures of eternal inflation~\cite{Aguirre:2007an}; it is these signatures which our analysis targets. If we are to detect such bubble collisions, their predicted signatures must be consistent with our observed cosmology, but sufficiently distinct to be differentiated from other possible signals in the CMB. In addition, the theory must predict that we expect to have causal access to bubble collisions. While these criteria are not met in every model of eternal inflation, recent work~\cite{Guth:1981uk,Guth:1982pn,Hawking:1982ga,Wu:1984eda,Gott:1984ps,Garriga:2006hw,Aguirre:2007an,Aguirre:2007wm,Aguirre:2008wy,Chang:2007eq,Chang:2008gj,Dahlen:2008rd,Freivogel:2009it,Easther:2009ft,Larjo:2009mt,Zhang:2010qg,Czech:2010rg} (for a review, see Ref.~\cite{Aguirre:2009ug}) has established that bubble collisions could in some theories be both expected and detectable. Bubble collisions produce a fairly characteristic set of inhomogeneities in the very early universe, which are processed into temperature anisotropies in the CMB. From the spherical symmetry of the colliding bubbles, the collision possess azimuthal symmetry, and by causality must be confined to a disc on the sky. The CMB temperature and its derivatives need not be constant across the causal boundary. Therefore, the signals we are searching for are localized, and because they are primordial, consist of a long-wavelength modulation of the standard inflationary density fluctuations inside the affected region~\cite{Chang:2008gj}. The amplitude and angular scale of the signal is dependent upon the underlying model and kinematics of the collision. These general features suggest a set of strategies for data analysis. The localization of the collision implies that wavelet analysis could be a sensitive tool for picking out both the location and angular scale of a candidate signal. The causal boundary, across which the temperature and its derivatives need not be constant, suggests the use of edge detection algorithms similar to those used in searches for cosmic strings \cite{Kaiser:1984iv,Lo:2005xt,Danos:2008fq,Amsel:2007ki}. Finally, the prediction that the temperature modulation induced by the collision is rather long-wavelength yields a sufficiently generic template to perform a full Bayesian parameter estimation and model selection analysis. In this paper, we describe a modular analysis algorithm designed to look for the signatures of eternal inflation, and apply it to the WMAP 7-year data~\cite{Jarosik:2010iu}. This algorithm can easily be adapted to test any model that predicts a population of spatially localized sources in addition to the standard fluctuations predicted by $\Lambda$CDM. A summary of our results was presented in Ref.~\cite{Feeney:2010jj}; in this paper we describe our analysis in detail. Currently available full sky CMB data are rather limited in their sensitivity to the signatures of bubble collisions listed above; the main current limitation is the low resolution. Therefore, we apply our algorithm to current data mainly as a validation exercise; to exploit its full power would require future high resolution data, e.g., from {\em Planck}. The individual steps of our analysis pipeline are calibrated using realistic simulations of the WMAP experiment with and without bubble collisions. The calibrated pipeline applied to data is fully automated, identifying the candidate signals and processing them without any human intervention. This removes any {\em a posteriori} choices from our analysis, which must be carefully avoided in any analysis of a large data-set such as the WMAP 7-year data~\cite{Bennett:2010jb}. The plan of the paper is as follows. In Sec.~\ref{sec:collisionintro}, we review some of the background on bubble collisions in eternal inflation, and outline the predicted observable signatures. Our analysis pipeline is summarized in Sec.~\ref{sec:summaryofpipeline}. We describe some properties of the WMAP experiment and our simulations in Sec.~\ref{sec:simulatedmaps}, and detail our analysis tools in Sec.~\ref{sec:analysistools}. Sec.~\ref{sec:WMAP7} summarizes the results of our analysis of the WMAP 7-year data, and we conclude in Sec.~\ref{sec:conclusions}.

\label{sec:conclusions} An exciting opportunity to confront the eternal inflation scenario with experiment lies in the observation of collisions between other bubble universes and our own. In this paper, we have described an algorithm to search for the imprint of bubble collisions on the cosmic microwave background, and applied it to the WMAP 7-year data. Our search algorithm targets the generic signatures expected from bubble collisions: azimuthal symmetry, long-wavelength modulation of the temperature confined to discs on the sky, and circular temperature discontinuities. For this reason, we expect our analysis to be fairly robust under changing assumptions about the underlying theory, which is presently rather poorly understood. The analysis pipeline we have developed takes a two-pronged approach, applied in parallel. The first uses heuristic techniques to test for the presence of features specific to bubble collisions. The second is a fully Bayesian algorithm for the general problem of non-Gaussian source detection, implemented as a patch-wise approximation to the full-sky model selection and parameter estimation problem. The data set is segmented in a completely automated way, allowing us to avoid {\em a posteriori} selection effects associated with choosing the most ``interesting" features on the CMB sky by hand. The algorithm is tested and thresholds at each step are calibrated using extensive simulations, and then frozen before ever looking at the data, to follow as much as possible the philosophy of a blind analysis. Candidate collisions are identified from an input temperature map based on the response to a suite of needlet transforms (calibrated using simulations with and without bubble collisions), and grouped into ``blobs." These blobs are scrutinized for circular temperature discontinuities using an edge detection algorithm. The quantitative significance of an edge is characterized using the Circular Hough Transform (CHT). The blobs are also used to construct an approximation to the full-sky Bayesian parameter estimation and model selection problem for bubble collisions. The posterior probability distribution over the expectation value for the number of detectable collisions, $\nsavge$, is then obtained. This allows us to quantify which of the two models -- a theory which predicts on average $\nsavge$ bubble collision signatures described by temperature modulations of the form given in Eq.~\ref{eq:collfluct}, or else the standard model (specified by $\nsavge=0$) with CMB plus realistic noise and beam effects -- better explains the data. Applying our analysis pipeline to simulations, we have found that a circular temperature discontinuity at the causal boundary is a clear signature of bubble collisions.\footnote{The observational detection of a circular temperature discontinuity is so unlikely to arise spuriously that it provides conclusive evidence of a detection.} Although our analysis can identify collisions without temperature discontinuities, their presence greatly increases our ability to make a conclusive detection. Both the edge-detection and Bayesian model selection steps have the ability to identify a causal boundary in the patches of the sky that are highlighted as candidate collisions by the blob detection step of our analysis pipeline. We have found no evidence for circular temperature discontinuities in the WMAP 7-year data using either method. Based on our analysis of simulations, this allows us to rule out the presence of collisions in the exclusion region of Fig.~\ref{fig-chtexclusion}. For collisions larger than $\theta_{\rm crit} \agt 10^{\circ}$, this corresponds to $10^5 |z_{\rm crit}| \alt 3$--$6$ for the amplitude of the circular temperature discontinuity defined in Eq.~\ref{eq:collfluct}. For collisions on smaller scales, the CHT step loses sensitivity due to the proliferation of degree-scale blobs in the background CMB fluctuations. The posterior evaluated using the WMAP 7-year data is maximized at $\nsavge=0$, and constrains $\nsavge < 1.6$ at $68 \%$ confidence. We therefore conclude that this data set does not favor the bubble collision hypothesis for any value of $\nsavge$. In light of this null detection, comparing with the simulated bubble collisions, we can constrain the central amplitude of the temperature modulation caused by the collision (defined in Eq.~\ref{eq:collfluct}) to be $z_0 \alt 1 \times 10^{-4}$ over the range of scales $\theta_{\rm crit} \agt 5^{\circ}$ we have simulated. If the collision is described by a single super-Hubble wavelength mode confined to a disc on the sky, from Eq.~\ref{eq:griszeleffect} we can use these bounds (with the largest collision size we have simulated: $\theta_{\rm crit}=25^{\circ}$) to constrain $\Omega_k^{1/2} \Phi(0) \alt 7 \times 10^{-4}$ (where $\Omega_k$ is the present component in curvature and $\Phi(0)$ is the initial magnitude of the Newtonian potential caused by the collision). More generally, Eq.~\ref{eq:collbound} bounds the nucleation rate of bubbles in our parent vacuum, provided gravitational waves and negative curvature are observed with future experiments. Although we have obtained a null result, our analysis pipeline has identified four features in the WMAP 7-year data that have Bayesian evidence ratios that are significantly larger than expected for false detections from an end-to-end simulation of the WMAP experiment. Two of these features (features 2 and 3) have been noted in previous literature. Feature 2 corresponds to the WMAP Cold Spot~\cite{Cruz:2004ce} (see Ref.~\cite{Cruz:2009nd} for a review of its properties), and feature 3 was identified using standard needlets in Ref.~\cite{Pietrobon:2008rf}. All four features are far from the Galactic cut of the KQ75 7-year mask (see Fig.~\ref{fig-detection_locations}), and none appear to be responses to the point source components of the mask (see Fig.~\ref{fig-detection_closeups}). We have confirmed that the signal in each case is not strongly dependent on the frequency band used (see Fig.~\ref{fig-detection_freqdep}), providing evidence that these features are not due to astrophysical foregrounds. A number of analyses, most recently the redshift analysis of Ref.~\cite{2010arXiv1004.1178B}, suggest that the Cold Spot is primordial and not associated with the integrated Sachs-Wolfe effect of a large void. Further studies of the other three features would be needed to confirm that they are truly primordial. Our ability to detect bubble collisions will improve greatly with data from the {\em Planck} satellite. Decreased instrumental noise will enlarge the exclusion and sensitivity regions in parameter space for the needlet step of the analysis, as evidenced by our ability to detect more simulated collisions in low-noise regions of the WMAP data. The threefold increase in resolution will greatly improve our ability to detect circular edges. In addition, the polarization data from {\em Planck} will be of sufficient resolution to look for complementary signatures of bubble collisions~\cite{Czech:2010rg,Dvorkin:2007jp}. Such an analysis should be able to confirm if the features we have identified are in fact bubble collisions. It is also important to determine if other theories predicting azimuthally symmetric features in the CMB~\cite{Cornish:1997ab,2008MNRAS.390..913C,Afshordi:2010wn,Kovetz:2010kv} are better fits to the data. The blob and edge detection steps in our analysis pipeline are sensitive to a variety of possible signatures, and given a model, the Bayesian model comparison step could be easily tailored to accommodate different forms of the temperature modulation. Because our pipeline is automated, we can compare the evidence ratios obtained for different models to decide which is a better fit, without recourse to {\em a posteriori} choices of which features to analyze. In conclusion, we have presented a powerful algorithm for analyzing CMB data for signatures of bubble collisions. Applying this pipeline to the WMAP 7-year data, we have constrained the possible parameter space of bubble collisions, as well as identifying interesting candidate signatures in the data for further investigation. Future data from the {\em Planck} experiment will allow us to greatly improve on these results. If confirmed, the presence of bubble collisions in the CMB would be an extraordinary insight into the origins of our universe.

1012.3667

In the picture of eternal inflation, our observable universe resides inside a single bubble nucleated from an inflating false vacuum. Many of the theories giving rise to eternal inflation predict that we have causal access to collisions with other bubble universes, providing an opportunity to confront these theories with observation. We present the results from the first observational search for the effects of bubble collisions, using cosmic microwave background data from the WMAP satellite. Our search targets a generic set of properties associated with a bubble-collision spacetime, which we describe in detail. We use a modular algorithm that is designed to avoid a posteriori selection effects, automatically picking out the most promising signals, performing a search for causal boundaries, and conducting a full Bayesian parameter estimation and model selection analysis. We outline each component of this algorithm, describing its response to simulated CMB skies with and without bubble collisions. Comparing the results for simulated bubble collisions to the results from an analysis of the WMAP 7-year data, we rule out bubble collisions over a range of parameter space. Our model selection results based on WMAP 7-year data do not warrant augmenting ΛCDM with bubble collisions. Data from the Planck satellite can be used to more definitively test the bubble-collision hypothesis.

