ReadingTimeMachine/rtm-sgt-ocr-v1

1source,target2 The Parkes radio telescope is part of the Australia Telescope.which is funded by the Conunomvealth of Australia for operation as a National Facility managed by CSIRO.," The Parkes radio telescope is part of the Australia Telescope,which is funded by the Commonwealth of Australia for operation as a National Facility managed by CSIRO."3"luminosity distance d;(z) depends on the geometry of the universe, i.e. the sign of Q,, and is given by where and Dark energy parameterization schemes enter through f(z).","luminosity distance $d_L(z)$ depends on the geometry of the universe, i.e. the sign of $\Omega_k$, and is given by where and Dark energy parameterization schemes enter through $f(z)$."4" For the case where EOS is piecewise constant in redshift, f(z) can be rewritten as (?) where w; is the EOS parameter in the i redshift bin defined by an upper boundary at z;, and the zeroth bin is defined as zo=0."," For the case where EOS is piecewise constant in redshift, $f(z)$ can be rewritten as \citep{Sullivan:2007pd}5 where $w_i$ is the EOS parameter in the $i^{\mathrm{th}}$ redshift bin defined by an upper boundary at $z_i$ , and the zeroth bin is defined as $z_0=0$."6" In order to compare with previous analysis (?),, we define the first three redshift bins to be the same as those used by ? by setting z;=0.2, 2=0.5, and z=1.8."," In order to compare with previous analysis \citep{Sullivan:2007pd}, we define the first three redshift bins to be the same as those used by \citet{Sullivan:2007pd} by setting $z_1=0.2$, $z_2=0.5$, and $z_3=1.8$."7 The fourth bin is defined by z4=7 to include GRBs., The fourth bin is defined by $z_4=7$ to include GRBs.8" We carry out our analyses under two different assumptions about the high redshift (redshift greater than zg=7 in our case) behavior of dark energy, i.e. the so-called (see?) “weak” prior, which makes no assumptions about w(z) at z>7 and the “strong” prior, which assumes w(z)=-1 at z> 7."," We carry out our analyses under two different assumptions about the high redshift (redshift greater than $z_4=7$ in our case) behavior of dark energy, i.e. the so-called \citep[see][]{Riess:2006fw}9 “weak” prior, which makes no assumptions about $w(z)$ at $z>7$ and the “strong” prior, which assumes $w(z)=-1$ at $z>7$ ."10 In this paper we adopt x? statistic to estimate parameters., In this paper we adopt $\chi^2$ statistic to estimate parameters.11" For a physical quantity € with experimentally measured value €,, standard deviation σε,and theoretically predicted value &,(0), where 0 is a collection of parameters needed to calculate the theoretical value, the Y? value is given by and the total y? is the sum of all X258. ie. The likelihood function is then proportional to expwhich produces the posterior probability when (2X6)multiplied2), by the prior probability of 0."," For a physical quantity $\xi$ with experimentally measured value $\xi_o$, standard deviation $\sigma_{\xi}$,and theoretically predicted value $\xi_t(\theta)$ , where $\theta$ is a collection of parameters needed to calculate the theoretical value, the $\chi^2$ value is given by and the total $\chi^2$ is the sum of all $\chi_{\xi}^2$ s, i.e. The likelihood function is then proportional to $\exp\left(-\chi^2(\theta)/2\right)$,which produces the posterior probability when multiplied by the prior probability of $\theta$."12" In the case of our analysis, the calculation of x?s for different observational data is described in section 3.."," In the case of our analysis, the calculation of $\chi^2$ s for different observational data is described in section \ref{sec:observational_data}."13" According to the posterior probability derived in this way, Markov chains are generated through the Monte-Carlo algorithm to study the statistical properties of the parameters."," According to the posterior probability derived in this way, Markov chains are generated through the Monte-Carlo algorithm to study the statistical properties of the parameters."14" In this paper, we focus on the EOS parameters by marginalizing the others."," In this paper, we focus on the EOS parameters by marginalizing the others."15" As mentioned above, in the process of constraining cosmological parameters, standard candles play this roleby providing the luminosity distances at certain redshifts."," As mentioned above, in the process of constraining cosmological parameters, standard candles play this roleby providing the luminosity distances at certain redshifts."16" However, the luminosity distance depends on the integration of the behavior of the dark energy over redshift, so the estimates of the dark energy EOS parameters w; at high redshift depend on those at low redshift."," However, the luminosity distance depends on the integration of the behavior of the dark energy over redshift, so the estimates of the dark energy EOS parameters $w_i$ at high redshift depend on those at low redshift."17" In other words, the EOS parameters w; are correlated in the sense thatthe covariance matrix, is not diagonal."," In other words, the EOS parameters $w_i$ are correlated in the sense thatthe covariance matrix, is not diagonal."18" In the above equation, the w is a vector with components w; and the average is calculated by letting w run over the Markov chain."," In the above equation, the $\boldsymbol{w}$ is a vector with components $w_i$ and the average is calculated by letting $\boldsymbol{w}$ run over the Markov chain."19 We can obtain a set of decorrelated parameters w; through diagonalization of the covariance matrix by choosing an appropriate transformation There can be different choices for T., We can obtain a set of decorrelated parameters $\widetilde{w}_i$ through diagonalization of the covariance matrix by choosing an appropriate transformation There can be different choices for $\boldsymbol{T}$.20 In this paper we use the transformation advocated by ? (see below)., In this paper we use the transformation advocated by \citet{Huterer:2004ch} (see below).21" First we define the Fisher matrix and then the transformation matrix T is given by except that the rows of the matrix T are normalized such that The advantage of this transformation is that the weights (rows of T) are positive almost everywhere and localized in redshift fairly well, so the uncorrelated EOS parameters w; are easy to interpret intuitively (?).."," First we define the Fisher matrix and then the transformation matrix $\boldsymbol{T}$ is given by except that the rows of the matrix $\boldsymbol{T}$ are normalized such that The advantage of this transformation is that the weights (rows of $\boldsymbol{T}$ ) are positive almost everywhere and localized in redshift fairly well, so the uncorrelated EOS parameters $\widetilde{w}_i$ are easy to interpret intuitively \citep{Huterer:2004ch}."22" To constrain the dark energy EOS, we have made useof observational data described below."," To constrain the dark energy EOS, we have made useof observational data described below."23" Recently compiled SN Ia data (???) include 45nearby supernovae (???),, 60 ESSENCE supernovae (?),, 57 SNLS supernovae (?),, and 30 HST supernovae (?).."," Recently compiled SN Ia data \citep{Riess:2006fw,WoodVasey:2007jb,Davis:2007na} include 45nearby supernovae \citep{Hamuy:1996su, Riess:1998dv, Jha:2005jg}, 60 ESSENCE supernovae \citep{WoodVasey:2007jb}, , 57 SNLS supernovae \citep{Astier:2005qq}, , and 30 HST supernovae \citep{Riess:2006fw}. ."24 Figure 1 shows the distribution of these SN Ia samples versus redshift., Figure \ref{fig:SNIa_distr} shows the distribution of these SN Ia samples versus redshift.25αἱ the inner edge of the eric begin to interact dynamically.,at the inner edge of the grid begin to interact dynamically.26 A wave of strong dvnamieal interactions then moves out through the grid., A wave of strong dynamical interactions then moves out through the grid.27 It takes ~ 1 Myr for the wave to move from ~ 0.85 AU to e 1.15 AU., It takes $\sim$ 1 Myr for the wave to move from $\sim$ 0.85 AU to $\sim$ 1.15 AU.28 During this period. some oligarchs merge.," During this period, some oligarchs merge."29 Others migrate through the grid on highly eccentricity orbits., Others migrate through the grid on highly eccentricity orbits.30" From ~ 1 Myr onward. mergers slowly reduce AN,."," From $\sim$ 1 Myr onward, mergers slowly reduce $N_o$."31 It takes ~ 1 Myr lor the first 2 mergers and another ~ 2 Alyy for the second 2 mergers., It takes $\sim$ 1 Myr for the first 2 mergers and another $\sim$ 2 Myr for the second 2 mergers.32 After LOO AIve. only 3 oligarchs remain.," After 100 Myr, only 3 oligarchs remain."33 One of these has m~ 0.42 m.2a ~ 0.9 AU. and e& 0.08.," One of these has $m \sim$ 0.43 $m_{\oplus}$, $a \sim$ 0.9 AU, and $e \sim$ 0.08."34 The other two oligarchs have m.e 0.05 mi. and e 0.1 (Figure 7)., The other two oligarchs have $m \sim$ 0.05 $m_{\oplus}$ and $e \sim$ 0.1 (Figure 7).35 Aside from the eccentricity of the more massive planet. the properties of (hese objects are reasonably close to those of the Earth ancl Mars.," Aside from the eccentricity of the more massive planet, the properties of these objects are reasonably close to those of the Earth and Mars."36 Fragmentation and interactions with the gas probably promote smaller eccentricities for the largest objects (e.g..Wetherill&Stewart1993:AenorWard2002 ).," Fragmentation and interactions with the gas probably promote smaller eccentricities for the largest objects \citep[e.g.,][]{ws93,agn02,kom02}."37".Figure 7 illustrates the evolution of the semimajor axes of the oligarchs for three calculations with Mj = 2 g 7 (lower panel). Xj = 4 g ? (middle panel). and. X, =S8eem ? (upper panel)."," Figure 7 illustrates the evolution of the semimajor axes of the oligarchs for three calculations with $\Sigma_0$ = 2 g $^{-2}$ (lower panel), $\Sigma_0$ = 4 g $^{-2}$ (middle panel), and $\Sigma_0$ = 8 g $^{-2}$ (upper panel)."38 As in Figure IL. the (racks change color when two oligarchs collide and merge.," As in Figure 1, the tracks change color when two oligarchs collide and merge."39 The labels indicate the final mass. in Earth masses. of the largest oligarchs at 100 Myr.," The labels indicate the final mass, in Earth masses, of the largest oligarchs at 100 Myr."40 The Gmescale for the onset of chaotic growth depends on (he initial surface density., The timescale for the onset of chaotic growth depends on the initial surface density.41 More massive disks reach the transition first. (see.forexample.Lissaner1987).," More massive disks reach the transition first. \citep[see, for example,][]{lis87}."42". For M, = Secm 7 the transition begins al ~ a lew x10? vr."," For $\Sigma_0$ = 8 g $^{-2}$, the transition begins at $\sim$ a few $\times ~ 10^5$ yr."43" The transition is delaved to ~ 1 Myr for. E ""T M22gcem2-.", The transition is delayed to $\sim$ 1 Myr for $\Sigma_0$ = 2 g $^{-2}$.44 The character of the transition to chaotic growth also depends on (he initial surface densitv., The character of the transition to chaotic growth also depends on the initial surface density.45 In relatively massive disks with Xj ~ 8 g 7. many oligarchs develop highly eccentric orbits and exhibit large variations in their semimajor axes.," In relatively massive disks with $\Sigma_0$ $\sim$ 8 g $^{-2}$ , many oligarchs develop highly eccentric orbits and exhibit large variations in their semimajor axes."46 These large excursions result in many mergers and a rapid reduction in AN., These large excursions result in many mergers and a rapid reduction in $N_o$.47" In less massive disks with X,< 24 ο em7. only I or 2 oligarchs develop highly eccentric orbits."," In less massive disks with $\Sigma_0 \lesssim$ 2–4 g $^{-2}$, only 1 or 2 oligarchs develop highly eccentric orbits."48 Most mergers are caused by two-body interactions. instead of large-scale dynamical interactions throughout the grid.," Most mergers are caused by two-body interactions, instead of large-scale dynamical interactions throughout the grid."49 Figures 8I0 illustrate (hese general conclusions., Figures 8–10 illustrate these general conclusions.50 In Figure 8. (he orbit crossing paraimeler rapidly approaches zero for caleulations with Xj = 8 g 7.," In Figure 8, the orbit crossing parameter rapidly approaches zero for calculations with $\Sigma_0$ = 8 g $^{-2}$."51" At 0.11 Myr. p; has a long plateau: close approaches between oligarchs cause p, lo fall below zero: mergers cause p, (o jump above zero."," At 0.1–1 Myr, $p_o$ has a long plateau; close approaches between oligarchs cause $p_o$ to fall below zero; mergers cause $p_o$ to jump above zero."52" During a series of 4 mergers αἱ 10 Myr. p, remains below 0 [or a long period."," During a series of 4 mergers at 10 Myr, $p_o$ remains below 0 for a long period."53" After the final merger. p, jumps to 40. where it remains for many Myr."," After the final merger, $p_o$ jumps to 40, where it remains for many Myr."54 For smaller Mua. Po Yeluains well above zero until one or (wo close pairwise interactions pushes it below zero.," For smaller $\Sigma_0$ , $p_o$ remains well above zero until one or two close pairwise interactions pushes it below zero."55" Once(hese interactions produce a merger. (he svstems stabilize and p, moves. well above zero."," Oncethese interactions produce a merger, the systems stabilize and $p_o$ moves well above zero."56"plane is populated with an array of fiber positioners, distributed in a hexagonal pattern, so that a single device is used to position each fiber head in the desired location within its patrol disc.","plane is populated with an array of fiber positioners, distributed in a hexagonal pattern, so that a single device is used to position each fiber head in the desired location within its patrol disc."57 Each positioner is therefore devoted to observing a single target., Each positioner is therefore devoted to observing a single target.58" In a few words, the fiber positioning robot is a collection of positioners, all identical, distributed over an array which covers the entire focal plane."," In a few words, the fiber positioning robot is a collection of positioners, all identical, distributed over an array which covers the entire focal plane."59 The focal plane is therefore covered by these patrol discs so that all possible positions can be reached by at least one positioner (see Fig. 1))., The focal plane is therefore covered by these patrol discs so that all possible positions can be reached by at least one positioner (see Fig. \ref{fig:focalplane}) ).60" In order to cover the whole focal plane, a certain degree of overlap between patrol discs is needed so some regions of the focal plane can be reached by more than one positioner (two or maximum three), hence rising the possibility of fiber collisions."," In order to cover the whole focal plane, a certain degree of overlap between patrol discs is needed so some regions of the focal plane can be reached by more than one positioner (two or maximum three), hence rising the possibility of fiber collisions."61 In Fig., In Fig.62" 2 we show a view of a real subset of a fiber positioner robot, with 19 positioners in hexagonal pattern (?).."," \ref{fig:19_positioners} we show a view of a real subset of a fiber positioner robot, with $19$ positioners in hexagonal pattern \citep{Azzaro2010}."63 This concept offers a number of advantages as compared to others., This concept offers a number of advantages as compared to others.64" It is robust, scalable and easy to service and maintain (failure of one positioner causes the loss of one target only)."," It is robust, scalable and easy to service and maintain (failure of one positioner causes the loss of one target only)."65" Operationally, it provides extremely short reconfiguration times and allows an efficient real-time correction of differential atmospheric dispersion."," Operationally, it provides extremely short reconfiguration times and allows an efficient real-time correction of differential atmospheric dispersion."66" It is, however, specifically conceived for wide-field surveys."," It is, however, specifically conceived for wide-field surveys."67 The reason is that positioners cannot be densely packed onto a small portion of the focal plane., The reason is that positioners cannot be densely packed onto a small portion of the focal plane.68" Consequently, the system is efficient for rather uniform distributions of targets or, alternatively, for large areas where any possible bias is smoothed."," Consequently, the system is efficient for rather uniform distributions of targets or, alternatively, for large areas where any possible bias is smoothed."69 Find a complete discussion on this fiber positioning concept in ?.., Find a complete discussion on this fiber positioning concept in \cite{Azzaro2010}.70" The results presented in this work on the optimization of the fiber assignment process are based on a complete simulation of the focal plane of the SIDE spectrograph, which adopted a fiber positioning robot as that described above."," The results presented in this work on the optimization of the fiber assignment process are based on a complete simulation of the focal plane of the SIDE spectrograph, which adopted a fiber positioning robot as that described above."71 SIDE is an example of a state-of-the-art fiber-fed instrument capable of efficiently undergoing next-generation large-scale spectroscopic surveys., SIDE is an example of a state-of-the-art fiber-fed instrument capable of efficiently undergoing next-generation large-scale spectroscopic surveys.72 The focal plane of the SIDE spectrograph was designed as an array of 1003 positioners covering the 20-arcmin field of view at GTC., The focal plane of the SIDE spectrograph was designed as an array of 1003 positioners covering the 20-arcmin field of view at GTC.73" Important for this work, the results obtained with this simulation are valid for any focal plane consisting in an array of positioners as that outlined in this section."," Important for this work, the results obtained with this simulation are valid for any focal plane consisting in an array of positioners as that outlined in this section."74" The key parameter to describe the efficiency of the fiber assignment process, as we will see below, is the target-to-positioner ratio, η (instead of other parameters such as the size of the focal plane or the number of positioners)."," The key parameter to describe the efficiency of the fiber assignment process, as we will see below, is the target-to-positioner ratio, $\eta$ (instead of other parameters such as the size of the focal plane or the number of positioners)."75 In order to better illustrate our results we will also discuss another real-life example: BigBOSS., In order to better illustrate our results we will also discuss another real-life example: BigBOSS.76 In Table 1 we list some of the relevant parameters for both SIDE and In this section we briefly describe the philosophy behind our optimized fiber positioning algorithm: the draining algorithm., In Table \ref{tab:BB_SIDE} we list some of the relevant parameters for both SIDE and In this section we briefly describe the philosophy behind our optimized fiber positioning algorithm: the draining algorithm.77 We also discuss its performance as compared to a simple random approach., We also discuss its performance as compared to a simple random approach.78 Let us consider a generic instrument with a fiber positioning robot like that described in the previous section covering the entire focal plane with N positioners., Let us consider a generic instrument with a fiber positioning robot like that described in the previous section covering the entire focal plane with $N$ positioners.79" A set of targets is used to populate the focal plane according to a given target-to-positioner ratio, 7."," A set of targets is used to populate the focal plane according to a given target-to-positioner ratio, $\eta$."80 We will first assume that targets are randomly distributed in the focal plane., We will first assume that targets are randomly distributed in the focal plane.81" The design of the system considered implies that each target is reachable by one, two or maximum three"," The design of the system considered implies that each target is reachable by one, two or maximum three"82Since the discovery of the high-redshift Lava forest over 25 vears ago. these abundant absorption features in the spectra of QSOs have been used as evolutionary probes of the intergalactic medimm (IGMD). galactic halos. and now larec-scale structure and chemical evolution.,"Since the discovery of the high-redshift $\alpha$ forest over 25 years ago, these abundant absorption features in the spectra of QSOs have been used as evolutionary probes of the intergalactic medium (IGM), galactic halos, and now large-scale structure and chemical evolution."83 The rapid evolution iu the distribution of lines per uuit redshift. ΑλάνX(1τὸ (552.5 for c> 1.6) was consistent with a picture of these features as highly ionized “clouds” whose munbers aud sizes were controlled by the evolution of the IGM pressure. the mctagalactic ionizing radiation field. aud galaxy formation.," The rapid evolution in the distribution of lines per unit redshift, $d{\cal N}/dz \propto (1+z)^{\gamma}$ $\gamma \approx 2.5$ for $z \geq 1.6$ ), was consistent with a picture of these features as highly ionized “clouds” whose numbers and sizes were controlled by the evolution of the IGM pressure, the metagalactic ionizing radiation field, and galaxy formation."84 Early observations also suggested that Lya clouds had characteristic sizes 10 kpe. were much more abundant than (£.) galaxies. aud showed Little clusteriug iu velocity space.," Early observations also suggested that $\alpha$ clouds had characteristic sizes $\sim10$ kpc, were much more abundant than $L_*$ ) galaxies, and showed little clustering in velocity space."85 They were interpreted as pristine. zeroauctallicitv eas left over frou the recombination cra.," They were interpreted as pristine, zero-metallicity gas left over from the recombination era."86 One therefore expected low-redshift (2< 1) absorption clouds to show only traces of IT 1. due to photoionization aud evaporation in a lower pressure IGAL," One therefore expected low-redshift $z < 1$ ) absorption clouds to show only traces of H I, due to photoionization and evaporation in a lower pressure IGM."87 All these ideas have now changed with new data., All these ideas have now changed with new data.88 One of the delightful spectroscopic surprises from the (UST) was the discovery of Lava absorption lines toward the quasar 3€ 273 at ten0458 by both the Faint Object Spectrograph (FOS. (Baheall et al.," One of the delightful spectroscopic surprises from the (HST) was the discovery of $\alpha$ absorption lines toward the quasar 3C 273 at $z_{\rm em} = 0.158$ by both the Faint Object Spectrograph (FOS, (Bahcall et al."89 1991) and the Goddard Tigh Resolution Spectrograph (GITRS. Morris et al.," 1991) and the Goddard High Resolution Spectrograph (GHRS, Morris et al."90 1991)., 1991).91 Iu this review. I will describe (822) the current status of our eroups program with the IIST aud VLA to define the parameters and nature of the low-redshift Lya forest.," In this review, I will describe 2) the current status of our group's long-term program with the HST and VLA to define the parameters and nature of the low-redshift $\alpha$ forest."92 In 8323. I discuss related theoretical work ou the," In 3, I discuss related theoretical work on the"93solution and is seen to agree very well with the established codes.,solution and is seen to agree very well with the established codes.94" In our 3D code, two points may be at the same distance from the center, but separated by a large angle and they may not be radiatively connected, especially not in high opacity models."," In our 3D code, two points may be at the same distance from the center, but separated by a large angle and they may not be radiatively connected, especially not in high opacity models."95 These grid points may not see exactly the same radiation field because of the random nature of the grid and photon transport and therefore the level populations may not necessarily be exactly the same., These grid points may not see exactly the same radiation field because of the random nature of the grid and photon transport and therefore the level populations may not necessarily be exactly the same.96 There are no fluctuations in the results of the 1D codes because only a single cell with a single solution exist at a given radius., There are no fluctuations in the results of the 1D codes because only a single cell with a single solution exist at a given radius.97" If we take the solution to be the expected solution we can calculate reduced X? values, =>for cach of the 6 levels“of Pesta) the solution which gives us X20.59(1.08,1.11,1.02,1.17,1.00,1.00] for the optically thin case and Xo.s=(1.06,3.12,1.03,2.08,1.28,1.15} for the optically thick case."," If we take the solution to be the expected solution we can calculate reduced $\chi^2$ values, for each of the 6 levels of the solution which gives us $\chi^2_{l=0...5}\approx\{1.08, 1.11, 1.02, 1.17, 1.00, 1.00\}$ for the optically thin case and $\chi^2_{l=0...5}\approx\{1.06, 3.12, 1.03, 2.08, 1.28, 1.15\}$ for the optically thick case."98" The comparison is slightly worse for the optically thick case, but here we also see greater variation between the established codes."," The comparison is slightly worse for the optically thick case, but here we also see greater variation between the established codes."99 The example we present is a typical 2D hydrostatic protoplanetary disk model with a cold and dense mid-plane., The example we present is a typical 2D hydrostatic protoplanetary disk model with a cold and dense mid-plane.100 Such models are numerously found in the literature but here we use a simple analytic toy model for illustrative purposes.," Such models are numerously found in the literature \citep[e.g.,][]{chiang1997, dullemond2004, robitaille2006} but here we use a simple analytic toy model for illustrative purposes."101" The density structure is given by where We consider HCO*, H5O, and CH30H gas ata fractional abundance of 2x10? with respect to the Η2 density."," The density structure is given by where We consider $^+$, $_2$ O, and $_3$ OH gas ata fractional abundance of $\times 10^{-9}$ with respect to the $_2$ density."102" With these three species we illustrate the possibility of utilizing a large dynamic range in scales (HCO*), very opaque models (H20), and multiple overlapping lines (CH3OH)."," With these three species we illustrate the possibility of utilizing a large dynamic range in scales $^+$ ), very opaque models $_2$ O), and multiple overlapping lines $_3$ OH)."103 The temperature is given by a power-law In a more realistic model the temperature would be calculated self-consistently based on the radiation properties of the central source and the temperature would drop toward the mid-plane because this region would be shielded from stellar radiation by the upper layers of the disk., The temperature is given by a power-law In a more realistic model the temperature would be calculated self-consistently based on the radiation properties of the central source and the temperature would drop toward the mid-plane because this region would be shielded from stellar radiation by the upper layers of the disk.104" In our example we mimic this effect by lowering the temperature in a wedge shaped region around the mid-plane to 20 K. By letting water freeze out at temperatures below 90 K, we can simulate a complex abundance structure often used in protoplanetary disks (??).."," In our example we mimic this effect by lowering the temperature in a wedge shaped region around the mid-plane to 20 K. By letting water freeze out at temperatures below 90 K, we can simulate a complex abundance structure often used in protoplanetary disks \citep{jonkheid2007,woitke2009}."105 The disk extends to 500 AU and the values for no and Τρ are 10$ cm? and 90 K at the radius of 100 AU., The disk extends to 500 AU and the values for $n_0$ and $T_0$ are $10^8$ $^{-3}$ and 90 K at the radius of 100 AU.106 In addition we have added a 2 AU wide gap around a radius of 5 AU from the center., In addition we have added a 2 AU wide gap around a radius of 5 AU from the center.107 The disk is in Keplerian rotation., The disk is in Keplerian rotation.108" Figure shows the density, temperature, and H2O density of our disk."," Figure \ref{diskmod} shows the density, temperature, and $_2$ O density of our disk."109" Of other parameters that describes the disk are the turbulent velocity dispersion set to 150 ms-!, stellar mass of 1 Mo, and a gas-to-dust ratio of 100."," Of other parameters that describes the disk are the turbulent velocity dispersion set to 150 $^{-1}$, stellar mass of 1 $_\odot$ , and a gas-to-dust ratio of 100."110 We use thin mantled grains with 107 years of coagulation and the resulting disk mass is 0.02 Mo., We use thin mantled grains with $^7$ years of coagulation and the resulting disk mass is 0.02 $_\odot$.111" To break the azimuthal symmetry and make the model fully 3D, we have placed a protoplanetary condensation in the gap."," To break the azimuthal symmetry and make the model fully 3D, we have placed a protoplanetary condensation in the gap."112 The protoplanet has the same qualitative properties as the one described in ?.., The protoplanet has the same qualitative properties as the one described in \citet{narayanan2006}. .113 The protoplanet has been modeled by placing a spherical distribution of gridpoints at the desired spot and, The protoplanet has been modeled by placing a spherical distribution of gridpoints at the desired spot and114aabundance pattern.,abundance pattern.115 NUV spectra of aand wwere obtained using the Space Telescope Imaging Spectrograph (STIS) on the (HST))., NUV spectra of and were obtained using the Space Telescope Imaging Spectrograph (STIS) on the ).116" These spectra cover a wavelength region from aat R=A/AA~ 30,000."," These spectra cover a wavelength region from at $R \equiv \lambda/\Delta\lambda \sim$ 30,000."117" The optical spectrum of wwas obtained using the High Resolution Echelle Spectrograph (HIRES; Vogtetal. 1994)) on Keck I, and this spectrum covers a wavelength region from aat R~ 45,000."," The optical spectrum of was obtained using the High Resolution Echelle Spectrograph (HIRES; \citealt{vogt94}) ) on Keck I, and this spectrum covers a wavelength region from at $R \sim$ 45,000."118 See Cowanetal.(2005) for further details., See \citet{cowan05} for further details.119" In Figure 1, we show segments of the STIS spectra surrounding the Os transition at aand the the Cd transition at iin aand122563,, as well as A strong absorption feature is clearly identified at these wavelengths in bbut not in122563."," In Figure \ref{overplot}, we show segments of the STIS spectra surrounding the Os transition at and the the Cd transition at in and, as well as A strong absorption feature is clearly identified at these wavelengths in but not in."120. iis warmer (Το= 5200 K) and more metal-rich ([Fe/H] = —2.1) than (Te= 4570 K and [Fe/H] = —2.7)., is warmer $T_{\rm eff} =$ 5200 K) and more metal-rich ([Fe/H] $= -$ 2.1) than $T_{\rm eff} =$ 4570 K and [Fe/H] $= -$ 2.7).121" has a temperature (Το= 4720 K), metallicity ([Fe/H] = —2.9), and overall light element abundance distribution (i.e., 6 €Z 40) that closely resembles citepwestin00.."," has a temperature $T_{\rm eff} =$ 4720 K), metallicity ([Fe/H] $= -$ 2.9), and overall light element abundance distribution (i.e., 6 $\leq Z \leq$ 40) that closely resembles \\citep{westin00}."122 is overabundant in the heavy eelements ([Eu/Fe] = +0.7) relative to (([Eu/Fe] = —0.5)., is overabundant in the heavy elements ([Eu/Fe] $= +$ 0.7) relative to ([Eu/Fe] $= -$ 0.5).123 Therefore the only significant difference between the spectra of and sshould be the stronger heavy aabsorption lines in115444., Therefore the only significant difference between the spectra of and should be the stronger heavy absorption lines in.124". In Figure 1, we see that115444,, like3248,, also exhibits strong absorption features at 2282.28 and2288.02A,, but ddoes not."," In Figure \ref{overplot}, we see that, like, also exhibits strong absorption features at 2282.28 and, but does not."125 Thus heavy sspecies must be producing this absorption., Thus heavy species must be producing this absorption.126 We find no transitions of heavy sspecies at these wavelengths—or the Lu line at 2615.41À-—in the Kurucz or NIST line databases that could plausibly account for this absorption other than the species of interest., We find no transitions of heavy species at these wavelengths—or the Lu line at —in the Kurucz or NIST line databases that could plausibly account for this absorption other than the species of interest.127 References for published transition probabilities of the lines used in this analysis are given in Table 1.., References for published transition probabilities of the lines used in this analysis are given in Table \ref{abundtab}.128 We determined the transition probability of the Lu rresonance line to be log(gf) = +0.11 + 0.04 based on a laser-induced fluorescence lifetime measurement of its upper level (Fedchaketal.2000) and a branching fraction calculation of 0.971 (Quinetetal.1999).. (, We determined the transition probability of the Lu resonance line to be $gf$ ) $ = +$ 0.11 $\pm$ 0.04 based on a laser-induced fluorescence lifetime measurement of its upper level \citep{fedchak00} and a branching fraction calculation of 0.971 \citep{quinet99}. (129See also Lawleretal. 2009..),See also \citealt{lawler09}. .)130 The isotope is dominant of S. LLu; Lodders2003))., The isotope is dominant of S. Lu; \citealt{lodders03}) ).131" The isotope is blocked from production by the stable isotope, so it is expected to be entirely absent from 3248."," The isotope is blocked from production by the stable isotope, so it is expected to be entirely absent from ."132. The odd-Z isotope!?Lu has nonzero nuclear spin I= 7/2., The $Z$ isotope has nonzero nuclear spin $I =$ 7/2.133" Hyperfine structure (hfs) and an accurate line position are based on new laboratory measurements of the 6s6p 1Ρ0 level energy, 38223.406(8) cm"", hfs A, —0.03731(10) em~!, and hfs B, 0.0811(15) απ’."," Hyperfine structure (hfs) and an accurate line position are based on new laboratory measurements of the 6s6p $^{1}$ $^{0}$ level energy, 38223.406(8) $^{-1}$, hfs A, $-$ 0.03731(10) $^{-1}$, and hfs B, 0.0811(15) $^{-1}$."134" The naturally occurring iisotopes of Cd and Os are predominantly even-Z N isotopes with zero nuclear spin, thus we are justified in ignoring the hfs from their minority isotopes."," The naturally occurring isotopes of Cd and Os are predominantly $Z$ $N$ isotopes with zero nuclear spin, thus we are justified in ignoring the hfs from their minority isotopes."135 We use the current version of the LTE spectral analysis code MOOG (Sneden1973) to perform the abundance analysis., We use the current version of the LTE spectral analysis code MOOG \citep{sneden73} to perform the abundance analysis.136" We adopt the atmospheric parameters for aand dderived by Cowanetal.(2002) and Simmereretal.(2004) (T.g/log g/[M/H]/w= 5200 aand 4570 K/1.80/—2.08/1.9K/1.35/—2.50/2.9s~1,, respectively) and interpolate model atmospheres from the Kurucz grids (Castellietal.1997)."," We adopt the atmospheric parameters for and derived by \citet{cowan02} and \citet{simmerer04}137 $T_{\rm eff}$ /log $g$ $v_{t} =$ 5200 $-$ 2.08/1.9 and 4570 $-$ 2.50/2.9, respectively) and interpolate model atmospheres from the Kurucz grids \citep{castelli97}."138. We compare our results to abundances of other species derived from lines in the optical spectral range., We compare our results to abundances of other species derived from lines in the optical spectral range.139" In the NUV, bound-free continuous opacity from metals may be comparable to or greater than the bound-free continuous opacity from H that dominates in the optical spectral range for metal-poor stars (e.g., Travis&Matsushima 1968))."," In the NUV, bound-free continuous opacity from metals may be comparable to or greater than the bound-free continuous opacity from $^{-}$ that dominates in the optical spectral range for metal-poor stars (e.g., \citealt{travis68}) )."140" To compensate for deficiencies in our ability to model the continuousopacity in this spectral range, we have derived abundances of relatively clean, unsaturated, and unblended Fe and Zr lines across the NUV."," To compensate for deficiencies in our ability to model the continuousopacity in this spectral range, we have derived abundances of relatively clean, unsaturated, and unblended Fe and Zr lines across the NUV."141 We require that these lines have reliable log(g/) values, We require that these lines have reliable $gf$ ) values142dusty circumstellar envelopes of evolved AGB stars may be the source of the bulk of M33's diffuse 8um and 24μπι emission.,dusty circumstellar envelopes of evolved AGB stars may be the source of the bulk of M33's diffuse $\mu$ m and $\mu$ m emission.143 The contribution of the TP-AGB to the K-band has been combined with with the mean IR colors of Galactic TP-AGB and stars in order to estimate the contributions of bothC youngM and old stellar populations to mid-IR observations of galaxies., The contribution of the TP-AGB to the $K$ -band has been combined with with the mean IR colors of Galactic TP-AGB C and M stars in order to estimate the contributions of both young and old stellar populations to mid-IR observations of galaxies.144" Without tuning, we find that the resulting mid-IR luminosities of the TP- can reproduce the MIPS 24m fluxes for galaxies back to at least 2= in a manner consistent with restframe optical colors."," Without tuning, we find that the resulting mid-IR luminosities of the TP-AGB can reproduce the MIPS $24\mu$ m fluxes for galaxies back to at least $z=2$ in a manner consistent with restframe optical colors."145 We have also tested the validity of the model on local scales in the galaxy M81 and find reasonable agreement., We have also tested the validity of the model on local scales in the galaxy M81 and find reasonable agreement.146" The origins of correlations between optical colors and mid-IR luminosities seen by others, such as Salimetal.(2009),, can now be understood."," The origins of correlations between optical colors and mid-IR luminosities seen by others, such as \cite{salim2009}, can now be understood."147" With careful modeling of SEDs from the UV through the mid-IR, more detailed histories of star formation should be possible."," With careful modeling of SEDs from the UV through the mid-IR, more detailed histories of star formation should be possible."148" Unfortunately, stellar spectral libraries and theoretical modeling are neither sufficient for verifying nor reducing the uncertainties our models (Conroyetal.2009)."," Unfortunately, stellar spectral libraries and theoretical modeling are neither sufficient for verifying nor reducing the uncertainties our models \citep{conroy2009}."149". This is largely due to the great difficulty in modeling post-main-sequence evolution, including the envelopes of TP-AGB stars, though the UV may provide further constraints2008)."," This is largely due to the great difficulty in modeling post-main-sequence evolution, including the envelopes of TP-AGB stars, though the UV may provide further constraints."150" We are optimistic that improved characterization of the mid-IR colors of the 'TP-AGB can be incorporated into SED fitting, though our calculations have uncertainties perhaps on the order of a factor of two due to uncertainties in the ensemble colors of the TP-AGB populations at different ages."," We are optimistic that improved characterization of the mid-IR colors of the TP-AGB can be incorporated into SED fitting, though our calculations have uncertainties perhaps on the order of a factor of two due to uncertainties in the ensemble colors of the TP-AGB populations at different ages."151 With refinement we anticipate that incorporating the mid-IR into multiwavelength analysis of SEDs will provide the strongest constraints on the star formation histories of galaxies., With refinement we anticipate that incorporating the mid-IR into multiwavelength analysis of SEDs will provide the strongest constraints on the star formation histories of galaxies.152" There is little doubt that star formation and the growth of stellar mass was occurring more rapidly in the distant universe than today, but the nature of that growth has remained largely unknown."," There is little doubt that star formation and the growth of stellar mass was occurring more rapidly in the distant universe than today, but the nature of that growth has remained largely unknown."153 Earlier results (e.g.LeFloc’hetal.2005) had implied that ~1/3 of the stellar mass at the present epoch was formed after z=] — a result that appears to be at odds with the evolution in the stellar mass function to z—1 (e.g.Cirasuoloetal. 2007)., Earlier results \citep[e.g.][]{lefloch2005} had implied that $\sim 1/3$ of the stellar mass at the present epoch was formed after $z=1$ — a result that appears to be at odds with the evolution in the stellar mass function to $z=1$ \citep[e.g.][]{cirasuolo2007}.154. But the model presented here implies that the mid-IR provides the total mass in stars formed in windows stretching back 1.5 Gyr in cosmic time., But the model presented here implies that the mid-IR provides the total mass in stars formed in windows stretching back 1.5 Gyr in cosmic time.155" As a result, such observations must be used with care when constraining the star formation rate density of the universe at z« 2, or when considering whether variations in the initial mass function are warranted by the data (e.g.Davé2008;Wilkinsetal.2008)."," As a result, such observations must be used with care when constraining the star formation rate density of the universe at $z < 2$ , or when considering whether variations in the initial mass function are warranted by the data \citep[e.g.][]{dave2008,wilkins2008}."156. 'The detection of galaxies in the mid-infrared over most of a Hubble time has helped changeour view of galaxy, The detection of galaxies in the mid-infrared over most of a Hubble time has helped changeour view of galaxy157Ow work and data sharine-parallel tree (WDSILEPT) code is principally aimed at ruuniug LSS cosmological simmlations with a umber of particles as high as possible using supercomputers sucli as ταν T3E svstems.,Our work and data sharing-parallel tree (WDSH-PT) code is principally aimed at running LSS cosmological simulations with a number of particles as high as possible using supercomputers such as Cray T3E systems.158 In order to increase the code efficiency. we adopt initially the erouping iiethod proposed by Barues aud introduce a modified implementation of lis grouping policy vielding very high gains in the code performances with the same accuracy.," In order to increase the code efficiency, we adopt initially the grouping method proposed by Barnes and introduce a modified implementation of his grouping policy yielding very high gains in the code performances with the same accuracy."159 To compute the force ou a body. the BIL algorithm needs to build au interaction list (£L) for cach particle p.," To compute the force on a body, the BH algorithm needs to build an interaction list $IL$ ) for each particle $p$."160 Starting from the root cell. a tree inspection is done aud the opening angle parameter Ó is used to evaluate whether a cell ust be opened or closed as meutioned above.," Starting from the root cell, a tree inspection is done and the opening angle parameter $\theta$ is used to evaluate whether a cell must be opened or closed as mentioned above."161 If a cell las dimension C; aud distance d from the particle p so that Eq. (1)), If a cell has dimension $C_l$ and distance $d$ from the particle $p$ so that Eq. \ref{eq:un}) )162 is verified. the cell is closed. it is added to the ££. aud its subcells are not investigated further.," is verified, the cell is closed, it is added to the $IL$, and its subcells are not investigated further."163 Otherwise the cell is opened aud its subeells are investigated in the same wav., Otherwise the cell is opened and its subcells are investigated in the same way.164 Bodies belonging to au opened cell are added to the ZZ. Next. the force ou the body is computed using the monopole and quadrupole momenta for all the cells in the The tree inspection phase represeuts a sizeable task to compute the force because the cell opening criterion is applied may times for each particle.," Bodies belonging to an opened cell are added to the $IL$ Next, the force on the body is computed using the monopole and quadrupole momenta for all the cells in the The tree inspection phase represents a sizeable task to compute the force because the cell opening criterion is applied many times for each particle."165" The CPU time 2, to compute the force im a time-step for all the N particles is where (5j is the average time to build au JL aud (Ty) is the average time to compute the force on cach particle using the interaction", The CPU time $T_o$ to compute the force in a time-step for all the $N$ particles is where $\langle T_{l} \rangle$ is the average time to build an $IL$ and $\langle T_f \rangle$ is the average time to compute the force on each particle using the interaction166just a few cycles.,just a few cycles.167 Ideally. an image should ouly © conrpressed once. aud then all the subsequeut data analysis should be performed directlv on he tile-compressed FITS file.," Ideally, an image should only be compressed once, and then all the subsequent data analysis should be performed directly on the tile-compressed FITS file."168 If the software cannot read the compressed format directly. then it should operate on an uncompressed version of he original compressed fle. which then should rot be reconirpressed.," If the software cannot read the compressed format directly, then it should operate on an uncompressed version of the original compressed file, which then should not be recompressed."169 Tn order to make the nuage quautization aud colmpression techuiques that ave described in the previous sections more widely available to the astronomical conuuuuity. we have developed a pair : eeneral purpose utility programs. called. xd (Scamanctal.2007).. which can be used to compress and uuconpress any FITS nuage in integer or floating-point format.," In order to make the image quantization and compression techniques that are described in the previous sections more widely available to the astronomical community, we have developed a pair of general purpose utility programs, called and \citep{seaman2007}, which can be used to compress and uncompress any FITS image in integer or floating-point format."170 These utilities rely «i the underlying CFITSIO ibrary (Pence1999) to perform the quantization and colupression operations., These utilities rely on the underlying CFITSIO library \citep{pence1999} to perform the quantization and compression operations.171 Thefpack and utility programs were used in the experiments that are described in the following sectious to quantize id compress the images., The and utility programs were used in the experiments that are described in the following sections to quantize and compress the images.172 Further information yout andfunpachk is available frou the HEASARCweb siteat, Further information about and is available from the HEASARC web site at.173 Iu this section we present the results of experiments designed to show how the measurements ofobjects in an image are affected as the image is quantized by varviug degrees., In this section we present the results of experiments designed to show how the measurements of objects in an image are affected as the image is quantized by varying degrees.174 In particular. we will verity that he noise iu the inaege5 mereases as a fiction of q by the amount predicted bw equation 10.. and more siguificautly. that the statistical errors on the magnitude aud position measurcments of faint objects in an nuage. which are hnuited mainly by the backeround noise. also increase by a simular factor.," In particular, we will verify that the noise in the image increases as a function of q by the amount predicted by equation \ref{eq:fractionalnoise}, and more significantly, that the statistical errors on the magnitude and position measurements of faint objects in an image, which are limited mainly by the background noise, also increase by a similar factor."175 These results will provide ecucral euidelimes Του achieving the ercatest aout of mage compression while still preserviug the required level of scientific precision in the age., These results will provide general guidelines for achieving the greatest amount of image compression while still preserving the required level of scientific precision in the image.176 Section 3.1 describes the method of constructingc» the simulated CCD images that are used du the first 2 experiments., Section \ref{s:setup} describes the method of constructing the simulated CCD images that are used in the first 2 experiments.177 The first experiment. ii 83.2.. examines how quantization affects the uncertainties of measurements of single star inages. and the second experiuenut. iu 3.3.. examines the case where many quantized nuages are added together to detect sources far below the detection threshold of a single nuage.," The first experiment, in \ref{s:exp1}, examines how quantization affects the uncertainties of measurements of single star images, and the second experiment, in \ref{s:exp2}, examines the case where many quantized images are added together to detect sources far below the detection threshold of a single image."178" Finally. the third experiment. in §3.L.. is performed ou a set of actual astronomical images to verity the results obtained from the svuthetic Huages,"," Finally, the third experiment, in \ref{s:exp3}, is performed on a set of actual astronomical images to verify the results obtained from the synthetic images."179 Tn order to determine how quautization affects the precision of measurements of objects in an nuage. we eenecrated a large sample of realistic CCD star images with kuown iuput positious and maeguitudes.," In order to determine how quantization affects the precision of measurements of objects in an image, we generated a large sample of realistic CCD star images with known input positions and magnitudes."180 This allows us to xeciselv calculate the errors ou the measured positions and magnitudes in the quantized images., This allows us to precisely calculate the errors on the measured positions and magnitudes in the quantized images.181 All the stars have circular Gaussian profiles with o=1.0 aud FWIIM = 2.35 pixels., All the stars have circular Gaussian profiles with $\sigma = 1.0$ and FWHM = 2.35 pixels.182 This is typical of the spatial resolution couuuonlv found in astronomical CCD nuages and is adequate to avoid he difficulties when analyzing spatially undersampled iuages., This is typical of the spatial resolution commonly found in astronomical CCD images and is adequate to avoid the difficulties when analyzing spatially undersampled images.183 The central location of the star images. relative to the pixel grid. was varied so as to average out any subtle biases iu the subsequent star detectiou and measurement steps that might depend ou the //heasarc.gsfc.nasa.gov/fiusiosepaicw.," The central location of the star images, relative to the pixel grid, was varied so as to average out any subtle biases in the subsequent star detection and measurement steps that might depend on the exact position."184 The total integrated flux in the stars covered a range of LO magnitudes (a factor of 10000 in intensity) in 0.5 maenituce increments., The total integrated flux in the stars covered a range of 10 magnitudes (a factor of 10000 in intensity) in 0.5 magnitude increments.185 Finally. the skv backeround was simulated by adding 1000 counts ο cach pixel.," Finally, the sky background was simulated by adding 1000 counts to each pixel."186" Poissomau-cistributed shot noise aud Cassia distributed ""remd-out noise was then added to cach of these star ages to simulate real CCD nuages.", Poissonian-distributed shot noise and Gaussian distributed “read-out” noise was then added to each of these star images to simulate real CCD images.187" The shot noise iu each pixel was ταςοτι] calculated using a a equal to the square root of that pixel value (Gwhich implicitly assumes that the ""exin of the simulated CCD has been set to l electron per analog-to-digital readout count). and the readout noise was calculated using σ = LO."," The shot noise in each pixel was randomly calculated using a $\sigma$ equal to the square root of that pixel value (which implicitly assumes that the “gain” of the simulated CCD has been set to 1 electron per analog-to-digital readout count), and the readout noise was calculated using $\sigma$ = 10."188 The read-out noise iu these diuages is relatively siuall compared to the shot noise iu he sky backeround.g which is usually the case or real astronomical CCD nuages that lave a uoderatelv xieht vackeround level.," The read-out noise in these images is relatively small compared to the shot noise in the sky background, which is usually the case for real astronomical CCD images that have a moderately bright background level."189 The total roise in the background areas of these images as 8=ντους|10233.2.," The total noise in the background areas of these images has $\sigma =190\sqrt{1000 + 10^2} = 33.2$."191 Differeut starting random seed values were used so that the actual roise distribution varics in every nuage., Different starting random seed values were used so that the actual noise distribution varies in every image.192 The widely used SExtractor source extraction, The widely used SExtractor source extraction193superlliicl to the normal matter is likely to be responsible for pulsar elitches ancl postelitch relaxation.,superfluid to the normal matter is likely to be responsible for pulsar glitches and postglitch relaxation.194 A rotating neutron superlluid is thireaded by quantized vortex lines., A rotating neutron superfluid is threaded by quantized vortex lines.195 The superíIuid can alter its angular velocity by the radial motion of vortices., The superfluid can alter its angular velocity by the radial motion of vortices.196 In the inner crust of a neutron star neutron rich nuclei (hat form the bee lattice coexist with neutron superlluid., In the inner crust of a neutron star neutron rich nuclei that form the bcc lattice coexist with neutron superfluid.197 The vortex lines in (he inner crust are normally pinnecl to the lattice nuclei., The vortex lines in the inner crust are normally pinned to the lattice nuclei.198 Difference in velocity between superfIuid ancl nuclear lattice builds up since magnetic braking slows clown the nuclear latlice and pinning prevents (he vortex lines from moving., Difference in velocity between superfluid and nuclear lattice builds up since magnetic braking slows down the nuclear lattice and pinning prevents the vortex lines from moving.199 Then. (he neutron star undergoes sudden unpinning of a large number of vortex lines lollowed by Che outward motion.," Then, the neutron star undergoes sudden unpinning of a large number of vortex lines followed by the outward motion."200 This catastrophic unpinning has long been considered as a promising cause for glitehes Epstein.&Link 1992).," This catastrophic unpinning has long been considered as a promising cause for glitches \citep{and75, rud76, ala84, pin85, elb92, bel92}."201. A vortex line is subjected to the Magnus force when it moves relative to the superíIuid., A vortex line is subjected to the Magnus force when it moves relative to the superfluid.202 The catastrophic unpiniing model assumes that the pinnine force is strong enough to sustain the vortices to the nuclear lattice against the Magnus force until the moment just. before elitches., The catastrophic unpinning model assumes that the pinning force is strong enough to sustain the vortices to the nuclear lattice against the Magnus force until the moment just before glitches.203 Dased on the condensational ancl kinetic energies. (he pinning energy is estimated. ranging approximately 1—10 MeV. per nucleus (Alparetal.1984:Epstein&Bayvin.1938:pizzochero.Viverit.&Broglia 1997).," Based on the condensational and kinetic energies, the pinning energy is estimated, ranging approximately $1-10$ MeV per nucleus \citep{ala84, eps88, piz97}."204. This magnitude itself is strong enough against the Magnus force expected just before elitches., This magnitude itself is strong enough against the Magnus force expected just before glitches.205 Jones (1992) mentions. however. that (he pinnine forces on a randomly oriented rigid vortex line largely cancel since there are nearly an equal number of pinning sites on either side of the vortex line.," Jones (1992) mentions, however, that the pinning forces on a randomly oriented rigid vortex line largely cancel since there are nearly an equal number of pinning sites on either side of the vortex line."206" Link. Epstein. Davi (1993) lind that for a vortex line of finite tension. pinning becomes much more efficient by slightly bending aud forming kinks,"," Link, Epstein, Baym (1993) find that for a vortex line of finite tension, pinning becomes much more efficient by slightly bending and forming kinks."207 Recently. Jones (1997. 1993. 1999) argues that vortex interaction wilh a polvervstalline structure does not provide pinning strong enough (o explain the large elitches observed in the Vela pulsar (see Section ?? lor detail).," Recently, Jones (1997, 1998, 1999) argues that vortex interaction with a polycrystalline structure does not provide pinning strong enough to explain the large glitches observed in the Vela pulsar (see Section \ref{pinning} for detail)."208 llere we study the vortex configurations. oscillations aud pinning in (he inner crust οἱ a neutron star.," Here we study the vortex configurations, oscillations and pinning in the inner crust of a neutron star."209 Ii Section ?? we derive the equation of motion of vortex lines., In Section \ref{equation} we derive the equation of motion of vortex lines.210 In Section TT we present the equilibrium configurations of vortex lines., In Section \ref{configuration} we present the equilibrium configurations of vortex lines.211 In. Section ?? we examine the stability of equilibrium configurations ancl (he oscillations excited on a vortex line., In Section \ref{oscillation} we examine the stability of equilibrium configurations and the oscillations excited on a vortex line.212 In section ?? we discuss the vortex pinning based on (he results of the previous sections., In Section \ref{pinning} we discuss the vortex pinning based on the results of the previous sections.213 In the last section we summarize the results and mention the conclusions., In the last section we summarize the results and mention the conclusions.214relation.,relation.215 For their sample of acio-Ioud UDs. Berecr find /og(Ly/Ljg)~LI for late-M cwarfs. and fog(Ly/Lg)~12 for cooler dwarfs.," For their sample of radio-loud UDs, \citet{bbf+10} find $log(L_{X}/L_{R})\sim14$ for late-M dwarfs, and $log(L_{X}/L_{R})\sim12$ for cooler dwarfs."216 No clear trends jiwe been identified between the radio hunünosities of UDs aud stellar properties. such as rotation aud imiagnuetic Ποια streneth (Berecretal.2010).," No clear trends have been identified between the radio luminosities of UDs and stellar properties, such as rotation and magnetic field strength \citep{bbf+10}."217 Radio cussion fron UDs is variable on timescales [oveurss hours auc minutes.," Radio emission from UDs is variable on timescales of years, hours and minutes."218 Some UDs have radio icehteurves that are periodic ou the rotation periods of afew hours (Jallinanetal.2006.2007.2008:Bereeral. 2009).," Some UDs have radio lightcurves that are periodic on the rotation periods of a few hours \citep{had+06,hbl+07,had+08,brp+09}."219". These lighteurves are either characterized by uodulatious (οιοι,Tallinanetal.2006) or by short evcle peaks lasting for a few minutes (c.¢..Hallinauetal.Ww 07)."," These lightcurves are either characterized by modulations \citep[e.g.,][]{had+06} or by short duty-cycle peaks lasting for a few minutes \citep[e.g.,][]{hbl+07}."220. Other UDs exhibit isolated flares when otherwise radio-loud. also ou fes-uinute timescales (e...Durgasser&Putinan2005.hereafter DP05)..," Other UDs exhibit isolated flares when otherwise radio-loud, also on few-minute timescales \citep[e.g.,][hereafter BP05]{bp05}."221 The quickly-varvine cluission is generally 10054 circnlarly-polarized (e...ITal- 2008).," The quickly-varying emission is generally $\%$ circularly-polarized \citep[e.g.,][]{had+08}."222. In addition. the tvpes of radio ciission observed from UDs changes ou timescales of vears (Ostenetal. 2009).. varving between undetectable. quicsceut and periodic.," In addition, the types of radio emission observed from UDs changes on timescales of years \citep{oph+09}, varying between undetectable, quiescent and periodic."223 The variability of UD radio emission characteristics makes it hard to identify mubiased oud samples for population studies., The variability of UD radio emission characteristics makes it hard to identify unbiased radio-loud samples for population studies.224 Radio observations of UDs provide sguificaut insight iuto conditions in UD maeuctospheres., Radio observations of UDs provide significant insight into conditions in UD magnetospheres.225 Iu this Letter. we xeseut the widest-banud radio observations vet reported or the UD 395601. (hereafter DENISIO18). using the new Compact Array. Droadbaud Backend (Ferris&Wilsou2002) at the Australia Telescope Compact Array (ATCA.Fraterctal.1992)..," In this Letter, we present the widest-band radio observations yet reported for the UD $-$ 395604 (hereafter DENIS1048), using the new Compact Array Broadband Backend \citep{fw02} at the Australia Telescope Compact Array \citep[ATCA,][]{fbw92}."226 We propose a magnetospheric model which accounts for he violation of the CaiddelBeuz relation., We propose a magnetospheric model which accounts for the violation of the Güddel-Benz relation.227 This model could provide interesting insights into the magnetic field and plasina euviromneuts of these cuigliatic stars., This model could provide interesting insights into the magnetic field and plasma environments of these enigmatic stars.228 The tarect source. DENTSLOLS. was oue of seven Southern late-Al aud L dvarfs observed by BPOS with the ATCA in the ccm and ccm bands.," The target source, DENIS1048, was one of seven Southern late-M and L dwarfs observed by BP05 with the ATCA in the cm and cm bands."229 BPO5 reported a quicscent fux density of 1250.0 Limdy at 601. as well as a [5 nüuute flare in each baud. separated by 10 nüuutes with a peal flux density of 30nunuJw at Sec.," BP05 reported a quiescent flux density of $\pm$ mJy at cm, as well as a $-$ 5 minute flare in each band, separated by $\sim$ 10 minutes with a peak flux density of mJy at cm."230 DENISIOLS (spectralclassificationADS8.5.Πανal.2001). was identified as à UD in the DENTS survey (Epchteiuetal.1997). by Doelfosseetal.(2001).. ancl. at a distance of L.002E0.03 ppe (Costaetal.2005)... is one of the closest known stars.," DENIS1048 \citep[spectral classification M8.5,][]{hsb+04} was identified as a UD in the DENIS survey \citep{edc+97} by \citet{dfm+01}, and, at a distance of $\pm$ pc \citep{cmj+05}, is one of the closest known stars."231 A recent spectroscopic study w Martinetal.(2010). shows that it is uulikelv to be a own cawart., A recent spectroscopic study by \citet{mpb+10} shows that it is unlikely to be a brown dwarf.232 Επιστ:&Schinitt(2001). reported a aree optical flare aud a fast projected rotation velocity of esin/=2542 kkinss+., \citet{fs04} reported a large optical flare and a fast projected rotation velocity of $v\sin i=25\pm2$ $^{-1}$.233 While Ho eission was detected x Delfosseetal.(2001).. Schnutt&Liefke(2001). found 10 N-vav cuiission with an upper limit of241079 ss.! yon the ROSAT ΑΠδν Survey catalogue.," While $\alpha$ emission was detected by \citet{dfm+01}, \citet{sl04} found no X-ray emission with an upper limit of $2\times10^{26}$ $^{-1}$ from the ROSAT All-Sky Survey catalogue."234 Reiners&Das(2010) found au average lince-of-sielt magnetic field streugth of 23004LOO GG using measurements of Zeca xoadeniue in Fell absorption lines., \citet{rb10} found an average line-of-sight magnetic field strength of $2300\pm400$ G using measurements of Zeeman broadening in FeH absorption lines.235 We observed DENISIOI8 with the ATCA on 2000 August 10 and 11 in the 1.2ece02 band. aud sauultzueouslv im the 23 and Geem hands on 2009 August 15.," We observed DENIS1048 with the ATCA on 2009 August 10 and 11 in the cm band, and simultaneously in the 3 and cm bands on 2009 August 15."236 The six 22-metre ATCA antennas were placed in au exteuded configuration in order to niaxinise point-source sensitivity., The six 22-metre ATCA antennas were placed in an extended configuration in order to maximise point-source sensitivity.237 Bascline leneths ranged between 300nui aud 600012. corresponding to resolutions between approximately and at οσα.," Baseline lengths ranged between m and m, corresponding to resolutions between approximately and at cm."238 Visibility nieasurenüents for all basclines were recorded in two CCIz bands per Stokes polarization with NMMITz frequency resolution., Visibility measurements for all baselines were recorded in two GHz bands per Stokes polarization with MHz frequency resolution.239 The visibilitics were iutegrated over LOss intervals., The visibilities were integrated over s intervals.240 Details of the observations are even iu Table 1., Details of the observations are given in Table 1.241 We reduced the data using the MIRIAD software package (Saultetal.1995)., We reduced the data using the MIRIAD software package \citep{stw95}.242. Staudard calibrations were performed using observations of the ATCA primary calibrator 638 on cach dav. and frequeut observations of a radio galaxy 115 4.separated. bv 67 from DENISI0I8.," Standard calibrations were performed using observations of the ATCA primary calibrator $-$ 638 on each day, and frequent observations of a radio galaxy $-$ $-$ 445 $-$ separated by $^{\circ}$ from DENIS1048."243 Multi-fequeucy svuthesis totaliutensitv nuages were produced in sub-bands of MMIIz for the Geom aud 3ecu observations. and iu cach 2€GGIIzsub-baud for the L2ccm observations.," Multi-frequency synthesis total-intensity images were produced in sub-bands of MHz for the cm and cm observations, and in each GHzsub-band for the cm observations."244 We detected DENISLOLS as a point source iu all images. except that formed frou the 21GCGIIz data.," We detected DENIS1048 as a point source in all images, except that formed from the GHz data."245 The measured position of right- LOh. [812. 13.588 (40.038). declination: -397. 56. (40.55) is offset from the 2MASS position of DENISIOLS (Cutrietal2003). bv15.67.. which corresponds to the kuown proper motion (Deaconctal.20053.," The measured position of right-ascension: 10h, 48m, 13.58s $\pm0.03$ s), declination: $^{\circ}$, 56', $\pm0.5$ ) is offset from the 2MASS position of DENIS1048 \citep{c+03} by, which corresponds to the known proper motion \citep{dhc05}."246. The flux cdeusity of DENISLOLS was measured iu cach sub-hand by fitting the restoring Gaussian beans to the images., The flux density of DENIS1048 was measured in each sub-band by fitting the restoring Gaussian beams to the images.247 The beams were at position angles of 19° for the Gecin and 2e data. and 2° for the cc data. The rius noise levels. 0. in images made from cach sub-baud ranged between 2050 jy for the ccm data. aud between 25.35 Jy for the 3ccn data.," The beams were at position angles of $19^{\circ}$ for the cm and cm data, and $2^{\circ}$ for the cm data, The rms noise levels, $\sigma$, in images made from each sub-band ranged between $30-50$ $\mu$ Jy for the cm data, and between $25-35$ $\mu$ Jy for the cm data."248 For both the GGIIz aud CGCGIIz. 0=8 yr Jy.," For both the GHz and GHz, $\sigma=8$ $\mu$ Jy."249 We present the resulting spectrmm of DENISLOLS in Figure 1., We present the resulting spectrum of DENIS1048 in Figure 1.250" The Stokes I iicasureimoeuts of the flux deusityv of DENISI01ISs. Sv). at various frequencies r are au excellent fif to a power-law. ο)x»r"". where l.rlx 0.09."," The Stokes I measurements of the flux density of DENIS1048, $S(\nu)$, at various frequencies $\nu$ are an excellent fit to a power-law, $S(\nu)\propto\nu^{-\alpha}$, where $\alpha=1.71\pm0.09$ ."251 A similar process was also applied to the Stokes Q. U and V data.," A similar process was also applied to the Stokes Q, U and V data."252 While DENISIOIS was not found to have any detectable Stokes Q or U cinission. we detected Stokes Vo cussion atf frequencies up to GGIIz bv combining data iu imultiple sub-bauds.," While DENIS1048 was not found to have any detectable Stokes Q or U emission, we detected Stokes V emission at frequencies up to GHz by combining data in multiple sub-bands."253 Circular polarization fractions rangiug between 0.25 and 0.1 were found iu the cci band. aud 206 upper lanits of 0.2 were placed on the linear polarization fractious.," Circular polarization fractions ranging between 0.25 and 0.4 were found in the cm band, and $\sigma$ upper limits of 0.2 were placed on the linear polarization fractions."254 The Stokes V flux clensity measurenieuts are also plotted iu Figure 1l., The Stokes V flux density measurements are also plotted in Figure 1.255 An nuage of the data recorded between GGIIz aud [988CGGIIz. with Stokes V. contours overlayed ou a Stokes I erevscale nuage. is shown in Figure 2.," An image of the data recorded between GHz and GHz, with Stokes V contours overlayed on a Stokes I greyscale image, is shown in Figure 2."256 No sieuificaut short-timescale amplitude excursions or periodicities were detected iu any παςαλά in either the Stokes Tor V data., No significant short-timescale amplitude excursions or periodicities were detected in any sub-band in either the Stokes I or V data.257 We consider two possible radio emission mechanisms for DENISI0LS: evrosvuchrotron emission aud clectrou-evclotron maser (ECAL) eiission., We consider two possible radio emission mechanisms for DENIS1048: gyrosynchrotron emission and electron-cyclotron maser (ECM) emission.258 The spectral shape aud auuplitude are clearly inconsistent with thermal emission., The spectral shape and amplitude are clearly inconsistent with thermal emission.259 Cyrosvuchrotron and ECAD imechanisiis are both associated with muldly relativistic. nou-thermal electron populations. with cnergics >~— 20kkeV. Whereas eyvrosvuchrotron clissiou ds caused by incohereutlv radiating uou-thermal electrons propagating along magnetic field lines. ΤΟΝΤ cinission is coherent. and requires these electrons to lave au anisotropic pitcli-anele distribution (Dulk 1985)..," Gyrosynchrotron and ECM mechanisms are both associated with mildly relativistic, non-thermal electron populations, with energies $>\sim20$ keV. Whereas gyrosynchrotron emission is caused by incoherently radiating non-thermal electrons propagating along magnetic field lines, ECM emission is coherent, and requires these electrons to have an anisotropic pitch-angle distribution \citep{d85}. ."260 Gvrosvuchrotron, Gyrosynchrotron261"due to the fact that Sandell(1994) used the G93 model in order to determine his absolute calibration, which we have found to overestimate the brightness of Uranus by ~ 7%)).","due to the fact that \citet{sandell94}262 used the G93 model in order to determine his absolute calibration, which we have found to overestimate the brightness of Uranus by $\simeq 7$ )."263" It is not clear why our measured peak flux densities are higher for (94.5 and K3-50A, although the known extended emission in both sources, coupled with the fact that Sandell(1994) made single-pixel chopped photometry measurements at 27 arcsec resolution, may be the cause."," It is not clear why our measured peak flux densities are higher for G34.3 and K3-50A, although the known extended emission in both sources, coupled with the fact that \citet{sandell94}264 made single-pixel chopped photometry measurements at 27 arcsec resolution, may be the cause."265" We find that Uranus and Neptune behave as ideal sources for flux calibration at 143 GHz, with no evidence for temporal brightness variations."," We find that Uranus and Neptune behave as ideal sources for flux calibration at 143 GHz, with no evidence for temporal brightness variations."266" For Uranus, these results are in contrast to the lower frequency measurements of Kl06 and Kr08, who find -0.5 percent/year variations in the brightness temperature of Uranus at 8.6 and 90 GHz."," For Uranus, these results are in contrast to the lower frequency measurements of Kl06 and Kr08, who find $\simeq 0.5$ percent/year variations in the brightness temperature of Uranus at 8.6 and 90 GHz."267" Our data, combined with the measurements of O86 and G93, place a confidence level upper limit of 0.19 on the magnitude of variations in the brightness percent/yeartemperature of Uranus at c150 GHz over the same period."," Our data, combined with the measurements of O86 and G93, place a confidence level upper limit of 0.19 percent/year on the magnitude of variations in the brightness temperature of Uranus at $\simeq 150$ GHz over the same period."268 See Figure 2.., See Figure \ref{fig:f2}.269" A physical interpretation of the temporal variations in the brightness of Uranus seen at lower frequencies by K106 and Kr08, in combination with our static 143 GHz results, is beyond the scope of this manuscript, which is intended to quantify the magnitude and stability of the brightness of Uranus for the purposes of using it as a calibrator at 143 GHz."," A physical interpretation of the temporal variations in the brightness of Uranus seen at lower frequencies by Kl06 and Kr08, in combination with our static 143 GHz results, is beyond the scope of this manuscript, which is intended to quantify the magnitude and stability of the brightness of Uranus for the purposes of using it as a calibrator at 143 GHz."270" However, we do note that the combined results are not necessarily inconsistent, given that higher frequency observations of Uranus probe higher altitudes in the atmosphere (Kr08)."," However, we do note that the combined results are not necessarily inconsistent, given that higher frequency observations of Uranus probe higher altitudes in the atmosphere (Kr08)."271" Using Bolocam data collected between 2003 and 2010 we have tightly constrained the 143 GHz brightness ratio of Uranus and Neptune (1.027+0.006), and we find no evidence for temporal variations in the 143 GHz brightness temperature of either planet over that period."," Using Bolocam data collected between 2003 and 2010 we have tightly constrained the 143 GHz brightness ratio of Uranus and Neptune $1.027 \pm 0.006$ ), and we find no evidence for temporal variations in the 143 GHz brightness temperature of either planet over that period."272" Combining our results with those of O86 and G93, we find no evidence for 143 GHz brightness variations in either planet over the period from 1983—2010, and place a confidence level upper limit of on the magnitude of brightness variations over the 28 year period from 1983 to 2010."," Combining our results with those of O86 and G93, we find no evidence for 143 GHz brightness variations in either planet over the period from $1983-2010$, and place a confidence level upper limit of on the magnitude of brightness variations over the 28 year period from 1983 to 2010."273" By extrapolating the WMAP 94 GHz results given in Weilandetal.