12163166

2011A&A...526A..41O

1012

1012.1348_arXiv.txt

As part of a program with the InfraRed Spectrograph (IRS, \citealt{HO04}) onboard the {\it Spitzer Space Telescope} (SST, \citealt{WE04}) aimed at characterizing the circumstellar disks of a flux-limited sample of infrared-excess young stellar objects (YSOs) in the Serpens Molecular Cloud ($\alpha_{J2000}$=18$^h$29$^m$49$^s$, $\delta_{J2000}$=+01$^d$14$^m$48$^s$), an interesting object of unknown nature was discovered (SSTc2dJ18282720+0044450, or \#17 in \citealt{OL10}, hereafter OL17). The very bright, high ionization emission lines seen in the mid-IR spectrum of this object are not consistent with it being a YSO. Two types of galactic objects show such high excitation lines: dusty planetary nebulae (PNe, \citealt{BS09,ST07,GU07}) and supernova remnants (SNRs, \citealt{SA09,GA09}). SNRs typically show very broad emission lines, produced by the high velocity shock waves \citep{FE85,FH96,SP07}. PNe, on the other hand, are characterized by narrow emission lines arising from the low velocity expanding outer shells \citep{BF02,GO09}. Both classes of objects have been extensively studied by several authors, although just a few are so dusty that they were initially discovered only at mid-infrared wavelengths. To distinguish between these two possibilities, further spectroscopy on OL17 was needed. In this research note, we present the original IRS spectrum (\S~\ref{s_irs}) as well as follow-up VLT/X-shooter spectra obtained as part of the instrument science verification phase (\S~\ref{s_xshoot}) and report on our identification of this object as a PN (\S~\ref{s_res}).

1012.1348

As part of a mid-infrared spectroscopic survey of young stars with the Spitzer Space Telescope, an unclassified red emission line object was discovered. Based on its high ionization state indicated by the Spitzer spectrum, this object could either be a dusty supernova remnant (SNR) or a planetary nebula (PN). In this research note, the object is classified and the available spectroscopic data are presented to the community for further analysis. UV/optical/NIR spectra were obtained during the science verification run of the VLT/X-shooter. A large number of emission lines are identified allowing the determination of the nature of this object. The presence of strong, narrow (Δv ~8 - 74 km s<SUP>-1</SUP>) emission lines, combined with very low line ratios of, e.g., [N ii]/Hα and [S ii]/Hα show that the object is a PN that lies at an undetermined distance behind the Serpens Molecular Cloud. This illustrates the potential of X-shooter as an efficient tool for constraining the nature of faint sources with unknown spectral properties or colors.

12207259

2011JCAP...09..038P

1012

1012.1662_arXiv.txt

\label{sec:intro} The current accelerated expansion of the universe has been theoretically explained by either the cosmological constant or some kind of dynamical dark energy in the context of field or modified gravity (see Ref.\ \cite{DE-review-2010} for recent reviews). In the modified gravity such as $f(R)$ gravity \cite{fR-gravity-DE} (see Refs.\ \cite{fR-gravity-review} for reviews and \cite{fR-gravity-recent,Amendola-etal-2007,Li-Barrow-2007, Amendola-Tsujikawa-2008,Hu-Sawicki-2007,Starobinsky-2007,Appleby-Battye-2007, Tsujikawa-2008,Amendola-etal-2007-PRL,Carloni-etal-2008, Clifton-2008} for recent works), background evolution sometimes confronts with a numerical difficulty. The background evolution is not easily handled numerically during the early radiation dominated era in the modified gravity with a functional form $f(R)=R+f_\textrm{DE}(R)$, where $f_\textrm{DE}$ drives the late-time acceleration. The cosmologically viable forms of $f_\textrm{DE}$ proposed so far are $qR^{-n}$ \cite{Amendola-etal-2007,Li-Barrow-2007,Amendola-Tsujikawa-2008} and forms given by Hu \& Sawicki \cite{Hu-Sawicki-2007}, Starobinsky \cite{Starobinsky-2007}, and so on \cite{Appleby-Battye-2007,Tsujikawa-2008}. In all models, the second term $f_\textrm{DE}$ becomes extremely subdominant compared with $R$ in the early radiation dominated era so that the $f(R)$ gravity effectively goes over into the Einstein gravity; in general, however, the evolution could be more complicated, see \cite{Clifton-2008}. Since the quantity $F \equiv df/dR$ becomes extremely close to unity in such a situation, evolving a differential equation like Eq.\ (\ref{eq:ddF}) or (\ref{eq:ddR}) below is sometimes not numerically feasible. Adopting the modified form $f(R)=R^{1+\epsilon}+f_\textrm{DE}$ with small positive $\epsilon$ together with appropriate initial conditions we can evade this numerical problem (here we use the Planck unit with $8\pi G\equiv 1 \equiv c$). It is known that the first term $R^{1+\epsilon}$ which is dominant in the early epoch allows the density of gravity sector to follow that of dominant fluid (scaling evolution) \cite{Amendola-etal-2007}. We are motivated to study the case in order to investigate the observationally allowed regions with qualitatively different evolution available in our case of $f(R)$ gravity. By considering $R^{1+\epsilon}$ term, however, the gravity with $f_\textrm{DE}=-2\Lambda$ does not go over into the Einstein gravity in recent era. We will still consider values of $\epsilon$ which is likely to be excluded by the solar-system test because with vanishingly small $\epsilon$ the system of equations cannot be handled numerically due to limited numerical precision in the early era. Though, we will show that for smallest value of $\epsilon$ we considered, the cosmological evolution we study is numerically similar to the evolution in Einstein's gravity with cosmological constant. In this paper, we present initial conditions of background and perturbed variables during the scaling regime in this gravity. Using the initial conditions for scaling evolution we present background evolution, matter (density) and cosmic microwave background (CMB) anisotropy power spectra, and perturbation growth in the gravity with $f_\textrm{DE}=qR^{-n}$. We show that the CMB power spectrum is not sensitive to the model parameters, and explore the viable parameter space constrained by the type Ia supernova (SNIa), matter power spectrum, and the future perturbation growth factor observation. Throughout this paper we assume spatial flatness ($K\equiv 0$). Notations and the basic set of equations in $f(R)$ gravity are summarized in Ref.\ \cite{Hwang-etal-2010}.

In this paper we have studied a $f(R)$-gravity based dark energy model with early scaling era. We have presented initial conditions of background and perturbed variables during the early scaling evolution regime in the modified gravity with a pure power-law form $f(R)=R^{1+\epsilon}$ in the early era. With these initial conditions, the modified gravity with a form $f(R)=R^{1+\epsilon}+f_\textrm{DE}(R)$ where the second term drives the late-time acceleration becomes free from the numerical difficulty that is usually confronted during the background evolution in the early radiation dominated era for $f(R)=R+f_\textrm{DE}(R)$ gravity. Our initial conditions are general so that the scaling density evolution of the $X$-component is assured for any dominant fluid with a constant equation of state parameter $w$. As a possible dark energy model we have considered the gravity with a form $f(R)=R^{1+\epsilon}+qR^{-n}$ and compared the evolution of the background and perturbation variables in this gravity with the recent observational data and the $\Lambda\textrm{CDM}$ mock data. The present observational data already severely constrain our model parameters $n$ and $\epsilon$ so that only parameters extremely close to the $\Lambda\textrm{CDM}$ model is allowed. We found that the power spectrum of baryon component and the perturbation growth factor at small scales are more sensitive to the $f(R)$ gravity parameters than the SNIa distance modulus and the CMB anisotropy power spectra (see Figs.\ \ref{fig:fR_pow}--\ref{fig:fR_like}). Therefore, precise measurement of the perturbation growth is essential to tightly constrain our $f(R)$ gravity.

1012.1662

The modified gravity with f(R) = R<SUP>1+epsilon</SUP> (epsilon &gt; 0) allows a scaling solution where the energy density of gravity sector follows the energy density of the dominant fluid. We present initial conditions of background and perturbation variables during the scaling evolution regime in the modified gravity. As a possible dark energy model we consider a gravity with a form f(R) = R<SUP>1+epsilon</SUP>+qR<SUP>-n</SUP> (-1 &lt; n &lt;= 0) where the second term drives the late-time acceleration. We show that our f(R) gravity parameters are very sensitive to the baryon perturbation growth and baryon density power spectrum, and present observational constraints on the model parameters. We consider full perturbations of f(R) gravity. Our analysis suggests that only the parameter space extremely close to the ΛCDM model is allowed with epsilonlesssim5 × 10<SUP>-6</SUP> and ngtrsim-10<SUP>-4</SUP>.