(2011) to our observing band using the brightness models presented in G93, we are able to constrain the absolute 143 GHz brightness temperature of each planet to ~3%."," By extrapolating the WMAP 94 GHz results given in \citet{weiland11} to our observing band using the brightness models presented in G93, we are able to constrain the absolute 143 GHz brightness temperature of each planet to $\simeq 3$."274". Additionally, we determine ~3% absolute 143 GHz peak flux densities for the ultracompact HII regions G34.3 and K3-50A and the protostellar source NGC 2071IR."," Additionally, we determine $\simeq 3$ absolute 143 GHz peak flux densities for the ultracompact HII regions G34.3 and K3-50A and the protostellar source NGC 2071IR."275" We acknowledge the assistance of: the Bolocam instrument team: Peter Ade, James Aguirre, James Bock, Sam Edgington, Jason Glenn, Alexy Goldin, Sunil Golwala, Douglas Haig, Andrew Lange, Glenn Laurent, Phil Mauskopf, Hien Nguyen, Philippe Rossinot, and Jack Sayers; Matt Ferry, Matt Hollister, Patrick Koch, Kai-Yang Lin, Sandor Molnar, Seth Siegel, and Keiichi Umetsu who, in addition to the Bolocam instrument team, helped collect the data presented in this manuscript; the day crew and Hilo staff of the Caltech Submillimeter Observatory, who provided invaluable assistance during commissioning and data-taking for this data set; Kathy Deniston, Barbara Wertz, and Diana Bisel, who provided effective administrative support at Caltech and in Hilo; and the referee, who provided helpful comments and suggestions."," We acknowledge the assistance of: the Bolocam instrument team: Peter Ade, James Aguirre, James Bock, Sam Edgington, Jason Glenn, Alexy Goldin, Sunil Golwala, Douglas Haig, Andrew Lange, Glenn Laurent, Phil Mauskopf, Hien Nguyen, Philippe Rossinot, and Jack Sayers; Matt Ferry, Matt Hollister, Patrick Koch, Kai-Yang Lin, Sandor Molnar, Seth Siegel, and Keiichi Umetsu who, in addition to the Bolocam instrument team, helped collect the data presented in this manuscript; the day crew and Hilo staff of the Caltech Submillimeter Observatory, who provided invaluable assistance during commissioning and data-taking for this data set; Kathy Deniston, Barbara Wertz, and Diana Bisel, who provided effective administrative support at Caltech and in Hilo; and the referee, who provided helpful comments and suggestions."276" Bolocam was constructed and commissioned using funds from NSF/AST-0098737, NSF/AST-9980846, NSF/AST-9618798,NSF/AST-0229008, and NSF/AST-0206158."," Bolocam was constructed and commissioned using funds from NSF/AST-9618798, NSF/AST-0098737, NSF/AST-9980846, NSF/AST-0229008, and NSF/AST-0206158."277" JS was partially supported by a NASA Graduate Student Research Fellowship, a NASA Postdoctoral Program fellowship, NSF/AST-0838261, and NASA/NNXI1ABOT7G; NC was partially supported by NASA Graduate Student Research Fellowship; SG acknowledges an Alfred P. Sloan Foundation fellowship."," JS was partially supported by a NASA Graduate Student Research Fellowship, a NASA Postdoctoral Program fellowship, NSF/AST-0838261, and NASA/NNX11AB07G; NC was partially supported by NASA Graduate Student Research Fellowship; SG acknowledges an Alfred P. Sloan Foundation fellowship."278The class of stars stil preseuts poorlv defined characteristics. aud this more than 50 vears after the discovery of the prototype member by Morgan et al (1913) who noted the abnormally weak metal lines of this AO dwarf star.,"The class of stars still presents poorly defined characteristics, and this more than 50 years after the discovery of the prototype member by Morgan et al (1943) who noted the abnormally weak metal lines of this A0 dwarf star."279 The properties that should define a star are not clearly established: the proposed spectroscopic criteria are usually. based on the weakness of metal lines. especially of the Me II. List. compared with what is expectec from the lvdrogen line type. while C. N. O and S have nearly solar abundances.," The properties that should define a star are not clearly established; the proposed spectroscopic criteria are usually based on the weakness of metal lines, especially of the Mg II 4481, compared with what is expected from the hydrogen line type, while C, N, O and S have nearly solar abundances."280 The kincmatic behaviour shoul allow to distinguish these stars from the metal poor A-type IHorizoutal Branch stars., The kinematic behaviour should allow to distinguish these stars from the metal poor A-type Horizontal Branch stars.281 Moderate to high projectec rotational velocities are usually found απο stars. although some exceptions have been recently ideutifie (e.g. IID 61191 and TD 7£872 selected by Paunzen Gray L997).," Moderate to high projected rotational velocities are usually found among stars, although some exceptions have been recently identified (e.g. HD 64491 and HD 74873 selected by Paunzen Gray 1997)."282 The result of he vaeueOo definitions of these non-evolvec metal uucderabunudaut stars is well reflected by the variety of opinions existing at present about the members of this class., The result of the vague definitions of these non-evolved metal underabundant stars is well reflected by the variety of opinions existing at present about the members of this class.283 The metal abundances obtained up to now reveal a hieh scatter from star to star., The metal abundances obtained up to now reveal a high scatter from star to star.284 Details on the evolution with time of the definition are sununuarized in Farageiana Gerbalci (1998): the not clearly defined properties of these stars are. at least xuwtiall. responsible of the various hypotheses xoposed to explain the plienomenou as well as of he uncertainty on the age attributed to these objects. which spans from that of stars not vet on the Main Sequence to that of old objects descending from coutact inary svstenas.," Details on the evolution with time of the definition are summarized in Faraggiana Gerbaldi (1998); the not clearly defined properties of these stars are, at least partially, responsible of the various hypotheses proposed to explain the phenomenon as well as of the uncertainty on the age attributed to these objects, which spans from that of stars not yet on the Main Sequence to that of old objects descending from contact binary systems."285 The preseut paper reviews the characteristics of the uembers of this class according to recent compilations aud discusses the effect of duplicity ou a composite spectrum as source of nusclassification for some of these camucidatecs., The present paper reviews the characteristics of the members of this class according to recent compilations and discusses the effect of duplicity on a composite spectrum as source of misclassification for some of these candidates.286 Iu a moderu astrophysical perspective. age. positio- in the IIR. diagram aud chemical abuudances are the kev quantities which describe a class of stars.," In a modern astrophysical perspective, age, position in the HR diagram and chemical abundances are the key quantities which describe a class of stars."287 The purpose of introducing a class of stars is to help identify a comune- nuderling phenomenology., The purpose of introducing a class of stars is to help identify a common underlying phenomenology.288 For stars the ain is f find the common factor which can explain the observea chemical peculiarities: to be iieaniueful this iust explai- a statistically significant sample of stars., For stars the aim is to find the common factor which can explain the observed chemical peculiarities; to be meaningful this must explain a statistically significant sample of stars.289 Bearing thiρα in nind. it is clear that anv classification scheme whic[um does not rely on abundauce criteria is unlikely to be helpful.," Bearing this in mind, it is clear that any classification scheme which does not rely on abundance criteria is unlikely to be helpful."290 It is probable that as high accuracy abundance data acciunulate. we will have to revise our concept of stars and probably reach a more plivsical definition.," It is probable that as high accuracy abundance data accumulate, we will have to revise our concept of stars and probably reach a more physical definition."291 May stars were cassified as iu the past., Many stars were classified as in the past.292 The catalogue of Ποιοι ot al (1990) includes over 100 candidates;, The catalogue of Renson et al (1990) includes over 100 candidates.293 Mauv of these turned out to be misclassified and the whole sample results too heterogeneous., Many of these turned out to be misclassified and the whole sample results too heterogeneous.294 We selected. stars classified as iu recent papers based on modern data. loping to extract a more homogeneous sample.," We selected stars classified as in recent papers based on modern data, hoping to extract a more homogeneous sample."295 They should be considered candidates. since for many of them further analysis to check whether they match auv given definition is still required.," They should be considered candidates, since for many of them further analysis to check whether they match any given definition is still required."296 The caudidates we selected. with the exception of three of them. have been listed in at least oue of the following papers: Abt Morrell(1995: hereafter: AM). Paunzen ct al (1997: hereafter CC. Cray (1999: hereafter €).," The candidates we selected, with the exception of three of them, have been listed in at least one of the following papers: Abt Morrell(1995; hereafter AM), Paunzen et al (1997; hereafter CC), Gray (1999; hereafter G)."297 The exceptions are: TID 290192 aud TD 90821 which were classified as by Pauuzeu Cray (1997): ΠΟ 105759 for which the CG classification is unpublished., The exceptions are: HD 290492 and HD 90821 which were classified as by Paunzen Gray (1997); HD 105759 for which the G classification is unpublished.298 Both iethods aud scope differ :unong the three papers (ie. AALCCO.C).," Both methods and scope differ among the three papers (i.e. AM,CC,G)."299 However. the deeree of reliability of cach of them is difficult to clefine.," However, the degree of reliability of each of them is difficult to define."300" Promising candidates are found iun alb lists. although ταν is probably the most reliable source. because the author (ταν Carrison 1051, 1L989a.L989b: Carrison ααν 1991: CC im Table 1) has classified a large suuple of ποια aud “standard” stars using the same methods."," Promising candidates are found in all lists, although Gray is probably the most reliable source, because the author (Gray Garrison 1987, 1989a,1989b; Garrison Gray 1994; GG in Table 1) has classified a large sample of “normal” and “standard” stars using the same methods."301 AAD is a study of stellar v sn of 1700 A-type stars of the Bright Star Catalogue (IILoffiei Warren 1991) (BSC)., AM is a study of stellar v $ \sin i$ of 1700 A-type stars of the Bright Star Catalogue (Hoffleit Warren 1994) (BSC).302 Ou the basis of their available spectra (photographic spectra of dispersion 39 Amun1 ) they eive a classification for cach star., On the basis of their available spectra (photographic spectra of dispersion 39 $\rm mm^{-1}$ ) they give a classification for each star.303 Some are classified as A Doo., Some are classified as $\lambda$ Boo.304kinematics typical of a halo cluster.,kinematics typical of a halo cluster.305 It is also highly reddened. and its CMD is affected by some differential reddening.," It is also highly reddened, and its CMD is affected by some differential reddening."306 Alonso et al. (1997)), Alonso et al. \cite{alonso97}) )307 present the only other CCD 7 and T CMD existent for this cluster., present the only other CCD $B$ and $V$ CMD existent for this cluster.308 NGC 6366. together with NGC 5053. are the only two clusters that were observed under not exceptionally good seeing conditions.," NGC 6366, together with NGC 5053, are the only two clusters that were observed under not exceptionally good seeing conditions."309 Still. all the sequences in the CMD can be identified (apart from the upper RGB). including what seems to be a well populated blue straggler sequence.," Still, all the sequences in the CMD can be identified (apart from the upper RGB), including what seems to be a well populated blue straggler sequence."310 The HB is very red. as expected on the basis of the metallicity. and tilted.," The HB is very red, as expected on the basis of the metallicity, and tilted."311 We measured a total of ~5500 stars for this cluster. (, We measured a total of $\sim 5500$ stars for this cluster. (312Fig. 21) ,Fig. \ref{ngc6535}) )313To our knowledge. Sarajedini (1994)) has published the only previous CCD study of this cluster: a B and V. CMD down to V~21.," To our knowledge, Sarajedini \cite{sarajedini94}) ) has published the only previous CCD study of this cluster: a $B$ and $V$ CMD down to $V\sim21$."314 His stellar population is slightly smaller than ours for this range of magnitudes (we reach V ~ 23)., His stellar population is slightly smaller than ours for this range of magnitudes (we reach V $\sim23$ ).315 NGC 6535 is the least luminous object of our northern sample. and probably the one with the smallest number of stars.," NGC 6535 is the least luminous object of our northern sample, and probably the one with the smallest number of stars."316 We measured ~7500 stars for this cluster., We measured $\sim7800$ stars for this cluster.317 Its RGB is identifiable. but not clearly defined. due also to field star contamination.," Its RGB is identifiable, but not clearly defined, due also to field star contamination."318 Its CMD somehow resembles the CMD of NGC 6717 (Paper D). (, Its CMD somehow resembles the CMD of NGC 6717 (Paper I). (319Fig. 22) ,Fig. \ref{ngc6779}) )320We have not found any previous CCD study on this cluster., We have not found any previous CCD study on this cluster.321 Our CMD is well defined. though it is slightly contaminated by foreground/background stars.," Our CMD is well defined, though it is slightly contaminated by foreground/background stars."322 The broadening of the SGB-RGB might suggest the existence of some differential reddening., The broadening of the SGB-RGB might suggest the existence of some differential reddening.323 The distribution of the stars along the BHB seems to be not homogeneous. with the possible presence of a gap.," The distribution of the stars along the BHB seems to be not homogeneous, with the possible presence of a gap."324 The total of measured stars was of ~11300. (, The total of measured stars was of $\sim11300$. (325Fig. 23) ,Fig. \ref{ngc6838}) )326As suggested also by its CMD. M71 ts a metal rich cluster. similar to 47 Tue (Paper D.," As suggested also by its CMD, M71 is a metal rich cluster, similar to 47 Tuc (Paper I)."327 Our CMD is well defined and extends for more than 4 magnitudes below the TO. covering a total of ~12500 stars.," Our CMD is well defined and extends for more than 4 magnitudes below the TO, covering a total of $\sim12500$ stars."328 The cluster has only a RHB. and the upper part of the RGB is not very well defined.," The cluster has only a RHB, and the upper part of the RGB is not very well defined."329 This cluster is located close to the Galactic plane. and this explains the contamination by disk stars clearly visible in the CMD.," This cluster is located close to the Galactic plane, and this explains the contamination by disk stars clearly visible in the CMD."330 It is very bright and relatively nearby., It is very bright and relatively nearby.331 Despite this. there is no CMD in the literature after Hodder et al. (1992)).," Despite this, there is no CMD in the literature after Hodder et al. \cite{hodder92}) )."332 They present a good B and V diagram. less populated than ours. reaching V=22.," They present a good $B$ and $V$ diagram, less populated than ours, reaching $V=22$."333 Previous CCD studies are in Richer Fahlman (1988)). who present (.D.V. photometry for the main sequence. down to V—22 (0= 25).," Previous CCD studies are in Richer Fahlman \cite{richer88}) ), who present $U,B,V$ photometry for the main sequence, down to $V=22$ $U=25$ )."334 No evolved stars are present in this work. (, No evolved stars are present in this work. (335Fig. 24)),Fig. \ref{ngc7078}) )336 This cluster has been extensively studied in the past. both with groundbased facilities and a large number of HST observations.," This cluster has been extensively studied in the past, both with groundbased facilities and a large number of HST observations."337 HST studies include Stetson 1994 and Yanny et al 1994... were a CMD of the cluster center is presented.," HST studies include Stetson \cite{stetson94} and Yanny et al \cite{yanny94}, , were a CMD of the cluster center is presented."338 The CMD does not arrive to the MS TO. and is quite disperse.," The CMD does not arrive to the MS TO, and is quite disperse."339 Conversely. Sosin King 1997. and Piotto et al. 1997))," Conversely, Sosin King \cite{sosin97} and Piotto et al. \cite{piotto97}) )"340 present and extraordinarily well defined. MS. but no evolved stars are present.," present and extraordinarily well defined MS, but no evolved stars are present."341 The most recent ground-based study is the composite CMD of Durrell Harris (1993)) based on CCD data from two telescopes., The most recent ground-based study is the composite CMD of Durrell Harris \cite{durrell93}) ) based on CCD data from two telescopes.342 This is the kind of problem that we try to avoid with the present catalog., This is the kind of problem that we try to avoid with the present catalog.343 Our diagram is well populated (~27000 stars) from the RGB tip down to V=22.5., Our diagram is well populated $\sim27000$ stars) from the RGB tip down to $V=22.5$.344 The CMD features are better identifiable when a radial selection. avoiding the clusters center. is done.," The CMD features are better identifiable when a radial selection, avoiding the clusters center, is done."345 The CMD in Fig., The CMD in Fig.346 24 gives the visual impression that there are three distinct groups of stars in the HB., \ref{ngc7078} gives the visual impression that there are three distinct groups of stars in the HB.347 The third possible group. on the red side of the RR Lyrae gap is surely a statistical fluctuation in the distribution of the RR Lyrae magnitudes and colors at random phase.," The third possible group, on the red side of the RR Lyrae gap is surely a statistical fluctuation in the distribution of the RR Lyrae magnitudes and colors at random phase."348 It 1s present neither in the CMDs of M15 in the above quoted works nor in our CMD of a larger stellar sample. with more accurate photometry from our HSTdata base.," It is present neither in the CMDs of M15 in the above quoted works nor in our CMD of a larger stellar sample, with more accurate photometry from our HSTdata base."349the magnetic field is initialized to be purely azimutha in the x y plane at all radii such that r<0Mee.,the magnetic field is initialized to be purely azimuthal in the $x$ $y$ plane at all radii such that $r\le r_{max}$.350 Essentially. this confguration is a sphere of neste circular inaenetic feld loops of constant strength.," Essentially, this configuration is a sphere of nested circular magnetic field loops of constant strength."351 Muuevically this configuration is realized by utilizing the vector potential. where B—WVxA.," Numerically this configuration is realized by utilizing the vector potential, where $\boldsymbol{B} = \boldsymbol{\nabla}\times \boldsymbol{A}$."352 For a coustau magnetic field of streneth Bo defined as described. the vector poteutial is given as The vector potential is caleulated at cell corners aud then differenced to calculate the face-ceutered magnetic fields.," For a constant magnetic field of strength $B_0$ defined as described, the vector potential is given as The vector potential is calculated at cell corners and then differenced to calculate the face-centered magnetic fields."353 The advantage of this maeuetic field geometry is that the leat flux is zero in the initial configuration: thus. all subsequent evolution is due to the MTI.," The advantage of this magnetic field geometry is that the heat flux is zero in the initial configuration; thus, all subsequent evolution is due to the MTI."354 The fiducial magnetic field streugth is chosen to be 1 uc. possibly moderately higher than the poorly constrained primordial feld but substautially less than today's observed feld.," The fiducial magnetic field strength is chosen to be 1 nG, possibly moderately higher than the poorly constrained primordial field but substantially less than today's observed field."355 A second magnetic field geometry is chosen to be a inore realistic represcutation for the tangled fields observed in clusters today., A second magnetic field geometry is chosen to be a more realistic representation for the tangled fields observed in clusters today.356 There is little theoretical euidanuce for the magnetic field power spectrmm in the ICAL so we choose a I&oliuogorov power spectrum for the magnetic field. initialized in Fourier space as where is chosen as the waveunuber correspouding μαςto 21 times the grid scale.," There is little theoretical guidance for the magnetic field power spectrum in the ICM, so we choose a Kolmogorov power spectrum for the magnetic field, initialized in Fourier space as where $k_{\textrm{peak}}$ is chosen as the wavenumber corresponding to 2–4 times the grid scale."357| We utilize the Fast Fourier Transform (FFT) as was done iu ? to caleulate the vector potentials iu real space., We utilize the Fast Fourier Transform (FFT) as was done in \citet{rs99} to calculate the vector potentials in real space.358 We also raucdomize the phase to avoid correlating the modes., We also randomize the phase to avoid correlating the modes.359" The magnetic enerev scales as B2οςμασ),", The magnetic energy scales as $B^2 \propto k^{-2(\alpha-1)}$.360" We initialize cach Gaussian component separately. so for the componentwise [xoluosorov spectrum. the appropriate choice is @=—17/6 to give the familiar sealing for energv of &-11/2,"," We initialize each Gaussian component separately, so for the componentwise Kolmogorov spectrum, the appropriate choice is $\alpha= -17/6$ to give the familiar scaling for energy of $k^{-11/3}$."361 This Respace scaling for lohuogorov turbulence is appropriate for initializing the 1D power spectrum (as opposed to the 3D power spectruni) in each direction separately., This $k$ -space scaling for Kolmogorov turbulence is appropriate for initializing the 1D power spectrum (as opposed to the 3D power spectrum) in each direction separately.362 We numerically check that the divergence of the maguetic field) remains zero and renormalize the maguetic energev such that (B*)=(Bj)., We numerically check that the divergence of the magnetic field remains zero and renormalize the magnetic energy such that $\langle B^2\rangle = \langle B_0^2\rangle$.363 Iu this section we discuss several of the characteristic timescales for a cluster as was outlined in ?.., In this section we discuss several of the characteristic timescales for a cluster as was outlined in \citet{ps07b}.364 The Bruut- frequency can be written in a more useful form for the cluster problem as The iiaxinma erowth rate is given bv the isothermal Iuuit of the Drunt-Viuisallà frequency. iuncly Table 2 examines these timescales in the galaxy cluster at a fiducial radius of 150 kpc.," The Brunt-V\""aiis\""all\""a frequency can be written in a more useful form for the cluster problem as The maximum growth rate is given by the isothermal limit of the Brunt-Väiisällä frequency, namely Table \ref{tab:clust:timescales} examines these timescales in the galaxy cluster at a fiducial radius of 450 kpc."365 Several of these timescales are dependent on the leneth scale of interest., Several of these timescales are dependent on the length scale of interest.366 Thus. they are cousidered on elobal scales. roughly a waveleneth of l1 Alpe. and the scale of the mean free path. roughly 30 kpc.," Thus, they are considered on global scales, roughly a wavelength of 1 Mpc, and the scale of the mean free path, roughly 30 kpc."367 The maenetic field is assumed to be 1 nC for these calculations., The magnetic field is assumed to be 1 nG for these calculations.368 As can be secu by exiuuination. the MTI has teus of erowtl times during a IIubble time. allowing significant rearranecinent of the atimosphere.," As can be seen by examination, the MTI has tens of growth times during a Hubble time, allowing significant rearrangement of the atmosphere."369 The magnetic field estimated using primordial values clearly plays no role. even on the leusth scale of the mean free path.," The magnetic field estimated using primordial values clearly plays no role, even on the length scale of the mean free path."370 Ou mean free path timescales. the conduction time is faster than the sound crossing time.," On mean free path timescales, the conduction time is faster than the sound crossing time."371 The potential worry is due to the thought that the heat fiux is saturated in the sense of ?.., The potential worry is due to the thought that the heat flux is saturated in the sense of \citet{cm77}.372 The right comparison there is nof the couduction time to the mode sound crossing finie but the electron thermal velocity to the velocity needed for electrous to transport the heat., The right comparison there is not the conduction time to the mode sound crossing time but the electron thermal velocity to the velocity needed for electrons to transport the heat.373 For the atinosphliere considered here. the saturated heat fiux is several orders of magnitude larger than the actual leat fiux.," For the atmosphere considered here, the saturated heat flux is several orders of magnitude larger than the actual heat flux."374 We use the 3D version of the Athena ATID code (C?) to simulate the AITT far into the nonlinear regine.," We use the 3D version of the Athena MHD code \citep{gs08,sg08} to simulate the MTI far into the nonlinear regime."375 Both the MIID aud the heat transport methodology and tests have been described in these references aud T. respectively.," Both the MHD and the heat transport methodology and tests have been described in these references and \citet{ps05}, respectively."376 The thermal diffusivity is either set as a constant or pernutted to vary spatially according to the standard Spitzer prescription., The thermal diffusivity is either set as a constant or permitted to vary spatially according to the standard Spitzer prescription.377 Tere we adopt the adopt the method of ? for handling variable thermal conductivities., Here we adopt the adopt the method of \citet{sh07} for handling variable thermal conductivities.378" Namely. the product of deusity and thermal diffusivity is interpolated to the cell faces using a harmonic average. mm one dimension as This harmonic averaging preveuts the Couraut condition frou, becoming severe due to discoutiuuous diffusivities aud densities at interfaces. should they develop."," Namely, the product of density and thermal diffusivity is interpolated to the cell faces using a harmonic average, in one dimension as This harmonic averaging prevents the Courant condition from becoming severe due to discontinuous diffusivities and densities at interfaces, should they develop."379 The timestep is determined with respect to the maxim thermal diffusivity ou the erid and the conduction module is sub-cvcled., The timestep is determined with respect to the maximum thermal diffusivity on the grid and the conduction module is sub-cycled.380 All the runs described in this chapter are performed onu a uniform Cartesian threc-dineusional grid with the initial cluster atmosphere described iu refsecnmodel.., All the runs described in this chapter are performed on a uniform Cartesian three-dimensional grid with the initial cluster atmosphere described in \\ref{sec:model}.381 The domain extends frou. —3.0x1074 cii to +3.0x1074 cin (968 kpe) in cach direction., The domain extends from $-3.0\times 10^{24}$ cm to $+3.0\times 10^{24}$ cm (968 kpc) in each direction.382 Most of the rus presented in this paper are snulated at a resolution of (128)? or correspouding to a zone size of 15 kpe respectively., Most of the runs presented in this paper are simulated at a resolution of $(128)^3$ or corresponding to a zone size of 15 kpc respectively.383 The simulations are typically runi for 3x101 s or 9.5 Car. a substantial fraction of cosmic time.," The simulations are typically run for $3\times 10^{17}$ s or 9.5 Gyr, a substantial fraction of cosmic time."384 The magnetic field is initialized in a region Tmax Which is slightly less than the leneth of the cube., The magnetic field is initialized in a region $r_{max}$ which is slightly less than the length of the cube.385 Typically for our siuulatious rj44=2.5x1074 (806 kpc)., Typically for our simulations $r_{\textrm{max}} = 2.5\times 10^{24}$ (806 kpc).386 Thus. the initial magnetic field is coufined to a spherical region.," Thus, the initial magnetic field is confined to a spherical region."387 Since there is no magnetic field outside of Pax. there is initially no conduction in this region.," Since there is no magnetic field outside of $r_{\textrm{max}}$, there is initially no conduction in this region."388 We use modified reflecting boundary conditions for all the MIID variables. in which the pressure aud deusitv are extrapolated iu the ghost zones.," We use modified reflecting boundary conditions for all the MHD variables, in which the pressure and density are extrapolated in the ghost zones."389 In order to prevent any uceative values for pressure or density iu the ghost zones. we introduce pressure and deusitv floors applied oulv in the ghost zones.," In order to prevent any negative values for pressure or density in the ghost zones, we introduce pressure and density floors applied only in the ghost zones."390 The maguetic field aud velocity field componcuts are simply reflected at the boundary., The magnetic field and velocity field components are simply reflected at the boundary.391 The heat fluxes are forced to be aciabatic by setting —0ou the boundary., The heat fluxes are forced to be adiabatic by setting on the boundary.392region of 20 pixel radius.,region of 20 pixel radius.393 The channels in the source spectrum were then binned in groups of 20 or more photons., The channels in the source spectrum were then binned in groups of 20 or more photons.394 Spectral fitting was done in NSPEC 12.5.1 (see Arnaud 1996 for a description of an earlier version ofλ, Spectral fitting was done in XSPEC 12.5.1 (see Arnaud 1996 for a description of an earlier version of.395ος Bebinned channels including photons below 0.5 keV and above S keV are ignored. because the response matrix lor Chandra. is better calibrated. within this energv range than outside of it.," Rebinned channels including photons below 0.5 keV and above 8 keV are ignored, because the response matrix for Chandra is better calibrated within this energy range than outside of it."396 We report the Lo errors on source parameters., We report the $\sigma$ errors on source parameters.397 In observation 321. taken 12 June 2000. the cluster's X-ray spectrum was well fitted (Cvo= 4.74/4) with a disk blackbody model with foreground. neutral hydrogen column density of 1.6«107 7. &T—O88(415 keV and normalization in NSPEC of L2.10 corresponcding toa best-fitting inner disk racius of 2 km assuming a lace- disk and no spectral hardening correction. (Mitsuda. et al.," In observation 321, taken 12 June 2000, the cluster's X-ray spectrum was well fitted $\chi^2/\nu=4.74/4$ ) with a disk blackbody model with foreground neutral hydrogen column density of $1.6\times10^{20}$ $^{-2}$, $kT = 0.88^{+0.16}_{-0.13}$ keV and normalization in XSPEC of $^{+1.6}_{-0.9} \times10^{-3}$, corresponding to a best-fitting inner disk radius of 72 km assuming a face-on disk and no spectral hardening correction (Mitsuda et al."398 1984)., 1984).399 The source Dux from 0.5-8.0 keV is um10+! fem /sec., The source flux from 0.5-8.0 keV is $^{+0.1}_{-0.9}\times10^{-14}$ $^2$ /sec.400 Phe best fitting value corresponds to à luminosity of 6.5«1077 eres/sec using a distance of 16 \Ipe to NGC 4472. based on the distance to the Virgo Cluster in which it is contained (Macri οἱ al.," The best fitting value corresponds to a luminosity of $6.5\times10^{38}$ ergs/sec using a distance of 16 Mpc to NGC 4472, based on the distance to the Virgo Cluster in which it is contained (Macri et al."401 1999)., 1999).402 1n observation 11274. taken on 27 February 2010. the source spectrum was well fitted. (i.c. v v=lO.5/19) with AL=152008 keV and normalization in NSPEC of 300$.]0 corresponding to a best fitting inner disk radius of 34 km assuming a face-on disk and no spectral hardening correction.," In observation 11274, taken on 27 February 2010, the source spectrum was well fitted (i.e. $\chi^2/\nu$ =10.5/19) with $kT =4031.52^{+0.16}_{-0.13}$ keV and normalization in XSPEC of $^{+3.3}_{-2.5}\times10^{-4}$, corresponding to a best fitting inner disk radius of 34 km assuming a face-on disk and no spectral hardening correction."404 The source [lux from 0.5-S.0. keV is m10H7 ergs/em?/sec., The source flux from 0.5-8.0 keV is $^{+0.2}_{-1.4}\times10^{-14}$ $^2$ /sec.405 The best fitting. value corresponds to a luminosity of 2.7/510° ergs/scc., The best fitting value corresponds to a luminosity of $2.7\times10^{39}$ ergs/sec.406 The temperatures of the disk ave thus different at nearly. the 3a level. and both spectra are. consistent with standard phenomenology that the inner disk radius will vary little in high/solt states.," The temperatures of the disk are thus different at nearly the $3\sigma$ level, and both spectra are consistent with standard phenomenology that the inner disk radius will vary little in high/soft states."407 Phe Iuminosity. difference is significant at more than deo. and the 3e. lower limit to the luminosity dillerence is about 55.107 ergs/sec. above the Exddington luminosity for a single neutron star.," The luminosity difference is significant at more than $4\sigma$, and the $3\sigma$ lower limit to the luminosity difference is about $5\times10^{38}$ ergs/sec, above the Eddington luminosity for a single neutron star."408 The Chandra spectra are shown in figure 1.., The Chandra spectra are shown in figure \ref{spectra}. .409 We have also triedἱ to [fit power law models to. the ata., We have also tried to fit power law models to the data.410 For observation 321. a power law model with the foreground Ny gives an acceptable fit. with ασE= 4176/4. with a power law index of 2.004O23US). and the 1-0 confidence interval for the [lux ranging from 10Ll ergs/sec/em. Doa," For observation 321, a power law model with the foreground $N_H$ gives an acceptable fit, with $\chi^2/\nu=4.76/4$ , with a power law index of $2.00\pm^{+0.24}_{-0.23}$, and the $\sigma$ confidence interval for the flux ranging from $\times10^{-14}$ $^2$."411nThe spectral shape is: thus muweinallv consistent with expectations for à. ονπαν state. but the luminosity is well above the few percent of 1e Exdldington luminosity in which hard states are typically 'ound (Maccarone 2003).," The spectral shape is thus marginally consistent with expectations for a low/hard state, but the luminosity is well above the few percent of the Eddington luminosity in which hard states are typically found (Maccarone 2003)."412 We thus favor the diskbb mocel fit as providing parameter values more likely to be indicative ‘the real physical state of the system. but we do note hat the spectral fits do not distinguish between the two scenarios.," We thus favor the diskbb model fit as providing parameter values more likely to be indicative of the real physical state of the system, but we do note that the spectral fits do not distinguish between the two scenarios."413 For observation 11274. a power law mocdel with the orceround Ny is formally a good fit. with αν=23/19. out with a spectral index of 1.3023:0.07. considerably harder han is ever seen in à low hard state from a Galactic black role X-ray. binary.," For observation 11274, a power law model with the foreground $N_H$ is formally a good fit, with $\chi^2/\nu=23/19$, but with a spectral index of $1.30\pm0.07$, considerably harder than is ever seen in a low hard state from a Galactic black hole X-ray binary."414 Since in the former case. the power law nmocel provides a poor fit to the data. and in the latter case. he best fitting value of the power law index lies outside the range expected. from. phenomenology. there is a strong case o be mace that the data are genuinely better explained with a strong thermal component than a pure power law spectrum.," Since in the former case, the power law model provides a poor fit to the data, and in the latter case, the best fitting value of the power law index lies outside the range expected from phenomenology, there is a strong case to be made that the data are genuinely better explained with a strong thermal component than a pure power law spectrum."415 The two other Chandra observations of this field of view. observation 322 (made 19 March. 2000). ancl observation SO95 (made 23 February 2008). have much shorter integration times.," The two other Chandra observations of this field of view, observation 322 (made 19 March 2000), and observation 8095 (made 23 February 2008), have much shorter integration times."416 For CXOU 1229410|075744. observation 322 vields 40 counts in 10 kiloseconcds with ACLIS-L. and observation NOO5 vields 24 counts in 5 kiloseconds. with ACIS-S. Both observations viele [lux levels of 5. 107 ergs/sec/eni.," For CXOU 1229410+075744, observation 322 yields 40 counts in 10 kiloseconds with ACIS-I, and observation 8095 yields 24 counts in 5 kiloseconds with ACIS-S. Both observations yield flux levels of $\sim$ 5 $\times10^{38}$ $^2$."417 Because there are not enough. counts [or detailed spectral fitting. there is a considerable uncertainty on the counts-to-energy conversion.," Because there are not enough counts for detailed spectral fitting, there is a considerable uncertainty on the counts-to-energy conversion."418 The Poisson errors are also substantial., The Poisson errors are also substantial.419 As a result. it is cdillicult to determine whether the flux levels during the two short. observations were higher or lower than those curing the longer observations.," As a result, it is difficult to determine whether the flux levels during the two short observations were higher or lower than those during the longer observations."420 “The X-ray detections ancl upper limits are summarized in Table 1.., The X-ray detections and upper limits are summarized in Table \ref{xrays}.421 Two deep NMM-Nexwton observations of this source have been mace as well., Two deep XMM-Newton observations of this source have been made as well.422 Llowever. this source is close cnough to the center of NCC 4472 that the clilfuse gas emission significantly alfects NMM'S sensitivity.," However, this source is close enough to the center of NGC 4472 that the diffuse gas emission significantly affects XMM's sensitivity."423 Phe 2XMM catalog (Watson et al., The 2XMM catalog (Watson et al.424 2009) reports a source S away [rom CXOU 1229410|075744. with a positional error of 4477 at 12h29m40.40s. |1 57m4.Is on 5 June 2002.," 2009) reports a source 8” away from CXOU 1229410+075744, with a positional error of 4.47” at 12h29m40.49s, $+7^{\circ}57$ m47.1s on 5 June 2002."425 The source is given a quality Hag (the EELACG parameter value) of 4. indicating that it is located within a region where spurious detections are likely. and that the source itself may be a spurious source.," The source is given a quality flag (the FLAG parameter value) of 4, indicating that it is located within a region where spurious detections are likely, and that the source itself may be a spurious source."426 Formally. 4o upper limits can be obtained from the ELIX tool. using data corresponding to the 2XMMI-DIU data release.," Formally, $4\sigma$ upper limits can be obtained from the FLIX tool, using data corresponding to the 2XMMi-DR3 data release."427 ELIN finds that the source was no brighter than about τι107 ergs/sec on 5 June 2002. and 2Q7 cres/sce on 1 January 2004.," FLIX finds that the source was no brighter than about $7\times10^{38}$ ergs/sec on 5 June 2002, and $2\times10^{38}$ ergs/sec on 1 January 2004."428 While the upper limits from FLIN appear to indicate that the source faded sometime after 2001. and re-brightened sometime between 2004 and 2008. we have also looked. at the aperture photometry from the FLIX tool.," While the upper limits from FLIX appear to indicate that the source faded sometime after 2001, and re-brightened sometime between 2004 and 2008, we have also looked at the aperture photometry from the FLIX tool."429 We set the extraction region to 5... in order to limit the cllects of confusion from nearby eas emission and other point sources.," We set the extraction region to 5”, in order to limit the effects of confusion from nearby gas emission and other point sources."430 We find that in the obseravations made on | January 2004. all three NAIAL instruments show a flux more than 3.86 above background in the 0.2-12 keV band. with the most sensitive PN detection above 5e.," We find that in the obseravations made on 1 January 2004, all three XMM instruments show a flux more than $3.8\sigma$ above background in the 0.2-12 keV band, with the most sensitive PN detection above $5\sigma$."431 The flux within 5° is 19+0.3.101 eres in the EPIC-PN.2.140.6.10H ergs/sec in MOSI. and /see3.0E0.6«1044 orgs/sec in MOS2.," The flux within 5” is $1.9\pm0.3\times10^{-14}$ ergs/sec in the EPIC-PN, $2.1\pm0.6\times10^{-14}$ ergs/sec in MOS1, and $3.0\pm0.6\times10^{-14}$ ergs/sec in MOS2."432 The encircled energy fraction at 57 is about, The encircled energy fraction at 5” is about.433 Taking the aperture photometry at face value. we estimate that the source was at about L2«10° cres/sec.," Taking the aperture photometry at face value, we estimate that the source was at about $1-2\times10^{39}$ ergs/sec."434 The aperture photometry from. FLEX for the 5 June 2002 observation gives a Lux level similar to that in the 1January 2004 observation. but the source was only about 20 above the background on 5 June 2002.," The aperture photometry from FLIX for the 5 June 2002 observation gives a flux level similar to that in the 1January 2004 observation, but the source was only about $2\sigma$ above the background on 5 June 2002."435 We tentatively trust. the aperture photometry results. in. part because they indicate," We tentatively trust the aperture photometry results, in part because they indicate"436these features is that in this part of the profile the signal is dominated by more highly delayed. components. making the periodic correlation C'(ó.7) contain signal power at larger |r| values than are present at earlier phases.,"these features is that in this part of the profile the signal is dominated by more highly delayed components, making the periodic correlation $C(\phi,\tau)$ contain signal power at larger $|\tau|$ values than are present at earlier phases."437 When transformed to the periodic spectrum domain. this results in [iner frequeney structure appearing at later pulse phases.," When transformed to the periodic spectrum domain, this results in finer frequency structure appearing at later pulse phases."438 Given the signal model presented in I5qn. 11..," Given the signal model presented in Eqn. \ref{eqn:model},"439 it is possible to determine both the SM. response and intrinsic pulse profile. directly from a single evelic spectrum., it is possible to determine both the ISM response and intrinsic pulse profile directly from a single cyclic spectrum.440" Phe two-dimensional evelie spectrum σταν) contains NeierNL cata values. while {νο(7) and Z(n) ave described bv only Nous|NS, model parameters."," The two-dimensional cyclic spectrum $S(\nu;\alpha_n)$ contains $N_{chan} \times N_{lag}^\prime$ data values, while $H_{ISM}(\nu)$ and $I(n)$ are described by only $N_{chan}+N_{lag}^\prime$ model parameters."441 This provides sullicient constraints for both to be determined via iterative least-squares minimization., This provides sufficient constraints for both to be determined via iterative least-squares minimization.442 X detailed. analysis of this method will be presented. in a separate paper (7).., A detailed analysis of this method will be presented in a separate paper \citep{walker:cyc}.443 One previous measurement. method for ο) has been published (?).. based on a dynamic spectrum. phase retrieval. procedure.," One previous measurement method for $h_{ISM}(t)$ has been published \citep{walker:hol}, based on a dynamic spectrum phase retrieval procedure."444 The evelic spectrum method is much simpler since in this case the wave phases can be measured. directly. giving an estimate of hs; from a single. “snapshot” observation.," The cyclic spectrum method is much simpler since in this case the wave phases can be measured directly, giving an estimate of $h_{ISM}$ from a single “snapshot” observation."445 The evelic spectrum also incorporates pulse. profile shape information. which is lost in standard dynamic spectra.," The cyclic spectrum also incorporates pulse profile shape information, which is lost in standard dynamic spectra."446 ‘This naturally lacis. for the first time. to a true coherently descattered pulse profile shape (Figure 3)).," This naturally leads, for the first time, to a true coherently descattered pulse profile shape (Figure \ref{fig:profs}) )."447 In contrast with previous intensitv-based profile deconvolution methods (?).. this requires no assumption of a specific functional form for hist ond is not alfected by ambiguity between intrinsic and ISM-induced: profile features.," In contrast with previous intensity-based profile deconvolution methods \citep{bhat:pbf}, this requires no assumption of a specific functional form for $h_{ISM}$ and is not affected by ambiguity between intrinsic and ISM-induced profile features."448 “Phe ESAL response shown in ligure 3 has an initial exponential decay followed by a more slowlv-decaving tail., The ISM response shown in Figure \ref{fig:profs} has an initial exponential decay followed by a more slowly-decaying tail.449" This will be interesting to compare in detail with the predictions of the standard. Ixolmogorov scattering mocel (o.g..ο),"," This will be interesting to compare in detail with the predictions of the standard Kolmogorov scattering model \citep[e.g.,][]{rjt+09}."450 There are several degeneracies that the descattering process alone can not resolve., There are several degeneracies that the descattering process alone can not resolve.451 As is ‘lear from Eqn. 1...," As is clear from Eqn. \ref{eqn:model},"452 multiplving {say by a constant. phase factor will produce no change in the observed eveclic spectrum., multiplying $H_{ISM}$ by a constant phase factor will produce no change in the observed cyclic spectrum.453 Similarly. without additional assumptions. the pulsar's intrinsic [ux Sy is degenerate. with the magnitude of £4;53j..," Similarly, without additional assumptions, the pulsar's intrinsic flux $S_0$ is degenerate with the magnitude of $H_{ISM}$."454 Most critically for pulsar timing. an arbitrary rotation can be applied to Z(ó). and absorbed into £s.," Most critically for pulsar timing, an arbitrary rotation can be applied to $I(\phi)$, and absorbed into $H_{ISM}$."455 That is. at this level of analysis it is impossible to distinguish an LSAI-induced delay. from a pulsar spin deviation or other timing ellect.," That is, at this level of analysis it is impossible to distinguish an ISM-induced delay from a pulsar spin deviation or other timing effect."456 Resolving this situation to obtain properly scattering-corrected timing will require the development of additional analysis techniques., Resolving this situation to obtain properly scattering-corrected timing will require the development of additional analysis techniques.457 This could range from assuming a constrained form. or applying moment analvsis to. fysay to physical models of the spatial distribution of scattering material. and is an active topic for further study.," This could range from assuming a constrained form or applying moment analysis to $h_{ISM}$ to physical models of the spatial distribution of scattering material, and is an active topic for further study."458 Compared with previous methods. the evelic spectrum. provides a qualitatively new way to measure the ISM response. and this new information dramatically expands the possibilities [or scattering corrections to timing data.," Compared with previous methods, the cyclic spectrum provides a qualitatively new way to measure the ISM response, and this new information dramatically expands the possibilities for scattering corrections to timing data."459 In addition to the determination of /;s. Eqn.," In addition to the determination of $H_{ISM}$, Eqn."460 9 allows for the appleation of other phase-coherent filters to the final integrated evelie spectra.," \ref{eqn:inout}461 allows for the applcation of other phase-coherent filters to the final integrated cyclic spectra."462 This technique can be used (ο perform coherent. dispersion corrections. within the limits of the evelie spectrum resolution.," This technique can be used to perform coherent dispersion corrections, within the limits of the cyclic spectrum resolution."463 Phe high frequency resolution of the pulsar cvelic spectrum. enables. precise »ost-detection. removal of narrow-band. radio. frequency interference. without sacrificing pulse phase resolution.," The high frequency resolution of the pulsar cyclic spectrum enables precise post-detection removal of narrow-band radio frequency interference, without sacrificing pulse phase resolution."464 More sophisticated approaches may also incorporate information eained [rom the evelostationaritv of the interference (7).., More sophisticated approaches may also incorporate information gained from the cyclostationarity of the interference \citep{feliachi:phd}.465 Although the discussion in refsec:analysis focused on the analysis of a single stochastic signal. pairs of correlated. evclostationary signals can be analyzed via evelic erossespectra. in complete analogy with standard. cross-spectra (2)...," Although the discussion in \\ref{sec:analysis} focused on the analysis of a single stochastic signal, pairs of correlated cyclostationary signals can be analyzed via cyclic cross-spectra, in complete analogy with standard cross-spectra \citep{gardner:book}."466 For. dual-polarization radio data. this results in the creation of evelic Stokes parameters. and allows for the application of phase-coherent matrix convolution (7)— directly. to the evelie spectra.," For dual-polarization radio data, this results in the creation of cyclic Stokes parameters, and allows for the application of phase-coherent matrix convolution \citep{straten:phase} directly to the cyclic spectra."467 For radio interferometers. analvzing cata from antenna pairs will produce evelic visibilities.," For radio interferometers, analyzing data from antenna pairs will produce cyclic visibilities."468 Along with recently developed VLBI imaging techniques for investigating pulsar scintillation (?7).. this could. prove to be an extremely powerful tool for exploring the ΙΔ.," Along with recently developed VLBI imaging techniques for investigating pulsar scintillation \citep{brisken:vlbi_arcs}, this could prove to be an extremely powerful tool for exploring the ISM."469 Cyclic spectral analysis is a powerful new observational echnique for studying radio pulsars., Cyclic spectral analysis is a powerful new observational technique for studying radio pulsars.470 Lt provides a clata representation that simultaneously preserves both high pulse shase resolution and high frequency resolution information about the signal., It provides a data representation that simultaneously preserves both high pulse phase resolution and high frequency resolution information about the signal.471 With the accompanying preservation of signal phase content. this allows fundamentally new analysis echniques not possible with standard filterbank data.," With the accompanying preservation of signal phase content, this allows fundamentally new analysis techniques not possible with standard filterbank data."472 This olds promise both for increasing our understanding of the ionized ΔΙ. and eventually for removing ISM scattering as an obstacle to achieving the highest. possible pulsar timing oxrecision.," This holds promise both for increasing our understanding of the ionized ISM, and eventually for removing ISM scattering as an obstacle to achieving the highest possible pulsar timing precision."473Quasars often exhibit narrow heavy element absorption lines near their emission redshift.,Quasars often exhibit narrow heavy element absorption lines near their emission redshift.474 One possible explanation for the origin of these systems is that they arise in clouds of matter associated with galaxies in the clusters surrounding the quasars., One possible explanation for the origin of these systems is that they arise in clouds of matter associated with galaxies in the clusters surrounding the quasars.475 Alternate possibility is that they originate in the clouds that are physically associated with the quasars themselves., Alternate possibility is that they originate in the clouds that are physically associated with the quasars themselves.476 It has been shown that both these scenarios may be true EEllingson et citeell1994:; Hamann et citeham 19972)., It has been shown that both these scenarios may be true Ellingson et \\cite{ell1994}; Hamann et \\cite{ham1997a}) ).477 Clusters of absorbers of any type are of particular interest., Clusters of absorbers of any type are of particular interest.478 On one hand. clustering. properties. of intervening. systems depend strongly on the type of absorbers. while on the other hand complexes of associated absorbers give a unique opportunity to analyze the properties of their hosts and of the quasar emission.," On one hand, clustering properties of intervening systems depend strongly on the type of absorbers, while on the other hand complexes of associated absorbers give a unique opportunity to analyze the properties of their hosts and of the quasar emission."479 The stumbling block in such studies is the fact that bright high redshift quasars with rich absorption are rare (for examples of such systems see. e.g.. Morris et citemor1986:: Foltz et citefol 987:: Petitjean et citepet1994:;; Hamann et citeham1997a:; Lespine Petitjean 1997:: Petitjean Srianand 1999)).," The stumbling block in such studies is the fact that bright high redshift quasars with rich absorption are rare (for examples of such systems see, e.g., Morris et \\cite{mor1986}; ; Foltz et \\cite{fol1987}; Petitjean et \\cite{pet1994}; Hamann et \\cite{ham1997a}; Lespine Petitjean \cite{les1997}; Petitjean Srianand \cite{pet1999}) )."480 There are only a few quasars known with more than just a handful of associated absorption systems which are bright enough to perform high resolution spectral analysis., There are only a few quasars known with more than just a handful of associated absorption systems which are bright enough to perform high resolution spectral analysis.481 In this paper we present a bright τμ=2.51 quasar 1160343820 (eyosy=16:03:07.7.04193) +38:20:07) with a very rich metal absorption spectrum.," In this paper we present a bright $z_{\rm em}=2.51$ quasar 1603+3820 $\alpha_{\rm 1950} = \mbox{16:03:07.7}, \delta_{\rm4821950} = \mbox{+38:20:07}$ ) with a very rich metal absorption spectrum."483 It was discovered during the course of the Hamburg/CfA Bright Quasar Survey (Hagen et 1995;; Dobrzycki et citedob1996)).Quasar candidates in the survey are selected, It was discovered during the course of the Hamburg/CfA Bright Quasar Survey (Hagen et \\cite{hag1995}; ; Dobrzycki et \\cite{dob1996}) ).Quasar candidates in the survey are selected484The research work of ο. D. is supported by the University. Grants Comunission. Government of India Grant No.,"The research work of S. D. is supported by the University Grants Commission, Government of India Grant No."485 3432/2008 (SR)., 34-32/2008 (SR).486" 5. G. and ο, V. acknowledges (he financial support provided by Council for Scientific and Industrial Research (CSIR) and University Grants Commission (UGC). Government of India. respectively."," S. G. and S. V. acknowledges the financial support provided by Council for Scientific and Industrial Research (CSIR) and University Grants Commission (UGC), Government of India, respectively."487from the emergent fluxes computed with the SYNSPEC code. applying Eq.,"from the emergent fluxes computed with the SYNSPEC code, applying Eq."488 5 for individual rotational phases., \ref{velik} for individual rotational phases.489 To study the influence of individual elements separately. we first calculated the light variations with the abundance map of one element only (Fig. 4)).," To study the influence of individual elements separately, we first calculated the light variations with the abundance map of one element only (Fig. \ref{prv_hvvel}) ),"490 assuming a fixed abundance of other elements (eq.= —1.0. eg=-3.75. ος=-5.9. εις= —44)," assuming a fixed abundance of other elements $\varepsilon_\text{He}=-1.0$ , $\varepsilon_\text{Si}=-3.75$, $\varepsilon_\text{Cr}=-5.9$, $\varepsilon_\text{Fe}=-4.4$ )."491 From Fig., From Fig.492 4. it follows that iron. silicon. and chromium contribute most to the light variations. while the contribution of helium is only marginal.," \ref{prv_hvvel} it follows that iron, silicon, and chromium contribute most to the light variations, while the contribution of helium is only marginal."493 This ts because of the large overabundance of these elements in the spots and by their large abundance variations on the stellar surface., This is because of the large overabundance of these elements in the spots and by their large abundance variations on the stellar surface.494" The amplitude of the light variations increases with decreasing wavelength. as can be expected from the plot of the magnitude difference Aun, in Fig. 3.."," The amplitude of the light variations increases with decreasing wavelength, as can be expected from the plot of the magnitude difference $\Delta m_\lambda$ in Fig. \ref{magtoky}. ."495 Because the overabundant regions are brighter in the (ον colours. the predicted light variations. reflect the equivalent width variations1).," Because the overabundant regions are brighter in the $uvby$ colours, the predicted light variations reflect the equivalent width variations."496. The light maximum occurs at the same phase at which the equivalent widths of a given element are the largest., The light maximum occurs at the same phase at which the equivalent widths of a given element are the largest.497 Because this happens at slightly different phases for individual elements. the light curves in Fig.," Because this happens at slightly different phases for individual elements, the light curves in Fig."498 4+ are slightly shifted., \ref{prv_hvvel} are slightly shifted.499 Taking into account the surface distribution of helium. silicon. chromium. and iron m the calculation. of the light curves (Fig. 5)).," Taking into account the surface distribution of helium, silicon, chromium, and iron in the calculation of the light curves (Fig. \ref{cuvir_hvvel}) ),"500 we obtained a good agreement between the observed and predicted light curves in the v. 6 and y bands of the Strómmgren photometric system.," we obtained a good agreement between the observed and predicted light curves in the $v$, $b$ and $y$ bands of the Strömmgren photometric system."501 On the other hand. our models are able to explain only about half of the amplitude in the 4 filter.," On the other hand, our models are able to explain only about half of the amplitude in the $u$ filter."502 The disagreement between the predicted and observed light curves is mostly apparent around phase @=0.6., The disagreement between the predicted and observed light curves is mostly apparent around phase $\phi=0.6$.503 ote also that a similar disagreement visible in 4 can be also found in other filters. but to a much smaller extent.," Note also that a similar disagreement visible in $u$ can be also found in other filters, but to a much smaller extent."504 These differences clearly point to an existence of some additional. unknown mechanism working especially in the violet band that still needs to be investigated (Fig. 5..," These differences clearly point to an existence of some additional, unknown mechanism working especially in the violet band that still needs to be investigated (Fig. \ref{cuvir_hvvel},"505 and see also Sect. 7))., and see also Sect. \ref{kecame}) ).506 The discrepancies between the predicted ad observed light curves increase when comparing the predicted and observed colour indices (see Fig. 6))., The discrepancies between the predicted and observed light curves increase when comparing the predicted and observed colour indices (see Fig. \ref{cuvir_uvby}) ).507 While the (5—v) data agree reasonably well. the (v—5) curves are mutually shifted. and the predicted metallic index gj=(v—5)—(b—xy) shows a significantly lower amplitude than the observed one.," While the $(b-y)$ data agree reasonably well, the $(v-b)$ curves are mutually shifted, and the predicted metallic index $m_{\text{1}}=(v-b)-(b-y)$ shows a significantly lower amplitude than the observed one."508 The inhomogeneous surface distribution of individual elements causes bright spots on the stellar surface., The inhomogeneous surface distribution of individual elements causes bright spots on the stellar surface.509 The spots. whose surface distribution can be derived using abundance maps and model atmospheres (see Fig. 7)).," The spots, whose surface distribution can be derived using abundance maps and model atmospheres (see Fig. \ref{cuvir_povrch}) ),"510 cause the light variability., cause the light variability.511 We have shown that the light variability of iis caused by the redistribution of flux from the far UV to the near UV and visible regions., We have shown that the light variability of is caused by the redistribution of flux from the far UV to the near UV and visible regions.512 Consequently. the light variability in the farUV regionshould be in antiphase with the visual one.," Consequently, the light variability in the farUV regionshould be in antiphase with the visual one."513 This behaviour was indeed found in a detailed analysis of IUE, This behaviour was indeed found in a detailed analysis of IUE514disces are derived by 2.. assuming saturation docs not occur vet.,"discs are derived by \citet{pbck09}, assuming saturation does not occur yet."515 7. reexamined the corotation torque in acliabatic disces. taking into account the effects. of pressure.," \citet{mc09} reexamined the corotation torque in adiabatic discs, taking into account the effects of pressure."516 Γον showed that the. entropy-relatecl torque does. not arise. from. the overdense ancl underdense regions. as had. been previously thought.," They showed that the entropy-related torque does not arise from the overdense and underdense regions, as had been previously thought."517 Instead. they showed. that entropy. gradients make the horseshoe region considerably asymmetric. and this asymmetry. exerts the entropy-related torque by. the excitation of evanescent waves at the horseshoe sparatrices.," Instead, they showed that entropy gradients make the horseshoe region considerably asymmetric, and this asymmetry exerts the entropy-related torque by the excitation of evanescent waves at the horseshoe sparatrices."518 Since these evanescent waves are excited by pressure disturbance that are a result. of horseshoe orbits. the magnitude of the entropy-relatecl torque depends on the perturbed pressure.," Since these evanescent waves are excited by pressure disturbance that are a result of horseshoe orbits, the magnitude of the entropy-related torque depends on the perturbed pressure."519 This gives an explanation of the results of? that the corotation torque in a very cold. acliabatic disc is identical to that in an isothermal disc.," This gives an explanation of the results of \citet{bm08} that the corotation torque in a very cold, adiabatic disc is identical to that in an isothermal disc."520 Thus. the corotation torque or horseshoe drag can be important only for disces assumed to be adiabatic with high emperatures. that is. the radiative cooling timescale in disces is longer than the planet's orbital period.," Thus, the corotation torque or horseshoe drag can be important only for discs assumed to be adiabatic with high temperatures, that is, the radiative cooling timescale in discs is longer than the planet's orbital period."521 Otherwise. he direction of migration is determined by the Lindblad orque.," Otherwise, the direction of migration is determined by the Lindblad torque."522 The horseshoe drag. however. has some problems.," The horseshoe drag, however, has some problems."523 One is the width of the horseshoe region. and the other is saturation.," One is the width of the horseshoe region, and the other is saturation."524 The former is characterised by the mass of the janets. but it is parameterized in numerical simulations xuwily because a planet is treated: as a point mass object (77) and partly because any simulation cannot infinitely resolve the horseshoe region (??)..," The former is characterised by the mass of the planets, but it is parameterized in numerical simulations partly because a planet is treated as a point mass object \citep{dkh03,pbck09} and partly because any simulation cannot infinitely resolve the horseshoe region \citep{mdk06,pp09}."525 For the latter problem. he horseshoe drag can be saturated (ic. negligible) due to acdiabatie invariance.," For the latter problem, the horseshoe drag can be saturated (i.e. negligible) due to adiabatic invariance."526 We emphasise that the problem of saturation is essentially how the disc viscosity connects the horseshoe region to the rest of disc by transferring angular momentunm. between them (e.g.2).., We emphasise that the problem of saturation is essentially how the disc viscosity connects the horseshoe region to the rest of disc by transferring angular momentum between them \citep[e.g.][]{mc10}.527 This is crucial for understanding horseshoe drags., This is crucial for understanding horseshoe drags.528 As a example. we summarise the work done bv ? who considered the saturation of the vortensitv-related and. entropy-related: horseshoc drags.," As a example, we summarise the work done by \citet{pbk10} who considered the saturation of the vortensity-related and entropy-related horseshoe drags."529 For the vortensity-related. one. the saturation is controlled. only bv. viscous cillusion (alsosee2?7).," For the vortensity-related one, the saturation is controlled only by viscous diffusion \citep[also see][]{m01,m02,ward07}."530. For the disces with a~LO he horseshoe drag becomes almost zero while it is unsaturated for the disc with a~0.01 (?)..," For the discs with $\alpha \sim 10^{-5}$, the horseshoe drag becomes almost zero while it is unsaturated for the disc with $\alpha \sim 0.01$ \citep{pp08}."531 ? (7)..," \citet{pbk10} \citep{pbk10},"532"tendency to overestimate the error as v becomes low; the ση, bias has no dependence on the pixel sampling.",tendency to overestimate the error as $\nu$ becomes low; the $\tilde{\sigma}_{\eta_i}$ bias has no dependence on the pixel sampling.533 The v dependence is approximately πω..., The $\nu$ dependence is approximately $\sqrt{\langle\Delta\eta_i^2/\tilde{\sigma}_{\eta_i}^2\rangle}\simeq $ $1-\frac{0.6}{\nu}$.534" Figure 7 shows the accuracy of a deconvolution fit when a symmetric, exponential galaxy is convolved with an Airy PSF."," Figure \ref{fig:dcvlshapeerror100} shows the accuracy of a deconvolution fit when a symmetric, exponential galaxy is convolved with an Airy PSF."535" We first discuss a high-S/N case, v=100."," We first discuss a $S/N$ case, $\nu=100$."536" The measurements were done at various minor axis resolutions, where rp=0.5, 1, 2, and 5 each correspond to the columns from left to right."," The measurements were done at various minor axis resolutions, where $r_b=0.5$, 1, 2, and 5 each correspond to the columns from left to right."537 The panels in the top two rows show the shape error (An;) as a function of the input galaxy ellipticity., The panels in the top two rows show the shape error $\langle\Delta\eta_i\rangle$ as a function of the input galaxy ellipticity.538" In the top row, the Airy PSF is circular (epsr= 0.0), while in the middle row the Airy PSF is anisotropic, with an ellipticity of"," In the top row, the Airy PSF is circular $e_{\rm PSF}=0.0$ ), while in the middle row the Airy PSF is anisotropic, with an ellipticity of"539he aperture used. to avoid confusion we do not give a ucasured excess value. oulv the plotospheric fux which can be subtracted from azuv later measurcieuts.,"the aperture used, to avoid confusion we do not give a measured excess value, only the photospheric flux which can be subtracted from any later measurements."540 The shotospheric flux given in Table 1 does not include the coutributiou from the C dwarf (90 aud 10τιν at 51 andl TOyan. respectively).," The photospheric flux given in Table \ref{tab:par} does not include the contribution from the G dwarf $90$ and $10 ~{\rm mJy}$ at $24$ and $70 ~\micron$, respectively)."541 The top left panel iu Figure 1 shows he sumuned iuaege from epochs 2 aud 3. to demoustrate he asviunietrv sueeested even before PSF subtraction.," The top left panel in Figure 1 shows the summed image from epochs 2 and 3, to demonstrate the asymmetry suggested even before PSF subtraction."542 For the fist epoch 21jan nuage. the refercuce star lage was subtracted from the image of 6 Velormu. with a scale factor chosen as the masxinuun value that would completely remove the nuage core without creating siguificant negative flux residuals.," For the first epoch $24 ~\micron$ image, the reference star image was subtracted from the image of $\delta$ Velorum, with a scale factor chosen as the maximum value that would completely remove the image core without creating significant negative flux residuals."543 The deeper exposures from the secoud and third epochs were designed to reveal faint structures far from the star. where the observed PSF is difficult to extract accurately.," The deeper exposures from the second and third epochs were designed to reveal faint structures far from the star, where the observed PSF is difficult to extract accurately."544 Therefore. we used simulated PSFs (from STinwTim 2002))) aud the AQPS simulator7.," Therefore, we used simulated PSFs (from STinyTim ) and the MIPS simulator."545. Because bright structures nearly iu he PSF contribute to the residuals at large distances. we oversubtracted the PSF to compensate.," Because bright structures nearly in the PSF contribute to the residuals at large distances, we oversubtracted the PSF to compensate."546 The first epoch PSF subtracted 21jan image is shown iu the bottom xuels of Figure 1. aud the composite from epochs 2 aud 3 in the upper right., The first epoch PSF subtracted $24 ~\micron$ image is shown in the bottom panels of Figure \ref{fig:im} and the composite from epochs 2 and 3 in the upper right.547 The PSF subtracted nuages in Figure l1. show that he asvuuuetrv is caused bv a bow shock., The PSF subtracted images in Figure \ref{fig:im} show that the asymmetry is caused by a bow shock.548 As shown in the lower left. the head of the bow shock points approximately toward the direction of the stellar proper notion.," As shown in the lower left, the head of the bow shock points approximately toward the direction of the stellar proper motion."549 The bottom right paucl shows the excess fiux contours and that it consists of incomplete spherical shells centered ou à. Velorzuu., The bottom right panel shows the excess flux contours and that it consists of incomplete spherical shells centered on $\delta$ Velorum.550 Combined with the upper right damage. there is also a parabolic cavity. as expected for a bow shock.," Combined with the upper right image, there is also a parabolic cavity, as expected for a bow shock."551 The stagnation points (where photon pressure equals gravitational force) of the grains in the bow shock are within ~200AU of the star. according to the observations.," The stagnation points (where photon pressure equals gravitational force) of the grains in the bow shock are within $\sim 200 ~{\rm AU}$ of the star, according to the observations."552 A notable feature in the upper right is the wines of the bow shock. which are detectable to 1500AU.," A notable feature in the upper right is the wings of the bow shock, which are detectable to $\sim 1500 ~{\rm AU}$."553 The 70ji observation is shown iu Figure 2.., The $70 ~\micron$ observation is shown in Figure \ref{fig:70}.554 The PSF subtraction (scaled to the point source flux of 125 αι) does uot reveal the bow shock structure at this wavelength. only that there is extended excess.," The PSF subtraction (scaled to the point source flux of $125 ~{\rm mJy}$ ) does not reveal the bow shock structure at this wavelength, only that there is extended excess."555 The total flux of the residual of the PSF subtracted inage is 119wy., The total flux of the residual of the PSF subtracted image is $119 ~{\rm mJy}$.556 The iuteusitv coutowurs panel) sugeest that the 70pin excess fades at the cavity behind the star. but the effect is small.," The intensity contours ) suggest that the $70 ~\micron$ excess fades at the cavity behind the star, but the effect is small."557 The ecometry aud direction of the bow shock are discussed in more detail iu 83.2., The geometry and direction of the bow shock are discussed in more detail in 3.2.558 Based on a previous sugeestion bv(1990).. proposed a physical model to explain the abundance— pattern of A Boottis stars through πανΔΙ interaction and the diffusion/accretion hvpothesis.," Based on a previous suggestion by, proposed a physical model to explain the abundance pattern of $\lambda$ Boöttis stars through star-ISM interaction and the diffusion/accretion hypothesis."559 Then imeocel is based on a lunimnous nmnaiu-sequeuce star passing through a diffuse ISM. cloud., Their model is based on a luminous main-sequence star passing through a diffuse ISM cloud.560 The star blows the interstellar dust eraius away by its radiation pressure. but accretes the interstellar gas onto its surface. thus establishing a thin surface laver with abundance anomalics.," The star blows the interstellar dust grains away by its radiation pressure, but accretes the interstellar gas onto its surface, thus establishing a thin surface layer with abundance anomalies."561 So long as the star is inside the cloud. the dust erains are heated to produce excess in the infrared above the photospheric radiation of the star.," So long as the star is inside the cloud, the dust grains are heated to produce excess in the infrared above the photospheric radiation of the star."562 Martfunez-CGalarza et ((2007. in prep.)," Martínnez-Galarza et (2007, in prep.)"563 have developed a model of this process aud show that the elobal spectral energy. distributions of a eroup of A Boottis type stars that have infrared) excesses are consistent with the emission from the hivpothesized ISM cloud., have developed a model of this process and show that the global spectral energy distributions of a group of $\lambda$ Boöttis type stars that have infrared excesses are consistent with the emission from the hypothesized ISM cloud.564 Details of the mocel can be found in their paper., Details of the model can be found in their paper.565" Tere we adapt their model anc improve its fidelity (c.e.. with higher resolution itceratious). aud also model the surface brightuess distribution to describe the observed bow shock seen around 6 ολοι,"," Here we adapt their model and improve its fidelity (e.g., with higher resolution integrations), and also model the surface brightness distribution to describe the observed bow shock seen around $\delta$ Velorum."566 The phenomenon of stzr-ISM interactions generating vow shocks was first studied by., The phenomenon of star-ISM interactions generating bow shocks was first studied by.567(1997)... They showed that the radiative pressure force on a subanicron dust erain can be many times that of he eravitational force as it approaches the star., They showed that the radiative pressure force on a sub-micron dust grain can be many times that of the gravitational force as it approaches the star.568 The scattering surface will be a parabola with the star at the focus point of the parabolic shaped dust cavity., The scattering surface will be a parabola with the star at the focus point of the parabolic shaped dust cavity.569 Since he star heats the erains outside of the cavity and close o the parabolic surface. an infrared-emitting bow shock cature is expected.," Since the star heats the grains outside of the cavity and close to the parabolic surface, an infrared-emitting bow shock feature is expected."570 The shape of the parabola (for cach erain size) cau ο given in terms of the distance between the star (focus) aud the vertex., The shape of the parabola (for each grain size) can be given in terms of the distance between the star (focus) and the vertex.571 This so-called avoidance radius (or the p/2 parameter of the scattering parabola) can be calculated from euerev. conservation to be1997): where 6 is the radius of the particle. AZ is the umass of the star ancl ey is the relative velocity between the star and the dust eraius.," This so-called avoidance radius (or the $p/2$ parameter of the scattering parabola) can be calculated from energy conservation to be: where $a$ is the radius of the particle, $M$ is the mass of the star and $v_{\rm rel}$ is the relative velocity between the star and the dust grains."572" “as the ratio of photon pressure to eravitatioual force on a erain and it is eiven by1979): where à is the bulk density of the erain material and QV, is the radiation pressure oficieucy averaged over tlie stellar spectrum.", $\beta^a$ is the ratio of photon pressure to gravitational force on a grain and it is given by: where $\delta$ is the bulk density of the grain material and $Q_{\rm pr}^a$ is the radiation pressure efficiency averaged over the stellar spectrum.573" QUU) can be expressed in terns of erain. properties. PRSEscattering. cocfiicicnt QU.[aan(A) and the scattering asviunietry.(Qu, factor"," $Q_{\rm pr}^a(\lambda)$ can be expressed in terms of grain properties )$, scattering coefficient $Q_{\rm sca}^a(\lambda)$ and the scattering asymmetry factor"574where the longitudinal and latitucinal components of j£ are proportional to the acceleration of the solar svstem projected on the celestial plane assumes that the acceleration of the solar svstem barvcenter has only dy = 4 component. (he equations (5)) (7)) for the elective proper motion caused by the secular aberration are simplified and reduced to Proper motion vectors j£ of a given number of objects represent a discrete vector field on the sphere which can be decomposed in a set of vector spherical harmonics (Thorne.1980)..,"where the longitudinal and latitudinal components of $\vec\mu$ are proportional to the acceleration of the solar system projected on the celestial plane If one assumes that the acceleration of the solar system barycenter has only $A_X=A$ component, the equations \ref{2a}) \ref{3a}) ) for the effective proper motion caused by the secular aberration are simplified and reduced to Proper motion vectors $\vec\mu$ of a given number of objects represent a discrete vector field on the sphere which can be decomposed in a set of vector spherical harmonics \citep{thorne}."575 The largest galactocentric component of (he secular aberration can be determined [rom global astromeltric observations as à svsteniatic dipole component of this vector field., The largest galactocentric component of the secular aberration can be determined from global astrometric observations as a systematic dipole component of this vector field.576 For quasars the problem of its determination is simpler (han for stars since {μον have negligible small proper motions caused by their peculiar velocities with respect to the Hubble flow., For quasars the problem of its determination is simpler than for stars since they have negligibly small proper motions caused by their peculiar velocities with respect to the Hubble flow.577 Therefore. the secular aberration can be directly measured Irom the observed proper motions of quasars.," Therefore, the secular aberration can be directly measured from the observed proper motions of quasars."578 The magnitude of the secular aberration effect in (he case givenby Eq. (8)), The magnitude of the secular aberration effect in the case givenby Eq. \ref{4}) )579 is where sin€=V1—cos?/b and ὁ is the angle between the direction towards the galactic center and (hat to the stir.," is \citep{gaia-report}580 where $\sin\zeta=\sqrt{1-\cos^2l\cos^2b}$ and $\zeta$ is the angle between the direction towards the galactic center and that to the star."581 In what follows. we investigate a more accurate approximation of Eq. (8))," In what follows, we investigate a more accurate approximation of Eq. \ref{4}) )"582 that includes all three components of the Sun's acceleration. ancl evaluate the effect of the secular aberration more adequately in terms of vector harmonics.," that includes all three components of the Sun's acceleration, and evaluate the effect of the secular aberration more adequately in terms of vector harmonics."583 The velocity vector of the Sun in the galaxy is commonly considered (o be the sum of two components (Dinnev&Merrifield1998)., The velocity vector of the Sun in the galaxy is commonly considered to be the sum of two components \citep{binney}.584. The first (and largest) component is the motion ol the so-called. Local Standard of Rest (LR).," The first (and largest) component is the motion of the so-called, Local Standard of Rest (LSR)."585 By definition (Dinnev&Merrifield1993).. the LSR. is involved in a cireular planar motion around the center of mass of the galaxy. at a constant rate wilh a period £7.," By definition \citep{binney}, the LSR is involved in a circular planar motion around the center of mass of the galaxy at a constant rate with a period $P_0$ ."586 The second component of the solar velocity is the differential, The second component of the solar velocity is the differential587he discrepauey. between the two colors is only at the 20 evel.,the discrepancy between the two colors is only at the $\sigma$ level.588 The colors predicted from a 0.2 keV. metallicity ISM are (C21.032 =0.51.0.22). whereas the colors for a 13 keV. motallicity ISM are (021.032 =1.17.0.16).," The colors predicted from a 0.2 keV, metallicity ISM are (C21,C32 $= 0.54, 0.22)$, whereas the colors for a 0.3 keV, metallicity ISM are (C21,C32 $= 1.17, 0.46)$."589" ""Thus. if any ISM is preseut in NGC 1697. its temperature uust be below 0.3 keV and at a low metallicity. or else he C21 color of the LAINBs|ISM would be higher than he LAINBs alone. which is not observed."," Thus, if any ISM is present in NGC 4697, its temperature must be below 0.3 keV and at a low metallicity, or else the C21 color of the LMXBs+ISM would be higher than the LMXBs alone, which is not observed."590 In conclusion. we cannot rule out the prescuce of some low temperature ISM in NGC L697. although it is certain that an ISAL cannot constitute a majority of the cussion.," In conclusion, we cannot rule out the presence of some low temperature ISM in NGC 4697, although it is certain that an ISM cannot constitute a majority of the emission."591 Lwiu Breeian (19995) found that Galactic and MOI elobular cluster LAINB ταν colors were correlated with the aetallicity of the globular cluster. in the sense that lugher metallicity globular clusters had LATINBs with softer N-vav colors.," Irwin Bregman (1999b) found that Galactic and M31 globular cluster LMXB X-ray colors were correlated with the metallicity of the globular cluster, in the sense that higher metallicity globular clusters had LMXBs with softer X-ray colors."592 If this correlation extended to all LAINBs. it would predict that the N-vav colors of NGC L697 should harden with increasing radius. since metallicity decreases with radius iu elliptical galaxies.," If this correlation extended to all LMXBs, it would predict that the X-ray colors of NGC 4697 should harden with increasing radius, since metallicity decreases with radius in elliptical galaxies."593 Furthermore. the metallicity of NCC 1697 is rather low.," Furthermore, the metallicity of NGC 4697 is rather low."594 Within half an effective radius. the average ietallicity is onlv solar (Trager et 22000).," Within half an effective radius, the average metallicity is only solar (Trager et 2000)."595 The metallicity-color relation of Irwin Bregman (1999b) would predict a C32 color of about 1.5 within half au effective radius. aud au increase with increasing radius.," The metallicity-color relation of Irwin Bregman (1999b) would predict a C32 color of about 1.5 within half an effective radius, and an increase with increasing radius."596 The colors of the four resolved sources as well as the unresolved cuiissiou are at odds with this inetallicitv-color relation., The colors of the four resolved sources as well as the unresolved emission are at odds with this metallicity-color relation.597 Appaveuthy. if such a metallicity-color relation truly exists for LAINBs. it oulv applies to LAINBs that reside iuglobular clusters.," Apparently, if such a metallicity-color relation truly exists for LMXBs, it only applies to LMXBs that reside inglobular clusters."598 The X-ray source in NGC 1697 associated with a globular cluster (Source 11) has a C32 color of 1.11250.30. which would be consistent with the metallicity-color relation ifthe elobular cluster has a high metallicity.," The X-ray source in NGC 4697 associated with a globular cluster (Source 11) has a C32 color of $1.14 \pm 0.30$, which would be consistent with the metallicity-color relation if the globular cluster has a high metallicity."599" We also investigated the possibility that LAINBs below the detection threshold of the IIRI could account for the ""resolved eimissiou eiven a reasonable LAINB Iuninosityv distribution function for NGC 1697.", We also investigated the possibility that LMXBs below the detection threshold of the HRI could account for the unresolved emission given a reasonable LMXB luminosity distribution function for NGC 4697.600 Wo assuued a Dhuunositv distribution function ΑςLy)xLU. which is consistent within the errors with the bhuuinositv distribution function of pointsources in MOI with huuinositics greater than 2«LO cress3| (Primini ct 11993).," We assumed a luminosity distribution function $N(>L_X) \propto L_X^{-1.3}$, which is consistent within the errors with the luminosity distribution function of pointsources in M31 with luminosities greater than $2 \times 10^{37}$ ergs $^{-1}$ (Primini et 1993)."601 The function was normalized to vield the ohservec N-rav luminosity of NGC 1697 when inteerated over al LMXD hiuuinosities., The function was normalized to yield the observed X-ray luminosity of NGC 4697 when integrated over all LMXB luminosities.602 This model predicted 11 sources with uuivositics over 3«LOPS eyes 1; which coutributed o the total N-rav enissiou from LAINBs.," This model predicted 11 sources with luminosities over $3 \times 10^{38}$ ergs $^{-1}$, which contributed to the total X-ray emission from LMXBs."603 This agrees wel with what was observed with the URI: neglecting the poiut source associated with au unidentified optical counterpart. he remaining 11 detected sources comprised of he total emission.," This agrees well with what was observed with the HRI; neglecting the point source associated with an unidentified optical counterpart, the remaining 11 detected sources comprised of the total emission."604 Thus. if the luminosity distribution muction of LAINBs of NGC. L697 is similar to that of he brighter LAINBs in M2. the integrated eimissiou frou LAINBs below the detection threshold of the IIRI cau account for most of the uuresolveck emission.," Thus, if the luminosity distribution function of LMXBs of NGC 4697 is similar to that of the brighter LMXBs in M31, the integrated emission from LMXBs below the detection threshold of the HRI can account for most of the unresolved emission."605 Tuterestingly. the detection limit of the IIRI observation of 3«1075 eres 1 Bes above the Eddinetou luminosity laut for a 1.1AZ. ueutrou star.," Interestingly, the detection limit of the HRI observation of $3 \times 10^{38}$ ergs $^{-1}$ lies above the Eddington luminosity limit for a $1.4~M_{\odot}$ neutron star."606 LAINBs of simular huninosities as the ones found here exist in our own Galaxy., LMXBs of similar luminosities as the ones found here exist in our own Galaxy.607"A compilation by Christian Swank (1997) found cight galactic LAINBs with Iuninosities ereater than Jos10"" ores + (we have converted their 0.71.5 keV huninosities to 0.25-10 keV. huninositics usine the spectral model of 3)).",A compilation by Christian Swank (1997) found eight galactic LMXBs with luminosities greater than $3 \times 10^{38}$ ergs $^{-1}$ (we have converted their 0.7–4.5 keV luminosities to 0.25-10 keV luminosities using the spectral model of \ref{sec:spectral}) ).608 These high Iuninosities imply either that he compact object within the binary is a black hole (with Aipyc6AL.. for the most huuiuous binaries) or that he huninositics truly exceed the Edcdiugton limit for a jeutron star., These high luminosities imply either that the compact object within the binary is a black hole (with $M_{BH} \ge 6~M_{\odot}$ for the most luminous binaries) or that the luminosities truly exceed the Eddington limit for a neutron star.609 The former would iuplv that active binaries with massive black holes are fairly couunon in galaxies., The former would imply that active binaries with massive black holes are fairly common in galaxies.610 It should be noted that the bulee of M31 lacks the very ugh huninosity LAINBs that NGC 1697 has: the brightest LMXDB in MO was oulv L8&1075 eres | (Supper et 11997)., It should be noted that the bulge of M31 lacks the very high luminosity LMXBs that NGC 4697 has; the brightest LMXB in M31 was only $1.8 \times 10^{38}$ ergs $^{-1}$ (Supper et 1997).611 Tlowever. this is likely the result of siuall πο statistics.," However, this is likely the result of small number statistics."612 Siniulatious of the bulge of M31 using a unnuimositv distribution fiction of the form INN(5Ly)x⊉∖⊳⋟∐↓≼∐↸⊳⋜↧↑↸∖≼↧↑∐⋜↧↑∪∐⋅↖↽⋅≩↴∖↴≺∏∐⋅↸⊳↸∖↴∖↴↖↖↽↕↑∐↕⋯⊔∐⋯∖↴↕⊓↸∖↴∖↴LBs ↽∙ ⋅ ⋅⋅⋅ ↸∖⊼↸⊳↸∖↸∖≼∐∐∶↴⋁∐∣⋮⋝↖↸∖↥⋅∶↴↜⋱∖↴↴∖↴↴∖↴↓∪∏↕≺∏⋝↸∖↕≯∪∏∐≼⊔∐⋀∖↕∶∐∙↕∐∐↖⇁↸, Simulations of the bulge of M31 using a luminosity distribution function of the form $N(>L_X) \propto L_X^{-1.3}$ indicated that only 1–3 sources with luminosities exceeding $10^{38}$ ergs $^{-1}$ should be found in M31.613∖ ↴∖↴↸∖↻⋜∐⋅⋜↧↑↸∖↴∖↴↕∐∐∏⋜↧↑↕∪∐↴∖↴∪↕⋟⋀∖↕∶≩↽∙↑∐↸∖↻↸∖⋜∐↘↽↕∏∐∏∐∪↴∖↴↕↑⋅↖↽↕⋟∪↥⋅⋜⋯ ∫⇀⋀∖↕⊸∖⊽↕≧≼∐≼∐∪↑↸∖⊼∩∖↸∖≼⇂∶≩∖↕∩⋮⋝⋉↸∖↥⋅∶↴∙∷∖↴↴∖↴↓∙↕∐⋜↧∶↴∙⊾↥⋅↸∖↸∖⋯↸∖∐↑ with observation.," In five separate simulations of M31, the peak luminosity for an LMXB did not exceed $3 \times 10^{38}$ ergs $^{-1}$, in agreement with observation."614 We cannot rule out the presence of at least some iutersteHar imuediuni in NCC 1697 aud X-rav faint type galaxies in general., We cannot rule out the presence of at least some interstellar medium in NGC 4697 and X-ray faint early-type galaxies in general.615 Using the X-ray temperatureoptical velocity. dispersion relation of Davis White (1996). anv ISAT present iun NCC L697 would be expected to have a temperature around 0.3 keV. This would be very difficult to distinguisli from the soft component from LAINBs on a spectroscopic basis alouc.," Using the X-ray temperature--optical velocity dispersion relation of Davis White (1996), any ISM present in NGC 4697 would be expected to have a temperature around 0.3 keV. This would be very difficult to distinguish from the soft component from LMXBs on a spectroscopic basis alone."616 What is needed to separate the ISAL component frou the LAINB cutissiou is the lieh spatial resolution that cau be afforded byChandra., What is needed to separate the ISM component from the LMXB emission is the high spatial resolution that can be afforded by.617 Delow. we preseut a simulation of what we expect the emission from NGC 1697 to look like in the event that the emission is composed solely of LAINBs.," Below, we present a simulation of what we expect the emission from NGC 4697 to look like in the event that the emission is composed solely of LMXBs."618 We have an approved Cwele 1 10.000. 