3747456

2011ApJ...732...74G

1012

1012.2382_arXiv.txt

The search for planets about other stars has led to the discovery of dozens of planets with unusual properties. As both radial velocity and transit surveys are biased towards planets that are both massive and close to their parent stars, the region of parameter space corresponding to planets with a mass larger than that of Neptune and a semi-major axis $< 0.1$ au is particularly well-explored \citep{Ida:2004p4433, Shen:2008p4684, Zakamska:2010p5800}. This has led to the discovery of many giant planets known colloquially as hot Jupiters and Neptunes, which are thought to have formed far away from their parent stars but were then later transplanted to their observed positions by currently undetermined means. Many of these exoplanets come so close to their parent stars that they toe the line between destruction and survival, with some observed exoplanets in danger of being destroyed on a relatively short timescale \citep{Li:2010p4958}. Additionally, the inclination distribution of the hot Jupiters seems to demonstrate significant misalignment between the planet's orbit and the stellar spin axis \citep{Triaud:2010p4772, Schlaufman:2010p4689}, a surprising result that may require a dynamical process that acts after the protoplanetary disk dissipates. There are three primary physical processes that can deposit a planet on an orbit that is very close to its parent star: disk migration, the Kozai mechanism, and planet-planet scattering. Disk migration can yield hot Jupiters, but as the star collapses from the same cloud as the protoplanetary disk that encircles it, it is difficult to explain the observed orbit misalignments using this mechanism alone \citep[though see][for discussions regarding star-disk interactions]{Foucart:2010p4925,Watson:2010p4924}. The Kozai mechanism \citep{Kozai:1962p5046} can lead to the large eccentricities required to produce close-in planets, but it can only operate in systems in which a massive planet or secondary star are present, and the mechanism may be mitigated by general relativistic effects that become important before tidal dissipation is large enough to circularize the orbit \citep{Takeda:2005p4492, Fabrycky:2007p5801}. Planet-planet scattering can produce both the observed semi-major axis and inclination distributions, and can deposit planets close enough such that tides can circularize the orbits in a time that is less than the system age. Additionally, the object that acts as a scatterer can have approximately the same mass as the scattered object itself and still yield a hot Jupiter in a significant fraction of systems \citep{Ford:2008p4328}, negating the need for a non-planetary companion in the system. Previous hydrodynamical work has only focused on the planet's first close fly-by \citep[hereafter FRW]{Faber:2005p4315}, and does not investigate how prolonged tidal forcing over many orbits affects a planet's chances for survival. In this paper we have performed hydrodynamical simulations of multiple passages of a Jupiter-like planet by a Sun-like star, bridging the gap between numerical and analytical work that have focused on extremely close and extremely grazing encounters respectively. We find that scattering planets into star-grazing orbits is more destructive than previously thought, with Jupiter-like planets being destroyed or ejected at distances no smaller than 2.7 times the tidal radius \(r_{\rm t} \equiv R_{\rm P} (M_\ast/M_{\rm P})^{1/3}\). As some exoplanets are currently observed to have semi-major axes less than twice this critical value, their initial eccentricities may be required to have been substantially smaller than unity if planet-planet scattering is the mechanism responsible for bringing them so close to their host stars. This strongly suggests that planet-planet scattering alone cannot explain the complete observed population of close-in Jupiter-like exoplanets, and that the process must operate along with one of either the Kozai mechanism, disk migration, or both. These three processes likely act in concert to produce the observed population of hot gas giants, with the relative importance of each process being a function of the system's initial conditions. If planet-planet scattering is common enough to explain the existence of hot Jupiters, we predict that there should be two signatures of disruption that are readily detectable with today's instruments. Firstly, we find that the parent star can have its spin significantly altered by the accretion of material removed from the planet as a result of the disruption, producing a star that can be significantly misaligned relative to any remaining planets. Secondly, we find that most planet disruption events lead to the planet's ejection from the host system prior to the planet being completely destroyed, and that this ejected planet can remain almost as bright as its host star for centuries. In this paper we focus on the results of numerical hydrodynamical simulations that have been used to attempt to ascertain the true radii of destruction and ejection for Jupiter-like exoplanets, and the consequences of these planet-removing processes on their stellar hosts. In Section \ref{sec:modeling} we review the history of the analytical and numerical work done to characterize the orbital evolution of a planet that comes within a few tidal radii of its host star, and then we detail our particular numerical approach to modeling tidal disruption. We report the results of our simulations in Section \ref{sec:simres}. In Section \ref{sec:disc} we discuss the implications of our results, with special attention paid to the viability of various mechanisms for producing hot Jupiters, and the observational signatures of planetary disruption and ejection. We summarize the shortcomings of our models and the possible fates of a Jupiter-like exoplanet in Section \ref{sec:conclusions}. Appendix \ref{sec:modgrav} is provided to detail our algorithm used to simulate multiple orbits and for presenting tests of the algorithm's conservative properties.

\label{sec:conclusions} \subsection{Limitations and Future Directions} The principle assumptions that we have made in this paper is that Jupiter-like planets are represented accurately by a polytropic model of its structure. One advantage of this model is that disruption simulations are trivially scalable to planets of a different size by a simple correction to $\beta$, assuming that the planet's mass interior to a given radius $M_{\rm P}(<r)$ scales self-similarly and that the fluid $\gamma$ remains unchanged. An $n = 1$ polytrope reproduces the mass profile of coreless 1 $M_{\rm J}$ planet relatively well, with the difference in $M_{\rm P}(<r)$ never exceeding 10\% throughout the planet's interior (N. Miller, private communication). The inclusion of a core of a few tens of $M_\oplus$ affects $M_{\rm P}(<r)$ out to a few times the core radius, for which $M_{\rm P}(<r) \sim 0.4 M(R_{\rm P})$. Beyond this radius, the structure of the planet is nearly identical to the coreless/polytropic models. This means that our simulations should be an accurate representation for disruptions where $\lesssim 70\%$ of the planet's mass is removed for Jupiter-like planets. For Neptune-like planets, where the core mass can be larger than the gas mass, the difference in $M_{\rm P}(<r)$ is substantial all the way to the planet's outermost layers, and thus our simulation results should not be directly applied. As the average densities of Neptune-like planets is larger than Jupiter-like planets, $r_\tau$ for Neptunes should assume a smaller value. Additionally, we assume that $\gamma = 2$ throughout the simulation volume, even for regions of very low density where the fluid is completely ionized and should behave as an ideal gas ($\gamma = 5/3$) or even a radiation pressure dominated fluid ($\gamma = 4/3$) in the lowest-density regions. This transition to different values of $\gamma$ should affect the structure of the hot envelope that forms from the re-accreted debris that surrounds a partially disrupted planet, which is dynamically unimportant but may affect the planet's observable signature. This is not to say that a more realistic equation of state would not affect the mass loss itself. As the process of ripping material from the planet involves rapid fluid decompression, a decrease in $\gamma$ may result in slightly altered disruption dynamics. Ideally, one would like to extend the models we have presented here to include a more physical equation of state that can treat all components of the pre- and post-disrupted planet realistically. As the resolution required to determine $r_\tau$ for multiple-orbit encounters beyond what we have presented here is prohibitive, it seems that the exploring the affects of using a more-complete equation of state with a realistic initial planet model is the next natural step for future studies. In the case of planets with a substantial core, these modifications are necessary to determine $r_\tau$ with any confidence. \subsection{The Fates of Scattered Jupiters} The fate of a Jupiter-like planet after a strong scattering event is a function of the strength of the tidal forces it experiences at periastron. In this paper, we have determined the disruption radius $r_\tau$ for Jupiter-like planets which sets the boundary between long-term survival and rapid tidal disruption. Below, we summarize the various post-scattering outcomes in order of decreasing distance, using $r_\tau$ and the tidal radius $r_{\rm t}$ as points of reference. Stalled $\left(r_{\rm p} \gtrsim 6 r_{\rm t}\right)$: The planet is deposited into an orbit where the rate of tidal dissipation is too small to result in a change in semi-major axis over the lifetime of the system. This planet may be in a Kozai state driven by a third body in the system, or could experience another strong scattering event, which may lead to an increase of eccentricity and subsequent circularization. Circularization/Migration $\left(r_\tau < r_{\rm p} < 6 r_{\rm t}, e \lesssim 0.9\right)$: In this region, the planet is close enough to its parent star that tidal dissipation is effective, and the planet can circularize in $10^9$ yr or less for moderate values of $e$. For near-radial orbits, circularization may still be longer than the stellar age, but again the Kozai mechanism or scattering could lead to a more rapid orbital evolution. All currently observed hot Jupiters are either stalled, in the process of circularizing, or already have circular orbits. If the planet is close enough to its parent star and $Q_\ast \lesssim 10^7$, the planet will raise a tide on the star and migrate inwards due to the transfer of angular momentum. Ejection $\left(r_{\rm p} < r_\tau, e \gtrsim 0.97\right)$: A planet that passes within the exclusion zone will be ejected from the system if its initial orbit is radial enough such that its orbital energy $E_{\rm orb}$ is significantly smaller than the self-binding energy $E_{\rm p}$ of the planet. Slightly less than half of the planet's initial mass remains bound to the central star, carrying with it a large reservoir of angular momentum that can significantly alter the host star's spin rate and axis of rotation. The ejected planet will remain as bright or brighter than its host star for a few years, eventually plateauing via hydrogen recombination as an object with a fraction of solar luminosity for a century. All Jupiter-like planets that scatter in from beyond $a_{\rm ice}$ such that $r_{\rm p} < r_\tau$ will be ejected from the system. Complete disruption $\left(r_{\rm p} < r_\tau, e \lesssim 0.97\right)$: For planets that are deep within their parent star's potential well, the planet cannot soak the change in energy required to significantly alter the orbit, which eventually leads to its complete disruption. Approximately half of the planet's mass accretes onto the stellar host, carrying the same specific angular momentum as in the ejection case, leading to even more pronounced effects on the stellar spin. Only planets that have migrated close to their stars prior to being scattered are destroyed before they are ejected. Collision with the central star $\left(r_{\rm p} < R_\ast\right)$: The planet strikes the surface of the star directly. Anywhere from half to all of the planet's mass is absorbed by the star, with the angular momentum being carried by this material potentially being smaller than that carried by the debris from a disruption, depending on how direct the impact is. These events are approximately twice as uncommon as ejections/disruptions.

1012.2382

The discovery of Jupiter-mass planets in close orbits about their parent stars has challenged models of planet formation. Recent observations have shown that a number of these planets have highly inclined, sometimes retrograde orbits about their parent stars, prompting much speculation as to their origin. It is known that migration alone cannot account for the observed population of these misaligned hot Jupiters, which suggests that dynamical processes after the gas disk dissipates play a substantial role in yielding the observed inclination and eccentricity distributions. One particularly promising candidate is planet-planet scattering, which is not very well understood in the nonlinear regime of tides. Through three-dimensional hydrodynamical simulations of multi-orbit encounters, we show that planets that are scattered into an orbit about their parent stars with closest approach distance being less than approximately three times the tidal radius are either destroyed or completely ejected from the system. We find that as few as 9 and as many as 12 of the currently known hot Jupiters have a maximum initial apastron for scattering that lies well within the ice line, implying that these planets must have migrated either before or after the scattering event that brought them to their current positions. If stellar tides are unimportant (Q <SUB>*</SUB> &gt;~ 10<SUP>7</SUP>), disk migration is required to explain the existence of the hot Jupiters present in these systems. Additionally, we find that the disruption and/or ejection of Jupiter-mass planets deposits a Sun's worth of angular momentum onto the host star. For systems in which planet-planet scattering is common, we predict that planetary hosts have up to a 35% chance of possessing an obliquity relative to the invariable plane of greater than 90°.

12224836

2011PhRvD..83l3517L

1012

1012.2511_arXiv.txt

\label{intro} The accumulation of WMAP seven year measurement on cosmic microwave background radiation (CMB) \cite{wmap7}, associated with observations of SDSS \cite{sdsslrg7}, provide wealthy information on the anisotropies and inhomogeneities of our universe, in light of which the perturbation theory has been tested in certain level. Currently, various observational data favor the simplest concordance cosmological model, which has six free parameters and the pure adiabatic initial condition \cite{wmap7,Xia:2008ex,Li:2006ev,Li:2008vf,Li:2010ac}. Although the concordance model fits the data quite well, it is always worthy to consider alternative candidates. Namely, it is important to study observational constraints on initial states of cosmological perturbations at reheating surface. Generically, there exists two classes of modes of cosmological perturbations, with one being adiabatic of which the trajectory is parallel to the background evolution, while the other isocurvature of which the trajectory is orthogonal to the background evolution\cite{Linde:1985yf, Kofman:1986wm, Mollerach:1990ue, Kawasaki:1995ta, Polarski:1994rz, Sasaki:1995aw, GarciaBellido:1995qq, Gordon:2000hv, Bartolo:2001rt}. However, a first lesson from observational data is that the primordial fluctuations are nearly adiabatic; additionally, an isocurvature mode is also expected to be negligible as is predicted by the simplest inflation model in terms of a single inflaton field and a rapidly reheating process\cite{Kofman:1997yn}. Therefore, one usually makes the data fitting without considering the isocurvature mode. Accompanied with developments of inflationary cosmology, models of multiple field inflation were extensively studied in the literature\cite{Liddle:1998jc, Kanti:1999vt, Copeland:1999cs, Langlois:1999dw, Wands:2002bn, Piao:2002vf, Dimopoulos:2005ac, Langlois:2008wt, Cai:2009hw, Cai:2008if, Cai:2010wt}, which predicted an existence of primordial isocurvature fluctuations. These primordial isocurvature modes could be transferred into cosmological perturbations after reheating, such as Bayon, CDM, DE, and neutrino respectively\cite{Seckel:1985tj, Linde:1991km, Linde:1996gt,Liu:2010ba}. Consequently, the attention on isocurvature modes has been awaken in recent years\cite{Bucher:1999re, Challinor:1998xk}. One may notice that, although the pure isocurvature primordial perturbation has been ruled out by the Boomerang and MAXIMA-1 data \cite{Enqvist:2000hp} already, a mixture of adiabatic and isocurvature modes can be in agreement with the current data fortunately. In the literature \cite{Gordon:2002gv, Crotty:2003rz, Bucher:2004an, Moodley:2004nz, KurkiSuonio:2004mn, Beltran:2005xd, Bean:2006qz, Trotta:2006ww, Sollom:2009vd, Valiviita:2009bp,Beltran:2005gr,Mangilli:2010ut}, the studies on constraining the isocurvature fluctuation have been performed extensively using various observational data, such as CMB, LSS, integrated Sachs-Wolfe effect or Lyman-$\alpha$ forest data. With the WMAP7 data, the totally un-correlated and anti-correlated adiabatic, non-adiabatic perturbations are constrained \cite{wmap7}, which are performed by fixing the correlation coefficients to be $0$ or $-1$, respectively. In this paper, we study the constraints on the isocurvature modes of cosmological perturbations in light of the latest observational data. We consider a generic scenario that the adiabatic modes and isocurvature ones are allowed mixed through a correlation matrix which could be arising from a time-varying background trajectory. In the detailed analysis, we treat the coefficients of correlation matrix as free parameter and make the data fitting. Comparing with the previous results in the literature, our result shows a slight improvement on the final constraints of the mixed initial condition parameters. Namely, CDM isocurvature components are stringently limited, the contribution from the isocurvature modes are only allowed in small scales, also the error bars of initial condition parameters become smaller than the past works in the literature, which can be observed from the one dimensional probability distribution of the fraction parameter $\alpha$ as analyzed in the main context. The WAMP normalization priors, $R$, $l_A$ and $z_{*}$ encoding the information of background cosmic distances, can be applied to greatly simplify the numerical calculations of determining cosmological parameters, such as the EoS of dark energy. It was found that these priors could be sensitive to the peak locations and local structures of the CMB temperature power spectrum \cite{Li:2008cj}. However, it is well-known that these quantities can be affected by CDM isocurvature perturbation. Therefore, we study the effects on the WMAP normalization priors given by the WMAP group from the isocurvature mode perturbation. The outline of this paper is as follows. In Section II, we describe the parametrization of cosmological perturbations with adiabatic and isocurvature modes mixed, and then we consider a specific example to illustrate the effects of isocurvature perturbation imprinted on the CMB temperature power spectrum and LSS matter power spectrum respectively. In Section III, we perform a global analysis and introduce the data we applied and the parameters used in the data fitting. The constraints on these parameters are present in detail in Section IV, and the corresponding effects on reduced distance parameters are discussed. Finally, Section V includes the conclusions and discussion.

\label{summary} In this paper, we study the constraints on mixed adiabatic and isocurvature modes of cosmological perturbations. Using the current observational data, such as WMAP7 CMB power spectrum, matter power spectrum of SDSS data released seven LRG data and SNIa ``union2" sample, we find an adiabatic initial condition with the presence of certain isocurvature modes can explain the experiments better than a pure adiabatic initial condition. Moreover, we obtained more stringent constraints on the parameters of cosmological perturbations by virtue of the improvement of accuracy of observational data in recent years. Our result shows that the spectral index of isocurvature perturbation has a very blue tilt. This quantity may lead to either enhancement or depression of power spectrum on small scales, which depends on the correlation angle of isocurvature and adiabatic modes. Given the WMAP normalization priors are widely used in performing testing of cosmological models, we provide the comparison on constraints of $R$, $l_A$ and $z_{\ast}$ from different initial conditions. Since WMAP normalization priors are derived parameters from the CMB power spectra, the constraints are changed obviously when the isocurvature mode modifies of the shape of the peaks and troughs in CMB power spectra. As an end, we would like to mention that, from the point of view of information criteria, detailed constraints on cosmological parameters depend on the set of parameters of the model which is compared with observational data\cite{Liddle:2004nh}. The simplest model with purely adiabatic scale-invariant primordial power spectrum is able to capture the most relevant clues of early universe physics, however, more information deserves to be explored along with the improvement of observational data from forthcoming experiments, we expect the global analysis on both CMB power spectrum and matter power spectrum on small scales will be crucial to constrain the isocurvature perturbation.

1012.2511

In this paper we present the constraints on cold dark matter (CDM) isocurvature contributions to the cosmological perturbations. By employing Markov chain Monte Carlo method, we perform a global analysis for cosmological parameters using the latest astronomical data, such as 7-year Wilkinson Microwave Anisotropy Probe (WMAP7) observations, matter power spectrum from the Sloan Digital Sky Survey, luminous red galaxies, and the Union2 type Ia supernovae sample. We find that the correlated mixture of adiabatic and isocurvature modes is slightly favored by the current observational data. We also obtain a tight limit on the fraction of the CDM isocurvature contributions, which should be less than 14.6% at 95% confidence level. With the presence of the isocurvature modes, the adiabatic spectral index becomes slightly bigger, n<SUB>s</SUB><SUP>adi</SUP>=0.972±0.014 (1σ), and the tilt for isocurvature spectrum could be large, namely, the best fit value is n<SUB>s</SUB><SUP>iso</SUP>=3.020. Finally, we discuss the effect on WMAP normalization priors, shift parameter R, acoustic scale l<SUB>A</SUB>, and z<SUB>*</SUB>, from the CDM isocurvature perturbation. By fitting the mixed initial condition to the combined data, we find the mean values of R, l<SUB>A</SUB>, and z<SUB>*</SUB> can be changed about 2.9σ, 2.8σ, and 1.5σ respectively, comparing with those obtained in the pure adiabatic condition.

12298870

2012PhRvD..85d3011K

1012

1012.1274_arXiv.txt

\label{intro} Neutrinos are unique probes of the physics of collapsing stars (supernovae). Diffusing from the dense region surrounding the collapsed stellar core, they can deliver first hand information on the collapse of their stars, on the physics of matter near nuclear density and on the propagation of \ns\ from such high densities to the interstellar space and to detectors on Earth. The physics potential of \ns\ from \sne\ has been studied only minimally due to the scarcity of data. These are limited to handful of events from SN1987A \cite{Hirata:1987hu,Bionta:1987qt}, which was the only recent \sn\ close enough for its neutrino flux to be detectable. While the rarity of nearby \sne\ seems an insurmountable problem, a new phase of data taking is expected to begin with the detection of the {\it diffuse \sn\ \n\ flux} (or background, \df), on which only upper limits exist \cite{Malek:2002ns,Eguchi:2003gg,Aharmim:2006wq,Lunardini:2008xd}. Tiny but continuous in time, the diffuse flux will give tens to hundreds of events in a few years at future Mt scale detectors, ensuring constant progress for decades. Besides the practical advantages, the diffuse flux has a theoretical value of its own: indeed, it has the unique potential to probe the entire \sn\ population of the universe in its diversity. An important advancement in this direction is the study of the \n\ flux from {\it failed \sne}, stars that collapse directly into a black hole with no explosion and no significant emissions other than \ns\ and gravitational waves. These direct \bh\ collapses are rare; they are estimated to account for less than $\sim 22\%$ of all collapses \cite{Sumiyoshi:2006id,Lunardini:2009ya}. The physics of failed \sne\ were modeled numerically in a number of works, including \cite{Liebendoerfer:2002xn,Sumiyoshi:2006id,Sumiyoshi:2007pp,Fischer:2008rh,Sumiyoshi:2008zw,Nakazato:2008vj,Nakazato:2010ue}, which predicted the emission of a \n\ flux with a higher luminosity and average energy compared to the flux from regular, \nts\ collapses. In \cite{Lunardini:2009ya} this result was used to make the first calculation of the diffuse \n\ flux from failed \sne. The main result was the possibility that, due to their higher energetics, failed \sne\ might contribute substantially to the \df, with an enhancement of the total flux and event rate in water detectors of up to $\sim 100\%$. The possibility to detect \ns\ from failed \sne\ in the form of a diffuse flux has several interesting implications. Experimentally, the enhancement of the total flux is attractive because it means that a detection might be closer in time and within the reach of the next phase of \sk, especially in the configuration with Gadolinium \cite{Beacom:2003nk}. Theoretically, detecting the diffuse flux would make it possible to learn about direct \bh\ collapses, specifically by constraining their energetics and cosmological rate. This opportunity is especially precious, considering that failed \sne\ are virtually invisible to telescopes \footnote{An interesting possibility has been suggested recently \cite{Kochanek:2008mp}: observing the disappearance of stars rather than the appearance of \sne. In principle, a ``before vs after" comparison, together with the non-observation of a \sn\ explosion, could reveal a direct \bh\ collapse.}. The published \sk\ \n\ data already constrain the rate of failed \sne\ \cite{Lien:2010yb}. A new, preliminary, analysis from the \sk\ collaboration \cite{iida,iidathesis} considers the \n\ flux from failed \sne, and limits it to about a factor of two from the most optimistic predictions. Neutrinos from failed \sne\ can also increase the amount of Technetium 97 ($^{97}$Tc) that accumulates in Molybdenum ores over millions of years due to solar and galactic supernova \n\ irradiation \cite{Lazauskas:2009yh}. It was observed that they also enhance the proposed neutrino-based mechanisms to create amino acid enantiomerism \cite{Boyd:2010ak}. In this paper we elaborate further on the theme of the diffuse \n\ flux from failed \sne, with a focus on its dependence on the relevant parameters and on its signatures at the next generation of \n\ detectors with 0.1 - 1 Mt masses. Specifically, we consider a Mt water \ck\ detector and a 0.1 Mt liquid argon (\lar) experiment. Our results for water \ck\ detectors elaborate on those of ref. \cite{Lunardini:2009ya}, while the discussion of the potential of liquid argon detectors is presented here for the first time. The advent of liquid argon technology will be a revolution for the study of \sn\ \ns due to its strong sensitivity to electron \ns, which complements the sensitivity of water detectors to antineutrinos. A mass of 0.1 Mt is considered to be the minimum mass required to have any sensitivity to diffuse \sn\ \ns. We will show that this configuration might be particularly suited for \ns\ from failed supernovae: their higher energies imply a larger detection cross section compared to \ns\ from neutron-star forming collapses, and their event energy spectrum might peak above the background of solar \ns. The enhancement of the event rate due to failed \sne\ increases the potential of discovery of the \df\ during the earliest phase of the liquid argon technology development. The paper opens with generalities on \ns\ from failed \sne, their expected flux at Earth and basics of their detection and relevant backgrounds (sec. \ref{general}). We then give results for fluxes and event rates in the antineutrino channel (sec. \ref{antinu}) and the \n\ channel (sec. \ref{nu}). In sec. \ref{disc} the results are discussed and summarized.