8Chendre Qsorvation of NGC 1697: here we show that this observation should resolve the hard and soft N-raw Cluission iuto individual sources. assunüug that the wission ds frou LAINBs.," We have an approved Cycle 1 40,000 s observation of NGC 4697; here we show that this observation should resolve the hard and soft X-ray emission into individual sources, assuming that the emission is from LMXBs."619 Couversely. the observation should cleanly seurate a truly diffuse cussion from tla of LAINBs.," Conversely, the observation should cleanly separate a truly diffuse emission from that of LMXBs."620 Since the main goal is to resolve the issue of the very soft componcut. the soft N-ray sensitive backside-Uhuninated (BI) S3 chip of the ACIS-S array will be usec for the observation.," Since the main goal is to resolve the issue of the very soft component, the soft X-ray sensitive backside-illuminated (BI) S3 chip of the ACIS-S array will be used for the observation."621 We have used the MARX (Mode ofANVAF Respouse to N-ravs: Wise. IIeuenmoerder. Davis 1997) Simulator to generate a svuthetic image of NGC 1697.," We have used the MARX (Model of Response to X-rays; Wise, Huenemoerder, Davis 1997) Simulator to generate a synthetic image of NGC 4697."622 The MARX Simulator takes as input the desired. spectral model aud spatial distribution model of an N-ray source and creates an image of the source as it would appear ouce having passed through the optics ofChandra., The MARX Simulator takes as input the desired spectral model and spatial distribution model of an X-ray source and creates an image of the source as it would appear once having passed through the optics of.623 The spectral and spatial distribution models described below were fed into ATARN using the ACTIS-S BI response to produce an image of NCC 1697 for a 10.000 x observation.," The spectral and spatial distribution models described below were fed into MARX using the ACIS-S BI response to produce an image of NGC 4697 for a 40,000 s observation."624 For the spectra of the LAINBs. we assume a model that best fit the jointROSAT PSPC | spectiiii of NGC 1697 discussed in 23.," For the spectra of the LMXBs, we assume a model that best fit the joint PSPC + spectrum of NGC 4697 discussed in \ref{sec:spectral}."625 For the spatial distribution of the N-rav enussion. we assume that the LAINBs have the same spatial distribution as the stellaz light.," For the spatial distribution of the X-ray emission, we assume that the LMXBs have the same spatial distribution as the stellar light."626" We taxe the optical distribution to be a de Vaucouleurs profile (de Vaucouleurs et 11991) with a mean haltlieht radius of 72"". an effective senminnajor axis of 95"". an effective scluimuuor axis of 55"". resulting in an ellipticity of 0.12. and elougated at a position augle of 677. ("," We take the optical distribution to be a de Vaucouleurs profile (de Vaucouleurs et 1991) with a mean half-light radius of $72^{\prime\prime}$, an effective semimajor axis of $95^{\prime\prime}$ , an effective semiminor axis of $55^{\prime\prime}$ , resulting in an ellipticity of 0.42, and elongated at a position angle of $^\circ$ . ("627Jedrzejewsld et 11987: Faber et 11989: Peleticer et 11990).,Jedrzejewski et 1987; Faber et 1989; Peletier et 1990).628where p and e are respectively (he pressure and the macroscopic energy density measured in proper coordinates.,where $p$ and $\epsilon $ are respectively the pressure and the macroscopic energy density measured in proper coordinates.629 Einsteins field equations without the cosmological constant reduce to denote differentiation with respect (o r., Einstein's field equations without the cosmological constant reduce to where primes denote differentiation with respect to $r$.630" These three equations together with the equation of stateof the material e = p(e) determine the mechanical equilibrium of the matter distribution as well as the dependence of the metric g,,'s on r."," These three equations together with the equation of stateof the material $\epsilon $ = $631p(\epsilon )$ determine the mechanical equilibrium of the matter distribution as well as the dependence of the metric $g_{\mu \nu }$ 's on $r$."632 The boundary of the matter distribution is the value of 7=ry for which p= 0. aud such (hat for r«rg.p>0.," The boundary of the matter distribution is the value of $r=r_{b}$ for which $%633p= 0, and such that for $r<r_{b},p>0$."634 For r«ry the solution depends on the equation of state of the malerial connecting p aud e., For $r<r_{b}$ the solution depends on the equation of state of the material connecting $p$ and $\epsilon $.635 For many equations of state a sharp boundary. exists with a finite value of rj., For many equations of state a sharp boundary exists with a finite value of $r_{b}$.636 In ihe empty space. p=€ 0. surrounding the spherically svinnmetric distribution of matter. the Sehwarzschild’s exterior solution is obtained: where m. is the Newlonian mass of the matter as calculated by a distant observer.," In the empty space, $p = \epsilon = 0$ , surrounding the spherically symmetric distribution of matter, the Schwarzschild's exterior solution is obtained: where $m$ is the Newtonian mass of the matter as calculated by a distant observer."637 Inside the boundary. Eqs. (5)). (6)).," Inside the boundary, Eqs. \ref{fieldeq1}) ), \ref{fieldeq2}) ),"638 and (7)) may be rewritten as follows., and \ref{fieldeq3}) ) may be rewritten as follows.639 Using the equation of state e= e(p). Eq. (7))," Using the equation of state $\epsilon = \epsilon(p)$ , Eq. \ref{fieldeq3}) )"640" max be immediately integrated The constant e""!) is determinedby making e continuous across the boundary.", may be immediately integrated The constant $e^{\nu(r_b)}$ is determinedby making $e^\nu$ continuous across the boundary.641the New Luvytei Two-Tenths Catalogue (Luyteo1. 1980 - NLTT).,"the New Luyten Two-Tenths Catalogue (Luyteon, 1980 - NLTT)."642" Once correlated. we can use the Gu,— Ks) colo was a crude plotometric parallax estimator. with the long baseline compeusatiug lo sone extent or the uncertainties in the red 1magnitudes."," Once correlated, we can use the $m_r-K_S$ ) colour as a crude photometric parallax estimator, with the long baseline compensating to some extent for the uncertainties in the red magnitudes."643 Paper I describes the definitiou of au initial sample o. uearby-star candidates. drawu rom NLTT sources which have potential 2ALASS counterparts wihin a 10-arcsecoud search racius.," Paper I describes the definition of an initial sample of nearby-star candidates, drawn from NLTT sources which have potential 2MASS counterparts within a 10-arcsecond search radius."644 By applying a series of cuts in colour-maguituce ancl colour-colow planes. we reduced the list of 23795 optical/IR matches to 1215 sources with yhotometric pre)perties consistent. with their beiig late-type clwarls within 20 parsecs of the Suu.," By applying a series of cuts in colour-magnitude and colour-colour planes, we reduced the list of 23795 optical/IR matches to 1245 sources with photometric properties consistent with their being late-type dwarfs within 20 parsecs of the Sun."645 These stars coistitute NLTT Sample 1., These stars constitute NLTT Sample 1.646 Paper I compiles photometric data from the literature or 169 st:ars. and uses colour-maguitude relationsat (Ny. (V-Ix). (Na. (V-I)) and (Aly. (1-0) to estimate ¢istauces to those stars.," Paper I compiles photometric data from the literature for 469 stars, and uses colour-magnitude relationsat $_V$ , (V-K), $_V$ , (V-I)) and $_I$, (I-J)) to estimate distances to those stars."647 Three hundred of those stars. aud a further 39 ultracool (spectral ype >M6) cdwarls. have formal distauces of less than 20 parsecs. including 76 stars not previously included i1 nearby star catalogues.," Three hundred of those stars, and a further 39 ultracool (spectral type $>$ M6) dwarfs, have formal distances of less than 20 parsecs, including 76 stars not previously included in nearby star catalogues."648 The current paper coutinues analysis of the NLTT Sample 1. presenting optical photometry ol a sampe of 180 relatively-bright southeru stars.," The current paper continues analysis of the NLTT Sample 1, presenting optical photometry of a sample of 180 relatively-bright southern stars."649 The following section outlines the sample and p'esents the observations., The following section outlines the sample and presents the observations.650 Section 3 describes our procedures for estimating distauces to these stars. ud Sectiou f discusses some of the more interesting stars in the sample.," Section 3 describes our procedures for estimating distances to these stars, and Section 4 discusses some of the more interesting stars in the sample."651" Our results are πα, in the final section.", Our results are summarised in the final section.652 As descried iu Paper L tie 1215 stars in NLTT Sample 1 were selected ou he basis of their having locatious iu tlie Gn. Crp) and (J-H)/(H-Ix)) planes consistent with ukl- or late-type M dwar witlin 20 parsecs of he Suu.," As described in Paper I, the 1245 stars in NLTT Sample 1 were selected on the basis of their having locations in the $m_r$, $m_r$ $_S$ )) and ((J-H)/(H-K)) planes consistent with mid- or late-type M dwarfs within 20 parsecs of the Sun."653 Regious within —1¢ of he Galactic Plane were excludedpriori. since tle NLTT catalogte has :i significantly brighter linitiug maguitu« ea ilose latitudes.," Regions within $\pm10^o$ of the Galactic Plane were excluded, since the NLTT catalogue has a significantly brighter limiting magnitude at those latitudes."654 The selected stars have maguituces iu he range 8«n<2Q. wil hover Lyi 1g|jetween Lith aud 16th imagnituce.," The selected stars have magnitudes in the range $8 < m_r < 20$, with over lying between 11th and 16th magnitude."655 They span the full raee of Right Ascension. alhough the majOriv lie at northieru Decliuation. reflectiug both the areal coverage of the second iiicremental release of 2NLASS «ala auc incompleteness ii the NLTT sou rol dé=—30°.," They span the full range of Right Ascension, although the majority lie at northern Declination, reflecting both the areal coverage of the second incremental release of 2MASS data and incompleteness in the NLTT south of $\delta = -30^o$."656 Several huuclrecl stars. howeve'. lie sout rol the equator.," Several hundred stars, however, lie south of the equator."657 Southern hemisplere proper-1uojon stars have generally. received less attention than tjelr northern counterparts. aud. as a result. even relatively bright objects in the current. sample |ave uo previous detailed 1jeasureiments.," Southern hemisphere proper-motion stars have generally received less attention than their northern counterparts, and, as a result, even relatively bright objects in the current sample have no previous detailed measurements."658 Figure 1 shows the distribution ou tlie celestial sphere aud in the (ny. Qny-Ixs)) coour-magnitude plane of the 180 NLTT cwarls targeted here.," Figure 1 shows the distribution on the celestial sphere and in the $m_r$ , $m_r$$_S$)) colour-magnitude plane of the 180 NLTT dwarfs targeted here."659For an observer at X=(X.Y.Z). the condition IX- should be satisfied in order to detect a holonomy y in a shell region ofrj«rrs.,"For an observer at $\vec X=(X, Y, Z)$, the condition $|\vec X - \gamma \vec X| /2< r_2$ should be satisfied in order to detect a holonomy $\gamma$ in a shell region of $r_1<r<r_2$."660 The lefthand side is the distance from the observer to the face of his Dirichlet domain. which corresponds to y.," The lefthand side is the distance from the observer to the face of his Dirichlet domain, which corresponds to $\gamma$ ."661 ↑↴⋯⊲⋗∕⋮⊺⋮≺↕∠∕∖⋒⊺∖≺∠∕∖⋒∪⋮≺∕⊓⋅⋯⋯ for y= ⊺⋮≺↕∠∕∖∕⋮⋟⊺∖≼−∠∕∖∕≣⋟∪⋮≺∕⊓⋅∖∖⇁∣↴⊜∣⊲⊜∠⋮∣≺↴⊂↿∣⊃∁ and rs»Ξ8.2 Gpe.," for $\gamma= T_z(\pm L/\sqrt{2})T_x(L/\sqrt{2})O_z(\pi)$ , and for $\gamma=T_z(\pm L/\sqrt{2})T_x(-L/\sqrt{2})O_z(\pi) $, where $L=16$ Gpc and $r_2=8.2$ Gpc."662 An observer satisfying each condition can detect each par of holonomies., An observer satisfying each condition can detect each pair of holonomies.663 It can be seen that the Z- of the observer affects nothing., It can be seen that the $Z$ -location of the observer affects nothing.664 However. the bigger the Y-location becomes. the farther away all faces are located. hence the harder to detect all holonomies.," However, the bigger the $Y$ -location becomes, the farther away all faces are located, hence the harder to detect all holonomies."665 The X-location has an effect of pushing two faces away. while drawing the other two. which is due to the translational direction that does not accord with the rotational axis.," The $X$ -location has an effect of pushing two faces away, while drawing the other two, which is due to the translational direction that does not accord with the rotational axis."666 The last effect cannot be seen in spaces without half-turn corkscrew motion (type IL) or glide reflection (type ID., The last effect cannot be seen in spaces without half-turn corkscrew motion (type II) or glide reflection (type II).667 According to these conditions. we chose six locations of (X.Y.Z)= (0. 0. 0). (1. 0. 0). (3. 0. 0). (5. 0. 0). (0. 0.5. 0). and (0. I. 0). in units of Gpe.," According to these conditions, we chose six locations of $(X,Y,Z)=$ (0, 0, 0), (1, 0, 0), (3, 0, 0), (5, 0, 0), (0, 0.5, 0), and (0, 1, 0), in units of Gpc."668 We prepared catalogs of toy quasars seen from these observers. and applied our method to them.," We prepared catalogs of toy quasars seen from these observers, and applied our method to them."669 Thecoordinate axes were chosen correctly here., Thecoordinate axes were chosen correctly here.670 Results for these observers are given in Table 4. and Figure 6.., Results for these observers are given in Table \ref{table4} and Figure \ref{figure6}.671 In this table. it can be seen that the signal gets weaker às the observer moves along the y-axis. since the faces get farther and the ghosts decrease in number.," In this table, it can be seen that the signal gets weaker as the observer moves along the $y$ -axis, since the faces get farther and the ghosts decrease in number."672 This is also seen in the histograms. where the hills constituted by ghosts disappear.," This is also seen in the histograms, where the hills constituted by ghosts disappear."673 It is necessary to use a larger shell region. if possible. to detect the ghosts.," It is necessary to use a larger shell region, if possible, to detect the ghosts."674 As the observer moves along the x-axis. on the other hand. the signal and the hills remain like the initial one in which the observer is located at the center. since one pair of faces gets farther. but the others get closer.," As the observer moves along the $x$ -axis, on the other hand, the signal and the hills remain like the initial one in which the observer is located at the center, since one pair of faces gets farther, but the others get closer."675 As long as the coordinate axes are chosen correctly. our method can detect holonomies that are close enough to the observer.," As long as the coordinate axes are chosen correctly, our method can detect holonomies that are close enough to the observer."676 Our method ts suited to obtaining a lower limit to the size of the Universe., Our method is suited to obtaining a lower limit to the size of the .677 In this paper we have developed a cosmic crystallography method that significantly extends previous methods., In this paper we have developed a cosmic crystallography method that significantly extends previous methods.678 A thin. shell-Hke part (rj«r 7: of the full catalog is used. similar to the circles-in-the-sky method.," A thin, shell-like part $r_1<r<r_2$ ) of the full catalog is used, similar to the circles-in-the-sky method."679 This region should be as large às possible. but the comoving density of objects should be simultaneously high enough there.," This region should be as large as possible, but the comoving density of objects should be simultaneously high enough there."680 The quadruplets of these objects are filtered three times: separation. vectorial condition. and lifetime of objects.," The quadruplets of these objects are filtered three times; separation, vectorial condition, and lifetime of objects."681 These filters drop the false stochastic ones. while keeping the real topological ones.," These filters drop the false stochastic ones, while keeping the real topological ones."682 Flat spaces described by Euclidean geometry are assumed. and the second filter. vectorial condition. i$ specialized to detect their holonomies.," Flat spaces described by Euclidean geometry are assumed, and the second filter, vectorial condition, is specialized to detect their holonomies."683 This assumption ts not fundamental. since similar filters for the holonomies in spherical or hyperbolic spaces can also be constructed.," This assumption is not fundamental, since similar filters for the holonomies in spherical or hyperbolic spaces can also be constructed."684 The number of quadruplets having passed these filters is translated into an index of multiconnectednessof theuniverse. and then the objects are classifiedby 5;. the number of such," The number of quadruplets having passed these filters is translated into an index of multiconnectednessof theuniverse, and then the objects are classifiedby $s_i$ , the number of such"685ol dark matter clearing by black hole mergers of Merrittetal.(2002).,of dark matter clearing by black hole mergers of \citet{Meretal02}.686". To ealeulate the effect of gravitational focusing byΑι, a black hole mass of 3x105 is asstuned. and the velocily distribution is assigned to be isotropic wilh a dispersion of σι=155 km/s. The black hole mass. the assumption of an isotropic velocity distribution. andthe value of σι used here are all approximately. consistent will observed stellar kinematics in the (Ghezetal.1993:Genzel2000)."," To calculate the effect of gravitational focusing by, a black hole mass of $3\times 10^6$ is assumed, and the velocity distribution is assigned to be isotropic with a dispersion of $\sigma_v=155$ km/s. The black hole mass, the assumption of an isotropic velocity distribution, andthe value of $\sigma_v$ used here are all approximately consistent with observed stellar kinematics in the \citep{Gheetal98, Genetal00}."687". The observed dark matter sell-aunihilation flux depend on the local emissivity. where Pam/y; is the dark matter number density. and (0,,,0) gives the sell-annihilation rate per unit densitv."," The observed dark matter self-annihilation flux depend on the local emissivity, where $\rho_{\rm dm}/m_{\rm dm}$ is the dark matter number density, and $\left<\sigma_{ann} v\right>$ gives the self-annihilation rate per unit density."688" For the neutralino. (wpical values from the literature (e.g.DergstrómGondolo1996) are my=may100 GeV. and (0,,,0)=107"" cm/s. independent of pairwise closing speed v."," For the neutralino, typical values from the literature \citep[e.g.][]{BerGon96} are $m_{\chi}689\equiv m_{\rm dm} = 100$ GeV, and $\left<\sigma_{ann} v\right> =69010^{-26}$ $^3$ /s, independent of pairwise closing speed $v$."691 The quantity Y specifies the. vield of decay by-products: for example. the bolometric vield corresponds to Y—mic.," The quantity $Y$ specifies the yield of decay by-products; for example, the bolometric yield corresponds to $Y=m_\chi c^2$."692 The emissivity per unit frequency ol photons produced by electrons in a magnetic field depends on electron-positron production channels. as well as svnchrotron radiative efficiencies Tvler(e.g..2002).," The emissivity per unit frequency of photons produced by electrons in a magnetic field depends on electron-positron production channels, as well as synchrotron radiative efficiencies \citet[e.g.,][]{Tyl02}."693. In the case of certain neutralino decay products. namely neutrinos. the flux is a straightlorwarc line-o[-sight integral over (he emissivity. since sell-absorption and diffusion do not occur 1999).," In the case of certain neutralino decay products, namely neutrinos, the flux is a straightforward line-of-sight integral over the emissivity, since self-absorption and diffusion do not occur \citep{GonSil99}."694. The line of sight integral along some sky direction » is conveniently expressed in dimensionless form as (Bergstromοἱal.1998:Merritt2002) Figure 2 gives J. averaged inside a circular aperture centered onAx.. as a function of aperture radius.," The line of sight integral along some sky direction $\hat{n}$ is conveniently expressed in dimensionless form as \citep{BerUllBuc98, Meretal02}695 Figure 2 gives $J$, averaged inside a circular aperture centered on, as a function of aperture radius."696 The density profile is (he same as in Figure 1., The density profile is the same as in Figure 1.697 The enhancement from eravitational focusing is significant insile small apertures., The enhancement from gravitational focusing is significant inside small apertures.698 Even so. Bertoneetal.(2004) point out that (he neutrino flux from (the Galactic Center will be undetectable if (he current eamnin-ray constraints are anv indication of the annihilation rate.," Even so, \citet{Beretal04} point out that the neutrino flux from the Galactic Center will be undetectable if the current gamma-ray constraints are any indication of the annihilation rate."699 Erkocaetal.(2010) are more hopeful [rom a theoretical perspective. while (he observations are providing limits to the neutrino flux (e.g..fromIeeCubeAbbasietal. 2011).. but no Galactic Center signal at (his point.," \citet{Erketal10} are more hopeful from a theoretical perspective, while the observations are providing limits to the neutrino flux \citep[e.g., from IceCube][]{Abbetal11}, , but no Galactic Center signal at this point."700 The strength. of a gravitationallv focused. censitw profile around a compact object is, The strength of a gravitationally focused density profile around a compact object is701Previous ROSAT observations have shown that red. giants are not substantial X-ray emittors.,Previous ROSAT observations have shown that red giants are not substantial X-ray emittors.702 Only one late-type giant was deteced in the ROSAT all-skv survey (2:?:2). and pointed observations placed an extremely tight upper [init of 310aeeres ton the X-ray lux of the LELE red. giant Arcturus (2)..," Only one late-type giant was detected in the ROSAT all-sky survey \cite{Haisch91,Haisch92,Huensch96}, and pointed observations placed an extremely tight upper limit of $3\times10^{25}\rm\,erg\,s^{-1}$ on the X-ray flux of the III red giant Arcturus \cite{Ayers91}."703 Thus we can be confident that the X-ray emission from 4 Draconis reported in this paper originates on the ultraviolet companion. 4 DD. Our. ROSAT observations are consistent with this secondary containing an accreting white chwart.," Thus we can be confident that the X-ray emission from 4 Draconis reported in this paper originates on the ultraviolet companion, 4 B. Our ROSAT observations are consistent with this secondary containing an accreting white dwarf."704" The ~HkkeV temperature of the optically-thin X-ray spectrum is characteristic of non-magnetic cataclysmic variables (c.g.7) and of the ""bombardment solution"" for radial accretion onto a white chwarl (c.g.2)..", The $\sim$ keV temperature of the optically-thin X-ray spectrum is characteristic of non-magnetic cataclysmic variables \egcite{Wheatley96} and of the “bombardment solution” for radial accretion onto a white dwarf \egcite{Woelk95}.705 Phe bombardment solution applies when the mass accretion rate per unit area is too low for a stand-olI shock to form (im«]0+estem 7).," The bombardment solution applies when the mass accretion rate per unit area is too low for a stand-off shock to form $\rm\dot{m}<10^{-1}\,g\,s^{-1}\,cm^{-2}$ )."706 Our measured Luminosity of 6107ergs l-iniplies. an accretionoa rate of (24LS10bgs+ for white dwarf masses in the range 0.30LOALY.," Our measured luminosity of $6\times10^{31}\rm\,erg\,s^{-1}$ implies an accretion rate of $0.24-1.8\times10^{15}\rm\,g\,s^{-1}$ for white dwarf masses in the range $0.3-1.0\rm\,M_{\sun}$."707 For the bombardment solution to aplv his accretion rate must be spread over an area of at. least ηLs101em?. although this is a small fraction of the surface area of even a massive white ναι," For the bombardment solution to apply this accretion rate must be spread over an area of at least $0.24-1.8\times10^{16}\rm\,cm^2$, although this is a small fraction of the surface area of even a massive white dwarf."708 Although our observations are consistent with the yrescnee of an accreting white να they do not support the presence of an AAL Ller system.," Although our observations are consistent with the presence of an accreting white dwarf, they do not support the presence of an AM Her system."709 First. the τοΛΙ spectra of AAT ers are typically dominated: by intense optically-thick soft) emission. with characteristic temperatures of ~20ceV. We can rule out the presence of such a component in the 1993 spectrum of 4 BB reffig-spec)).," First, the ROSAT spectra of AM Hers are typically dominated by intense optically-thick soft emission, with characteristic temperatures of $\sim$ eV. We can rule out the presence of such a component in the 1993 spectrum of 4 B \\ref{fig-spec}) )."710 Second. it is clear from the LIBI lighteurve reffie-le}) that the X-ray emission is not strongly modulated abi a period of Shh. as it is for everv known hieh-state AM Ler system and most other magnetic cataclysmic variables.," Second, it is clear from the HRI lightcurve \\ref{fig-lc}) ) that the X-ray emission is not strongly modulated at a period of h, as it is for every known high-state AM Her system and most other magnetic cataclysmic variables."711 AM Her systems have shown spectra much like that of 4 Draconis during low accretion rate states (e.g.2).. but our measured Luminosity is rather high for a low-state ANI Ler. ancl we believe the lack of an A-ray orbital periodicity alone is sullicient evidence to rule out the presence of an AM Ler in the 4 Draconis system.," AM Her systems have shown spectra much like that of 4 Draconis during low accretion rate states \egcite{Ramsay95}, but our measured luminosity is rather high for a low-state AM Her, and we believe the lack of an X-ray orbital periodicity alone is sufficient evidence to rule out the presence of an AM Her in the 4 Draconis system."712 Non-magnetic cataclysmic variables chwarl novae) usually have no optically-thick component in the ROSAT bandpass (e.g.2).. have characteristic N-ray. temperatures lower than AM LHers. and do not exhibit strong orbital. X-rav modulation.," Non-magnetic cataclysmic variables dwarf novae) usually have no optically-thick component in the ROSAT bandpass \egcite{Wheatley96}, have characteristic X-ray temperatures lower than AM Hers, and do not exhibit strong orbital X-ray modulation."713 Pherefore we cannot rule out the presence of a non-magnetic cataclysmic variable., Therefore we cannot rule out the presence of a non-magnetic cataclysmic variable.714 However. the original case for the presence of a cataclysmic variable was based upon the claimed detection of a 4hh ultraviolet period (?)..," However, the original case for the presence of a cataclysmic variable was based upon the claimed detection of a h ultraviolet period \cite{Reimers88}."715 Reviewing the lighteurve in Fig.33 of Reimers ct we believe the case for a periodic moculation is not strong., Reviewing the lightcurve in 3 of Reimers et we believe the case for a periodic modulation is not strong.716 Also. more recent LIST. observations do not support the presence of a 4hh period Gaensicke. private communication).," Also, more recent HST observations do not support the presence of a h period Gaensicke, private communication)."717 ‘Thus. in the [ace [ni evidence clearly supporting the presence of an accreting white chwarl. but none requiring the," Thus, in the face of evidence clearly supporting the presence of an accreting white dwarf, but none requiring the"718deteriuuatiou of D4». since it relies ou cosmological siuulatious for inodeliug the temperature and density field.,"determination of $\Gamma_{-12}$, since it relies on cosmological simulations for modeling the temperature and density field."719"""Disks. Extrasolar Planets and Mown Dwarfs held at. the LAP in July. 2000 for useful discussions.","'Disks, Extrasolar Planets and Brown Dwarfs' held at the IAP in July 2000 for useful discussions."720"(109A7... (IKI&oriunendy kl~ (TTrinchiert. ZL,~LO 1 E03 (PPellegriui (BBlaudford (11952) DP-—0.61.5) (A (11996) 1) (SSlee (11999; ",$10^9 M_\odot$ \markcite{KoRi95}K $kT \sim$ \markcite{TrFa86}T $L_x \sim 10^{45}$ $^{-1}$ $< 10^{-3}$ \markcite{Pell99}P \markcite{NaYi95} \markcite{Abrm95} \markcite{DiMa00} \markcite{BlBg99}B \markcite{Bndi52}1 $\Gamma \sim 0.6 - 1.5$ \markcite{AlDF00}A \markcite{Reyn96}1 \markcite{Harr98}1 \markcite{Slee94}S \markcite{DiMa99}1 721"The basic requirement of the observed spectro-astrometry is that the line forming gas must be orbiting the central star with strongly sub-Keplerian azimuthal velocities in order to produce the single peak without requiring that the emission is extended at the spatial resolution of CRIRES (—0""115).",The basic requirement of the observed spectro-astrometry is that the line forming gas must be orbiting the central star with strongly sub-Keplerian azimuthal velocities in order to produce the single peak without requiring that the emission is extended at the spatial resolution of CRIRES $\sim$ 15).722" A wide angle wind provides a convenient physical way of accomplishing this through simple conservation of angular momentum — as a gas parcel is forced outwards due to the wind pressure, the azimuthal velocity decreases linearly with radius, in comparison with the underlying Keplerian disk in which the velocity experiences a shallower decrease as R-!/2."," A wide angle wind provides a convenient physical way of accomplishing this through simple conservation of angular momentum – as a gas parcel is forced outwards due to the wind pressure, the azimuthal velocity decreases linearly with radius, in comparison with the underlying Keplerian disk in which the velocity experiences a shallower decrease as $R^{-1/2}$."723 'This generates gas above the disk that is supported by wind pressure and orbits at low azimuthal velocities., This generates gas above the disk that is supported by wind pressure and orbits at low azimuthal velocities.724" Following?,, the wind is constructed as set of linear streamlines with a locus below the centrala star at a distance d in units of R,."," Following, the wind is constructed as a set of linear streamlines with a locus below the central star at a distance $d$ in units of $R_*$."725 This generates a conical wind with no flow along the disk axis., This generates a conical wind with no flow along the disk axis.726" Briefly, the wind is accelerated along the field lines as: where | is the coordinate alongthe stream line, c, is the sound speed, Όρες is the asymptotic velocity at the end of the stream line and Agcaie is the scale of the acceleration region of the wind."," Briefly, the wind is accelerated along the field lines as: where $l$ is the coordinate alongthe stream line, $c_s$ is the sound speed, $v_{\rm esc}$ is the asymptotic velocity at the end of the stream line and $A_{\rm scale}$ is the scale of the acceleration region of the wind."727 ϱ is the wind acceleration parameter., $\beta$ is the wind acceleration parameter.728" Requiring angular momentum conservation, the azimuthal velocity component is: where F is the radial disk coordinate."," Requiring angular momentum conservation, the azimuthal velocity component is: where $R$ is the radial disk coordinate."729" 'The density of the wind is calculated assuming mass conservation: Here, X(w)οςR? is the local mass-loss rate, ὃ is the angle between the stream line and the disk normal and S is the distance to the wind locus."," The density of the wind is calculated assuming mass conservation: Here, $\dot{\Sigma}(w) \propto R^{-p}$ is the local mass-loss rate, $\delta$ is the angle between the stream line and the disk normal and S is the distance to the wind locus."730 The exponent of the local mass loss rate is taken to be p=7/2(?)., The exponent of the local mass loss rate is taken to be $p=7/2$.731". The total wind mass loss rate can be calculated by integrating over the disk and multiplying by two to include the opposite surface: The raytracer RADLite is used to render model lines and spectro-astrometry (7)for the wind models, based on a generic model of a flared protoplanetary disk, and assuming level populations in LTE."," The total wind mass loss rate can be calculated by integrating over the disk and multiplying by two to include the opposite surface: The raytracer RADLite is used to render model lines and spectro-astrometry for the wind models, based on a generic model of a flared protoplanetary disk, and assuming level populations in LTE."732" Specifically, the temperature structure is assumed to be in equilibrium with the stellar radiation field and dominated by dust heating/cooling."," Specifically, the temperature structure is assumed to be in equilibrium with the stellar radiation field and dominated by dust heating/cooling."733" In reality, the heating of the wind is likely to be dominated by photo-electric heating similar to the heating of the disk atmosphere οι, perhaps, ambipolar diffusion(?)."," In reality, the heating of the wind is likely to be dominated by photo-electric heating similar to the heating of the disk atmosphere or, perhaps, ambipolar diffusion."734". The cooling(1173) may be dominated by adiabatic expansion and molecular cooling (e.g., partly via the observed CO and H20O lines)."," The cooling may be dominated by adiabatic expansion and molecular cooling (e.g., partly via the observed CO and $_2$ O lines)."735" However, we restrict ourselves to qualitative models in this paper (see also refCaveats)), since a detailed and appropriate physical treatment of the thermal wind structure required to match the observations will be likely be a significant study in its own right."," However, we restrict ourselves to qualitative models in this paper (see also \\ref{Caveats}) ), since a detailed and appropriate physical treatment of the thermal wind structure required to match the observations will be likely be a significant study in its own right."736" Figure 10 illustrates the wind geometry and compares the observables generated using the wind modelfor the spectro-astrometry of AS 205N. The total mass-loss rate is 9x107?Mo yr-!, assuming a CO abundance of 5x107? relative to H."," Figure \ref{Wind_sketch} illustrates the wind geometry and compares the observables generated using the wind modelfor the spectro-astrometry of AS 205N. The total mass-loss rate is $9\times 10^{-9}\,\rm M_{\odot}\,yr^{-1}$ , assuming a CO abundance of $5\times 10^{-5}$ relative to $\rm H$ ."737 This mass-loss rate is consistent, This mass-loss rate is consistent738fitting a gaussian function to their emission profile derived from the SH data.,fitting a gaussian function to their emission profile derived from the SH data.739" This is illustrated in reffig:hires,, which shows the SH IRS observations obtained for the [SIV], [NeII], [NeIII] and [SHI] ionic lines."," This is illustrated in \\ref{fig:hires}, which shows the SH IRS observations obtained for the [SIV], [NeII], [NeIII] and [SIII] ionic lines."740" Since no sky subtraction could be performed for these data though, the underlying continuum and the equivalent widths (EW) of these features were estimated from the continuum of the low-resolution spectrum as modeled with PAHFIT."," Since no sky subtraction could be performed for these data though, the underlying continuum and the equivalent widths (EW) of these features were estimated from the continuum of the low-resolution spectrum as modeled with PAHFIT."741 Our measurements are given in Table 2 along with the fluxes and the equivalent widths of the main PAHs and ionic lines measured in the low-resolution data., Our measurements are given in Table \ref{table:features} along with the fluxes and the equivalent widths of the main PAHs and ionic lines measured in the low-resolution data.742" For the strongest and isolated PAHs (i.e.,um,,uum,, and µπι)), we indicate the results obtained with the global PAHFIT decomposition but we also provide the measures that we derived with a local fit of the continuum underlying each individual feature using a spline function."," For the strongest and isolated PAHs (i.e., and ), we indicate the results obtained with the global PAHFIT decomposition but we also provide the measures that we derived with a local fit of the continuum underlying each individual feature using a spline function."743 The latter approach has been commonly used in the literature to characterize the mid-IR spectra of star-forming galaxies., The latter approach has been commonly used in the literature to characterize the mid-IR spectra of star-forming galaxies.744" It usually leads to lower values than obtained with PAHFIT, since PAHFIT accounts for the full extent of the PAH wings."," It usually leads to lower values than obtained with PAHFIT, since PAHFIT accounts for the full extent of the PAH wings."745" We constrained the spectral energy distributions of the GRB host galaxy and the WR region over the full infrared wavelength range by fitting the broad-band photometry with the empirical libraries of galaxy templates published by Chary&Elbaz(2001), Dale&Helou and Lagacheetal. as well as with the (2002)physical SEDs derived from (2004),radiative transfer modeling by Siebenmorgen&Kriigel(2007)."," We constrained the spectral energy distributions of the GRB host galaxy and the WR region over the full infrared wavelength range by fitting the broad-band photometry with the empirical libraries of galaxy templates published by \citet{Chary01}, \citet{Dale02} and \citet{Lagache04}, as well as with the physical SEDs derived from radiative transfer modeling by \citet{Siebenmorgen07}."746". Between these different libraries the SEDs mostly vary in the relative strength of the PAH features with respect to the hot dust continuum, as well as in the temperature and the emissivity of the cold dust component shaping the peak of the SED in the far-IR."," Between these different libraries the SEDs mostly vary in the relative strength of the PAH features with respect to the hot dust continuum, as well as in the temperature and the emissivity of the cold dust component shaping the peak of the SED in the far-IR."747" In the library of Siebenmorgen&Krügel the SEDs also depend on the size of the star-forming(2007) region responsible for the IR. emission, and they are given for radii of 0.35, 1, 3, 9 and kkpc."," In the library of \citet{Siebenmorgen07} the SEDs also depend on the size of the star-forming region responsible for the IR emission, and they are given for radii of 0.35, 1, 3, 9 and kpc."748" We only considered sizes of kkpc and kkpc for the WR region and the whole galaxy, respectively."," We only considered sizes of kpc and kpc for the WR region and the whole galaxy, respectively."749 'To obtain the best possible constraints we combined the MIPS fluxes presented in refsec:photo with the photometry already published by LeFloc’hetal., To obtain the best possible constraints we combined the MIPS fluxes presented in \\ref{sec:photo} with the photometry already published by \citet{LeFloch06}.750" The fitting was performed separately for each library,(2006).. using the code (Arnoutsetal.1999;Ibert2006)."," The fitting was performed separately for each library, using the code \citep{Arnouts99,Ilbert06}."751". Although most of the IR SED templates from the aforementioned libraries vary as a function of the total IR luminosity (but see Dale&Helou2002 for a dependence on dust temperature), their normalization was kept as a free parameter and the best templates were derived from a basic x? minimization of the fit."," Although most of the IR SED templates from the aforementioned libraries vary as a function of the total IR luminosity (but see \citealt{Dale02} for a dependence on dust temperature), their normalization was kept as a free parameter and the best templates were derived from a basic $\chi^2$ minimization of the fit."752" In the case of the WR region, we did not include the photometry at since at this wavelength we were unable to separate its contribution from the emission of the host (see refsec:mips160))."," In the case of the WR region, we did not include the photometry at since at this wavelength we were unable to separate its contribution from the emission of the host (see \\ref{sec:mips160}) )."753 We checked however that the best fits obtained for the WR region did not exceed the total flux measured for the GRB host galaxy atum., We checked however that the best fits obtained for the WR region did not exceed the total flux measured for the GRB host galaxy at.754". 'The results are illustrated in , whichshowsthemeasurements f romourbroad— bandphotometryandI RS spectroscopytogetherwiththeglobalrangeo; fittemplatesthatwereobtained f oreachoftheAlibraries."," The results are illustrated in \\ref{fig:ir_total_sed}, which shows the measurements from our broad-band photometry and IRS spectroscopy together with the global range of possible SED fits defined from the best-fit templates that were obtained for each of the 4 libraries."755 Asexpected f bandf luxesmeasuredwithI RACandMIP Swenotethatthemid— tofar—I Rspectralslopeo ftheW regionismuchsteeperthanobserve , As expected from the different broad-band fluxes measured with IRAC and MIPS we note that the mid- to far-IR spectral slope of the WR region is much steeper than observed for the whole GRB host.756"andpm,, we derivedR total IR luminosities of log(Lip/Lo) 88.6640.04 and log(Lirn/Leo) 99.0140.07 for the WR region and the whole GRB host galaxy, respectively."," By integrating the best fit SEDs between and, we derived total IR luminosities of $_{\rm IR}$ $_{\odot}$ $\pm$ 0.04 and $_{\rm IR}$ $_{\odot}$ $\pm$ 0.07 for the WR region and the whole GRB host galaxy, respectively."757 In these estimates the uncertainties were obtained by combining, In these estimates the uncertainties were obtained by combining758is ~3% of the total energy.,is only $\sim 3\%$ of the total magnetic energy.759 This value may increaseonly slightly on a much longer magnetictimescale., This value may increase slightly on a much longer timescale.760 Also in , Also reported in Fig.7612 as dashed lines are the corresponding reportedevolution Fig.of the magnetic energies when a smaller resistivity of ηο/Μο=0.06 is used., \ref{fig:fig2} as dashed lines are the corresponding evolution of the magnetic energies when a smaller resistivity of $\eta_0/M_{\odot}=0.06$ is used.762" Since the evolution of the instability in this case is qualitatively very similar the evolution of E, we have confidence that our prescription for the resistive tor),behaviour of the magnetic field near the stellar surface does not influence the dynamics of the instability."," Since the evolution of the instability in this case is qualitatively very similar the evolution of $E_{\mathrm{m,tor}}$ ), we have confidence that our prescription for the resistive behaviour of the magnetic field near the stellar surface does not influence the dynamics of the instability."763" At the same time, however, a smaller is also responsible for a smaller of the poloidal resistivitymagnetic field (see inset), which is decayconsiderably the end of the simulation."," At the same time, however, a smaller resistivity is also responsible for a smaller decay of the poloidal magnetic field (see inset), which is considerably dissipated by the end of the simulation."764" While this behaviour is dissipatedinevitable byin a resistive context and has been reported also by other authors (Braithwaite2007),, it an aspect of these evolutions which could be improved representswith a fully consistent resistive MHD approach (Palenzuelaetal.2009).."," While this behaviour is inevitable in a resistive context and has been reported also by other authors \citep{Braithwaite2007}, it represents an aspect of these evolutions which could be improved with a fully consistent resistive MHD approach \citep{Palenzuela:2008sf}."765" Another important confirmation of the perturbative analysis is offered in Fig. 3,,"," Another important confirmation of the perturbative analysis is offered in Fig. \ref{fig:fig3},"766" where we show the inverse of the growth- 7, defined through the exponential growth of the toroidal component, versus the initial magnetic-field strength (red empty circles)."," where we show the inverse of the growth-time $\tau$, defined through the exponential growth of the toroidal component, versus the initial magnetic-field strength (red empty circles)."767" Note that the scaling is essentially linear for Bo&€7x1019 G, deviating from this for higher values, because of the magnetic tension."," Note that the scaling is essentially linear for $B_0768\lesssim 7 \times 10^{16}\,$ G, deviating from this for higher values, because of the stronger magnetic tension."769" More specifically, the stronger Lorentz strongerforce will tend to oppose the fluid motions in the direction near the neutral line and which trigger the instability."," More specifically, the stronger Lorentz force will tend to oppose the fluid motions in the polar direction near the neutral line and which trigger the instability."770"polar The presence of a linear scaling is essential to extend our results to pulsar magnetic- strengths, thus estimating a growth-timetypical of ~10s for a neutron star with Bg=10!? G. Also marked in Fig."," The presence of a linear scaling is essential to extend our results to typical pulsar magnetic-field strengths, thus estimating a growth-time of $\sim 10\,$ s for a neutron star with $B_0=10^{12}\,$ G. Also marked in Fig."771" 3 (blue star) is the inverse growth-time for the fiducial star evolved with the smaller of jo/M;=0.06; again, the close similarity in the resistivitytimescales confirms our expectation that the is not influenced the choice of the The instability"," \ref{fig:fig3} (blue star) is the inverse growth-time for the fiducial star evolved with the smaller resistivity of $\eta_0/M_{\odot}=0.06$; again, the close similarity in the timescales confirms our expectation that the instability is not influenced by the choice of the resistivity."772final discussion is reservedby for the potential GW resistivity.signal emitted during the of the, The final discussion is reserved for the potential GW signal emitted during the development of the instability.773 In Fig., In Fig.774 we report the GW strain in developmentthe + and x instability.polarizations as computed from the Newtonian formula., \ref{fig:fig4} we report the GW strain in the $+$ and $\times$ polarizations as computed from the Newtonian quadrupole formula.775" It is quite apparent that the signal is not ofquadrupole a burst type but, rather, that the main effect of the is that of triggering large-amplitude oscillations ofinstability the star in its fundamental F-mode."," It is quite apparent that the signal is not of a burst type but, rather, that the main effect of the instability is that of triggering large-amplitude oscillations of the star in its fundamental $F$ -mode."776" These GWs start from the numerical noise already at ~3.5 ms, butemerging are associated to high-m oscillations and hence not efficient sources of GWs."," These GWs start emerging from the numerical noise already at $\sim7773.5\,$ ms, but are associated to $m$ oscillations and hence not efficient sources of GWs."778" However, as the magnetic field starts to approach the final m=2 configuration at ~ 7ms, the oscillations become more efficient in producing a GW signal (Note that a m=N pertubation in the magnetic field leads to am=2N perturbation in the density)."," However, as the magnetic field starts to approach the final $m=2$ configuration at $\sim 7\,$ ms, the oscillations become more efficient in producing a GW signal (Note that a $m=N$ pertubation in the magnetic field leads to a $m=2N$ perturbation in the density)."779" Because these oscillations will have a rather narrow spectral distribution peaked around the F'-mode frequency (which is not significantly affected by the presence of magnetic fields), they represent very good sources of a periodic signal, potentially detectable by future advanced detectors."," Because these oscillations will have a rather narrow spectral distribution peaked around the $F$ -mode frequency (which is not significantly affected by the presence of magnetic fields), they represent very good sources of a periodic signal, potentially detectable by future advanced detectors."780" Defining the root-sum-square amplitude of the cross polarization as h;,,=fredth2,1/2, and assuming that the oscillations will persist| undamped(| for ~0.1— 1s, we estimate 4,=(0.54—1.7)x107? for a source at 10 kpc."," Defining the root-sum-square amplitude of the cross polarization as $h_{\rm781 rss}=\left[\int_{-\infty}^{+\infty} dt \,782 h_{\times}^2(t)\right]^{1/2}$, and assuming that the oscillations will persist undamped for $\simeq 0.1-1\,$ s, we estimate $h_{\rm rss}783= (0.54-1.7)\times 10^{-22}$ for a source at $10\,$ kpc."784" The corresponding signal-to-noise ratio for a detector such as advanced-LIGO or advanced-Virgo is S/N~1.6—5, thus potentially observable."," The corresponding signal-to-noise ratio for a detector such as advanced-LIGO or advanced-Virgo is $S/N \simeq 1.6-5$, thus potentially observable."785 A more detailed analysis of the spectral properties of the GW signal will be in a future work., A more detailed analysis of the spectral properties of the GW signal will be presented in a future work.786 These waveforms represent the first presentedestimate of the conversion of the kinetic energy generated the instability into GWs., These waveforms represent the first estimate of the conversion of the kinetic energy generated through the instability into GWs.787" For weaker magnetic fields, throughperturbative analyses have suggested this coupling is much weaker (Levin&vanHoven2011),, but more work is needed to investigate nonlinearly this regime."," For weaker magnetic fields, perturbative analyses have suggested this coupling is much weaker \citep{Levin:2011}, but more work is needed to investigate nonlinearly this regime."788 We report on numerical evolutions of the instability of poloidal magnetic fields in relativistic stars and the subsequent of a mixed-field configuration in quasi- In generat, We report on numerical evolutions of the instability of poloidal magnetic fields in relativistic stars and the subsequent generation of a mixed-field configuration in quasi-equilibrium.789"ionagreement with the expectations from analytic studies (Markey&Tayler1973; 1973),, we show perturbativethat the instability appears after about an WrightAlfvénn"," In agreement with the expectations from analytic perturbative studies \citep{Markey1973,790 Wright1973}, , we show that the instability appears after about an Alfvénn"791budget for (he secondary. pairs is quite limited. unless a higher emission can be ellectively suppressed in the Fermi-LAT band.,"budget for the secondary pairs is quite limited, unless a higher emission can be effectively suppressed in the Fermi-LAT band."792 This could be achieved either by assuming a broad enerev distribution of target photons extending to X-ray. energies. so to provide a significant attenuation also in the GeV band. or by introducing a very high lower-energy eutoff in the proton distribution.," This could be achieved either by assuming a broad energy distribution of target photons extending to X-ray energies, so to provide a significant attenuation also in the GeV band, or by introducing a very high lower-energy cutoff in the proton distribution."793 In absence of (hese (wo conditions. the X-ray. svnchrotron flix of the secondary. pairs would be approximately an order of magnitude below the reported X-ray [luxes.," In absence of these two conditions, the X-ray synchrotron flux of the secondary pairs would be approximately an order of magnitude below the reported X-ray fluxes."794 Therefore. in (his specific case. the internal absorption scenario requires additional ad-hoc assumptions to provide a sell-consistent interpretation of the TeV and X-ray data.," Therefore, in this specific case, the internal absorption scenario requires additional ad-hoc assumptions to provide a self-consistent interpretation of the TeV and X-ray data."795 These additional assumptions mstead are not needed in (he case of a hard proton spectrum (p= —0.5)., These additional assumptions instead are not needed in the case of a hard proton spectrum $p=-0.5$ ).796 The latter can provide both the energv budget to explain the X-ray data and. GeV. fluxes below the Fermi-LAT limits. as shown in Fig.," The latter can provide both the energy budget to explain the X-ray data and GeV fluxes below the Fermi-LAT limits, as shown in Fig."797 3. (Fit 2. whose corresponding parameters are given in Table 1)).," \ref{fig:0229} (Fit 2, whose corresponding parameters are given in Table \ref{table:parameters}) )."798 In the case of high EBL flux (model F1.6). the de-absorbed VIE spectrum has a photon index close to Di21. ie. harder than the unabsorbed svnchrotron spectrum from a proton clistvibution with index p~2.," In the case of high EBL flux (model F1.6), the de-absorbed VHE spectrum has a photon index close to $\Gamma_{\rm int}\simeq1$, i.e. harder than the unabsorbed synchrotron spectrum from a proton distribution with index $p\sim2$."799 Internal absorption allows the hardening of the TeV spectrum to the required level. but in the case of a conventional proton distribution the diserepaney with the Fermi-LAT upper limits is very strong.," Internal absorption allows the hardening of the TeV spectrum to the required level, but in the case of a conventional proton distribution the discrepancy with the Fermi-LAT upper limits is very strong."800" To avoid the conllict with Fermi-LAT data we need to suppress the GeV emission. by introducing additional assumptions such as an effective absorption of GeV 5-ravs (e.g. by X-rays) or a very. high lower-energv cutolf (at. 10"" TeV in proton energv)."," To avoid the conflict with Fermi-LAT data we need to suppress the GeV emission, by introducing additional assumptions such as an effective absorption of GeV $\gamma$ -rays (e.g. by X-rays) or a very high lower-energy cutoff (at $10^{6}$ TeV in proton energy)."801 However. these assumptions can hardly be endorsed without an additional observational or theoretical justification.," However, these assumptions can hardly be endorsed without an additional observational or theoretical justification."802 , 803"There are two major mechanisms for gas-giant planet ormation: one is the core accretion mechanism in which a massive solic core forms first and the clisk gas acerctes onto he core (Saronov1969:Goldreich&Ware1973:Pollackο)αἱ, 1996). and the other is the gravitational instability (GI) mechanism in which the circumstellar disk. cirectly ragmoents ino gas-giant planets via GI (CameronLOTS).","There are two major mechanisms for gas-giant planet formation: one is the core accretion mechanism in which a massive solid core forms first and the disk gas accretes onto the core \citep{safronov69,goldreich_ward73,pollacketal96}, and the other is the gravitational instability (GI) mechanism in which the circumstellar disk directly fragments into gas-giant planets via GI \citep{cameron78}."804. tecent discovery. of extra-solar planets at a great distance rom the central star such as HIVSTOOb. c. d and e (Marolsοἱ2008.201) and (το CPhalmannefa£.2009). creates a new problem for the planet formation.," Recent discovery of extra-solar planets at a great distance from the central star such as HR8799b, c, d and e \citep{maroisetal08,maroisetal10} and GJ579b \citep{thalmannetal09} creates a new problem for the planet formation."805 Lt is dillicult to orm planets in the regions far from central stars according o the core accretion mechanism. because massive solid core ormation before the dissipation of gaseous disk seems to be caüllicult (Dodson-Robinsonefad.20090).," It is difficult to form planets in the regions far from central stars according to the core accretion mechanism, because massive solid core formation before the dissipation of gaseous disk seems to be difficult \citep{dodsonetal09}."806. ‘The gravitationally instability mechanism may be more plausible for the formation o ‘these planets., The gravitationally instability mechanism may be more plausible for the formation of these planets.807 Many studies of clisk fragment:uijon have been done using either an analytic approach (ltalikov2005) or numerical simulations (e.g..&Date 2010).," Many studies of disk fragmentation have been done using either an analytic approach \citep{rafikov05} or numerical simulations \citep[e.g.,][]{ stamatellos_whitworth08, caietal08, boleyetal06, mejiaetal05, pickettetal03,meru_bate10}."808". These elforts. however. μα... to build a Consensus tha the. planet ormation by Gl within ~50r AU is highly «illicult when the clisk-to-stellar mass ratio is MaafAda""Yo0.1 whose ratio is suggested by observations (see.e.g...Witamurae£αἱ.2)02)."," These efforts, however, seem to build a consensus that the planet formation by GI within $\sim 50$ AU is highly difficult when the disk-to-stellar mass ratio is $M_{\rm disk}/M_{\rm star}\lesssim 0.1$ whose ratio is suggested by observations \citep[see, e.g.,][]{kitamuraetal02}."809. On the otyer hand. IEnutsukaefαἱ.(2010). showed that the circumstelar disk is comparable to or more massive than the protostar. (AdagfA 1) during the (early) main accretior1 phase (i.e. Class 0 or Class | stages) and is hiehly σανίαionally unstable.," On the other hand, \citet{imm10} showed that the circumstellar disk is comparable to or more massive than the protostar $M_{\rm disk}/M_{\rm star}\gtrsim 1$ ) during the (early) main accretion phase (i.e., Class 0 or Class I stages) and is highly gravitationally unstable."810 Recently. such massive disks were observed around. very young protostars," Recently, such massive disks were observed around very young protostars"811proposed byDiStefano. Greiner. Garcia. Murray 2001.,"proposed by, Greiner, Garcia, Murray 2001."812 Given the expected rate of TDs. in a galaxy such as M3l. several of these remnants could be active at any given time.," Given the expected rate of TDs, in a galaxy such as M31, several of these remnants could be active at any given time."813 Some stripped cores are expected to be WDs or pre-WDs. and some are expected to be helium stars.," Some stripped cores are expected to be WDs or pre-WDs, and some are expected to be helium stars."814 The possibility that SSSs in the center of nearby galaxies could be signatures of TDs ts interesting. particularly because other signatures of TDs are so difficult to identify with confidence.," The possibility that SSSs in the center of nearby galaxies could be signatures of TDs is interesting, particularly because other signatures of TDs are so difficult to identify with confidence."815 The complementary signature most considered is an event due to the accretion of a portion of the disrupted star’s envelope by the BH (Hills 1975. Lidskit Ozernoi 1979. Gurzadyan Ozernoi 1980. Rees 1988).," The complementary signature most considered is an event due to the accretion of a portion of the disrupted star's envelope by the BH (Hills 1975, Lidskii Ozernoi 1979, Gurzadyan Ozernoi 1980, Rees 1988)."816 The associated accretion event can last for months or decades. with luminosities possibly as high as ~10—1075 erg/s. There is a growing body of data on UV and X-ray flares that may be consistent with these sorts of accretion events (see references in eet al.," The associated accretion event can last for months or decades, with luminosities possibly as high as $\sim 10^{44}-10^{46}$ erg/s. There is a growing body of data on UV and X-ray flares that may be consistent with these sorts of accretion events (see references in et al."817 2001)., 2001).818 It is nevertheless difficult to establish a definite link between observed flare events and accretion events. so information. about stripped cores in nearby galaxies would be important.," It is nevertheless difficult to establish a definite link between observed flare events and accretion events, so information about stripped cores in nearby galaxies would be important."819 The possibility of studying the stripped cores of disrupted stars is an important motivation of the search for SSSs in the central regions of galaxies., The possibility of studying the stripped cores of disrupted stars is an important motivation of the search for SSSs in the central regions of galaxies.820 In this section we have so far focused on models in which the SSSs we discover are luminous X-ray binaries., In this section we have so far focused on models in which the SSSs we discover are luminous X-ray binaries.821 We expect. however. that other types of objects will produce the same broadband X-ray signatures.," We expect, however, that other types of objects will produce the same broadband X-ray signatures."822" SNRs form the primary class of SSS ""contaminants"" that are luminous(Ly>1076 ere y. and which are actually members of the galaxy being observed."," SNRs form the primary class of SSS “contaminants"" that are luminous$L_X > 10^{36}$ erg $^{-1}$ ), and which are actually members of the galaxy being observed."823 In M31. 2 SSSs are SNRs.," In M31, $2$ SSSs are SNRs."824 Interestingly. one of these M31 SNRs is among the softest sources in M31.," Interestingly, one of these M31 SNRs is among the softest sources in M31."825 Most of the other 33 M31 SSSs we identified using the criteria presented in this paper. are highly variable: many are transients.," Most of the other $33$ M31 SSSs we identified using the criteria presented in this paper, are highly variable; many are transients."826 We therefore know that SNR contanimants form only à minor portion of the SSSs in M31., We therefore know that SNR contanimants form only a minor portion of the SSSs in M31.827 In Kong (2003 a). we studied the available data on the variability of SSSs in 4 more distant galaxies (MIOI. M83. M51. and NGC 4697; see also Kong 2003 b)," In Kong (2003 a), we studied the available data on the variability of SSSs in $4$ more distant galaxies (M101, M83, M51, and NGC 4697; see also Kong 2003 b)."828 Although the limited time coverage of the observations we studied allowed only the brightest sources to be checked for variability. we did find evidence of variability on time scales of a year.," Although the limited time coverage of the observations we studied allowed only the brightest sources to be checked for variability, we did find evidence of variability on time scales of a year."829 This is consistent with an X-ray binary nature for the majority of bright SSSs., This is consistent with an X-ray binary nature for the majority of bright SSSs.830 Knots in diffuse emission from the galaxy can also have very soft spectra. and they may be misidentified as SSSs when they cannot be spatially resolved.," Knots in diffuse emission from the galaxy can also have very soft spectra, and they may be misidentified as SSSs when they cannot be spatially resolved."831 This is most likely to occur near the centers of galaxies with a significant diffuse soft component. but can happen in any location in. which the X-ray emission appears to be dominated by diffuse emission.," This is most likely to occur near the centers of galaxies with a significant diffuse soft component, but can happen in any location in which the X-ray emission appears to be dominated by diffuse emission."832 If the sources are bright or the time sampling is good. time variability can help to identify which sources in regions of diffuse emission may be X-ray binaries: observations at other wavelengths may be helpful in finding counterparts to extended objects.," If the sources are bright or the time sampling is good, time variability can help to identify which sources in regions of diffuse emission may be X-ray binaries; observations at other wavelengths may be helpful in finding counterparts to extended objects."833 In the absence of such complementary information. however. SSSs discovered in regions of diffuse emission should not be assumed to be X-ray binaries.," In the absence of such complementary information, however, SSSs discovered in regions of diffuse emission should not be assumed to be X-ray binaries."834 Other systems identified by our algorithm are dim foreground objects or bright background objects., Other systems identified by our algorithm are dim foreground objects or bright background objects.835 Foreground stars can emit soft X-rays., Foreground stars can emit soft X-rays.836 In many cases such stars will have been identified by optical surveys and can be ruled out as luminous X-ray binaries., In many cases such stars will have been identified by optical surveys and can be ruled out as luminous X-ray binaries.837 In high surface brightness regions of the observed galaxy. however. the survey of foreground stars may be less complete. and we may not be able to identify which SSSs are foreground stars.," In high surface brightness regions of the observed galaxy, however, the survey of foreground stars may be less complete, and we may not be able to identify which SSSs are foreground stars."838 A further complication is that the soft X-ray emission from foreground stars can be highly variable. so variability cannot be taken as a signature that the SSS is an X-ray binary.," A further complication is that the soft X-ray emission from foreground stars can be highly variable, so variability cannot be taken as a signature that the SSS is an X-ray binary."839 Distant soft AGN can also be selected as SSSs: observations at other wavelengths can help to identify some. but probably not all of these.," Distant soft AGN can also be selected as SSSs; observations at other wavelengths can help to identify some, but probably not all of these."840 Further. some nearby magnetic CVs can also be selected as SSSs: it may be difficult to identify such sources at other wavelengths.," Further, some nearby magnetic CVs can also be selected as SSSs; it may be difficult to identify such sources at other wavelengths."841 The standard method to estimate the contribution of foreground and background sources is to use results derived from deep field surveys (Giaccont et 2001. Brandt et 2001).," The standard method to estimate the contribution of foreground and background sources is to use results derived from deep field surveys (Giacconi et 2001, Brandt et 2001)."842 Because. however. we are specifically interested 11 SSSs. which have not yet been studied in the deep fields. we have used another approach. sketched below. and discussed i more detail in Stefano et 2003).," Because, however, we are specifically interested in SSSs, which have not yet been studied in the deep fields, we have used another approach, sketched below, and discussed in more detail in Stefano et 2003)."843 Briefly. we have applied our algorithm to data from several fields analyzed by the ChAMP team.," Briefly, we have applied our algorithm to data from several fields analyzed by the ChAMP team."844 We consider only fields located away from the Galactic plane. and containing no clusters or galaxies.," We consider only fields located away from the Galactic plane, and containing no clusters or galaxies."845 [t such fields. we generally we find 1—3 VSSs in the 53 CCD.," In such fields, we generally we find $1-3$ VSSs in the S3 CCD."846 When. therefore. in observations of an external galaxy. we discover tens of VSSs in the S3 CCD. we can assume that the majority of them are associated with the galaxy.," When, therefore, in observations of an external galaxy, we discover tens of VSSs in the S3 CCD, we can assume that the majority of them are associated with the galaxy."847" Finally. we note that. although SNRs. foreground stars. and other ""contaminants"" do not dominate the VSSs we identify with galaxies. our algorithm does provide an efficient way to search for X-ray active SNRs and for a subset of foreground stars."," Finally, we note that, although SNRs, foreground stars, and other “contaminants"" do not dominate the VSSs we identify with galaxies, our algorithm does provide an efficient way to search for X-ray active SNRs and for a subset of foreground stars."848 Our phenomenological definition should. select. sources described by the physical models discussed above., Our phenomenological definition should select sources described by the physical models discussed above.849 The WD models alone define a broad range of temperatures. from 10 eV up to ~150 For example. while V751 Cyg had a best fit temperature with AT<10 eV. a L4M.. WD with Eddington-luminosity nuclear burning on its surface would have ΚΤ— 150eV. NBWD luminosities range from ~I0? erg s! up to the Eddington limit fora ΤΕΜ. object (~2«10? ere s! y.," The WD models alone define a broad range of temperatures, from $k\, T < 10$ eV up to $\sim 150$ For example, while V751 Cyg had a best fit temperature with $k\, T < 10$ eV, a $1.4\, M_\odot$ WD with Eddington-luminosity nuclear burning on its surface would have $k\, T \sim 150$ eV. NBWD luminosities range from $\sim 10^{35}$ erg $^{-1}$ up to the Eddington limit for a $1.4\, M_\odot$ object $\sim 2 \times 10^{38}$ erg $^{-1}$ )."850" The stripped core of à high-mass star might have a temperature near the low end of the temperature range. but a luminosity in excess of 10?? erg s!, Accreting BHs could have even higher luminosities. with temperatures in the SSS range or even higher."," The stripped core of a high-mass star might have a temperature near the low end of the temperature range, but a luminosity in excess of $10^{39}$ erg $^{-1}.$ Accreting BHs could have even higher luminosities, with temperatures in the SSS range or even higher."851 The sensitivities of the detectors used for X-ray astronomy tend to peak for photons with energies near or above | keV. Until the advent of and then ROSAT. it was difficult to study sources with energy distributions peaked significantly below | keV. It was also difficult to detect and study such," The sensitivities of the detectors used for X-ray astronomy tend to peak for photons with energies near or above $1$ keV. Until the advent of and then , it was difficult to study sources with energy distributions peaked significantly below $1$ keV. It was also difficult to detect and study such"852individually small contributions can be added together to produce a large overall result. if one has enough of them: and (hough Newtonian-evel perturbations weaken as ~1/7. the number of them in a spherical shell increases as 77? (easily overwhelming anv [factors of ο). creating a total perturbative effect that would formally be when integrated out to r=x. if not reined in by the finite causal horizon out to which an observer can ‘see clumped structure (hat has had sufficient time since the Die Bane to form: a situation reminiscent of Olbers: Paradox (Weinberg1972)..,"individually small contributions can be added together to produce a large overall result, if one has enough of them; and though Newtonian-level perturbations weaken as $\sim$$1/r$, the number of them in a spherical shell increases as $\sim$$r^{2}$ (easily overwhelming any factors of $v^{2}/c^{2}$ ), creating a total perturbative effect that would formally be when integrated out to $r = \infty$, if not reined in by the finite causal horizon out to which an observer can `see' clumped structure that has had sufficient time since the Big Bang to form; a situation reminiscent of Olbers' Paradox \citep{WeinbergGravCosmo}."853 For a causalitv-respecting approach. one must instead (as in electrodyvnanmics) use the full wave equation. lor special-relativistically consistent perturbation potential function $a: (Note that factors relating to the cosmic expansion are still neglected here. for simplicity.)," For a causality-respecting approach, one must instead (as in electrodynamics) use the full wave equation, for special-relativistically consistent perturbation potential function $\Phi _{\mathrm{SR}}$: (Note that factors relating to the cosmic expansion are still neglected here, for simplicity.)"854 Now. (he usual impulse is to immediately drop the extra term in Equation 1.. involving OpaΟΙ equivalent to dropping the eravilomagnetlic terms. as is done in the Buchert formalism — because of its resulting prefactor of 02/62: this factor would seem to make il verv small given (he assumption of nonrelativistic speeds for most matter flows. ancl thus (assumedlv) ensuring it to be negligible compared to the spatial variations (erm in any backreaction ealeulation.," Now, the usual impulse is to immediately drop the extra term in Equation \ref{EqnPoissonDynamic}, involving $\partial^2 \Phi_{\mathrm{SR}} / \partial t^2$ – equivalent to dropping the gravitomagnetic terms, as is done in the Buchert formalism – because of its resulting prefactor of $v^{2}/c^{2}$; this factor would seem to make it very small given the assumption of nonrelativistic speeds for most matter flows, and thus (assumedly) ensuring it to be negligible compared to the spatial variations term in any backreaction calculation."855 But this thinking is based only upon considerations of individual Fourier perturbation modes. not on the overall causal behavior of information flow in the siructure-forming universe.," But this thinking is based only upon considerations of individual Fourier perturbation modes, not on the overall causal behavior of information flow in the structure-forming universe."856 If we instead. all terms. ancl solve Equation 1. as-is. then one eels (adaptingfromJackson1915.eq. 6.69): where the bracketed numerator is always evaluated at theHime. l=1—|x—-x|/c.," If we instead all terms, and solve Equation \ref{EqnPoissonDynamic} as-is, then one gets \citep[adapting from][eq. 6.69]{JacksonEM}: : where the bracketed numerator is always evaluated at the, $t^{\prime} = 857t - \vert {\bf x} - {\bf x}^{\prime} \vert / c$."858 lt is Chis retarded-time condition which restores causalitv. allowing different regions of the universe (o communicate wilh (and gravitationally perturb) one another: and which provides (he escape route [rom the Buchert suppression of Newtonian-level backreaction. because such backreaction isnof (uly expressible as a total divergence.," It is this retarded-time condition which restores causality, allowing different regions of the universe to communicate with (and gravitationally perturb) one another; and which provides the escape route from the Buchert suppression of Newtonian-level backreaction, because such backreaction is truly expressible as a total divergence."859 We will refer to this propagation of gravitational perturbation information between distant (though communicating) regions as “Causal updating.," We will refer to this propagation of gravitational perturbation information between distant (though communicating) regions as “causal updating""."860 Given (he [act that the kev metric perturbation function. Φωνή). is predominantly affected by inhomogeneity information coming in from distant locations. the retarded-time condition of an integrated formula like Equation 2 (suitably modified for cosmological calculations) implicitly gives it the abilitv toexhibit relativistic behavior in what would," Given the fact that the key metric perturbation function, $\Phi _{\mathrm{SR}} (t)$, is predominantly affected by inhomogeneity information coming in from distant locations, the retarded-time condition of an integrated formula like Equation \ref{EqnSRpotential} (suitably modified for cosmological calculations) implicitly gives it the ability toexhibit relativistic behavior in what would"861overtone) for highest luminosities or smallest. &ravities in the figure (see Fig.,overtone) for highest luminosities or smallest gravities in the figure (see Fig.862 3 in Pamvatuykh 2000)., 3 in Pamyatnykh 2000).863 The Blue Edge for the radial fundamental mode lies approximately in the center of the 9 Scuti instability strip., The Blue Edge for the radial fundamental mode lies approximately in the center of the $\delta$ Scuti instability strip.864 An aclelitional study of the instability along the 2.5 M. evolutionary track shows that the best theoretical general Blue Edge for X=0.716. Y=0.26. Z=0.024 will be located. very. elose to the blue edge shown in Fig.," An additional study of the instability along the 2.5 $M_{\odot}$ evolutionary track shows that the best theoretical general Blue Edge for $X=0.716$, $Y=0.26$, $Z=0.024$ will be located very close to the blue edge shown in Fig."865 5. because the dillerences in the rotational velocity and in the overshooting eLlliclency do not inlluence the position of the Blue Eclees and because the dillerences in the helium abundance are small.," 5, because the differences in the rotational velocity and in the overshooting efficiency do not influence the position of the Blue Edges and because the differences in the helium abundance are small."866 Moreover. convection has only a minor influence on the position of this hot general Bluc Eclee (see Fig.," Moreover, convection has only a minor influence on the position of this hot general Blue Edge (see Fig."867 9 in Pamvatuykh 2000)., 9 in Pamyatnykh 2000).868 For the fundamental racial moodle the best Blue Edge will be hotter by 0.008—0.009 in logZ;47., For the fundamental radial mode the best Blue Edge will be hotter by $0.008-0.009$ in $\log T_{\rm{eff}}$.869 Phis is mainly due to a higher value of the mixing- parameter., This is mainly due to a higher value of the mixing-length parameter.870 From Fig., From Fig.871 5 we immediately obtain a strong constraint on the possible effective. temperature of the primary of 67 ‘Yau., 5 we immediately obtain a strong constraint on the possible effective temperature of the primary of $\theta^2$ Tau.872 All models with log(L/L.)