\label{disc} Let us summarize our results. \begin{itemize} \item The diffuse flux from \bhf\ reflects the features of the original \n\ flux from a failed \sn: it is more luminous and more energetic than the flux from \nts\ collapses, with the most energetic spectra being realized for the stiffer, Shen et al. equation of state. In energy windows relevant for detection, the $\barnue$ component of this flux is at a maximum for the largest $\barnue$ survival probability, due to the especially large flux of $\barnue$s originally produced in the star. An analogous result holds for the $\nue$ component as well. This contrasts with the case of \nts\ collapses, where the luminosity is roughly equipartitioned among the \n\ species. \item Because of its more energetic spectrum, the flux from \bhf\ has a cosmological component -- from stars with $z\gta 1$ -- as large as $\sim 40\%$ above 20 MeV. This is interestingly larger than the $\sim 10\%$ or less expected in the same interval for \nsf, for which the cosmological component largely accumulates below the experimental energy threshold. This could result in new possibilities to use \ns\ to test the rate of collapses at cosmological distances. \item The harder spectrum of the \bhf\ flux can result in a wider energy window of detection for the \df\ (defined as the energy interval where the core collapse flux exceeds the background fluxes of \ns\ of other origin). The window can be up to roughly 7 MeV wider than for \nsf\ only, depending on the magnitude of the atmospheric background relative to the signal (fig. \ref{backgroundsites}). \item The diffuse flux of \ns\ from failed \sne\ could be substantial, up to $\phi^{BH}_{\bar e}=0.38~{\rm ~s^{-1}cm^{-2}}$ ($\phi^{BH}_{ e}=0.28 ~ {\rm ~ s^{-1}cm^{-2}}$) for $\barnue$ ($\nue$) in the interval $19.3 - 29.3$ MeV, normalized to a local rate of core collapses of $R_{cc}(0)=10^{-4}~{\rm yr^{-1}Mpc^{-3}}$. This is only a factor of $\sim 4$ lower than the current sensitivity of \sk, indicating the possibility of detection in the near future. \item Depending on the parameters (the oscillation probabilities, the fraction of \bh\ collapses and the EoS), the flux from failed \sne\ ranges from 6-10\% to a dominant fraction of the total \df, for energies of experimental interest. The total flux is enhanced -- compared to \nts\ collapses only -- by up to a factor $\sim 2.3$, reaching $\phi_{\bar e} \simeq 0.67~{\rm ~s^{-1}cm^{-2}}$ in the $19.3 - 29.3$ MeV window, and $\phi^{BH}_{\bar e}=0.89~{\rm ~s^{-1}cm^{-2}}$ in the open interval $E>19.3 $ MeV. The latter estimate is only a factor of $\sim 2$ lower than of the current \sk\ limit, Eq. (\ref{sklim}), and therefore it is very promising for the next phase of experimental searches. \item The \sk\ limit constrains the multi-dimensional region of the parameter space. This loose constraint can be expressed in terms of conditional limits on the individual parameters: for example, the rate of core collapses is constrained to $R_{cc}(0) < 2.1 \cdot 10^{-4}~{\rm yr^{-1}Mpc^{-3}}$ when all the other parameters are fixed to maximize $\Phi^{BH}_{\bar e}$. Similarly, one gets a limit on the fraction of failed \sne, $f_{BH}=1-f_{NS} \lta 0.7 $, for the same set of parameters and $R_{cc}(0) = 10^{-4}~{\rm yr^{-1}Mpc^{-3}}$. \item in a detector, the most immediate effect of the \n\ flux from \bhf\ is an enhancement of the event rate, which reflects the enhancement of the flux compared to \nts\ collapses only. In a water \ck\ detector with a 2.25 Mt$\cdot$yr exposure (e.g., 0.45 Mt for 5 years) we expect $\sim 5 - 65$ events from failed \sne\ in the window 18-28 MeV of positron energy, out of a total of 63-113 events from all collapses. These represent an excess of $2.3 - 3.9\sigma$, after background rates have been included to calculate errors. For the extended window 10-38 MeV, relevant to water plus Gadolinium, we get 13-165 events from failed \sne\ and a total of $\sim 190-310$ events from all collapses, corresponding to an excess of 4 - 6$\sigma$ above background. \item in liquid argon the spectrum of events from \bhf\ peaks above $\sim 19$ MeV, where the solar \n\ flux terminates. This is a distinctive feature of liquid argon, and is due to the fast increase of the cross section with the \n\ energy. \item For a liquid argon detector with exposure of $0.5$ Mt$\cdot$yr, the larger energy window of 19-39 MeV is overall convenient to increase statistics at the price of a slightly worse signal-to-background ratio. For this window we predict 1-11 events from \bhf, and a total of 12-20 signal events, with 9 events from background. Statistical significance of $3\sigma$ is realized for the S EoS, in the absence of background systematics. \end{itemize} Our results show that, with an improvement by a factor of 2 (in flux) of its sensitivity, \sk\ can start to probe the parameter space of \ns\ from failed \sne\ at the basic level and that next generation detectors should cover a substantial portion of this space. Due to uncertainties in the normalizations, the most robust signature of failed \sne\ in the diffuse flux would be the harder spectrum, and possibly a deviation from the characteristic exponential shape of the spectrum of the \df\ \cite{Lunardini:2006pd} where the two contributions, from \nts\ and \bh\ collapses, are comparable. Such spectral distortion could be visible with the extended energy window of a water+Gd detector or with a liquid scintillator detector \cite{Wurm:2007cy}, which both have the advantage of a better energy resolution. To establish the presence of a flux from failed \sne\ would already be a fundamental result, being the first detection of new type of \n\ source. Beyond the discovery phase, with a high statistics signal it might be possible to distinguish between different models of \bh\ collapse, at the level of favouring one EoS over another, although a model-independent discrimination might not be possible due to the large errors. The cases with the largest $\Phi^{BH}$ -- maximum survival of the electron flavors and smallest $f_{NS}$ -- might be established or ruled out relatively easily, while other scenarios might be more difficult to probe because their lower flux is more shadowed by the atmospheric background. For a given model of \n\ spectra from \bh\ and \nts\ collapses, the position of the spectral distortion (with respect to an exponential spectrum) might be used to probe the relative frequency of the two types, in other words $f_{NS}$, and in turn the minimum progenitor mass required to produce a direct \bh\ collapse (sec. \ref{general}). It is likely that in the space of a few years the rate of \nts\ collapses will be known with good precision from astronomy \cite{snap,snls}, and this will allow translation of the information on $f_{NS}$ obtained from \ns\ into an absolute (as opposed to relative) rate of failed \sne. Data on \ns\ from failed \sne\ would also constitute a new ground to test \n\ oscillations, and therefore \n\ masses and mixings. Realistically, only average survival probabilities could be extracted from a fit to high statistics data. It would be especially interesting to look for differences in the oscillation patterns for \bhf\ and \nsf, as these could give insight on the different physics at play in the two types of collapses (e.g., different matter density profiles influencing the MSW resonances). To conclude, the detection of a diffuse \n\ flux from failed \sne\ is a realistic possibility. It would have profound implications on the study of these invisible objects, on which we have no data so far. The flux is uncertain by more than one order of magnitude, and therefore it remains to be established whether it dominates the total flux or just modifies it at the 10\% level. In the first case, a change of perspective in the field will be needed. In the latter, failed \sne\ would be an ingredient of precision modeling of the \df\ and their parameter space would be constrained experimentally. \subsection*

1012.1274

We discuss the diffuse flux of electron neutrinos and antineutrinos from cosmological failed supernovae, stars that collapse directly into a black hole with no explosion. This flux has a hotter energy spectrum compared to the flux from regular, neutron star-forming collapses and therefore it dominates the total diffuse flux from core collapses above 20-45 MeV of neutrino energy. Reflecting the features of the originally emitted neutrinos, the flux of ν<SUB>e</SUB> and ν¯<SUB>e</SUB> at Earth is larger when the survival probability of these species is larger, and also when the equations of state of nuclear matter are stiffer. In the 19-29 MeV energy window, the flux from failed supernovae is substantial, ranging from ∼7% to a dominant fraction of the total flux from all core collapses. It can be as large as ϕ<SUB>e¯</SUB><SUP>BH</SUP>=0.38s<SUP>-1</SUP>cm<SUP>-2</SUP> for ν¯<SUB>e</SUB> and as large as ϕ<SUB>e</SUB><SUP>BH</SUP>=0.28s<SUP>-1</SUP>cm<SUP>-2</SUP> for ν<SUB>e</SUB>, normalized to a local rate of core collapses of R<SUB>cc</SUB>(0)=10<SUP>-4</SUP>yr<SUP>-1</SUP>Mpc<SUP>-3</SUP>. In 5 years, a 0.45 Mt water Cherenkov detector should see ∼5-65 events from failed supernovae, while up to ∼160 events are expected for the same mass with Gadolinium added. A 0.1 Mt liquid argon experiment should record ∼1-11 events. Signatures of neutrinos from failed supernovae are the enhancement of the total rates of events from core collapses (up to a factor of ∼2) and the appearance of high energy tails in the event spectra.

12286068

2012MNRAS.425..862G

1012

1012.1873_arXiv.txt

Diffuse background light is very important in understanding conditions and classes of objects in the Universe. This is due to the fact that the spectral, spatial, and amplitude information in a diffuse background is linked to the properties of the otherwise unresolved contributing sources. For example, microwave background measurements include contributions of cosmic origin \citep{2009ApJS..180..330K}, as well as foregrounds of Galactic origin \citep{Bernardi:2005hx,2011MNRAS.412.2383C,2008ApJ...680.1222D,Dobler:2008av,2009ApJS..180..265G}. As another example, $\gamma$-ray background measurements include contributions from unresolved blazars \citep{1993MNRAS.260L..21P, 1996ApJ...464..600S, 2000MNRAS.312..177M, 1998ApJ...496..752C,2006ApJ...643...81N, 2009RAA.....9.1205B, 2010ApJ...710.1530V, 2009MNRAS.400.2122A, 2009ApJ...702..523I, Dodelson:2009ih}, inverse Compton scattering of CMB photons by electrons accelerated at shocks around galaxy clusters and cosmic filaments \citep{2000Natur.405..156L,2002MNRAS.337..199M,1998APh.....9..227C,2007ApJ...667L...1M}, starburst galaxies \citep{2002ApJ...575L...5P}, cosmic ray interactions with atomic and molecular gas in the Milky Way \citep{1986A&A...157..223D,2009PhRvL.103y1101A}, as well as the possible annihilation of dark matter \citep{Ando:2009fp,Ando:2006cr,2007JCAP...04..013C,Cuoco:2007sh,Fornasa:2009qh,Hooper:2007be,2009JCAP...07..007L,2010ASPC..426...87S,SiegalGaskins:2009ux,SiegalGaskins:2008ge,Taoso:2008qz,2010PhRvD..82l3511B}. Background events may be divided into two classes. Some events are generated by localized sources while others are generated by mechanisms which cannot be localized. In the first class the sources can be either spatially fixed (in celestial coordinates) or may exhibit proper motion (i.e. over a period of time their displacements are larger than the angular resolution of the detector). Using again the diffuse $\gamma$-ray background as an example, unresolved blazars, starburst galaxies, and emission from structure formation shocks would be considered spatially fixed sources of background. Cosmic ray events with interstellar gas would be considered a non-localized random process. Sources of background which will exhibit proper motion include the interaction of energetic cosmic rays with solar system bodies (e.g., small objects in the asteroid belt or objects in the Kuiper belt and the Oort cloud) \citep{1984JGR....8910685M,2008ApJ...681.1708M,2009ApJ...692L..54M}, dark matter annihilation around primordial black holes \citep{Mack:2006gz,2010ApJ...720L..67L}, and potentially nearby remnants of high density dark matter density peaks \citep{2006PhRvL..97s1301K,2008MNRAS.384.1627P,2008PhRvD..78j1301A,2009NJPh...11j5012K}. In all these cases, individual emission from any single object is not distinguishable, but the sum of these contributions may contribute to the diffuse $\gamma$-ray background. Correlations between individual events can help disentangle the contribution of various sources to the background. In this manuscript we present a formalism and a technique that can be used to identify the presence of background sources that exhibit spatial motion. In Sec.~\ref{sec:overview} we present an overview of the problem. In Sec.~\ref{sec:definitions} we detail definitions that are used in the statistical techniques that follow. This allows us to write down the formal definition of the spacetime 2-point correlation function, which can be used to extract the moving signal in the diffuse background. In Sec.~\ref{sec:2d} we derive the form of the spacetime correlation function in 2 dimensions. A quantitive account of the uncertainty in the method is found in Sec.~\ref{sec:errors}. In Sec.~\ref{sec:demos} we demonstrate the method's robustness in toy experiments and comment about the use of an instrumental point spread function. We generalize the formalism to realistic problems in 3 dimensions in Sec.~\ref{sec:3d}, discuss generalizations of the formalism in Sec.~\ref{sec:discussion} and discuss applications and conclude in Sec.~\ref{sec:conclusions}. \begin{figure*} \centering \resizebox{2.2in}{!} {\includegraphics{limitA.pdf}} \resizebox{2.2in}{!} {\includegraphics{limitB.pdf}} \resizebox{2.2in}{!} {\includegraphics{limitC.pdf}} \caption{Illustration of the two limits in the problem. The first figure contains 5 objects each with event rate 10 and the second contains 50,000 objects with event rate 0.01. The third figure contains the same number of events as the second but they are distributed randomly. Naively, it is impossible to tell which of the last two figures contains random events and which contains moving objects.} \label{fig:2limits} \end{figure*}