>167 and logT;3.907 CHapZ7SOTO WIS) are stable in all modes., All models with $\log (L/L_{\odot})>1.67$ and $\log T_{\rm{eff}}>3.907$ $T_{\rm{eff}}>8070$ K) are stable in all modes.873 A AIS moclel of 2.5 AJ. on the Blue Edge (τμ=SOTO WIS) is marginally unstable in radial mode ps with the frequency. 18.74 1 which is well outside the observed frequency range.," A MS model of 2.5 $M_{\odot}$ on the Blue Edge $T_{\rm{eff}} = 8070$ K) is marginally unstable in radial mode $p_6$ with the frequency 18.74 $^{-1}$, which is well outside the observed frequency range."874 We can conclude that only significantly. cooler models can pulsate with the observed frequencies in the 10.8 to 16 ed.+ range., We can conclude that only significantly cooler models can pulsate with the observed frequencies in the 10.8 to 14.6 $^{-1}$ range.875 This conclusion. is confirmed by computation of oscillations of the selected test models for the primary of 67 ‘Tau., This conclusion is confirmed by computation of oscillations of the selected test models for the primary of $\theta^2$ Tau.876 In Fig., In Fig.877 6. the normalized. growth rates of. radial and. nonraclial modes are plotted: against frequency for. all nine higher-mass models which are marked in Fig.," 6, the normalized growth rates of radial and nonradial modes are plotted against frequency for all nine higher-mass models which are marked in Fig."878 5., 5.879 Only axisvmmetrie modes (i= 0) are shown., Only axisymmetric modes $m = 0$ ) are shown.880 The independence of the growth rate on the spherical harmonic degree. f£. is a tvpical feature of modes excited by the & mechanism.," The independence of the growth rate on the spherical harmonic degree, $\ell$, is a typical feature of modes excited by the $\kappa$ mechanism."881 The rotational velocities of the models are SI to 86 km/s. The rotational splitting of the modes can extend the frequency range by approximately 0.5 c/d on both sides., The rotational velocities of the models are 81 to 86 km/s. The rotational splitting of the modes can extend the frequency range by approximately 0.5 c/d on both sides.882 We can see that excited. frequencies of the 245 AL. model with dig=TSOO WH are in excellent agreement with the observed frequeney range., We can see that excited frequencies of the 2.45 $M_{\odot}$ model with $T_{\rm{eff}} = 7800$ K are in excellent agreement with the observed frequency range.883 The frequeney range of unstable modes spans three racial orders from py to ps for radial. modes (mode pj is marginally unstable)., The frequency range of unstable modes spans three radial orders from $p_4$ to $p_6$ for radial modes (mode $p_4$ is marginally unstable).884 As was noted already. thje results are sensitive to the treatment of convection.," As was noted already, the results are sensitive to the treatment of convection."885 For example. if we use à mixing-length parameter a=2.0 insead of à=1.6. the frequency range of unstable moces [ου 245 AJ. model with Zip TSOOIx. is extended by 1 cf on both sides.," For example, if we use a mixing-length parameter $\alpha=2.0$ instead of $\alpha=1.6$, the frequency range of unstable modes for 2.45 $M_{\odot}$ model with $T_{\rm{eff}} = 7800$ K is extended by 1 c/d on both sides."886 Moreover. our assumption about the unperurbed convective [ux during an oscillation cvwcle is. not fulfilled) inside the hydrogen convective zone and may result in artificial acdcditiona driving in this zone.," Moreover, our assumption about the unperturbed convective flux during an oscillation cycle is not fulfilled inside the hydrogen convective zone and may result in artificial additional driving in this zone."887 Pheree. these preliminary results must be considered with caution.," Therefore, these preliminary results must be considered with caution."888 Similar results were obt:uned for the models. withou overshooting. the best fitting is achieved. in this case for 2.50 AJ. model with Yup= TSOOWKK. Also. for slightly more massive ALS models with overshooting. we obtainec a good agreement between the observed and the theoretica frequency ranges.," Similar results were obtained for the models without overshooting, the best fitting is achieved in this case for 2.50 $M_{\odot}$ model with $T_{\rm{eff}} = 7800$ K. Also, for slightly more massive MS models with overshooting, we obtained a good agreement between the observed and the theoretical frequency ranges."889 The parameters of some mocels are given in ‘Table 4 below., The parameters of some models are given in Table 4 below.890 In Fig., In Fig.891 7 we show the normalized. growth rates in test models of the secondary. component., 7 we show the normalized growth rates in test models of the secondary component.892 These models are also nmiwked in Fig., These models are also marked in Fig.893 5., 5.894 As for the primary. we used mocdoels with elective. temperatures from. 7800. to KIN. The," As for the primary, we used models with effective temperatures from 7800 to K. The"895At this frequency the resolving power of the telescope is 9.6 aresee (Πα Power Beam Width) aud the 1 CIIz bandwidth corresponds to 11705.,At this frequency the resolving power of the telescope is 9.6 arcsec (Half Power Beam Width) and the 1 GHz bandwidth corresponds to 1170.896 The resulting spectral resolution and the noise iu the coadded and rebiuued spectrini were. respectively. 56 aand 0.3? mds (2.8 wv). leading to a accuracy of the flux density scale.," The resulting spectral resolution and the noise in the coadded and rebinned spectrum were, respectively, 56 and 0.3 mK (2.8 mJy), leading to a accuracy of the flux density scale."897" The i| line was detected at a siguificance level of Sa for the total exposure time of 12.1 h. The redshift and the peak intensity of the [C1] line are ig=6.11589+0.0006 and τος=11.5 ταν,", The ] line was detected at a significance level of $8\sigma$ for the total exposure time of 12.4 h. The redshift and the peak intensity of the ] line are $z_{\rm fs} = 6.4189 \pm 0.0006$ and $I_{158} = 11.8$ mJy.898 The reported error σ.=0.0006 corresponds to the uncertainty of the line position mcasurement of a κιν. whic[um is about one bin size iu the 11] spectrum: at the Nyquis+ iuit of 2 resolution clemeuts.," The reported error $\sigma_z = 0.0006$ corresponds to the uncertainty of the line position measurement of $\sigma_v \sim$ 24, which is about one bin size in the ] spectrum at the Nyquist limit of 2 resolution elements."899" Observations of the CO 3G) and 35) omission ines were obtained with the IRAAL Plateau de Bure interferometer at the frequencies 108.721 GIIz (the total integration tiuie Dig,=22 h) aud 93.206 GIIz (Lig,11 1). respectively (Bertoldi et al."," Observations of the CO $\rightarrow$ 6) and $\rightarrow$ 5) emission lines were obtained with the IRAM Plateau de Bure interferometer at the frequencies 108.724 GHz (the total integration time $T_{\rm exp} = 22$ h) and 93.206 GHz $T_{\rm exp} = 14$ h), respectively (Bertoldi et al."900 2003)., 2003).901 At about 5 arcsec angular resolution (5.77«LA” at 3.2 mim) the CO emission ine is unresolved aud colucides within the astrometric uncertainties of £0.3 arcsec with the optical position of he quasar eiven by Fan ct al. (, At about 5 arcsec angular resolution $5.7''\times4.1''$ at 3.2 mm) the CO emission line is unresolved and coincides within the astrometric uncertainties of $\pm0.3$ arcsec with the optical position of the quasar given by Fan et al. (9022003).,2003).903 The coadded 3 nuu data were rebiuned to GL (27 »6)andSo ((7—6 »5) resulting in the accuracy of the me position lüueasurenmieuts of a. ~36 aud 21 respectively.," The coadded 3 mm data were rebinned to 64 $J=7\rightarrow6$ ) and 55 $J=6\rightarrow5$ ) resulting in the accuracy of the line position measurements of $\sigma_v \sim$ 36 and 24, respectively."904 These uncertainties are again of a bin size in the reduced spectra., These uncertainties are again of a bin size in the reduced spectra.905 The redshifts and the peak intensities of the CO ⋅ : ⋖∣≻↭⋜⋯≼⊔∩≻⋅↱⊐⋝↕∐∐∖↴∖↴⋜∐⋅↸∖∙↥⋅↸∖↴∖↴↻↸∖↸⊳⊓↖⇁↸∖↕⋅↖↽∙−∙↕∴⊺⇂∶∩⋅⊔≝∟≻∶≓: (T6)⋅↽⊲ ∩∙∩∩∩∩∙∫∩−⋯∶⊇∙↕↕⊔⋅↧⋅↖⇁∙⋜⋯≼↧ 27-6.4189dcOLQUQG. fig3)=2.15 wwdy.," The redshifts and the peak intensities of the CO $\rightarrow$ 6) and $\rightarrow$ 5) lines are, respectively, $z^{(7-6)}_{\rm rot} = 6.4192 \pm 0.0009$ , $I_{(7-6)} = 2.14$ mJy, and $z^{(6-5)}_{\rm rot} = 6.4189 \pm 0.0006$, $I_{(6-5)} = 2.45$ mJy."906" Weighting the reported rotational redshifts with these peak iutensitics. one obtains the mean typo,=6.1190+ 0.0005."," Weighting the reported rotational redshifts with these peak intensities, one obtains the mean $z_{\rm rot} = 6.4190 \pm 0.0005$ ."907 We will take this value for the quasars svstenic redshift 7., We will take this value for the quasar's systemic redshift $z$.908 Using the reported redshift τε aud the averaged noe Eq.(9)) viclds ΔΕΕΞ(0.141.0)«103.," Using the reported redshift $z_{\rm fs}$ and the averaged $z_{\rm rot}$, \ref{EQ8}) ) yields ${\Delta F}/{F} = (0.1 \pm 1.0)\times10^{-4}$."909 The second [Cu] line was detected at τε=16908 towards the northern componcut of the quasar (lone et al., The second ] line was detected at $z_{\rm fs} = 4.6908$ towards the northern component of the quasar (Iono et al.910 2006)., 2006).911 The profile of this line is similar to the CO »1) and +6) lines seen at tor=LG6916 from the same component (Oment et al., The profile of this line is similar to the CO $\rightarrow$ 4) and $\rightarrow$ 6) lines seen at $z_{\rm rot} = 4.6916$ from the same component (Omont et al.912 1999)., 1999).913 The [Cu] 158 pau cinission was observed with the Subuullimeter Arrav interferometer (SMA. Πο oet al.," The ] 158 $\mu$ m emission was observed with the Submillimeter Array interferometer (SMA, Ho et al."914 2001)., 2004).915" The total exposure fine at a redshifted [C1] frequeney of 333.969 GIIz was Τον=19.6 h. and the angular resolution was 3.17«2.7""."," The total exposure time at a redshifted ] frequency of 333.969 GHz was $T_{\rm exp} = 19.6$ h, and the angular resolution was $3.4''\times2.7''$."916 The coadded spectrun was averaged using 120 bbin size resulting in the rs noise of 7.5 mJy. or the signal-to-noise ratio S/N~3 (the peak flux deuxitv ~23 mJy as shown inFig.," The coadded spectrum was averaged using 120 bin size resulting in the rms noise of 7.5 mJy, or the signal-to-noise ratio $\sim 3$ (the peak flux density $\sim$ 23 mJy as shown inFig."917 b iu Ione et al)., 1 in Iono et al.).918" Assiuuine the uncertainty of the | line position as —1/1 bin size. 61ο ects the error o,230 "," Assuming the uncertainty of the ] line position as $\sim$ 1/4 bin size, one gets the error $\sigma_v \simeq 30$ ."919"The CO J=5l1 line observed with the IRAM iuterferoimneter is detected at the ~56 confidence level (Ting,~16 b). and the CO πλ>6 line observed with the IRAM. 30-1 telescope at the 30 confidence level."," The CO $J=5\rightarrow4$ line observed with the IRAM interferometer is detected at the $\sim 5\sigma$ confidence level $T_{\rm exp} \sim 16$ h), and the CO $J=7\rightarrow6$ line observed with the IRAM 30-m telescope – at the $3\sigma$ confidence level."920 The aneularOo resolution im the interferometric observations was 5442.5 arcsec. and velocity resolution of about GOἘ," The angular resolution in the interferometric observations was $5.0\times2.5$ arcsec, and velocity resolution of about 60."921"ν, The redshift of the northern compouent is τα2=L6016.", The redshift of the northern component is $z_{5-4} = 4.6916$.922 TIje uncertaintv of this value is. probably. 25-30Woede approximately one resolution clement. cousidering rather noisy line profiles shown in Fig.," The uncertainty of this value is, probably, 25-30, i.e. approximately one resolution element, considering rather noisy line profiles shown in Fig."923 2 in Omout et al., 2 in Omont et al.924 The aneular resolution of the 30-12 telescope for the 2-uuu beam dis 17 arcsec., The angular resolution of the 30-m telescope for the 2-mm beam is 17 arcsec.925 The error of the reporte redshift i£[46=|6915d:0.001 corresponds to the radia velocity uncertainty of 53, The error of the reported redshift $z_{7-6} = 4.6915\pm0.001$ corresponds to the radial velocity uncertainty of 53.926" Taking mto account that the angular resolutions are similar in observations of the [Cul and CO f=5)> enisson lines. we| en use their redshifts τε=L6908+0.0006 and s=1.691643:0.0006 (both errors correspon to a,=30 Ly) to calculate ΑΕΕ=(1.1510 "," Taking into account that the angular resolutions are similar in observations of the ] and CO $J=5\rightarrow4$ emission lines, we can use their redshifts $z_{\rm fs} = 4.6908\pm0.0006$ and $z_{\rm rot} = 4.6916\pm0.0006$ (both errors correspond to $\sigma_v = 30$ ) to calculate ${\Delta F}/{F} = (1.4 \pm 1.5)\times10^{-4}$ ."927While comparing the redshifted. frequencies. of. to castre hvwpothetical variations of physical constants. one niust account for random Doppler shifts of the line positions caused by non-identical spatial distributions of s)ecies (referred to as the Doppler noise hereinafter) whicl1 can niuie non-zero signals in oor oor Ina combinaion of these quantities (0... Levshakoy 1991: Caxilli et al.," While comparing the redshifted frequencies of to measure hypothetical variations of physical constants, one must account for random Doppler shifts of the line positions caused by non-identical spatial distributions of species (referred to as the Doppler noise hereinafter) which can mimic non-zero signals in or or in a combination of these quantities (e.g., Levshakov 1994; Carilli et al."928 2000: Bahcall et al., 2000; Bahcall et al.929 2001)., 2004).930 To quantity Πο iudiwed by the Doppler noise a sample of (Aqfa)(Ai) measurements is to be collected., To quantify uncertainties induced by the Doppler noise a sample of $(\Delta\alpha/\alpha)/(\Delta\mu/\mu)$ measurements is to be collected.931 The man pr‘oblem here is how to cstimate the dispersion of random velocity shifts 0. for a given system of spectral lines., The main problem here is how to estimate the dispersion of random velocity shifts $\sigma_{\rm v}$ for a given system of spectral lines.932" In case of a laree sample size the value of ao, can be found from the scatter of »outs.", In case of a large sample size the value of $\sigma_{\rm v}$ can be found from the scatter of points.933 For a sinele measurement. a guess for σς comes roni the comparison with data on velocity differences vetwween spectral lines of similar species in nearby clouds.," For a single measurement, a guess for $\sigma_{\rm v}$ comes from the comparison with data on velocity differences between spectral lines of similar species in nearby clouds."934" Observations of k(al galaxies show that the inteusitv of Ci] is strongly correlated with the intensities of the low-vine rotational |ines of CO. the fine-structure dines of Ci] AASTO0.609. , πα and [OL AAG3.116. pan (Malhotra ( al."," Observations of local galaxies show that the intensity of ] is strongly correlated with the intensities of the low-lying rotational lines of CO, the fine-structure lines of ] $\lambda\lambda370, 609$ $\mu$ m and ] $\lambda\lambda63, 146$ $\mu$ m (Malhotra et al."935 2001). ane the fine-structure line of [ΠΠ A205 pau (Petuchowski Bennett 1993: Abel 2006).," 2001), and the fine-structure line of ] $\lambda205$ $\mu$ m (Petuchowski Bennett 1993; Abel 2006)."936 However. 1ο surface distribution of the [Ci] emission may not xeciselv. follow t1ο actual D 2CO- contours (Stacey.+ et al.," However, the surface distribution of the ] emission may not precisely follow the actual $^{12}$ CO contours (Stacey et al."937 19855)., 1985).938 The CO rotational ues. if optically thin. are (uitted throughout the whole molecular cloud.," The CO rotational lines, if optically thin, are emitted throughout the whole molecular cloud."939 As for jio. r1]. endissio1 it is usually enhanced at the edges of the molecular cloud in the photodissociation regions (PDRs).," As for the ] emission, it is usually enhanced at the edges of the molecular cloud in the photodissociation regions (PDRs)."940 Additionally. diffuse eas from the regions cau coutributeto the intensity of the ΡΟ lines (Ixaufuia et al.," Additionally, diffuse gas from the regions can contributeto the intensity of the `PDR' lines (Kaufman et al."941 1999)., 1999).942 ILowever. the impact from the diffuse gas decreases with increasing gas densities anddrops from YA at nyp~locuni ," However, the impact from the diffuse gas decreases with increasing gas densities anddrops from $\sim$ at $n_{\rm H} \sim 1$ "943the uncertainties ou these parameters. which lie between ~1050%.,"the uncertainties on these parameters, which lie between $\sim$ 10–50."944. Therefore. in the general study above. this effect is not crucial in the determination of the elobal characteristics of the flares.," Therefore, in the general study above, this effect is not crucial in the determination of the global characteristics of the flares."945 Since the refined wwodel also leads to lower \2 values. the main source of uucertaiuty that now remains is the level of the background radio flix.," Since the refined model also leads to lower $\chi_\nu^2$ values, the main source of uncertainty that now remains is the level of the background radio flux."946 However. in the case of a strong backeround. this level is poorly constrained.," However, in the case of a strong background, this level is poorly constrained."947 Indeed. there is a degeneracy between the offset level Avy aud the other parameters: if [y is fixed to lower values. the fitting routine will chauge the other parameters to increase the overlap between flares. which will lead to an acceptable ft.," Indeed, there is a degeneracy between the offset level $K_{0}$ and the other parameters: if $K_{0}$ is fixed to lower values, the fitting routine will change the other parameters to increase the overlap between flares, which will lead to an acceptable fit."948 This effect is particularly true in the case of strong overlap., This effect is particularly true in the case of strong overlap.949" Iu order to quantify this issue. we attributed a ""confidence index” to our data."," In order to quantify this issue, we attributed a “confidence index"" to our data."950 This index is based onu two criteria: the level of backeround fiux. aud the uorphologv of the observed flares.," This index is based on two criteria: the level of background flux, and the morphology of the observed flares."951 Indeed. when a radio oscillation is fully shaped. im particular with a clear exponential tail we have good confidence that the vackeround subtraction is accurate.," Indeed, when a radio oscillation is fully shaped, in particular with a clear exponential tail, we have good confidence that the background subtraction is accurate."952 Ou the other haud. in the case when only the “tip” of the flare is visible. the vackerouud level caunot be precisely constrained. which is asource of higher uncertainty for the determination of flare properties.," On the other hand, in the case when only the “tip"" of the flare is visible, the background level cannot be precisely constrained, which is a source of higher uncertainty for the determination of flare properties."953 We attributed a confidence iudex of 1 to the most reliable data: observations with fully shaped flares. either isolated or ou top of a low background enisiou (= 15 mJy. Fig. 1)).," We attributed a confidence index of 1 to the most reliable data: observations with fully shaped flares, either isolated or on top of a low background emission $\la$ 15 mJy, Fig. \ref{exemples1}) )."954 Au iudex of 2 correspouds to observations where the expoucutial decrease is still clearly visible. witli inoderate background emission (<30 mJy. Fig. 2..," An index of 2 corresponds to observations where the exponential decrease is still clearly visible, with moderate background emission $\la$ 30 mJy, Fig. \ref{exemples2},"955 left)., left).956 Finally. observations with almost sinusoidal oscillations ou top of a strong backgrouud (up to 120 indy. Fig. 2..," Finally, observations with almost sinusoidal oscillations on top of a strong background (up to 120 mJy, Fig. \ref{exemples2},"957 rieht) correspoud to iudex 3., right) correspond to index 3.958 Thus. parameters deduced roni iudex 3 observations should be considered with care. While iudex 1 observations would produce reliable »uwanmeters.," Thus, parameters deduced from index 3 observations should be considered with care, while index 1 observations would produce reliable parameters."959 Among the observations listed iun. Table 1.. four observations had to be excluded. from our data set: on lhree observatious. two successive flares are visible but rot distiuguishable. while ou one more observation the radio data are too noisy to ect any constraint on the Hare.," Among the observations listed in Table \ref{logg}, four observations had to be excluded from our data set: on three observations, two successive flares are visible but not distinguishable, while on one more observation the radio data are too noisy to get any constraint on the flare."960 To measure the characteristics of the N-rayv dips. we used the following definition.," To measure the characteristics of the X-ray dips, we used the following definition."961 Time 0 is defined as the time when the phase of highly variable N-ray flux euds., Time 0 is defined as the time when the phase of highly variable X-ray flux ends.962 This time is also the first point of spectral hardening (visible in the ITR)., This time is also the first point of spectral hardening (visible in the HR).963 We used the position of the masini flux of the N-rav spike to mark the end of the dip. whose duration will be noted Af.," We used the position of the maximum flux of the X-ray spike to mark the end of the dip, whose duration will be noted $\Delta t$."964" At this time. the IIR has gone down to values close to the pre-dip ράσο,"," At this time, the HR has gone down to values close to the pre-dip phase."965Furthermore. it has recently been suggested that θά of the stellar mass at 2c5 is missed by traditional drop-out selection techniques (Stark et al.,"Furthermore, it has recently been suggested that $\simeq60\%$ of the stellar mass at $z\simeq5$ is missed by traditional drop-out selection techniques (Stark et al."966 2006). which exclude high-redshift galaxies which are too red in the rest-frame optical to satisfy LBG selection criteria (see discussion in Section 6).," 2006), which exclude high-redshift galaxies which are too red in the rest-frame optical to satisfy LBG selection criteria (see discussion in Section 6)."967 However. even if we increase our estimated number density of massive galaxies at 2o5 by a factor of two. and assume ax=0.75. XCDM can still produce the required number of dark matter halos if the stellar to dark matter ratio is 10.," However, even if we increase our estimated number density of massive galaxies at $z\simeq5$ by a factor of two, and assume $\sigma_{8}=0.75$, $\Lambda$ CDM can still produce the required number of dark matter halos if the stellar to dark matter ratio is $\simeq 10$."968 In Fig 6 we also show the estimated number densities of ogML galaxies at 2—3 derived from α ἐν band selectec sample of galaxies in the GOODS CDFS (Caputi et al., In Fig 6 we also show the estimated number densities of $\geq10^{11}\Msolar$ galaxies at $z\leq3$ derived from a $K-$ band selected sample of galaxies in the GOODS CDFS (Caputi et al.969 2006)., 2006).970 The Caputi et al., The Caputi et al.971 sample has a magnitude limit of ἐν<23.3. and is c85% complete for stellar masses >1075M. (Salpeter IMP) a redshifts 2.—3.," sample has a magnitude limit of $K\leq23.3$, and is $\geq85$ complete for stellar masses $\geq10^{11}\Msolar$ (Salpeter IMF) at redshifts $z\leq3$."972 As can be seen from Fig 6. the conclusions of Caputi et al.," As can be seen from Fig 6, the conclusions of Caputi et al."973 were that 204€ of the local mass density comprised by galaxies with stellar masses2104M. was already in place by z&2. and that —the majority was in place by z—1.," were that $\simeq20$ of the local mass density comprised by galaxies with stellar masses$\geq10^{11}\Msolar$ was already in place by $z\simeq2$, and that the majority was in place by $z\simeq1$."974 Similar conclusions have been reached by numerous studies of deep. smal area. near-infrared surveys (e.g. Drory et al.," Similar conclusions have been reached by numerous studies of deep, small area, near-infrared surveys (e.g. Drory et al."975 2005: Fontana 2004)., 2005; Fontana 2004).976 The results presented here strongly suggest that the increase in number density of massive galaxies within the redshift interva 2.5«25.5 closely traces the build-up of suitable dark matter halos and that. as a consequence. only ::14 of the local density of 2104}AL. galaxies was in place by zc5.," The results presented here strongly suggest that the increase in number density of massive galaxies within the redshift interval $2.5<z<5.5$ closely traces the build-up of suitable dark matter halos and that, as a consequence, only $\ltsim\,1\%$ of the local density of $\geq10^{11}\Msolar$ galaxies was in place by $z\simeq5$."977 It has recently become apparent that the number density of massive. red galaxies observed at high redshift is much larger than would have been predicted by semi-analytic galaxy formation models five years ago (e.g. Cimatti et al.," It has recently become apparent that the number density of massive, red galaxies observed at high redshift is much larger than would have been predicted by semi-analytic galaxy formation models five years ago (e.g. Cimatti et al."978 2002: Glazebrook et al., 2002; Glazebrook et al.979 2004)., 2004).980 In essence this is simply a result of the fact that galaxy formation models based on hierarchical assembly naturally produce the most massive galaxies at late times., In essence this is simply a result of the fact that galaxy formation models based on hierarchical assembly naturally produce the most massive galaxies at late times.981" This apparent discrepancy between the models and observations has been described variously as ""anti-hierarchical” behaviour. or cosmic downsizing."," This apparent discrepancy between the models and observations has been described variously as ``anti-hierarchical'' behaviour, or cosmic downsizing."982 It is clear therefore that important constraints can be placed on the latest generation of galaxy formation models via comparison with the observed number density of massive galaxies already in place at 25., It is clear therefore that important constraints can be placed on the latest generation of galaxy formation models via comparison with the observed number density of massive galaxies already in place at $z\geq5$.983 In Fig 7 we show the predicted number density of galaxies as a function of stellar mass at 2=5.3 (median. redshift of our LBG candidates) from the latest generation of Durham semi-analytic galaxy formation models. which incorporate AGN feedback (Bower et al.," In Fig 7 we show the predicted number density of galaxies as a function of stellar mass at $z=5.3$ (median redshift of our LBG candidates) from the latest generation of Durham semi-analytic galaxy formation models, which incorporate AGN feedback (Bower et al."984 2006)., 2006).985 Also shown in Fig 7 are the predictions (G.L. Granato. private communication) from the latest version of the Granato et al. (," Also shown in Fig 7 are the predictions (G.L. Granato, private communication) from the latest version of the Granato et al. ("9862004) galaxy formation model. which is based on an anti-hierarchieal baryon collapse scenario.,"2004) galaxy formation model, which is based on an anti-hierarchical baryon collapse scenario."987 The filled circle is our estimate for the number density of galaxies within a bin centred on 104XL.. based on the results of this study.," The filled circle is our estimate for the number density of galaxies within a bin centred on $10^{11}\Msolar$, based on the results of this study."988 When plotting the data-point in Fig 7. we have shifted the centre ofthe bin to lower mass by 0.15 dex. to account for the difference between the Salpeter IMF adopted in this paper. and the IMFs adopted in the Bower et al.," When plotting the data-point in Fig 7, we have shifted the centre of the bin to lower mass by 0.15 dex, to account for the difference between the Salpeter IMF adopted in this paper, and the IMFs adopted in the Bower et al."989 and Granato et al., and Granato et al.990 models (Kennicutt and Romano et al., models (Kennicutt and Romano et al.991 2002 respectively)., 2002 respectively).992 As can be seen from Fig 7. our estimated number density is clearly in good agreement with the predictions of the Bower et al.," As can be seen from Fig 7, our estimated number density is clearly in good agreement with the predictions of the Bower et al."993 and Granato et al., and Granato et al.994 models., models.995 Interestingly. given that our number density is almost certainly a lower limit. both model predictions can accommodate a number density which is a factor of &2—3 higher than our estimate. as required if LBGs constitute <50% of the total stellar mass at >25 (e.g. Stark et al.," Interestingly, given that our number density is almost certainly a lower limit, both model predictions can accommodate a number density which is a factor of $\simeq2-3$ higher than our estimate, as required if LBGs constitute $\leq50\%$ of the total stellar mass at $z\geq5$ (e.g. Stark et al."996 2006)., 2006).997 Finally. although not included in Fig 7. we also note that our estimated number density of 2°<25 LBGs at 2>5 is fully consistent with the hydro-dynamical simulations of Night et al. (," Finally, although not included in Fig 7, we also note that our estimated number density of $z^{\prime}\leq25$ LBGs at $z\geq5$ is fully consistent with the hydro-dynamical simulations of Night et al. ("9982006).,2006).999 In this study we have exploited the large co-moving volume covered by the SXDS/UDS data-set to identify a sample of bright (x 25) LBGs at 2>5., In this study we have exploited the large co-moving volume covered by the SXDS/UDS data-set to identify a sample of bright $z^{\prime}\leq25$ ) LBGs at $z\geq5$.1000 In this section we compare our results with those of three recent studies based on the extensive multi-wavelength data available in the GOODS CDFS (Grazian et al., In this section we compare our results with those of three recent studies based on the extensive multi-wavelength data available in the GOODS CDFS (Grazian et al.1001 2006: Stark et al., 2006; Stark et al.1002 2006: Yan et al., 2006; Yan et al.1003 2006)., 2006).1004 The GOODS-MUSIC sample (Grazian et al., The GOODS-MUSIC sample (Grazian et al.1005 2006) consists ofz bandand A/ band selected samples in the southern GOODS field. and provides a photometric redshift for each object calculated from SED fits to the extensive. high-quality. HST+VLT+Spitzer imaging available in the field.," 2006) consists of $z-$ band and $K-$ band selected samples in the southern GOODS field, and provides a photometric redshift for each object calculated from SED fits to the extensive, high-quality, HST+VLT+Spitzer imaging available in the field."1006 The + band catalogue is 100A complete to zs).=25 and. although it only covers 142. square aremin. it is obviously of some interest to compare the number of 25 galaxies found in the MUSIC sample to the results presented here.," The $z-$ band catalogue is $100\%$ complete to $z_{850}=25$ and, although it only covers 142 square arcmin, it is obviously of some interest to compare the number of $z\geq5$ galaxies found in the MUSIC sample to the results presented here."1007 Based on the results of this study. and correcting for the difference in area. the number of za5o25 LBG-type galaxies within the GOODS-MUSIC catalogue should be consistent with 0.732:0.30.," Based on the results of this study, and correcting for the difference in area, the number of $z_{850}\leq25$ LBG-type galaxies within the GOODS-MUSIC catalogue should be consistent with $0.73\pm0.30$."1008" Interestingly. the GOODS-MUSIC catalogue contains seven objects with 2.5525 and a photometric redshift estimate of 25, seemingly very inconsistent with the results presented here."," Interestingly, the GOODS-MUSIC catalogue contains seven objects with $z_{850}\leq25$ and a photometric redshift estimate of $z\geq5$, seemingly very inconsistent with the results presented here."1009 However. our own SED fits to the MUSIC photometry for these seven objects suggest that only two are believable high-redshift candidates.," However, our own SED fits to the MUSIC photometry for these seven objects suggest that only two are believable high-redshift candidates."1010 Three of the seven objects (MUSIC IDs 1133. 8316 12966) have SEDs which are clearly stellar in nature. and have SExtractor stellarity parameters of 0.98. 0.99 and 0.99 respectively.," Three of the seven objects (MUSIC IDs 1133, 8316 12966) have SEDs which are clearly stellar in nature, and have SExtractor stellarity parameters of 0.98, 0.99 and 0.99 respectively."1011" In addition. MUSIC ID27004 has a clear V— band detection which is inconsistent with the best-fitting photometric redshift of 2,,,,,= 6.91."," In addition, MUSIC ID=7004 has a clear $V-$ band detection which is inconsistent with the best-fitting photometric redshift of $z_{phot}=6.91$ ."1012 Finally. MUSIC ID=10140 has a very unusual SED which would appear to be due to severe blending with bright nearby objects in the near-infrared and Spitzer bands.," Finally, MUSIC ID=10140 has a very unusual SED which would appear to be due to severe blending with bright nearby objects in the near-infrared and Spitzer bands."1013 Of the original seven, Of the original seven1014"characteristic of this region, in order to reproduce the triangular shape of the observed feature.","characteristic of this region, in order to reproduce the triangular shape of the observed feature."1015 The two intense humps at both sides of the central maximum would come from the hollow shells., The two intense humps at both sides of the central maximum would come from the hollow shells.1016" However, the high-velocity wings are severely underestimated by our standard model, and we would have to significantly increase the excitation conditions of the very fast bipolar outflow of 6618 to reproduce their intensity."," However, the high-velocity wings are severely underestimated by our standard model, and we would have to significantly increase the excitation conditions of the very fast bipolar outflow of 618 to reproduce their intensity."1017" We can also see in 44 our predictions for a model similar to the previous one, but with Ty ~ 200 K in the regions that present expansion velocities  100Κπις-]."," We can also see in 4 our predictions for a model similar to the previous one, but with $T_{\rm k}$ $\sim$ 200 K in the regions that present expansion velocities $\sim$ 100."1018". Other parameters of the fast outflow, particularly its velocity and density distributions do not change with respect to the original model."," Other parameters of the fast outflow, particularly its velocity and density distributions do not change with respect to the original model."1019 The asymmetry between the red and blue line wings is not reproduced by our predictions also in this case., The asymmetry between the red and blue line wings is not reproduced by our predictions also in this case.1020 It is remarkable that the high-velocity outflow obviously contributes to the emission at the profile central features., It is remarkable that the high-velocity outflow obviously contributes to the emission at the profile central features.1021" Therefore, in this model the requirements to reproduce the secondary maxima are significantly weaker, and a temperature similar to or even lower than assumed in our original model for the empty shells would be compatible with the observations."," Therefore, in this model the requirements to reproduce the secondary maxima are significantly weaker, and a temperature similar to or even lower than assumed in our original model for the empty shells would be compatible with the observations."1022" The general properties of the central, dense component mentioned before, in particular its low velocity, are required in this case too."," The general properties of the central, dense component mentioned before, in particular its low velocity, are required in this case too."1023" The high velocity wings are also detected in emission from other molecules, such as H20, HCN, and CN 11), which must be significantly abundant in this recently shocked gas."," The high velocity wings are also detected in emission from other molecules, such as $_2$ O, HCN, and CN 1), which must be significantly abundant in this recently shocked gas."1024 The relatively high temperatures deduced here for the fast bipolar flow relax the discrepancy usually found between the high excitation of the shocked gas predicted by theoretical models and the observational results (anintricatetheoretical 2009).," The relatively high temperatures deduced here for the fast bipolar flow relax the discrepancy usually found between the high excitation of the shocked gas predicted by theoretical models and the observational results \citep[an intricate theoretical problem not discussed in this1025letter, see e.g.][]{lee09}."1026" However, the discrepancy persists."," However, the discrepancy persists."1027" Leeetal.(2009) predict temperatures of the high-velocity gas in 6618 over ~ 1000 K and too weak CO emission in all rotational lines, since shocks are expected to dissociate molecules."," \cite{lee09} predict temperatures of the high-velocity gas in 618 over $\sim$ 1000 K and too weak CO emission in all rotational lines, since shocks are expected to dissociate molecules."1028" The low-excitation component of the nebula model by SC04, the extended halo, has apparently no counterpart in the observations, because the model predicts a very low intensity from such cool gas and the high-J line profiles do not seem to require any contribution from it."," The low-excitation component of the nebula model by SC04, the extended halo, has apparently no counterpart in the observations, because the model predicts a very low intensity from such cool gas and the $J$ line profiles do not seem to require any contribution from it."1029 In 11 we also show our HIFI spectra of llines., In 1 we also show our HIFI spectra of lines.1030" The oobservations are not very sensitive, so the line is not detected with a limit Timp 00.2 K. As we can see, the contrast between aand llines is high (mainly in the line wings, about a factor ten for the highest transitions)."," The observations are not very sensitive, so the line is not detected with a limit $T_{\rm mb}$ 0.2 K. As we can see, the contrast between and lines is high (mainly in the line wings, about a factor ten for the highest transitions)."1031" This result is compatible with our calculations, which suggest moderate opacities in high-J llines from the main nebular components, in particular with 1(16-15) 11 for gas flowing at 1100s“!."," This result is compatible with our calculations, which suggest moderate opacities in $J$ lines from the main nebular components, in particular with $\tau$ $-$ 15) 1 for gas flowing at 100."1032", Our results can therefore be summarized as follows 1.", Our results can therefore be summarized as follows 1.1033 We detected high-J CO emission using Herschel/HIFI., We detected $J$ CO emission using Herschel/HIFI.1034" The high-velocity line wings characteristic of this source become progressively dominant as the level energies increase, with sshowing a spectacular composite profile."," The high-velocity line wings characteristic of this source become progressively dominant as the level energies increase, with showing a spectacular composite profile."1035 2., 2.1036" The temperature of the very fast bipolar outflow in CRL6618 is high, significantly higher than the previously adopted values."," The temperature of the very fast bipolar outflow in 618 is high, significantly higher than the previously adopted values."1037" SCO4 proposed a temperature for this component < 100 K, which, in view of the intense line wings seen in the ttransition, must be significantly increased."," SC04 proposed a temperature for this component $<$ 100 K, which, in view of the intense line wings seen in the transition, must be significantly increased."1038" From our calculations, we estimate that gas flowing at about 100 mmust have a temperature ~ 200 K. We suggest that this very fast outflow, with a kinematic age < 100 yr, was accelerated by a shock and has not yet fully cooled down."," From our calculations, we estimate that gas flowing at about 100 must have a temperature $\sim$ 200 K. We suggest that this very fast outflow, with a kinematic age $<$ 100 yr, was accelerated by a shock and has not yet fully cooled down."1039" The rest of the physical conditions are not significantly changed, therefore"," The rest of the physical conditions are not significantly changed, therefore"1040the number of stars in the second bin. the 50 returned. “detections” should include 4-5 real planetary transits if 47 Tuc were like the Solar Neighbourhood.,"the number of stars in the second bin, the 50 returned “detections” should include 4-5 real planetary transits if 47 Tuc were like the Solar Neighbourhood."1041 It is important to note that the proceedure used to weed out false positive detections could conceivably include clisreearcling real transits., It is important to note that the proceedure used to weed out false positive detections could conceivably include disregarding real transits.1042 It was found by visual inspection of randomly selected model transits that the signature is clearly visible on the data and has an extremely high. detection significance compared to the svstematic false positives., It was found by visual inspection of randomly selected model transits that the signature is clearly visible on the data and has an extremely high detection significance compared to the systematic false positives.1043 It is nol expected that the transit recoverability would be significantly inpinged by this visual inspection method., It is not expected that the transit recoverability would be significantly impinged by this visual inspection method.1044 The majority of false detections were caused by erowcding and are immediately identified by plotüng the lighteurve. which displays systematic features occuring al the same times Lor most false detections.," The majority of false detections were caused by crowding and are immediately identified by plotting the lightcurve, which displays systematic features occuring at the same times for most false detections."1045 The problem was less severe in (he uncrowcded outer regions of 47 Tuc. where the false detection rates were 1/3 the values presented above.," The problem was less severe in the uncrowded outer regions of 47 Tuc, where the false detection rates were 1/3 the values presented above."1046 Fig.12. shows the period distribution of all lighteurves that pass the detection criteria and as such is the period distribution of the false cleteetions that were common in the dataset., \ref{fdhist} shows the period distribution of all lightcurves that pass the detection criteria and as such is the period distribution of the false detections that were common in the dataset.1047 A few periods have a significant. excess of detections over (he general trend., A few periods have a significant excess of detections over the general trend.1048 In particular {ρω= 1.5 days and Prod= 448 days show a large excess of “detections”., In particular $\it{P_{mod}}=$ 1.5 days and $\it{P_{mod}}=$ 4.48 days show a large excess of “detections”.1049 As it is more difficult for svstematies to produce a strong detection when the number of in-(üransit data points is reduced. the nunbers of false detections decreases as period increases.," As it is more difficult for systematics to produce a strong detection when the number of in-transit data points is reduced, the numbers of false detections decreases as period increases."1050 The code is quick and easy to apply., The code is quick and easy to apply.1051 Application to our whole 47 Tue dataset was accomplished in 16 steps on four 3057MIILz 4096Mb RAAT. i86pe machines.," Application to our whole 47 Tuc dataset was accomplished in 16 steps on four 3057MHz 4096Mb RAM, i86pc machines."1052 The code completed the task in 12 hours. and is easily modifiable to run on existing datasets.," The code completed the task in 12 hours, and is easily modifiable to run on existing datasets."1053 We have presented and described a quick. efficient. and easy. to apply computational method for the detection of planetary transits in large photometric (me-series datasets.," We have presented and described a quick, efficient and easy to apply computational method for the detection of planetary transits in large photometric time-series datasets."1054 Using a cross correlation Function (tlie code compares each sampled lishteurve wilh a database ol model transits of appropriate transit depth and duration., Using a cross correlation function the code compares each sampled lightcurve with a database of model transits of appropriate transit depth and duration.1055 A detection is implied by a sienificantly high value of the correlation distribution., A detection is implied by a significantly high value of the correlation distribution.1056 Monte Carlo simulations using the actual temporal sampling ancl photometric characteristics of the data superimposed. on appropriately modeled transits. show an excellent weighted mean recoverability rate over the whole of the sampled period range with a relatively low false detection probability.," Monte Carlo simulations using the actual temporal sampling and photometric characteristics of the data superimposed on appropriately modeled transits, show an excellent weighted mean recoverability rate over the whole of the sampled period range with a relatively low false detection probability."1057 In parücular the code achieves à very good recoverability when searching lor transits at or just below (he photometric noise level of the data., In particular the code achieves a very good recoverability when searching for transits at or just below the photometric noise level of the data.1058 The code is easily adaptable to run on existing datasets to search for the same photometric signals. and is capable of testing 10.000 stars in," The code is easily adaptable to run on existing datasets to search for the same photometric signals, and is capable of testing 10,000 stars in"1059Bellaziuis surface brightuess estimates are expected to be slightly too faint in the outer parts 22119.,Bellazini's surface brightness estimates are expected to be slightly too faint in the outer parts 2419.1060 The radial surface brightuess profiles of four exemplary moclels projected onto the sky aud converted to units⋅ of⋅ mag > are shown in⋅ Figure⊲⋅ ye13., The radial surface brightness profiles of four exemplary models projected onto the sky and converted to units of mag $^{-2}$ are shown in Figure \ref{surfbright}.1061.The merger objects show a Ixing-like prolile out to radii of about, The merger objects show a King-like profile out to radii of about.1062The observed profile from Bellazzini(2007) is acided to Figure 13 to allow for a direct comparison., The observed profile from \cite{bellazzini} is added to Figure \ref{surfbright} to allow for a direct comparison.1063 The surface brightness profile of model 91060 agrees very well with the observed profile at all radii., The surface brightness profile of model 50 agrees very well with the observed profile at all radii.1064 The other three models in Figre 13. illustrate how he profiles chauge when oue of the parameter mass. size and initial distribution of star clusters is inodified.," The other three models in Figure \ref{surfbright} illustrate how the profiles change when one of the parameter mass, size and initial distribution of star clusters is modified."1065 Model M.11-11.5-5250 has a similar shape as 520.but is too bright at al adii due to the larger mass.," Model 50 has a similar shape as 50, but is too bright at all radii due to the larger mass."1066" Model 1100. which is a more extencec| version of mocle 1.11.11.5.220. agrees well with the observations between 'and"".. but it is cousicerably wiehter in the center aud the outer parts."," Model 100, which is a more extended version of model 50, agrees well with the observations between and, but it is considerably brighter in the center and the outer parts."1067 Model 1100. wh has a broader initia distribution of star clusters than mocel 1100. shows a smaller cdeviatic1 from the observec surface brightness profile than moclel 1100.," Model 100, which has a broader initial distribution of star clusters than model 100, shows a smaller deviation from the observed surface brightness profile than model 100."1068 Baungardtetal.(2009) observed the radial velocities of {0 stars within a srojected radius of 100 pe of 22119 aud derived a velocity dispersion of o; = 1240. 15 kin ., \cite{baumgardt} observed the radial velocities of 40 stars within a projected radius of 100 pc of 2419 and derived a velocity dispersion of $\sigma$ = $\pm$ 0.48 km $^{-1}$.1069 The liue-o[-sight velocity dispersious within a projected radius of 100 pe of the moclels are lised in Table 3.., The line-of-sight velocity dispersions within a projected radius of 100 pc of the models are listed in Table \ref{tbl-2}.1070 A number of models with masses A5 = 1.0. 1.5. and 2.0 x 105 M.. have velociv dispersious that are within the one sigma error of the observed velocity cdispersio1.," A number of models with masses $M^{\rm CC}$ = 1.0, 1.5, and 2.0 $\times$ $^{6}$ $_{\odot}$ have velocity dispersions that are within the one sigma error of the observed velocity dispersion."1071 Model 1.0.5250. which has au effective radius aud euclosed mass very close to tle newest observed values f'om aud a surface brightuess profile that is a good app'oximatiou of the observed profile. has a velocity dispersion of 6 = L13 kms |. Le. almost exactly the observed ole.," Model 50, which has an effective radius and enclosed mass very close to the newest observed values from \cite{baumgardt} and a surface brightness profile that is a good approximation of the observed profile, has a velocity dispersion of $\sigma$ = 4.13 km $^{-1}$, i.e. almost exactly the observed one."1072 C'ouskleriug masses. effective radii. surface brightness profies. aud velocity dispersions. model NLILII.0.550 provides the best representation of 22119.," Considering masses, effective radii, surface brightness profiles, and velocity dispersions, model 50 provides the best representation of 2419."1073 However. a uunber of models reproduce the observed structural parameters of 22119 cuite well with the observatioual uucertaiuties. demonstratiug that au object like NOC22119 can be formed from merged. CCs without the need of fine-tuning of the input. parameters.," However, a number of models reproduce the observed structural parameters of 2419 quite well within the observational uncertainties, demonstrating that an object like 2419 can be formed from merged CCs without the need of fine-tuning of the input parameters."1074 The proposed formation scenario for. NCC22119 starts with newly. born complexes of star clusters in the Galactic halo with orbital parameters allowing for au highly eccentric orbit., The proposed formation scenario for 2419 starts with newly born complexes of star clusters in the Galactic halo with orbital parameters allowing for an highly eccentric orbit.1075 We model the dynamical evolution of various CCs leading to merger objects., We model the dynamical evolution of various CCs leading to merger objects.1076 We do not. however. consider the process which formed the CCs in the first place as this wouldiucrease the complexity ol the sunulatious and add more degrees of Ireecdoum makiug the interpretation of the results difficult.," We do not, however, consider the process which formed the CCs in the first place as this would increase the complexity of the simulations and add more degrees of freedom making the interpretation of the results difficult."1077A starburst is an intense period of star formation within a galaxv that cannot be sustained over ils lifetime. due to the huge amount of gas needed to fuel the process.,"A starburst is an intense period of star formation within a galaxy that cannot be sustained over its lifetime, due to the huge amount of gas needed to fuel the process."1078 Typical starburst lifetimes are 105 to LO” vears., Typical starburst lifetimes are $^8$ to $^9$ years.1079 Starbursts can occur on (wo scales: centralized in the ealactic nucleus (a nuclear starburst) or throughout the entire galaxy (a global starburst)., Starbursts can occur on two scales: centralized in the galactic nucleus (a nuclear starburst) or throughout the entire galaxy (a global starburst).1080 In order to (rigger a starburst. there must be a plentiful supply of eas to the centre of the galaxy. which occurs when a large-scale disturbance to the angular momentum of the central region causes the infall of enough material.," In order to trigger a starburst, there must be a plentiful supply of gas to the centre of the galaxy, which occurs when a large-scale disturbance to the angular momentum of the central region causes the infall of enough material."1081 In a global starburst this disturbance is usually associated with the collision (and possible merger) of (wo galaxies., In a global starburst this disturbance is usually associated with the collision (and possible merger) of two galaxies.1082 li a nuclear starburst (such as M32) the disturbance may be due to the effects of shock Ivonis causecl by a stellar bar. a tidal interaction wilh another galaxy. or perhaps an active galactic nucleus (AGN).," In a nuclear starburst (such as M82) the disturbance may be due to the effects of shock fronts caused by a stellar bar, a tidal interaction with another galaxy, or perhaps an active galactic nucleus (AGN)."1083 The most luminous starbursts appear to all be triggered by interactions with other ealaxies., The most luminous starbursts appear to all be triggered by interactions with other galaxies.1084 Smith. Herter and Haynes (1993) studied a sample of 20 radio-Iuminous starburst ealaxies.," Smith, Herter and Haynes (1998) studied a sample of 20 radio-luminous starburst galaxies."1085 They found that all 20 galaxies were mergers. interacting pairs. or members of eroups or clusters. suggesting that all the starbursts in the sample were most likely triggered by interactions with other galaxies.," They found that all 20 galaxies were mergers, interacting pairs, or members of groups or clusters, suggesting that all the starbursts in the sample were most likely triggered by interactions with other galaxies."1086 It has become clear that a “pan-spectval” approach to observations is necessary (o comprehensively study. starburst galaxies. combining observations taken over as large a wavelength range as possible. and supported by theoretical advances and modeling.," It has become clear that a “pan-spectral"" approach to observations is necessary to comprehensively study starburst galaxies, combining observations taken over as large a wavelength range as possible, and supported by theoretical advances and modeling."1087 However. parüceularlv as the more compact and Iuminous nuclear starbursts are enshrouded in dust clouds which are optically thick at shorter wavelengths. long wavelength data are uniquely able to peer through the curtain of extinction and derive (he basic properties of the stellar population.," However, particularly as the more compact and luminous nuclear starbursts are enshrouded in dust clouds which are optically thick at shorter wavelengths, long wavelength data are uniquely able to peer through the curtain of extinction and derive the basic properties of the stellar population."1088 In some cases. IR. data is sufficient. but. where even IR radiation is strongly absorbed. it is necessary to (turn to the mm and radio domains.," In some cases, IR data is sufficient, but where even IR radiation is strongly absorbed, it is necessary to turn to the mm and radio domains."1089 The IR luminosity is a widely used diagnostic of starburst activity., The IR luminosity is a widely used diagnostic of starburst activity.1090 This strong emission comes Irom cust eraius presumed {ο be heated by photons from voung stars Iormed in the, This strong emission comes from dust grains presumed to be heated by photons from young stars formed in the10910.3in 0.211 Figure 1.,0.3in 0.2in Figure 1.1092 Iuteerated intensity maps of the six PO QSOs detected in CO(L>0)., Integrated intensity maps of the six PG QSOs detected in $1\to0$ ).1093 a} PG. 0838|770 - coutours are plotted as lo«(.1.6.1.6.2.6.3.6.16.5.6.6.6).," a) PG 0838+770 - contours are plotted as $\sigma1094\times (-1.6, 1.6, 2.6, 3.6, 4.6, 5.6, 6.6)$."1095 The peal intensity is (017. Jy. beam1 and corresponds to the position RAH=Os:1115.22 dec|76:53:08.96. (72000.0)., The peak intensity is 0.017 Jy $^{-1}$ and corresponds to the position RA=08:44:45.22 dec=+76:53:08.96 (J2000.0).1096 b) PG 11191120 - contours are plotted as lo«(2.3.2.3.3:3.1.3.5.3).," b) PG 1119+120 - contours are plotted as $\sigma \times1097(-2.3, 2.3, 3.3, 4.3, 5.3)$."1098 The peak intensity is 0.011 Iv beau1. and corresponds to the position RA-11:21:17.12 dec=|11:LET8.30. (12000.0).," The peak intensity is 0.014 Jy $^{-1}$, and corresponds to the position RA=11:21:47.12 dec=+11:44:18.30 (J2000.0)."1099 c) PG 13511610 - contours are plotted as lo<(1.6.1.6.2.6.3.6.1.6): the peak intensity is 0.0055 Jy beamLand corresponds to the position RA=13:53:15.62 dec=|63:15:15.72 (J2000.0).," c) PG 1351+640 - contours are plotted as $\sigma \times1100(-1.6, 1.6, 2.6, 3.6, 4.6)$; the peak intensity is 0.0055 Jy $^{-1}$, and corresponds to the position RA=13:53:15.62 dec=+63:45:45.72 (J2000.0)."1101 d) PC 11151151 - coutours are plotted as lo.(2.2.2.2.5.2.1.2.5.2): the peal intebsity is 0.051 Jv 5. and corresponds to the position RA=LL17:00.76 dec-|LLESG6.50 (J2000.03.," d) PG 1415+451 - contours are plotted as $\sigma \times1102(-2.2, 2.2, 3.2, 4.2, 5.2)$; the peak intensity is 0.051 Jy $^{-1}$, and corresponds to the position RA=14:17:00.76 dec=+44:56:06.50 (J2000.0)."1103 0) PG 111012356the - contours are plotted as lao«(.2.3.2.3.ολοι1.3.5.3.6.3): the peak tensity is (016 Jy J|. and corresponds to position RA=11:12:07.18 dec=|35:26:22.33 (J2000.0).," e) PG 1440+356 - contours are plotted as $\sigma \times1104(-2.3, 2.3, 3.3, 4.3, 5.3, 6.3)$; the peak intensity is 0.016 Jy $^{-1}$, and corresponds to the position RA=14:42:07.48 dec=+35:26:22.33 (J2000.0)."1105" Ὁ PC 16131658 - contours are plotted as lo«(/—1.6.2.66.6.7.6.8.6): the peak intensity is 0.012 Jy 1 aud corresponds to the position RA=16:13:57.15 dec=|65:13:09.623.0,1.0. ooeed(J2000.0)."," f) PG 1613+658 - contours are plotted as $\sigma \times1106(-1.6, 1.6, 2.6, 3.6, 4.6, 5.6, 6.6, 7.6, 8.6)$; the peak intensity is 0.012 Jy $^{-1}$, and corresponds to the position RA=16:13:57.15 dec=+65:43:09.62 (J2000.0)."1107 Olin Figure 2., 0.1in Figure 2.1108 COGL> 0) spectra of the SiN de PG QSOs., $1\to0$ ) spectra of the six detected PG QSOs.1109 For PG 0838|770 aud PC 11151151. the data have been smoothed to a resolution of 16 MITIz with a 8 MIIz sampling.," For PG 0838+770 and PG 1415+451, the data have been smoothed to a resolution of 16 MHz with a 8 MHz sampling."1110 For PC 11191120. the data have been siioothed. to a resolution of 32 MIIz with a 16 MITz sampling.," For PG 1119+120, the data have been smoothed to a resolution of 32 MHz with a 16 MHz sampling."1111 For PG 123511610. the data have been smoothed to a resolutiou of LO AIIIz with a 20 MITz sampling.," For PG 1351+640, the data have been smoothed to a resolution of 40 MHz with a 20 MHz sampling."1112 For PC 111012356 aud PG 1613|658. the data have been smoothed to a resolution of 21 Mz with a 12 MIIz sampling.," For PG 1440+356 and PG 1613+658, the data have been smoothed to a resolution of 24 MHz with a 12 MHz sampling."1113 O.lin Figure 3., 0.1in Figure 3.1114 A plot of the CO fiux vs. the 100 pan flux density for PG OSOs aud ULICs detected to date in CO., A plot of the CO flux vs. the 100 $\mu$ m flux density for PG QSOs and ULIGs detected to date in CO.1115 The CO data for IZwl iud Mrk 1014 were obtained from Barvainis et al. (, The CO data for IZw1 and Mrk 1014 were obtained from Barvainis et al. (11161989) and Solomon et al. (,1989) and Solomon et al. (11171997). respectively.,"1997), respectively."1118 The ULICs CO data were obtained from Sanders et al. (, The ULIGs CO data were obtained from Sanders et al. (11191989b). Sanders. Scoville. Soifer (1988. 1991). Solomon et al. (,"1989b), Sanders, Scoville, Soifer (1988, 1991), Solomon et al. ("11201997). and Evans et al. (,"1997), and Evans et al. ("11211999).,1999).1122 The 100722 flux deusities were obtained frou Saucers et al. (, The $\mu$ m flux densities were obtained from Sanders et al. (112319592. 1991) and Solomon et al. (,"1989a, 1991) and Solomon et al. ("11241997).,1997).1125 Arrows denote 36 upper limits on the CO luminosity of PC 1126-011. PG 1202|281. and PC 1102|261.," Arrows denote $\sigma$ upper limits on the CO luminosity of PG 1126-041, PG 1202+281, and PG 1402+261."1126 Ouly upper limits on Zt4 aud figo curently exist for PC 0007|106. aud thus PC 0007|106 is not included on the plot.," Only upper limits on $L'_{\rm CO}$ and $f_{100}$ currently exist for PG 0007+106, and thus PG 0007+106 is not included on the plot."1127 Adapted from Figure. lof Solomon et al. (, Adapted from Figure 4 of Solomon et al. (11281997).,1997).1129 O.lin Figure L., 0.1in Figure 4.1130 Α plot of logLog versus redshift (2) for the low-: PG QSO sample. a fiux-linited sample (1ου>5.21 Jv) of iufrared huninous galaxies aud a sample of ultrahuninous infrared galaxies.," A plot of $\log L'_{\rm CO}$ versus redshift $z$ ) for the $z$ PG QSO sample, a flux-limited sample $f_{60} > 5.24$ Jy) of infrared luminous galaxies and a sample of ultraluminous infrared galaxies."1131 The vertical dashed lines represeut the upper and lower 2 boundaries of the PO QSO sample., The vertical dashed lines represent the upper and lower $z$ boundaries of the PG QSO sample.1132 The LIC and ULICs data have been obtained from the same sources as in Fieure 3., The LIG and ULIGs data have been obtained from the same sources as in Figure 3.1133 Arrows denote 36 upper lanits on the CO Lhnunuinositv of PC 1126-011. PC 1202|281. and PC 1102|261.," Arrows denote $\sigma$ upper limits on the CO luminosity of PG 1126-041, PG 1202+281, and PG 1402+261."1134 Q.lin Figure 5., 0.1in Figure 5.1135 à) A plot of Zi vs. Lt; for the low-z: QSO sample. a fux-Iimited sample 275.21 Jy) of infrared huninous galaxies and a sample of ultralumiuous infrared galaxies.," a) A plot of $L_{\rm IR}$ vs. $L'_{\rm CO}$ for the $z$ QSO sample, a flux-limited sample $f_{60 \mu{\rm m}} > 5.24$ Jy) of infrared luminous galaxies and a sample of ultraluminous infrared galaxies."1136 Arrows denote 30 upper (foolimits on the CO huninosity of PG 1126-011. PC 1202|281. and PG ΙΟΣ|261.," Arrows denote $\sigma$ upper limits on the CO luminosity of PG 1126-041, PG 1202+281, and PG 1402+261."1137 b) A plot of Lin/Log vs. Eqs for the low-z QSO sample. a flux-limited saluple (fion>5.21 Jy) of infrared Iuninous galaxies and a sample of ültraluuiuous infrared. galaxies.," b) A plot of $L_{\rm IR}/L'_{\rm CO}$ vs. $L_{\rm IR}$ for the $z$ QSO sample, a flux-limited sample $f_{601138\mu{\rm m}} > 5.24$ Jy) of infrared luminous galaxies and a sample of ultraluminous infrared galaxies."1139 Arrowsdenote Jo lower limits on Lig[ο of PG 1126-011. PG 1202|281. and PG 1102|261.," Arrowsdenote $\sigma$ lower limits on $L_{\rm IR}/L'_{\rm CO}$ of PG 1126-041, PG 1202+281, and PG 1402+261."1140 For simplicity. all of the cool infrared huninous galaxies are plotted as plus sigue.," For simplicity, all of the cool infrared luminous galaxies are plotted as plus signs."1141 0.1iu Figure 6., 0.1in Figure 6.1142" A plot of logLey, versus redshift (+) for the low-. PG QSO sample and the high-redshift QSOs tected iu CO to date.", A plot of $\log L'_{\rm CO}$ versus redshift $z$ ) for the $z$ PG QSO sample and the high-redshift QSOs detected in CO to date.1143 The moderate and οντοdshift QSO data havebeen obtained from the following references: 3€ 18dct (2=0.37: Scoville et al., The moderate and high-redshift QSO data have been obtained from the following references: 3C 48 $z = 0.37$: Scoville et al.1144 1993). T1113|117 (2=¢2.56:n Barvainis1900). et al.," 1993), H1413+117 $z =11452.56$: Barvainis et al."1146 199D. ATC OLLEE)O531 (2=2.61: Barvainis et al.," 1994), MG 0414+0534 $z=2.64$: Barvainis et al."1147 1998). APM 08279|5255 (1=3.91: Downes sot ," 1998), APM 08279+5255 $z=3.91$: Downes et al."1148BRI 1335-0115 --Ll: Cuilloteau et al.," 1999), BRI 1335-0415 $z=4.41$: Guilloteau et al."1149 1997). and BR 1202-0725 (2=L69: Omonut et al.," 1997), and BR 1202-0725 $z=4.69$: Omont et al."1150 1996)., 1996).1151 For gravitationally lensed QSOs. Li is plotted in terms of both the observed. value aud the iutrinsic value.," For gravitationally lensed QSOs, $L'_{\rm CO}$ is plotted in terms of both the observed value and the intrinsic value."1152 For MG 0141110531. the amplification is unknown. thus MG O04L4LL110531 is plotted with a solid line extending downward from the observed Lig.," For MG 0414+0534, the amplification is unknown, thus MG 0414+0534 is plotted with a solid line extending downward from the observed $L'_{\rm CO}$ ."1153 Adapted from Figure 3 of Fraver ct al. (, Adapted from Figure 3 of Frayer et al. (11541999).,1999).1155where If the Iuminositv function defined in Eq.(12)) has slope a=—1 the total flux [rom resolved and unresolved PWNe divided over the HESS field of view if 2.5 per cent of radio loud pulsars enit 5-ravs is The total [lux from the Galactic SNR. population is For a=-1 the integral flux expected lor the Milagro field of view [rom IIESS-like sources assuming a pulsar shaped surface densitv if 2.5 per cent of radio loud pulsars emit -FAVs Is while the contribution Lom SNRs shaped surface density is The fluxes in Eq. (19)),where If the luminosity function defined in \ref{eqn:luminosity}) ) has slope $\alpha=-1$ the total flux from resolved and unresolved PWNe divided over the HESS field of view if 2.5 per cent of radio loud pulsars emit $\gamma$ -rays is The total flux from the Galactic SNR population is For $\alpha=-1$ the integral flux expected for the Milagro field of view from HESS-like sources assuming a pulsar shaped surface density if 2.5 per cent of radio loud pulsars emit $\gamma$ -rays is while the contribution from SNRs shaped surface density is The fluxes in Eq. \ref{eqn:terzabis}) )1156 and (20)) are lower than the HESS fluxes from resolved aud unresolved sources., and \ref{eqn:quarta}) ) are lower than the HESS fluxes from resolved and unresolved sources.1157 In fact. as shown in Fig.3.. ESS observes the inner region of the Galaxy. whereas Milagro field of view is more spread toward (he outer regions in (he Galaxy. where the number of sources is substantially lower. so a lower contribution from uiresolved sources is expected for Milagros region of the Galactic plane than for HESS.," In fact, as shown in \ref{fig3}, HESS observes the inner region of the Galaxy, whereas Milagro field of view is more spread toward the outer regions in the Galaxy, where the number of sources is substantially lower, so a lower contribution from unresolved sources is expected for Milagro's region of the Galactic plane than for HESS."1158cylindrical radius alouc. D.=(5). but the axial iaenetic field D. is constant iu our model.,"cylindrical radius alone, $B_{\varphi}= B_{\varphi}(s)$, but the axial magnetic field $B_z$ is constant in our model."1159 We study the behaviour of ΑΠΟ modes in the incompressible linut., We study the behaviour of MHD modes in the incompressible limit.1160 This approximation is well justified for modes with the characeristic time-scale longer than the ooriod. of sound. waves àuid for subsonic notions (see. e.g.àY Landau Lifshitz 198]j.," This approximation is well justified for modes with the characteristic time-scale longer than the period of sound waves and for subsonic motions (see, e.g., Landau Lifshitz 1981)."1161 Poerturbalous of the deusityv caused by these motions are ια. aud. cau be uceleced in MIID equatiois., Perturbations of the density caused by these motions are small and can be neglected in MHD equations.1162 By making use of the incompressile lut iua magnIDsed eas; one can cosider slow modes with the erowth rate (or frequency) lower than the frequeney| of fast maeuetosoniecS waves.," By making use of the incompressible limit in a magnetised gas, one can consider slow modes with the growth rate (or frequency) lower than the frequency of fast magnetosonic waves."1163 Since the typical time scale of nodes uuder study is of the order of the inverse Alfvénn frequency. our consideration applies if the Alfvéuu velocity is lower than the sound speed or. in terms of the plasma. ;-parameter. if.m1 )j is the ratio of the eas and maguctic pressures).," Since the typical time scale of modes under study is of the order of the inverse Alfvénn frequency, our consideration applies if the Alfvénn velocity is lower than the sound speed or, in terms of the plasma $\beta$ -parameter, if $\beta \gg 1$ $\beta$ is the ratio of the gas and magnetic pressures)."1164 Note that the velocity of jet V can be much higher than the Alfvéóunu aud sound speed., Note that the velocity of jet $V$ can be much higher than the Alfvénn and sound speed.1165 Iu the incompressible lait. the MIID equations read Iun the basic state. the eas is assumed to be iu hydrostatic equilibrium in the racial direction. theu Stability will VD-qANOBbe studied. by maine use of a linear perturbative analysis.," In the incompressible limit, the MHD equations read In the basic state, the gas is assumed to be in hydrostatic equilibrium in the radial direction, then Stability will be studied by making use of a linear perturbative analysis."1166 Because the basic state is stationary and axisvunnetric. the dependence of perturbations on f. y. and + can be taken in the form exp(otthotung) where fk. is the wavevector in the axial direction aud m is the azimuthal wavemuuber.," Because the basic state is stationary and axisymmetric, the dependence of perturbations on $t$, $\varphi$ , and $z$ can be taken in the form $\exp{(\sigma t - i k_z z - i m \varphi)}$ where $k_z$ is the wavevector in the axial direction and $m$ is the azimuthal wavenumber."1167 Small perturbations will be indicated by subscript 1. while unperturbed quautitics will have no subscript.," Small perturbations will be indicated by subscript 1, while unperturbed quantities will have no subscript."1168 Then. the liearized Eqs. (," Then, the linearized Eqs. ("11691)-(1) read Eliminating all variables in favour of the radial velocity perturbation ¢y.. we obtain This eqation was first derived by Bonanno Urpin (20081) in treir analysis of the non-axisvuuuaetrie stability of stellar uaenetic configurations.,"1)-(4) read Eliminating all variables in favour of the radial velocity perturbation $v_{1s}$, we obtain where This equation was first derived by Bonanno Urpin (2008b) in their analysis of the non-axisymmetric stability of stellar magnetic configurations."1170 For axisvuuuetric )orturbatiois (Q0=0). Eq. (," For axisymmetric perturbations $m=0$ ), Eq. ("117110) recovers Eq. (,10) recovers Eq. (117211) of the per by Bonanno Urpin (2008a).,11) of the paper by Bonanno Urpin (2008a).1173 Tustahiliics of the magnetic configurations associated o the electric current are basically absolue Iustabilities. i.c. they erow but do not propagate.," Instabilities of the magnetic configurations associated to the electric current are basically absolute instabilities, i.e. they grow but do not propagate."1174 We suppose tiat this is the case also in the rest frame of the jet. aud iustable )orturbations are therefore simply advectcc| with the flow att he jet velocity.," We suppose that this is the case also in the rest frame of the jet, and unstable perturbations are therefore simply advected with the flow at the jet velocity."1175 Equation (10) represeuts a lol-ear eleeuvalue problem for c. which cau be xved oX‘© the )nmdary conuditious are given.," Equation (10) represents a non-linear eigenvalue problem for $\sigma$, which can be solved once the boundary conditions are given."1176 We asse hat (44 s1o0uld o finite at the t ands., We assume that $v_{1s}$ should be finite at the jet axis.1177" As far as the outer bouudarv ds concerned, it i newhat difficult to formuate ap ausible )onndary co1 because :vetuallv there is no bouidirv )otween the j id the ailvent medium."," As far as the outer boundary is concerned, it is somewhat difficult to formulate a plausible boundary condition because actually there is no boundary between the jet and the ambient medium."1178 Likely. t1ο jet is separated he zuubieif plasma by he shear laver lia has a fi uckness.," Likely, the jet is separated from the ambient plasma by the shear layer that has a finite thickness."1179 The effect of this thickness ou he growth rate «justabilities las been studied by BaY (2005) who fornd hat the results are only slightly affect«xl wea change of ιο thickness., The effect of this thickness on the growth rate of instabilities has been studied by Baty (2005) who found that the results are only slightly affected by a change of the thickness.1180 For the sake of simplicitY. herefore. we assume that the shear laver is iufiuiteY hin (see aso Appl et al.," For the sake of simplicity, therefore, we assume that the shear layer is infinitely thin (see also Appl et al."1181 200D., 2000).1182 In asi perimagnuetoson jet. uo signal can propagate from the jet interior fo 1 sumroundings. aid we can expect that the iustabilitic )oliave as 1 the jet is bounded by a rig colleποπιο wa," In a supermagnetosonic jet, no signal can propagate from the jet interior to its surroundings, and we can expect that the instabilities behave as if the jet is bounded by a rigid conducting wall."1183" ""Therefore. we ca απλής the outer boundary by supposing CLs=0) at the jet radius s=,."," Therefore, we can mimic the outer boundary by supposing $v_{1s}=0$ at the jet radius $s=s_1$."1184 Note tthat we also ried other outer boundary conditions (or example. with ty. £0). but this docs not change the results qualitatively.," Note that we also tried other outer boundary conditions (for example, with $v_{1s} 1185\neq 0$ ), but this does not change the results qualitatively."1186 We can represent the azimuthal maeuetic field as where By is the characteristic field streugth aud c l., We can represent the azimuthal magnetic field as where $B_{\varphi 0}$ is the characteristic field strength and $\psi \sim 1$ .1187 To calculate the erowth rate of the instability. it is convenieut to introduce dimensionless (quantities Then. Eq. (," To calculate the growth rate of the instability, it is convenient to introduce dimensionless quantities Then, Eq. ("118810) reads where We thus solve Eq. (,10) reads where We thus solve Eq. (118913)for a wide rauge of the pariuneters and different functional dependence eC).,13)for a wide range of the parameters and different functional dependence $\psi(x)$ .1190racius of 0.2554. however. the and smooth matter clistributions are virtually identical. and for larger radii the smooth matter distribution is found to be narrower than the smooth matter cistribution.,"radius of $0.25\eta_0$, however, the and smooth matter distributions are virtually identical, and for larger radii the smooth matter distribution is found to be narrower than the smooth matter distribution."1191 We next compare the observed. elefel Dux ratio for Ale: 041410534. with our model probability distributions and construct an a posteriori. probability distribution for characteristic source radius and smooth matter percentage., We next compare the observed $A_2/A_1$ flux ratio for MG 0414+0534 with our model probability distributions and construct an a posteriori probability distribution for characteristic source radius and smooth matter percentage.1192 Following WMS95 and SWO2 we conduct. our analysis using an observed Dux ratio of Ron.=(Aofle.O45d0.06 (Schechter&Moore1993)., Following WMS95 and SW02 we conduct our analysis using an observed flux ratio of $R_{obs} = (A_2/A_1)_{obs} = 0.45 \pm 0.06$ \citep{sm93}.1193 ὃν comparing this observed Hux ratio with conditional probability distributions for the Hux ratio. we constructed likelihoods for the observed. ratio elven varius source raclii j. in units of the Einstein Racius gj. and smooth matter percentages s=K/h.," By comparing this observed flux ratio with conditional probability distributions for the flux ratio, we constructed likelihoods for the observed ratio given varius source radii $\eta$, in units of the Einstein Radius $\eta_0$, and smooth matter percentages $s = \kappa_c/\kappa_{tot}$."1194" Using Baves’ theorem. this likelihood LI,5.1) was converted to an a posteriori differential probability distribution for smooth matter percentage and source radius as a fuction of e."," Using Bayes' theorem, this likelihood $L(R_{obs}|s, \eta)$ was converted to an a posteriori differential probability distribution for smooth matter percentage and source radius as a fuction of $R_{obs}$."1195 These distributions are then mareinalised over the observed distribution for Dux ratio. where the error. in the [lux ratio was treated as a Gaussian with characteristic radius equal to. the observational error of £0.06.," These distributions are then marginalised over the observed distribution for flux ratio, where the error in the flux ratio was treated as a Gaussian with characteristic radius equal to the observational error of $\pm 0.06$."1196 We used a logarithmic Bayesian prior lor source radius. and a constant Bavesian prior for smooth matter percentage. Source size is à quantity with units (in this case. units of Einstein Iacius). and is therefore assumed to have à prior probability that is constant per unit logarithm.," We used a logarithmic Bayesian prior for source radius, and a constant Bayesian prior for smooth matter percentage, Source size is a quantity with units (in this case, units of Einstein Radius), and is therefore assumed to have a prior probability that is constant per unit logarithm."1197 The choice of a prior that is Dat in the logarithm ensures that the ratio of prior probability for two values of source size does. not depend on the units chosen., The choice of a prior that is flat in the logarithm ensures that the ratio of prior probability for two values of source size does not depend on the units chosen.1198 On the other hand. the value of the smooth matter fraction is a climensionless quantity and is therefore assumed to have a uniformi prior.," On the other hand, the value of the smooth matter fraction is a dimensionless quantity and is therefore assumed to have a uniform prior."1199 Probability contours were then drawn. through the resulting distribution and are plotted in Fig. 5.., Probability contours were then drawn through the resulting distribution and are plotted in Fig. \ref{contour}.1200 This shows hat our simulations do not constrain the smooth matter content of the lens., This shows that our simulations do not constrain the smooth matter content of the lens.

Showing the first 1,200 of 10357 lines. Download the file for the rest.