\label{sec:conclusions} We present a new tool, based on the familiar 2-point correlation function, which can be applied to astrophysical maps of diffuse emission. The measured quantity $\xi$ is designed to detect the presence of moving objects, each of which is too dim to be resolved individually. We derived the form of $\xi$ based on the physical parameters which describe the classes of objects which might be present in the sky (\ref{X3d}). A measurement of $\xi$ along with the theoretical prediction for $\xi$ can be used to find best-fit quantities for the physical parameters describing the populations of objects. We emphasize that all the technology invented to study the angular 2-point correlation function can be directly applied to the generalization to the spacetime 2-point correlation function. There are numerous applications of the derived formalism. An obvious place to start is the diffuse gamma-ray background measured by the Fermi-LAT instrument. The all-sky capabilities of LAT, coupled with its high angular resolution provide a convenient testbed where this technique can be applied. The interesting question is what kind of sources contribute to the gamma-ray background and also exhibit proper motion over the duration of observation. One potential source is the generation of gamma-rays from cosmic-ray interactions in rocky debris present in the solar system. Cosmic ray interactions with nuclei on a solar system body lead to hadronization, and the subsequent decay of neutral pions to a photon final state \citep{1970Ap&SS...6..377S,1981Ap&SS..76..213S,1986A&A...157..223D,2006PhRvD..74c4018K}. A detection of a large population of these sources is important as it provides information about the origin of the solar system and its evolution with time, as well as the energy spectrum and composition of the incident cosmic ray flux. The detection of gamma-rays from cosmic ray interactions with solar system bodies has been discussed in the context of past measurements by the Energetic Gamma Ray Experiment Telescope (EGRET) on board the Compton Gamma-ray Observatory, and measurements with Fermi-LAT \citep{2008ApJ...681.1708M,2009ApJ...692L..54M,2009arXiv0907.0541G}. Sources include small objects in the main asteroid belt, Trans-Neptunian objects in the Kuiper belt, as well as objects in the Oort cloud, including icy bodies such as comets. It was shown that for objects where the cosmic ray cascade fully develops (objects with size greater than $\sim 1$ m) it may be possible for Fermi to detect the cumulative gamma-ray emission from a collection of such bodies. These estimates are based on the distribution and composition of objects. Even though both of these quantities are partially constrained for objects in the main asteroid belt, large uncertainties are present for the populations in the Kuiper belt and the even more speculative Oort cloud. It is conceivable that a very large number of bodies may be present in the outskirts of the solar system. The proximity of these populations makes them ideal for an application of the spacetime correlation function, as each source will traverse an angular distance which is larger than the angular resolution limit of Fermi. Typical angular displacement (assuming Keplerian orbits) of an object at distance $d$ from the Sun is $\theta = 2 \pi \, {\mathrm{rad}} (\Delta T/\mathrm{yr}) (d / \mathrm{AU})^{-3/2}$ during the course of an integration for time $\Delta T$. The composition of these objects can be assumed to be similar to the composition of the Moon, though their mass density varies considerably. This similarity in composition is convenient as the gamma-ray flux due to cosmic interactions with the lunar rock is well understood \citep{2007ApJ...670.1467M,1984JGR....8910685M} (see also \citep{2009arXiv0907.0543G,2009arXiv0912.3734G}). If we assume that the spectral shape of the gamma-ray emission from solar system bodies is similar to that of the rim of the Moon (emission above 600 MeV is dominated by the rim of the Moon rather than the lunar disc) and we scale the flux from the object to the flux from the Moon ($\Phi_M = 1.1 \times 10^{-6} \mathrm{cm}^{-2} \mathrm{s}^{-1}$, \cite{2009arXiv0907.0543G}), the flux from an object of radius $r$ at distance $d$ would then be $ \Phi = \Phi_M ( r / r_M) (d_M / d)^2$. For a distance to the Moon of $d_M=0.0024$ AU and a lunar radius of $r_M = 1740$ km, the total number of photons per year detected by the Fermi-LAT instrument (with an orbit-averaged effective area of 2000 $\mathrm{cm}^2$) is $\Phi \approx 2 \times 10^{-4} \mathrm{yr}^{-1} ( r / \mathrm{km} ) ( d / \mathrm{1AU} )^{-2}$. Therefore, given this information, one can apply the spacetime correlation function to determine the abundance and radial distribution of solar system objects that contribute to the gamma-ray background \citep{GSK10b}. It is important to note that even though a theoretical estimate of $\xi$ requires knowledge of the objects one is searching for, the measurement of $\xi$ requires no such knowledge. Similar arguments can be used in search of the energetic neutrino signal from cosmic ray interactions with solar system bodies. The decay of kaons to charged pions leads to an energetic signal with a spectral signature that is different from the cosmic ray neutrino flux expected from spallation of nuclei. Therefore, energetic neutrinos from cosmic ray interactions with solar system bodies should be present in the signal measured by IceCube \citep{2006APh....26..155I}. The sources of these neutrinos will traverse an angular distance based on the distance of the source from the Sun, and therefore the spacetime correlation function derived here can be used in search of these sources. However, as in the case of gamma-rays, the uncertainties in the distribution and composition of small solar system bodies make predictions for such signal difficult. Nevertheless, a blind analysis of neutrino events from IceCube could place constraints on the parameters that describe the different populations of small bodies in the solar system. Another application is in the search for primordial black holes in the solar neighborhood. Primordial black holes may form in the early Universe through the collapse of large primordial fluctuations \citep{1971MNRAS.152...75H}. Current bounds on the abundance of such black holes are of order $\Omega_{\mathrm{PBH}} \sim 10^{-9}$ for most of the range of black hole masses \citep{2010ApJ...720L..67L}. If primordial black holes exist in an otherwise dark matter dominated Universe, they will acquire a dark matter halo \citep{Mack:2006gz,2007ApJ...662...53R}. Dark matter annihilation around primordial black holes and/or high density ultracompact halos will result in gamma-ray emission \citep{2009ApJ...707..979R,2009PhRvL.103u1301S}. Such objects with very small mass will in fact be very dense and survive in the Milky Way halo. If we assume that primordial black holes trace the distribution of dark matter in the Milky Way we can use their abundance to determine the angular distance that a black hole may traverse in a given time interval. For simplicity, let's assume that primordial black holes have mass $M_{\mathrm{PBH}} = 10^{-15} M_\odot$, $\Omega_{\mathrm{PBH}} = 10^{-9}$, and that the local dark matter density is 0.01 $M_\odot \mathrm{pc}^{-3}$. Then the mean distance between primordial black holes in the solar neighborhood is $\sim 10^{-2} \mathrm{pc}$. Assuming that this is the maximum distance to a primordial black hole, and that the mean velocity of primordial black holes is similar to the mean velocity of dark matter, i.e., 220 km/s, then the angular displacement of these gamma ray sources can be as large as 4.5 degrees in 10 years. As the angular resolution of Fermi is significantly less for energies greater than 1 GeV, constraints on the abundance and size of these black holes can be placed by applying the spacetime correlation function to the LAT all-sky map. A more speculative contribution to the gamma-ray background is from dark matter halos formed on scales close to the cutoff scale of the dark matter power spectrum. These objects typically have sub-solar masses \citep{1999PhRvD..59d3517S,2001PhRvD..64h3507H,2004MNRAS.353L..23G,2005JCAP...08..003G,2005PhRvD..71j3520L,2001PhRvD..64b1302C,2006PhRvL..97c1301P}. Even though their survival and abundance in the present-day Milky Way halo is unknown, it is possible that dark matter annihilation in these high-density objects may contribute to the gamma-ray background \citep{2008MNRAS.384.1627P,2008PhRvD..78j1301A}. The probability that such sources will exhibit spatial motion in the duration of the Fermi-LAT mission is directly linked to their abundance, and thus the use of the correlation function can provide information on the survival rate of these extremely early-forming objects. The spacetime correlation function can be applied to lensing surveys to search for compact objects in the Milky Way. Past studies suggest that up to 20\% of unseen matter is in the form of Massive Compact Halo Objects (MACHOs) \citep{2000ApJ...542..281A,2004ApJ...612..877U}. With the advent of dedicated surveys e.g., LSST, \citep{2009arXiv0912.0201L}, as well as astrometric missions such as SIM \citep{2008PASP..120...38U} and Gaia \citep{2008IAUS..248..217L}, it will be possible to generate time-domain maps of lensing events in dense stellar fields. Such information can be used to probe correlated events originating from the spatial translation of compact objects, thus probing the projected velocity distribution of the compact population in the Milky Way. In addition, it may also be possible to place constraints on the density, abundance and distribution of dark matter substructure \citep{2011ApJ...729...49E}. Throughout the development of the analysis we assumed that the event rate due to any source was constant in time. There are many classes of astrophysical objects with time-dependent emission. Most notably, unresolved pulsars are thought to contribute to the diffuse gamma-ray background (e.g. \cite{2011ApJ...727..123W,2010JCAP...01..005F}). While these sources will not exhibit proper motion over the course of observations the temporal correlations of their emitted photons may be discovered through techniques based on the ones presented here \citep{2012MNRAS.421.1813G}. Essentially, one chooses the volumes $V(p)$ according to (\ref{Vdefrect}) (illustrated in the left panel of Fig.~\ref{fig:Vdef}), but with a non-trivial slicing along the time axis. Such a $V(p)$ picks up on stationary objects which exhibit correlations within their photon time series. The power of this analysis for untangling the contribution of different classes of sources requires that each class have ``different enough'' velocity, luminosity, and spatial distributions. For example, if two classes have similar velocity and spatial distributions then one may as well just treat them as a single class with a modified luminosity function. This points to a problem that is likely to be encountered in many realistic astrophysical applications: the angular velocities of almost all objects will be much too small to be resolved by a detector. That is, when one combines the velocity distribution $f_i$ with the spatial distribution $n_i$ in (\ref{X3d}) it may be that $\xi=0$ at all angular velocities except in a tiny range near $\w=0$. This is because virtually all of the objects have distances and speeds such that their apparent proper motion is below the angular resolution of the detector. A large degeneracy is created and it will be impossible to pull out information about any specific class of objects. The fact that $\xi$ is not zero at $\w=0$ indicates the {\em existence} of objects. However, without being able to measure the shape of $\xi$ for different angular speeds $\w$ the 2-point function loses its value as a tool to untangle the contributions from different classes of objects. Of course, as the resolutions of detectors improve, the 2-point function becomes more useful. It is a straightforward task to calculate $\xi(\w_1,\w_2; t_1,t_2)$ for specific classes of objects and find out over what ranges of $\w$ and $t$ the correlation drops to zero. For example, if $\xi$ goes to zero around ($\w_1=\w', t_1=t'$) then a detector which has resolution better than $\T \approx \w' t'$ can measure the shape of $\xi(\w,t)$ as it goes from a maximum at ($\w_1=0,t_1=0$) to zero at ($\w_1=\w', t_1=t'$). In summary, we introduced the spacetime correlation function, a statistical tool that can be used to search for the presence of moving, flux-unresolved sources in a diffuse background. This formalism has numerous applications. With large area sky surveys and long duration baselines the spacetime correlation function can be used to disentangle the contributions from spatially moving sources, and may aid in the discovery of new sources. We thank the anonymous referee for comments and suggestions that improved the quality of this manuscript. We acknowledge useful conversations with John Beacom, Jacqueline Chen, Richard Cook, Ian Dell'Antonio, Scott Dodelson, Andrew Favaloro, Salman Habib, Arthur Kosowsky, David Laidlaw, Miguel Morales, Louie Strigari, Andrew Zentner. SMK and AGS are funded by NSF PHYS-0969853 and by Brown University.

1012.1873

We present a statistical technique which can be used to detect the presence and properties of moving sources contributing to a diffuse background. The method is a generalization of the two-point correlation function to include temporal as well as spatial information. We develop a formalism which allows for a derivation of the space-time two-point function in terms of the properties of the contributing sources. We test this technique in simulated sky maps, and demonstrate its robustness in identifying the presence of moving and stationary sources. Applications of this formalism to the diffuse gamma-ray background include searches for Solar system bodies, fast moving primordial black holes and dense cores of dark matter protohaloes in the solar neighbourhood. Other applications include detecting the contribution of energetic neutrinos originating in the Solar system, as well as probing compact objects in long-timeline lensing experiments.

12168394

2011ApJ...734...62H

1012

1012.2870_arXiv.txt

Dark matter (DM) dominated cosmological structures are a direct outcome of hierarchical structure formation with a subdominant baryon fraction. The baryons can either cool fast to form stars or end up as a hot virialized gas \citep{1977MNRAS.179..541R,1977ApJ...211..638S,2003MNRAS.345..349B,2005MNRAS.363....2K}. This hot gas is customarily observed in galaxies and galaxy clusters \citep{1986RvMP...58....1S}, and can be heated through adiabatic compression, shock heating or non-gravitational processes and cooled through radiative processes. The ability of the hot gas to shock heat to the virial temperature opens the possibility that one can combine the equations governing the dynamics of the gas and the DM, namely the equation of hydrostatic equilibrium and the Jeans equation. The simultaneous solution to these equations can in principle allow us to determine the DM velocity dispersion anisotropy, which holds information about the fundamental difference in the way DM and baryons equilibrate in cosmological structures. In fact, such a method has already been applied to galaxy clusters, where numerical simulations of cluster formation and evolution have been used to confirm that the gas equilibrium temperature is determined through the averaged velocity dispersion of the DM~\citep{hansenpiff,host2009}. The resolution of these simulations, however, did not allow us to probe radii smaller than a few hundred kpc. The temperature of the cooling gas remains to be determined at smaller radii and in the presence of an active AGN. Furthermore, it remains unknown if the gas temperature in galaxies will obey a similar simple relation. In this paper we conduct a dedicated comparison between the gas and DM in a range of simulated cosmological objects. The structure of the paper is as follows: first we discuss the relationship between the gas temperature and the DM velocity dispersion. In section \ref{sec:num} we present the results of numerical simulations of galaxy formation, and confirm that the gas temperature in galaxies exhibits a bimodal distribution --- the hot gas resides at the DM temperature and the cold gas cools below $10^4$~K and forms stars. In section \ref{sec:agn} we present the results of numerical simulations of AGN outflows, which demonstrate that even when the gas is heated episodically by a self-regulated AGN, it still moves towards the DM temperature, which it succeeds in reaching already beyond $\approx 30$~kpc. Finally, in section~\ref{sec:discussion} we explain how our findings open up for the possibility of measuring the DM velocity anisotropy in galaxies, and in section \ref{sec:conclusion} we briefly offer our conclusions.

\label{sec:conclusion} In summary, we have found a very simple relation between the temperature of the hot gas and the averaged dispersion of the dark matter, namely that their ratio is close to unity, $\kappa \approx 1$. We have demonstrated that this near equality of the gas and DM temperatures holds in the case of equilibrated and relaxed galaxies. This relation is not only conceptually important, but it will also allow an indirect determination of the dark matter velocity anisotropy in galaxies. Such an observation will provide an important and independent confirmation of all cosmological DM halo formation simulations. We have also studied this relation in galaxy clusters containing an active AGN, and our results confirm that one can indeed use galaxy clusters to infer the dark matter velocity anisotropy. \newpage \noindent It is a pleasure to thank Isaac Shlosman for insightful discussions. SHH thanks Ole Host, Marco Roncadelli and Darach Watson for discussions. Numerical simulations were performed on the PIA and on the PanStarrs2 clusters of the Max-Planck-Institut f\"ur Astronomie at the Rechenzentrum in Garching and on the zBox2 supercomputer at the University of Z\"urich. Special thanks to B. Moore, D. Potter and J. Stadel for bringing zBox2 to life. E.R.D. simulations have been run on a dedicated cluster. The AGN simulations were conducted on the "Saguaro" cluster operated by the Fulton School of Engineering at Arizona State University. Y.H. has been partially supported by the ISF (13/08). The Dark Cosmology Centre is funded by the Danish National Research Foundation.

1012.2870

The temperature profile of hot gas in galaxies and galaxy clusters is largely determined by the depth of the total gravitational potential and thereby by the dark matter (DM) distribution. We use high-resolution hydrodynamical simulations of galaxy formation to derive a surprisingly simple relation between the gas temperature and DM properties. We show that this relation holds not just for galaxy clusters but also for equilibrated and relaxed galaxies at radii beyond the central stellar-dominated region of typically a few kpc. It is then clarified how a measurement of the temperature and density of the hot gas component can lead to an indirect measurement of the DM velocity anisotropy in galaxies. We also study the temperature relation for galaxy clusters in the presence of self-regulated, recurrent active galactic nuclei (AGNs), and demonstrate that this temperature relation even holds outside the inner region of ≈30 kpc in clusters with an active AGN.

12230689

2011PNAS..10812641N

1012

1012.5661_arXiv.txt

Enormously powerful events illuminate the universe that challenge our understanding of the cosmos. Indeed, intense electromagnetic emissions of order $\gtrsim 10^{51}{\rm ergs}$~\cite{Hamuy:2003xv,Woosley:2006fn,2001ApJ...562L.149M} have been routinely observed in supernovae, gamma ray bursts~(GRBs), and active galactic nuclei~(AGNs), for example. However, despite important theoretical and observational advances, we still lack a thorough understanding of these systems (e.g.~\cite{Fender:2010tk}). Intense observational and theoretical efforts are ongoing in order to unravel these fascinating phenomena. While the full details remain elusive, one of the natural ingredients in theoretical models is the inclusion of a rotating black hole which serves to convert binding and rotational energy of the system to electromagnetic radiation in a highly efficient manner. The starting point for these theoretical models can be traced back to ideas laid out by Penrose~\cite{Penrose:1969pc} and Blandford and Znajek~(BZ)~\cite{Blandford:1977ds}, which explain the extraction of energy from a rotating black hole. These seminal studies, along with subsequent work (see references in e.g.~\cite{2000PhR...325...83L,2002luml.conf..381B,2008ASSL..355.....P}), have provided a basic understanding of highly energetic emissions from single black hole systems interacting with their surroundings. Recent work has indicated that related systems can also tap kinetic energy and lead to powerful jets~\cite{sciencejets}. This work concentrated on non-spinning black holes moving through a plasma and highlighted that relative black hole motion alone (with respect to a stationary electromagnetic field topology at far distances from it) can induce the production of jets. Furthermore, subsequent work~\cite{prdjets} demonstrated that even black holes with misaligned spins with respect to the asymptotic magnetic field direction induce strong emissions with power comparable to the aligned case. % These studies suggested that, independent of their inclination, astrophysical jets might be powered by the efficient extraction of both rotational and translational kinetic energy of black holes, and inducing even more powerful jets than the standard BZ mechanism would suggest. Galactic mergers provide a likely scenario for the production of the binary black hole systems considered here~\cite{1980Natur.287..307B,Milosavljevic:2004cg}. In such a merger, the supermassive black holes associated with each galaxy will ultimately form a binary in the merged galaxy. A variety of interactions will tighten the black hole binary. Eventually the dynamics of the system will be governed by gravitational radiation reaction which drives the binary to merge. The circumbinary disk will likely be magnetized and thereby anchor magnetic field lines, some of which will traverse the central region containing the binary. Preliminary observational evidence for supermassive black hole binaries resulting from galactic mergers has already been presented~\cite{Comerford:2008gm}. An ambient magnetic field threading a spinning black hole populates a low density plasma surrounding the black hole as explained by BZ~\cite{Blandford:1977ds}. Even for black holes with no spin, it was recently shown that the orbiting binary interacting with the surrounding plasma can lead to a collimated Poynting flux~\cite{sciencejets}. In this work, we consider this basic paradigm of energy extraction from black holes with the additional complexity of intrinsic black hole spin. Binaries consisting of spinning black holes demonstrate % similar, albeit energetically enhanced, phenomenology. Furthermore, we investigate the dependence of the energy flux on the black hole velocity and highlight a resulting strong boost in the emitted power. In addition to exquisite and powerful electromagnetic detectors, soon gravitational waves will be added to the arsenal of phenomena employed to understand our cosmos. These studies suggest excellent prospects for the coincident detection of both electromagnetic and gravitational signals from binary black hole systems. Certainly, dual detection of electromagnetic and gravitational wave signals would transform our understanding of these systems and lead to the refinement of theoretical models (e.g~\cite{2002luml.conf..207S,2003ApJ...591.1152S,2009arXiv0902.1527B,2009astro2010S.235P}). In the remainder of this paper, we elucidate the basic phenomenology arising from the interaction of binary black hole systems immersed within a plasma environment. We focus in particular on understanding the Poynting flux emissions from such binaries as well as single, possibly spinning, black holes moving through plasma. We describe our equations and assumptions employed together with some of the details of our numerical implementation in Sec.~\ref{sec:implementation}. Sec.~\ref{sec:results} describes our results for both single and binary black holes and we provide concluding comments in Sec.~\ref{sec:conclusions}.

1012.5661

The extraction of rotational energy from a spinning black hole via the Blandford-Znajek mechanism has long been understood as an important component in models to explain energetic jets from compact astrophysical sources. Here we show more generally that the kinetic energy of the black hole, both rotational and translational, can be tapped, thereby producing even more luminous jets powered by the interaction of the black hole with its surrounding plasma. We study the resulting Poynting jet that arises from single boosted black holes and binary black hole systems. In the latter case, we find that increasing the orbital angular momenta of the system and/or the spins of the individual black holes results in an enhanced Poynting flux.

12216300

2011NuPhB.849..410K

1012

1012.2107_arXiv.txt

The problem of cosmological constant presents a serious challenge to modern physics. Recently the following physical mechanism was proposed to overcome this problem \cite{AM-recent}. Let us imagine that the bare cosmological constant is present in the lagrangian. As is well known, it is the cause of gravitational repulsion, resulting in the accelerated expansion of the universe (the contraction is also possible, but we will discuss it later). According to the proposal, this acceleration leads to the explosive particle production. The gravitational attraction between the particles slows down the acceleration and thus reduces (asymptotically to zero) the effective cosmological constant. There are many puzzles associated with this proposal. Particles in the curved space are ill defined, does it make sense to ascribe to them a real physical effect? Even if it does, the Universe is exponentially expanding, so it may seem that the particles get diluted; isn't their back reaction negligible? How can massive particles considered in \cite{AM-recent} lead to large infrared effects? In this paper we will try to provide some clarifications, as well as present some new results. It is helpful to consider various cases of unstable vacua and to make their comparative studies. The basic origin of the difficulties lies in the non-equilibrium quantum field theory and are common to all the cases. Let us demonstrate this with the following example. Consider a~nucleus with a charge $Z$ and hit it with a $\gamma$-quantum of energy $\omega$ which produces a pair~$e^+e^-$. The electron forms a~bound state while the positron escapes to infinity. The~threshold singularity in~$% \omega$ is located at \begin{equation*} \omega=2m-|E_B| \end{equation*} where $|E_B(Z)|$ is the~binding energy. We see that when the~nucleus becomes supercharged, $|E_B(Z)|\!=\!2m$, we get a long-ranged correlations in time since the threshold is now at~$\omega\!=0$. Below we will find that such "long-memory" is a crucial factor in the non-equilibrium dynamics. In field theory it leads to the~"adiabatic catastrophe" \cite{Polyakov} and to the obstruction to Wick's rotation. In the next section we shall briefly summarize the situation. Another puzzle mentioned above is the dilution of particles in the expanding universe, the size of which grows as $a(t)\sim e^{t}$. However, in any physical quantity, this effect always cancels with the exponentially growing number of the comoving modes. The covariant cut-off for the comoving momentum $k$ is given by $k\leq k_{\max }\sim M_{planck}a(t)$, and this causes the growth. The naive reason for the above compensation is that the change $a\rightarrow \lambda a$ is a coordinate transformation and $a$ dependence can't be physical. As will be discussed below, there are caveats to this argument, but on a qualitative level they are unimportant. What about the infrared effects generated by the massive particles? They are not related to the interaction of these particles, which is short-ranged as usual. Their origin lies in the fact that the original vacuum is unstable with the non-zero decay rate. Therefore we get perturbative corrections containing secular terms, which represent the fact that, as the time goes, it is less and less likely for the vacuum to remain intact.

The physical interpretation of the above estimates is the following. We are considering a complete $dS$ space. All points of this space are equivalent, so that the statements that at a given point we have expansion or contraction are meaningless. However, if we fix the position of the observer, one can define domains, such that the signal sent from them will be either red shifted or blue shifted. The essence of the formula (\ref{phi^2 complete space}) can be grasped from the Fig.5. We integrate the interaction over the faraway past region. The size of the loop determines the interaction scale $\sim 1/m$, which is a large quantity. While the signal from the interaction region propagates along the geodesics to the observer, sitting at the point $\tau$, it is blue-shifted to the Planck scale $\sim 1/M_{pl}$. As a result we get a very curious UV/IR mixing. In the flat space we expect that UV and IR divergencies contribute to the physical quantities independently - we do not expect the terms, like (\ref{phi^2 complete space}), which are both UV and IR divergent at the same time. This is a specific feature of the curved space. The $\varepsilon $ dependence of the physical matrix element discussed above indicates a breakdown of the $dS$ symmetry; as always, spontaneous symmetry breaking manifests itself through the sensitive dependence on the boundary conditions. The logarithms will be present even for a patch of the $dS$ space, provided that it is "large", that is the past cone of the observer intersects a decent portion of the past infinity. As we saw from Fig.4, this is not the case for the Poincare patch; for it "the world is not enough". Let us explain our motives for using the global $dS$ space, while in the inflationary theories only a small part of it is usually present. Our goal is to resolve the puzzle of the cosmological constant by infrared means. We start with the Einstein action with the cosmological constant present. The standard procedure in field theory is to assume first that we can neglect quantum corrections at large distances, find a classical solution and then evaluate the corrections. It is this procedure which allows us to use classical Einstein or Navier - Stokes and forbids the similar use of the Yang - Mills equations (due to asymptotic freedom) and sometimes the diffusion equation (due to Anderson's localization). In such a setting we must consider the global $dS$ space as a first step. If a starting point were incomplete space, we would end up with the unitarity problem, since the particles can disappear from the space. Of course it is possible to have a Poincare patch glued to the Minkowski one in such a way that the result is geodesically complete. However this space will not be a solution of the Einstein equations with the cosmological constant only. It is also possible to modify the Einstein action so that we have a different background without IR divergences. This looks ambiguous and is far from our goal, which is to tame infrared divergencies. We should remember that to solve the $\Lambda$-problem one must be searching for the infrared effects and not running from them. IR divergence is not a problem but an opportunity. Another question is related to the choice of the Bunch-Davies vacuum in the above calculation. What is the reason for this (apart from the tradition)? It seems that the right starting point should be the state with the longest life time. We haven't proved that this is the case, but various estimates make us believe that the Bunch-Davies vacuum is the most stable one. In the appendix we present the propagators for the different possible ground states. It should not be difficult to extend our analysis to other vacua. Finally, there is a number of valuable papers \cite{intr-dev} intersecting with our work, but it seems that our approach brought some new and unusual results.

1012.2107

We discover that some unstable vacua have long memory. By that we mean that even in the theories containing only massive particles, there are correllators and expectation values which grow with time. We examine the cases of instabilities caused by the constant electric fields, expanding and contracting universes and, most importantly, the global de Sitter space. In the last case the interaction leads to a remarkable UV/IR mixing and to a large back reaction. This gives reasons to believe that the cosmological constant problem could be resolved by the infrared physics.

12134252

2010ASPC..438..236V

1012

1012.2936_arXiv.txt

The Galactic magnetic field plays critical roles in the interstellar medium, ranging from star formation to large-scale galactic dynamics. However, much remains unknown about how the field is generated or how it is evolving. The only way to make progress in addressing these questions is to fully understand the present overall structure of the field. This is an essential constraint to proposed evolutionary models of the field. One observation essential to the study of the Galactic magnetism is that of the rotation measure (RM). As a linearly polarised electromagnetic wave propagates through a region containing free thermal electrons and a magnetic field, such as the interstellar medium, the plane of polarisation will rotate through the process known as Faraday rotation. If we assume the polarised radiation emitted by a source is at a constant angle, $\tau_\circ$, and that this radiation is {\it{only}} affected by Faraday rotation, then the polarisation angle that we measure, $\tau$, will be given by \begin{equation} \tau = \tau_\circ + 0.812 \lambda^2 \int_{\rm source}^{\rm receiver} n_e \mathbf{B \cdot} {\rm{d}}\mathbf{l} =\tau_\circ+ \lambda^2 {\rm RM}, \end{equation} where $\lambda$ is the wavelength, $n_e$[cm$^{-3}$] is the electron density, $\bf{B}[\mu$G] is the magnetic field, d$\bf{l}$[pc] is the path length element. Consequently, measuring the polarisation angle at several wavelengths for a given source can provide a simple determination of the rotation measure for that line of sight. Compact polarised radio sources within the Galaxy (pulsars) and outside the Galaxy (extragalactic sources or EGS) subsequently act as line-of-sight probes for the magnetic field; the higher the projected spatial density of the observed probes, the easier it is to determine the field structure in the given region. Recent work has focused on developing and testing competing models and on determining the existence of large scale reversals in the magnetic field. Magnetic field reversals occur where the magnetic field direction completely reverses over a short change in radius and/or azimuth within the disk of the Galaxy. The number of reversals depends on the interpretation of the existing RM data and is presently a very controversial subject \citep[e.g.][]{bt01, Weisberg04, Han06,Brown07, Sun08, Vallee08, Men08, Ronnie, Katgert}. Most models are made to follow the spiral arm structure of the Galaxy since an approximate alignment of the regular magnetic fields and spiral arms is commonly observed in external galaxies \citep[e.g.][]{Fletcher}. For all of these models, sufficient numbers of low-latitude, high quality RM data in key regions have been lacking.

We have produced a catalog of 221 rotation measures of extragalactic sources to fill in critical gaps in the disk rotation measure coverage of the Canadian Galactic Plane Survey and the Southern Galactic Plane Survey. We then used these data, in conjunction with previously observed rotation measures, to propose a new magnetic field model, stemming from a new modeling strategy that studies the disk field in three different sectors. The division of sectors is roughly between the outer Galaxy (quadrants 2 and 3), and the two `inner Galaxy' quadrants: quadrant 1 and quadrant 4. Our modeling suggests that the outer Galaxy is dominated by an (almost) purely azimuthal field, whereas in the inner Galaxy magnetic lines have a spiral shape, and are likely to be aligned with the spiral arms. Furthermore, the model seems to indicate that the magnetic field in the Galaxy is predominantly clockwise, with a single reversed region that appears to spiral out from the center of the Galaxy. Additional observations, and more modeling work are needed to confirm our results. However, we believe that this multi-sector approach to empirical modeling yields a more reasonable model than attempting to model the entire Galaxy with the current status of both the electron density and available data.

1012.2936

Interstellar magnetic fields play critical roles in many astrophysical processes. Yet despite their importance, our knowledge about magnetic fields in our Galaxy remains limited. For the field within the Milky Way, much of what we do know comes from observations of polarisation and Faraday rotation measures (RMs) of extragalactic sources and pulsars. A high angular density of RM measurements in several critical areas of the Galaxy is needed to clarify the Galactic magnetic field structure. Using observations made with the VLA, we have determined RMs for sources in regions of the Galactic plane not covered by the Canadian Galactic Plane Survey (CGPS) and Southern Galactic Plane Survey (SGPS). We have combined these new RMs with those determined from the CGPS and SGPS and have produced a new model for the magnetic field of the Galactic disk.